{
    "0812/0812.2721_arXiv.txt": {
        "abstract": " ",
        "introduction": "Direct detection of gravitational waves (GWs) would be of huge importance to the physics and astrophysics communities. Detection would simultaneously verify our basic understanding of gravitational physics and usher in a new era of astronomy. Using models of galaxy merger rates and the properties of coalescing black holes, \\cite{svc+08} and references therein predict the existence of an isotropic, stochastic, low-frequency GW background.  Various models of cosmic strings \\citep{dv01,cbs96} and the inflationary era \\citep{tur97,bb08} also predict the existence of a low-frequency GW background \\citep{mag00}.  In the frequency range accessible by pulsar timing these backgrounds can be described by their characteristic strain spectrum, $h_c(f)$, which can be approximated as \\begin{equation} h_c(f) = A_g\\left(\\frac{f}{f_{\\rm 1 yr}} \\right)^\\alpha \\end{equation} where $f_{\\rm 1 yr} = 1/1$\\,yr.  The dimensionless amplitude of the background, $A_g$, and the spectral exponent, $\\alpha$, depend on the type of background \\citep{jhv+06}. \\cite{saz78} and \\cite{det79} showed that these GW signals could be identified through observations of pulsars.  Standard pulsar timing techniques \\citep[e.g.][]{ehm06,lk05} are used to search for the existence of GWs.  Timing software \\citep{hem06} is used to compare the predictions of a pulsar timing model with measured pulse times of arrival (TOAs).  The timing model for each pulsar contains information about the pulsar's astrometric, spin and orbital parameters. Any deviations of the actual arrival times from the  model predictions - the pulsar timing residuals - represent the presence of unmodelled effects which will include calibration errors, spin-down irregularities and the timing signal caused by GWs.  The signal from GWs can be disentangled from other unmodelled effects by looking for correlations in the timing residuals of a large number of pulsars that are distributed over the entire sky. Pulsar timing experiments are sensitive to GW signals in the ultra-low frequency (f $\\sim 10^{-9}$--$10^{-8}$\\,Hz) band making them complementary to the space- and ground-based interferometers such as the Laser Interferometer Space Antenna (LISA) and the Laser Interferometer Gravitational Wave Observatory (LIGO), which are sensitive to higher frequency GWs (see Figure~\\ref{fg:sens} and \\S\\ref{sec:band}).  \\begin{figure} \\includegraphics[width=6cm,angle=-90]{sens.ps} \\caption{Characteristic strain sensitivity for existing and proposed GW detectors as a function of GW frequency. Predicted signal levels from various astrophysical sources are shown. Figure from Hobbs (2008b).}\\label{fg:sens} \\end{figure}\\nocite{hob08b}  For a given pulsar and GW source, the induced timing residuals caused by a GW signal only depend on the GW characteristic strain at the pulsar and at the Earth.  The GW strains at the positions of well-separated pulsars will be uncorrelated, whereas the component at the Earth will lead to a correlated signal in the timing residuals of all pulsars \\citep{hd83}.  It is this correlated signal that the pulsar timing experiments aim to detect.  The first high-precision pulsar timing array experiments were carried out using the Arecibo and GreenBank radio telescopes \\citep[e.g.][]{fb90}.  Observations of PSRs~B1937+21 and B1855+09 were subsequently used to obtain a limit on the existence of a GW background \\citep[e.g.][]{srtr90,ktr94,mzvl96,lom02}.  Unfortunately, the timing precision presented in these early papers was too poor to attempt to \\emph{detect} the expected GW signals. New pulsar discoveries and better instrumentation has significantly improved the timing precision achievable.  Recently \\cite{vbv+08} showed that PSR~J0437$-$4715 could be timed using the Parkes radio telescope with an rms timing residual of only 200\\,ns over a 10\\,yr period. Timing precisions are continuing to improve.  \\cite{lb01} described how pulsar data sets could be used to constrain the existence of binary black holes in the centre of our Galaxy and in close-by galaxies. In 2003, a VLBI experiment detected motions in the radio galaxy 3C66B that were thought to be evidence of a binary supermassive black hole system \\citep{simt03}.  \\citet{jllw04} determined the expected PSR~B1855+09 timing residuals that would occur due to GW emission from the postulated system and compared them with archival Arecibo data. The pulsar timing observations allowed the existence of the binary system to be ruled out with 95\\% confidence.  The Parkes Pulsar Timing Array (PPTA) project began in the year 2003.  The three main aims of this project are 1) to detect GW signals, 2) to establish a pulsar-based time scale for comparison with terrestrial time standards and 3) to improve the Solar system planetary ephemerides.  In this paper, we describe the work that has already been carried out towards our first aim of GW detection. This includes determining the number of pulsars and observing parameters required to detect a GW signal, obtaining the most stringent constraint on the existence of a GW background and producing high quality data sets for use in GW astronomy. ",
        "conclusions": "Current pulsar timing array experiments have the potential to detect low-frequency GW signals within a decade.  Future experiments with the Square Kilometre Array telescope should allow detailed observations of both a stochastic GW background and individual GW sources.  The pulsar timing technique is complementary to other projects that are attempting to detect much higher frequency GWs."
    },
    "0812/0812.0809_arXiv.txt": {
        "abstract": "{To constrain Randall-Sundrum type braneworld gravity models and the expected rapid evaporation of astrophysical black holes due to the emission of gravitational modes in the extra dimension.} {It is argued that the black-hole X-ray binary XTE J1118+480 is suitable for a constraint on the asymptotic curvature radius of the extra dimension in such braneworld models. An upper limit on the rate of change of the orbital period of XTE J1118+480 is obtained.} {The expected black-hole evaporation in the extra dimension leads to a potentially observable rate of change of the orbital period in XTE J1118+480. The time-change of the orbital period is calculated from previous orbital period measurements from the literature. The lack of observed orbital period evolution is used to constrain the asymptotic curvature radius of the extra dimension.} {The asymptotic AdS radius of curvature is constrained to a value comparable to other limits from astrophysical sources. A unique property of XTE J1118+480 is that the expected rate of change of the orbital period due to magnetic braking alone is so large that only one additional measurement of the orbital period would lead to the first detection of orbital evolution in a black-hole binary and impose the tightest constraint to date on the size of one extra dimension of the order of $35{\\rm \\mu m}$.} {} ",
        "introduction": "The so-called hierarchy problem is one of the key open questions in the search for a grand theory beyond general relativity and the standard model. The fundamental scale of Einsteinian gravity, the Planck scale, differs from the electroweak scale by 16 orders of magnitude, thus making general relativity and the standard model hard to reconcile in a grand unification sceme. Braneworld Gravity, as a candidate for a unified theory, requires the existence of more than three spatial dimensions (see Maartens 2004 for a review), but the size of a potential extra dimension has already been constrained to the sub-mm range by precision tests of Newton's inverse square law (see, e.g., Adelberger, Heckel \\& Nelson 2003 for a review).  One possible solution (Arkani-Hamed, Dimopoulos \\& Dvali 1998) is to demand that extra dimensions be compactified at a scale smaller than those probed by experiment. All standard-model particles are bound to the known four-dimensional world (``the brane''). Only gravity can access the extra dimensions, which accounts for its apparent weakness, because it is spread out through higher-dimensional space (``the bulk''). This approach reconciles the gravitational and the electroweak scales if the extra dimensions are large enough, but it is beyond the scope of astrophysical tests.  A second scenario (Randall \\& Sundrum 1999) embeds the brane in a five-dimensional anti-de Sitter space, which allows for the extra dimension to be infinite. The bulk is filled with a negative cosmological constant such that the extra dimension only effects the brane on a length scale which is small enough and which is set by the asymptotic curvature radius $L$ of the bulk. This model has striking consequences for astrophysical black holes.  Perturbative solution of the classical bulk equations (Tanaka 2003) indicated that no stable black holes can exist on the brane. A treatment of higher-dimensional black holes via the AdS/CFT correspondence showed that black holes are indeed unstable and lose energy in the extra dimension through the emission of CFT modes (Emparan, Garc\\'{i}a-Bellido \\& Kaloper 2003; see, however, Fitzpatrick, Randall \\& Wiseman 2006). The lifetimes of astrophysical black holes are dramatically reduced and can be as short as a Megayear if the asymptotic curvature $L$ of the extra dimension turns out to be in the sub-mm range.  For a binary system consisting of a black hole and a companion star, this effect should be observable and lead to a measurable change of the orbital period (Johannsen, Psaltis \\& McClintock 2009). Several black-hole binaries were identified as candidates, and the system SXT A0620-00 yielded a constraint on the asymptotic curvature radius of $L<161~{\\rm \\mu m}$ at 3$\\sigma$ (Johannsen et al. 2009). Similar limits have been obtained from the age of the black hole XTE J1118+480 ($L<80~{\\rm\\mu m}$; Psaltis 2007) and from tabletop experiments of Newton's inverse square law (Adelberger et al. 2007; Geraci et al. 2008). The current 3$\\sigma$-upper limit on the AdS radius $L$ is of the order of $44~{\\rm\\mu m}$ (Kapner et al. 2007).  In this paper, I extend the analysis of Johannsen et al. (2009) to the black-hole binary XTE J1118+480 and compute an additional constraint on the asymptotic curvature radius $L$. In \\S 2, I briefly review the results of Johannsen et al. (2009) for the evolution of the orbital period of a black-hole binary system with non-conservative mass transfer and black-hole evaporation. In \\S 3, I apply this result to the black-hole binary J1118+480 and I obtain a constraint on the asymptotic curvature radius of $L<0.97~{\\rm mm}$ in \\S 4. ",
        "conclusions": "\\begin{enumerate} \\item The black-hole binary XTE J1118+480 is well-suited for a constraint on the asymptotic curvature radius of the extra dimension in Randall-Sundrum type braneworld gravity models. \\item An upper limit on the rate of change of the orbital period of XTE J1118+480 is calculated based on previous measurements of the orbital period from the literature. The lack of observed orbital evolution of this binary imposes a constraint on the asymptotic curvature radius of $L<0.97~{\\rm mm}$. \\item This constraint can be significantly improved by only one additional measurement of the orbital period of XTE J1118+480. This would be the first detection of orbital evolution in a black-hole binary. The expected predominance of magnetic braking would provide the best constraint to date on the asymptotic curvature radius of the extra dimension of the order of $35~{\\rm \\mu m}$. \\end{enumerate}"
    },
    "0812/0812.1514_arXiv.txt": {
        "abstract": "{We present the results of a 500 ksec observation of the Perseus cluster with \\intgr, with the aim of investigating the possible diffuse non-thermal component detected in a previous \\emph{Chandra} observation. In the 3-20 keV band with the JEM-X instrument, we detect the source with high significance and resolve it spatially. Above 20 keV with IBIS/ISGRI, we find that the source is point-like, and the cluster could be detected up to 120 keV. From the broad-band ISGRI/JEM-X spectrum, although we detect a non-thermal component, we find that the high-energy flux is variable and is consistent with the extrapolation of the 2-10 keV flux of the central AGN, NGC 1275. The extrapolation of the non-thermal component claimed from \\emph{Chandra} data exceeds the \\intgr\\ spectrum by a factor of 3.} ",
        "introduction": "The Perseus cluster Abell 426 ($z=0.0176$) is the brightest galaxy cluster in the X-ray band, and the prototype of cooling-core clusters \\citep{fabian}. Its central cD galaxy, NGC 1275, hosts a well-known narrow-line radio galaxy, Per A, which interacts with the intra-cluster gas through its jets and outflows \\citep{boehringer}. In the X-ray band, the cluster was the target of a very long (900 ksec) \\cra\\ observation \\citep{fabch1}, which revealed a very complex structure, with a temperature ranging from 2.5 keV in the central regions up to $\\sim8$ keV in the outskirts. Several structures associated with the propagation of high-energy particles injected by the central AGN in the thermal plasma (X-ray cavities, sound waves) were also detected.\\\\  Apart from the very deep \\emph{Chandra} observations, the cluster was also observed by \\xmm\\ \\citep{churazov}. Thanks to the larger FOV of \\xmm, it was also possible to observe the structure of the gas in the outer regions. An asymmetry of the surface brightness profile in the east-west direction was found, possibly corresponding to a small group of galaxies falling onto the main cluster. The contribution of the active nucleus in the center of NGC 1275 was also estimated. The spectrum of the AGN could be well-fitted by an absorbed power law with a photon index $\\Gamma=1.65$ typical of radio-loud AGN and a luminosity of $10^{43}$ ergs $s^{-1}$ in the 0.5-8 keV band.\\\\  In the radio domain, the bright radio source 3C 84 is consistent with the position of NGC 1275. While most of the radio emission comes from the radio galaxy, emission on a larger scale ($\\sim$10 arcmin) was also detected in the cluster (\\citet{burns}, \\citet{ferettiper}), designated as the Perseus ``mini-halo\". This implies the presence of relativistic electrons, whose origin is not yet fully understood. If the AGN jets are hadronic, the radio emission probably comes from secondary electrons produced by interactions between cosmic-ray protons and thermal ions \\citep{pfrommermh}. Alternatively, the electrons could also be directly accelerated through turbulence induced by the central nucleus \\citep{gitti}, although the shocks produced by the interaction of the AGN jets with the ICM do not seem to be strong enough.\\\\  Since the existence of relativistic electrons in the cluster is obvious because of the radio synchrotron emission, we expect that photons of the CMB as well as optical/IR photons coming from the central galaxy should be up-scattered to higher energies, in particular to the (hard) X-ray domain (see e.g. \\citet{sarazin}). An excess of emission compared to the thermal emission was probably detected in the Coma \\citep{fusco} and A2256 \\citep{fusco2256} clusters by \\bps\\ and in the Ophiuchus cluster \\citep{eckertoph} by \\intgr. Thanks to the very long \\emph{Chandra} observation of the cluster, such a component has been claimed to be detected in addition to the thermal emission (\\citet{sanders05}, \\citet{sanders07}). This allowed the authors to present a magnetic field map of the cluster and derive a steep magnetic field profile, ranging from $\\sim3\\,\\mu G$ in the center down to $\\sim0.1\\,\\mu G$ in the outskirts. However, a recent \\xmm\\ observation did not confirm the result \\citep{molendi}, so observations of the cluster above 10 keV where the thermal emission becomes weaker are required to confirm this result. \\\\  In the hard X-ray band, \\citet{nevalainen} presented \\bps/PDS observations of the cluster, and concluded that the emission was consistent with the extrapolation of the AGN flux. However, since the PDS instrument was non-imaging, it was not possible to put any constraints on the diffuse emission. Recently, \\citet{ajello} found the same result using data from the \\emph{Swift} satellite.\\\\  \\begin{figure*} \\centerline{\\hbox{\\includegraphics[width=8cm]{1154fig1a.eps} \\includegraphics[width=8cm]{1154fig1b.eps}}} \\caption{Left: ISGRI 20-30 keV significance image of the Perseus cluster region. For comparison, the contours show the JEM-X 3-7 keV emission. The position of the radio galaxy NGC 1275 is also displaid. Right: JEM-X significance image of the cluster in the 3-7 keV band with contours from \\emph{ROSAT}/PSPC. The green circle shows the position of NGC 1275 at the center of the cluster.} \\label{per_int} \\end{figure*}  In this paper, we present the results of a 500 ksec observation of the Perseus cluster with the \\intgr\\ satellite \\citep{Win03}. Thanks to the broad-band coverage and good sensitivity in the hard X-ray range, IBIS \\citep{Ube03} and JEM-X \\citep{lund} have the necessary capabilities to constrain a possible non-thermal component. Indeed, the extrapolation of the diffuse non-thermal flux claimed by \\citet{sanders05} within a radius of 3 arcmin from the center ($6.3\\times10^{-11}$ ergs cm$^{-2}$ $s^{-1}$ in the 2-10 keV band, and a power-law index in the range $1.4-2.2$) is clearly above the sensitivity of IBIS. Therefore, if it is present, \\intgr\\ has the capabilities to confirm or rule out the \\emph{Chandra} result. ",
        "conclusions": "In this paper, we presented the results of a 500 ksec \\intgr\\ observation of the Perseus cluster. We found that \\intgr\\ data can efficiently constrain a possible high-energy tail in the spectrum. Although we found clear evidence for a power-law component above 30 keV, the properties of this power-law are consistent with the extrapolation of the X-ray spectrum from the central nucleus, in agreement with the results from \\bps\\ \\citep{nevalainen} and \\emph{Swift} \\citep{ajello}. Moreover, we have found that the high-energy flux (30-120 keV) significantly varies between the different periods of the observation, covering $\\sim$ 1.5 yr. This clearly demonstrates that the high-energy emission is dominated by the AGN.\\\\  On the other hand, extrapolating the power-law component claimed by \\citet{sanders05} from \\emph{Chandra} data ($6.3\\times10^{-11}$ ergs cm$^{-2}$ $s^{-1}$ in the  2-10 keV band with a photon index $\\Gamma\\sim2.0$) to the ISGRI band, we found that this model clearly over-predicts the high-energy data by a factor of $\\sim$3 in the 40-60 keV band. Therefore, the presence of a flat power-law at such a high level is ruled out by ISGRI data. Overall, we find that the broad-band \\intgr\\ spectrum can be well-described by a two-temperature thermal model (a $\\sim$ 3 keV component from the cool core and a hotter component from the outskirts) plus a variable power-law from the central radio galaxy. The data can also be well-described by a model consisting of a single thermal component, the AGN contribution and a power-law with a high-energy cut-off at an energy $E_{cut}\\sim16$ keV in agreement with the discussion of \\citet{sanders07}. Given that deep X-ray observations clearly indicate the multi-temperature structure of the cluster, the two-temperature model is favoured. Moreover, a longer \\xmm\\ observation \\citep{molendi} did not confirm the \\emph{Chandra} result. The authors claim that the discrepancy between \\emph{Chandra} and \\xmm\\ results is due to uncertainties in the effective area calibration of \\emph{Chandra} in the highest band. \\intgr\\ data are in agreement with this result.\\\\  In conclusion, using \\intgr\\ data, we found no evidence for a diffuse power-law emission which would dominate the emission above 30 keV. Unfortunately, the angular resolution of IBIS/ISGRI is not sufficient to disentangle the point-like emission from the diffuse component, so it is not possible to set any upper limit on the diffuse non-thermal emission. Observations of the cluster in the hard X-ray band with better angular resolution and sensitivity would be very important to detect the diffuse non-thermal component and map the magnetic field strength over the cluster, which could bring a very important input to the understanding of the interactions between the AGN radio lobes and the thermal plasma. In this framework, the future focusing hard X-ray missions (\\emph{Simbol-X}, \\emph{NuSTAR}, \\emph{NeXT}) should be able to detect the diffuse non-thermal component."
    },
    "0812/0812.3511_arXiv.txt": {
        "abstract": "Recently, there are two hints arising from physics beyond the standard model. One is a possible energy loss mechanism due to emission of very weakly interacting light particles from white dwarf stars, with a coupling strength $\\sim 0.7\\times 10^{-13}$, and another is the high energy positrons observed by the PAMELA satellite experiment. We construct a supersymmetric flipped-SU(5) model, \\suf1 with appropriate additional symmetries,  [U(1)$_H]_{\\rm gauge}\\times$[U(1)$_R\\times$U(1)$ _\\Gamma]_{\\rm global}\\times Z_2$, such that these are explained by a very light {\\it electrophilic} axion of mass 0.5 meV from the spontaneously broken U(1)$_\\Gamma$ and two component cold dark matters from $Z_2$ parity. We show that in the flipped-SU(5) there exists a basic mechanism for allowing excess positrons through the charged SU(5) singlet leptons, but not allowing anti-proton excess due to the absence of the SU(5) singlet quarks. We show the discovery potential of the charged SU(5) singlet $E$ at the LHC experiments by observing the electron and positron spectrum. With these symmetries, we also comment on the mass hierarchy between the top and bottom quarks. ",
        "introduction": "The most looked-for particles on Earth at the present time are axions \\cite{KimCarosi} and weakly interacting massive particles(WIMPs) \\cite{highEpartDM}. Recently, hints on these particles have been reported from outer space sources. The 10 year old INTEGRAL data \\cite{INTEGRAL} gives the 511 keV line.  The white dwarf cooling has been suggested by the emission of very weakly interacting light particles(VWLP) (mass less than $\\sim 1$ eV) with a very small coupling strength to electron ($\\sim 0.7\\times 10^{-13}$) \\cite{Isern08}. At present, this white dwarf bound can be considered just as an upper bound, but in this paper we adopt a bold assumption that their best fit corresponds to the existence of a VWLP. More recently, the remarkable observation of high energy positron excess in the satellite PAMELA data \\cite{PAMELAe} hints heavy (100 GeV -- 10 TeV) cold dark matter(CDM) particles \\cite{% Bergstrom08,TwoDMs,Raidal08}. Since the 511 keV line can be explained by the astrophysical origin, or by the particles of mass ${\\cal O}(1-10)~{\\rm MeV}$ \\cite{Partgammaray}, or by an excited state almost degenerate to the WIMP \\cite{Weaner07}, in this paper we attempt to understand the latter two of these observations from a grand unification (GUT) viewpoint of the flipped-SU(5) \\cite{Barr82,DKN84flip}.  The PAMELA group \\cite{PAMELAe} reports the high energy positron excess, above 10 GeV up to 60 GeV, at the level  $e^+/(e^++e^-)\\sim O(0.1)$, but has not reported any noticeable excess of anti-protons \\cite{PAMELAp}. The charactersistic of the data between 10--60 GeV is a slightly rising positron flux, which is inconsistent to the lightest supersymmetric particle(LSP) in the minimal supersymmetric standard model(MSSM) with supersymmetric(SUSY) particles at ${\\cal O}(100~\\rm GeV)$ \\cite{TwoDMs}. Thus, the MSSM with ${\\cal O}(100~\\rm GeV)$  lightest neutralino DM component $\\chi$(\\LN) has to be modified if SUSY has to be kept for the gauge hierarchy solution. The minimal extension needs just three two component fermions $N_R, E_R$ and $E_R^c$ \\cite{TwoDMs}, with the coupling \\begin{equation} W=e_RE_R^cN_R.\\label{eEcN} \\end{equation} If one tries to introduce more fields around TeV scale, certainly the rising positron flux can be explained, but it is very difficult \\cite{Raidal08} to explain nonobservation of antiproton flux. This suggests a different treatment of electron from quarks. Indeed, the flipped-SU(5) GUT \\suf1 treats charged leptons differently from quarks \\cite{Barr82} such that the charged lepton is an \\suf1 singlet, \\begin{widetext} \\begin{equation} \\tent_1=\\left(\\begin{array}{ccccc} 0& d& -d& u^c& -d^c\\\\ -d& 0& d& -u^c& d^c\\\\ d& -d& 0& u^c& -d^c\\\\ -u^c& u^c& -u^c& 0& \\nu_0\\\\ d^c& -d^c& d^c& -\\nu_0& 0 \\end{array} \\right)_R,\\ \\fivebt_{-3}=\\left(\\begin{array}{c} u\\\\ u\\\\ u\\\\  \\overline{e}\\\\ -\\overline{\\nu}_e\\end{array}\\right)_R,\\ \\onet_{5}=e_R~ .\\label{flipRep} \\end{equation} \\end{widetext}  The specific form of the coupling (\\ref{eEcN}) is related to the problems of the Yukawa couplings. In this regard, we note the five old but fundamental problems on the masses of the standard model fermions. Firstly, there is a hierarchy of the top quark mass $m_t\\sim fv/\\sqrt2\\simeq 170$ GeV at the electroweak scale of $v\\simeq 250$ GeV, and the rest of the standard model particle masses are much smaller than the top quark mass, $m\\le m_b\\simeq 4.5~{\\rm GeV}\\ll m_t$. Second, there is a sizable Cabibbo mixing, especially between the first and the second family quarks. Third, the quark mass ratio is reversed in the first family compared to those of the second and the third family quarks, i.e. ${m_u}/{m_d}<1$ while ${m_c}/{m_s}>1$ and ${m_t}/{m_b}>1$. Even though it is not proper to define an inverted mass pattern for leptons, it looks like that the electron mass is also an inverted pattern because ${m_e}/{m_\\mu}\\simeq {m_u}/{4m_c}\\sim {\\cal O}(1/200)$. Fourth, there is a big hierarchy of singlet neutrino mass of $m_{\\nu_0}\\sim {\\cal O}(10^{14})$ GeV for the seesaw mechanism \\cite{Minkowski77} and the electroweak scale of $v\\simeq 250$ GeV. Finally, we point out that the SM neutrino mass pattern shows a tri-bi maximal mixing pattern \\cite{tribiPerkins}. Except the tri-bi maximal mixing, in this paper we try to understand the remaining four problems in a flipped-SU(5) GUT.  To have a very light axion implied by the white dwarf data, we introduce a Peccei-Quinn (PQ) symmetry U(1)$_\\Gamma$ so that the symmetry we consider is extended to \\suf1 times U(1)$_\\Gamma$. To have the inverted mass pattern of the quark mass ratio in the first family, we assign different $\\Gamma$ charges for the first family members and the second and the third family members. But, to have a large Cabibbo mixing, we require Q$_{\\rm em}=\\frac23$ quarks or Q$_{\\rm em}=-\\frac13$ quarks mix fully, restricted only by the strength of the corresponding Yukawa couplings but without any restriction from the symmetry arguments. In the flipped-SU(5), since both the quark singlet $d_R$ and the quark doublet $q_R^c$ appear in ${\\bf 10}_1$, viz. Eq. (\\ref{flipRep}), it is easy to mix Q$_{\\rm em}=-\\frac13$ quarks fully without restrictions from the symmetry arguments. So, we assign the same symmetry charges for three members ${\\bf 10}_1^{(i)}(i=1,2,3)$. Then, Q$_{\\rm em}=-\\frac13$ quarks can be mixed fully as desired. But the charges of $\\overline{5}_{-3}^{(i)}$ and ${1}_{5}^{(i)}$ can be different for different $i$ to allow the inverted mass pattern of the quark mass ratio for the first family. For the singlet neutrino mass problem, we need a large VEV, above the axion window of $10^9-10^{12}$ GeV.  If there is no other gauge symmetry, the PQ symmetry may be broken by this large VEV. Thus, we introduce another U(1) gauge symmetry which will be called a horizontal gauge symmetry U(1)$_H$, distinguishing the first family and the second and third families. Finally, to keep $E_R-E_R^c$ pair at the electroweak scale and to have proton stability and two dark matter components \\cite{TwoDMs}, we introduce the $R$ symmetry and $Z_2$ symmetry (the $R$-parity). Thus, the symmetry we consider is \\begin{equation} SU(5)\\times U(1)_X\\times U(1)_H\\times U(1)_R\\times U(1)_{\\Gamma}\\times Z_2.\\label{symmetry} \\end{equation}  In this model, the renormalizable couplings are gauge couplings and the top quark Yukawa coupling. All the other Yukawa couplings are non-renormalizable. The model we introduce does not have the $c_3$ term of Ref. \\cite{KimCarosi}, where the $c_3$ term is the axion and QCD anomaly coupling, and the PQ charges of Higgs doublets vanish. Therefore, the very light axion we obtain is not a proper Kim-Shifman-Vainstein-Zakharov (KSVZ) axion \\cite{Kim79},  or the Dine-Fischler-Srednicki-Zhitnitskii (DFSZ) axion \\cite{DFSZ81}, but can be called a `variant very light axion'. The electron--axion coupling is known to be small for the KSVZ axion, as calculated in Ref. \\cite{Sredecoupling85}. In our case, the axion arises from the $c_2$ term of Ref. \\cite{KimCarosi}, where the $c_2$ term is the light quark mass term, through a non-renormalizable interaction, and the electron--axion coupling turns out to be large. Since this `very light variant axion' is required to have a large electron coupling, it will be called a very light {\\it electrophilic} axion.  In Sec. \\ref{sec:WD}, we briefly review the white dwarf bound \\cite{Isern08}. In Sec. \\ref{sec:flip}, we present a flipped-SU(5) model, satisfying the criteria presented above. In Sec. \\ref{sec:phenomenology}, we discuss the SUSY phenomenology for the model presented in Sec. \\ref{sec:flip}. Sec. \\ref{sec:conclusion} is a conclusion. ",
        "conclusions": "\\label{sec:conclusion}  The recent PAMELA satellite data from the galactic source suggested the positron excess \\cite{PAMELAe} but no anti-proton excess \\cite{PAMELAp}. Regarding these observations, two dark matter components, the neutralino $\\chi$ and the singlet fermion $N$, were suggested in Ref. \\cite{TwoDMs}. The needed coupling has been given by Eq. (\\ref{eEcN}). Here, we studied the flipped-SU(5) where charged SU(2) singlet leptons such as $e_R, E_R$ and $E_R^c$ of  Eq. (\\ref{eEcN}) can be its allowed representations, but a similar coupling for an anti-proton excess is not allowed because quarks must be embedded in SU(5) non-singlet representations.  In addition, there exists an interesting suggestion of the possible existence of VWLP with the coupling strength (\\ref{aeWDcoupling}) from the study of white dwarf cooling \\cite{Isern08}. In this paper, we tried to interpret this VWLP as an  {\\it electrophilic} axion with the axion decay constant $F_a\\simeq 1.4\\times 10^{10}$ GeV. In the DFSZ model, the corresponding $F_a$ would be $\\simeq 1.2\\times 10^{9}$ GeV which would be barely consistent with the SN1987A bound. This kind of  {\\it electrophilic} axion can arise in `variant very light axion' models where different families possess different PQ quantum numbers $\\Gamma$. We argued that the inverted mass pattern of the first family from that of the second and third families is the logic for assigning different PQ quantum numbers for different family members. Allowing an effective domain wall number of $\\frac12$ [See Footnote \\ref{footDW}.], we constructed an {\\it electrophilic} axion. The symmetries we introduced in (\\ref{symmetry}) is bigger than \\suf1$\\times$U(1)$_\\Gamma$, to allow a GUT scale VEV with an intermediate scale $F_a$ and R-parity. With these symmetries, we explained the mass hierarchy between the top and bottom quarks. In this model, the gauge couplings and the top quark Yukawa coupling are the only renormalizable couplings.  This model needs a charged SU(2) singlet lepton $E$ at the electroweak scale. So, the existence of $E$ is absolutely needed for this suggestion of the positron excess to work. Here, the discovery potential of $E$, at the LHC experiments by observing the electron and positron spectrum, is presented. We calculated the $e^+e^-$ spectrum at the center of momentum frame of the quark and anti-quark pair and compared it with the expected background. We pointed out that the SU(2) singlet $E$ with the coupling (\\ref{eEcN}) is distinguishable from other possibilities we can imagine. Therefore, the discovery of $E$ at the LHC experiments will indirectly confirm the two dark matter component hypothesis of \\cite{TwoDMs}. Or it can be ruled out from the LHC experiments by the absence of the $e^+e^-$ excess above the background.   \\vskip 0.3cm"
    },
    "0812/0812.2101_arXiv.txt": {
        "abstract": "{We present some results on the study of stellar population properties and distances of galaxies using the SBF technique. The applications summarized here show that the Surface Brightness Fluctuations (SBF) method is able to $i)$ provide accurate distances of resolved and unresolved stellar systems from $\\sim$ 10~Kpc to $\\sim$ 150 Mpc, and $ii)$ to reliably constrain the physical properties (e.g. age and metallicity) of unresolved stellar systems. ",
        "introduction": "The knowledge of the complex evolution of the stellar component in high redshift galaxies relies on how detailed is our understanding of the history of stars in nearby galaxies. The best way to trace back the star formation episodes in a galaxy is to study resolved stars. However, this is possible only for few nearby objects, tipically using the best observational facilities to date, and only for the brightest stars in the population.  Given such limitations several techniques have been proposed to disentangle the properties of unresolved stellar systems, one of these is the Surface Brightness Fluctuations method \\citep[SBF hereafter]{ts88}. The SBF technique was introduced as a distance indicator for elliptical galaxies within $\\sim$1-20 Mpc. After more than two decades of systematic applications it is now recognized that the SBF method works in a much larger distance interval, and can be applied to a wider class of objects: ellipticals, bulges of spirals, dwarf ellipticals, globular clusters, etc. In addition, it is well accepted that SBF magnitudes and colors represent a potential tool to analyze in details the physical and chemical properties of unresolved and resolved stellar systems.  In this paper we will briefly discuss some applications of the SBF method done at the INAF-Observatory of Teramo by the SPoT\\footnote{Teramo--Stellar POpulations Tools group website: www.oa-teramo.inaf.it/SPoT} group. ",
        "conclusions": "The SBF technique is to date one of the most reliable distance indicators for elliptical galaxies. However, to our point of view, this technique is underestimated with respect to its real potentiality.  Concerning distance measurements, the SBF method is able to provide distances from few Kpc up to $\\sim$150 Mpc with present observing facilities, and possibly to much larger distances with future instrumentations. Such unique characteristic gives the SBF the potential to cover the distance scale ladder from local to low redshift ($z\\leq0.05$) distances. Thus, SBF measurements, coupled with a reliable calibration of absolute SBF magnitudes, provide a great opportunity to substantially reduce the systematic uncertainty that affects the cosmological distance scale.  With regard to stellar population analysis, there are only few studies dedicated to this topic based on the SBF method. However, it is now evident that SBF can greatly improve our understanding of the properties of unresolved stellar systems. We presented SPoT models in a specific SBF optical-to-near--IR color plane showing the potential of SBF colors to substantially remove the age-metallicity degeneracy. Future application of this technique - possibly coupled with the measure SBF-color gradients attainable with the next generation optical and near--IR large FoV detectors - will provide significant constraints to the knowledge on the formation and evolution of the stellar component in low-redshift galaxies."
    },
    "0812/0812.0797_arXiv.txt": {
        "abstract": "We improve earlier Galactic bounds that can be placed on the fraction of dark matter in charged elemental particles (CHAMPs). These constraints are of interest for CHAMPs whose mass is too large for them to have seen through their electromagnetic interaction with ordinary matter, and whose gyroradius in the galactic magnetic field is too small for halo CHAMPs to reach Earth.  If unneutralized CHAMPs in that mass range are well mixed in the halo, they can at most make up a fraction $\\lesssim (3-7)\\times 10^{-3}$ of the mass of the Galactic halo. CHAMPs might still be a solution to the cuspy halo problem if they decay to neutral dark matter but a fine-tuning is required. We also discuss the case where CHAMPs do not populate a spherical halo. ",
        "introduction": "Massive particles with integer electric charge (CHAMPs), denoted here by $X$, were considered as dark matter candidates in the late eighties (De R\\'ujula et al.~1990; Dimopoulos et al.~1990; Gould et al.~1990; Chivukula et al.~1990). Although stable CHAMPs are predicted in some extensions of the standard model, astrophysical constraints plus bounds from underground detectors, from balloon experiments and the lack of detection of anomalous hydrogen in the sea water, basically rule out CHAMPs as dark matter (see Perl et al.~2001 and Taoso et al.~2008). In particular, the non-detection of heavy water in the sea excludes CHAMPs with masses between $10$ and $10^{4}$ TeV (Verlerk et al.~1992). All the above constraints were derived for the standard flux of particles at Earth from the Galactic halo. If magnetic fields prevent the flux of CHAMPs to penetrate the Galactic disk, one must reevaluate earlier bounds (Chuzhoy \\& Kolb 2008).   Essentially all $X^{-}$ should have bound to protons, forming neutraCHAMPs, which decouple from the photon-baryon fluid and drive structure formation prior to recombination. NeutraCHAMPs reach Earth unimpeded. Searches for neutraCHAMPs in cosmic rays rule out particles with masses between $100$ and a few $10^{4}$ TeV (Barwick et al.~1990). Nevertheless, heavy-water searches, cosmic rays searches, and constraints from overproduction of $^{6}$Li (Berger et al.~2008) are only relevant if CHAMPs are singly charged, because for other charges, a CHAMP no longer behaves like a proton.  If a significant fraction of the mass of halos is made up by CHAMPs, it may have a strong impact on the observable Universe (e.g., Chuzhoy \\& Kolb 2008). It is therefore important to constrain the abundance of CHAMPs in galactic halos. After revising Galactic requirements for CHAMPs to be absent in the Galactic disk, we give an upper limit on the abundance of CHAMPs in the Galactic halo. ",
        "conclusions": "Whilst the common wisdom holds that dark matter is neutral and collisionless, it is important to explore the possibility of it having nonzero, not necessarily integer, charge. We have considered the pressure support of CHAMPs in our Galaxy to derive a simple, upper limit on the fraction of CHAMPs and milliCHAMPs, $f\\lesssim (2-7)\\times 10^{-3}$. If $f$ was roughly constant over time, this constraint rules out CHAMPs as the origin of the cores in LSB and dwarf galaxies. In the range of astrophysical interest, CHAMPs behave like strongly interacting (fluid-like) dark matter (SIDM). Thus, they face many of the problems attributed to SIDM. As some examples, we have discussed the survival of the Magellanic Stream and the mass distribution of the Bullet Cluster. Our constraint that the mass in CHAMPs in the Galaxy is not larger than the mass of coronal gas in the halo seems to apply also to galaxy clusters."
    },
    "0812/0812.0042_arXiv.txt": {
        "abstract": "{Waves connect all the layers of the Sun, from its interior to the upper atmosphere. It is becoming clear now the important role of magnetic field on the wave propagation. Magnetic field modifies propagation speed of waves, thus affecting the conclusions of helioseismological studies. It can change the direction of the wave propagation, help channeling them straight up to the corona, extending the dynamic and magnetic couplings between all the layers. Modern instruments provide measurements of complex patterns of oscillations in solar active regions and of tiny effects such as temporal oscillations of the magnetic field. The physics of oscillations in a variety of magnetic structures of the Sun is similar to that of pulsations of stars that posses strong magnetic fields, such as roAp stars. All these arguments point toward a need of systematic self-consistent modeling of waves in magnetic structures that is able to take into account the complexity of the magnetic field configurations. In this paper, we describe simulations of this kind, summarize our recent findings and bring together results from the theory and observations.} ",
        "introduction": "Any turbulent medium, as the interior of the Sun or stars, generates sound. The basics of the theory excitation of sound waves by the turbulent flow were developed by Lighthill in 1952 \\cite{Lighthill1952}. Since then, a vast amount of theoretical and numerical efforts has been dedicated to specify the properties of the spectrum of sound waves generated in a stratified stellar convection zone, by improving the description of the turbulent energy spectrum of the convective elements \\eg\\ \\cite{Bi+Li1998, Goldreich+Kumar1990, Goldreich+Murray+Kumar1994, Musielak+etal1994, Nordlund+Stein2001, Stein1967, Stein+Nordlund2001}. Without going into details of these works, the present knowledge can be summarized in the following way. The efficiency of the energy conversion from convective to acoustic is proportional to a high power of the Mach number of the convective motions ($M^{15/2}$ \\cite{Goldreich+Kumar1990}).  Most of the energy going into the $p$-modes, $f$-modes and propagating acoustic waves is emitted by convective eddies of size $h \\sim M^{3/2}H$ (where $H$ is the value of the pressure scale height), at frequencies close to the acoustic cut-off frequency  $ \\omega \\sim \\omega_c$ and wavelengths similar to $H$. This defines the frequency dependence of the oscillation spectrum observed in the Sun and stars \\cite{Goldreich+Kumar1990}. Since the Mach number is largest at the top part of the convection zone, the peak of the acoustic energy generation is located immediately below the photosphere \\cite{Nordlund+Stein2001, Stein+Nordlund2001}. Recent numerical simulations of magneto-convection have shown that the magnetic field modifies the spectrum of waves \\cite{Jacoutot+etal2008}. Apart from the power suppression in regions with enhanced magnetic field, these simulations suggest an increase of high-frequency power (above 5 mHz) for intermediate magnetic field strengths (of the order of 300--600 G) caused by changes of the spatial-temporal spectrum of turbulent convection in a magnetic field.    Waves generated in the convection zone resonate inside a cavity formed by the stellar interior and the photosphere and are used by helio- and astroseismology to derive its properties \\cite{Deubner1975, Ulrich1970}. The information contained in the frequencies of the trapped wave modes is used by the classical helioseismology. A relatively newer branch called local helioseismology uses the information contained in the velocity amplitudes of the propagating waves measured in a region of interest on the solar surface (Duvall \\etal\\ \\cite{Duvall+etal1993}). By inversion of these measurements, variations of the wave speed and velocities of mass flows can be recovered below the visible solar surface. The inversion results have been obtained for quiet Sun regions as well as for magnetic active regions including sunspots. It is known that sunspots possess strong magnetic fields with a complicated structure in the visible layers of the Sun where the Doppler measurements used by local helioseismology are taken \\cite{Solanki2003}. Consequently, such magnetic field can cause important effects on helioseismic waves, beyond the perturbation theories employed for helioseismic data analysis \\cite{Birch+Kosovichev2000, Gizon+Birch2002, Kosovichev+Duvall1997}. Recent numerical and analytical results demonstrate that the observed time-distance helioseismology signals in sunspot regions correspond to fast MHD waves \\cite{Khomenko+etal2008b, Moradi+etal2008}.   An estimation of the acoustic energy flux generated in the solar convection zone by the turbulent motions suggests that the it can be as large as, \\eg\\ $F_A = 5 \\times 10^7$ ergs/cm$^2$/s \\cite{Musielak+etal1994}. This is more than sufficient to maintain a hot chromosphere. It made the theory of acoustic heating of the upper atmosphere very attractive. However, soon after being proposed, the theory of acoustic heating encountered several major difficulties. It was found that both, low- and high-frequency waves are radiatively damped in the photosphere, reaching the upper layers with significantly less power \\cite{Ulmschneider1971}.  An additional difficulty comes from the fact that the measured high-frequency acoustic fluxes in the photosphere and chromosphere are uncertain and non-conclusive \\cite{Kalkofen2007, Wunnenberg+etal2002}. Low-frequency acoustic waves are reflected in the photosphere due to the effects of the cut-off frequency ( $\\sim 4$ mHz) before reaching chromospheric heights. Due to their long wavelengths they have too large shock formation distances. Despite that, the five-minute waves with enough energy were detected in the chromosphere and corona of the Sun  mainly above solar facular and network areas \\cite{Centeno+etal2006b, Krijer+etal2001, DeMoortel+etal2002, DePontieu+etal2003, Veccio+etal2007}. Several explanations of these waves involving the magnetic field have been recently proposed \\cite{Khomenko+etal2008, DePontieu+etal2004}. The wave energy can reach the upper layers not necessarily in the form of acoustic waves, but also in the form of other wave types, like magneto-acoustic waves, Alfv\\'en waves or a family of waves propagating in thin magnetic flux tubes (see a recent example in \\eg\\ \\cite{Hindman+Jain2008}). All of them are related to the magnetic field structure. Osterbrock in 1961 \\cite{Osterbrock1961} was the first to incorporate magnetic field into the theory of wave heating and to point out its importance on the propagation and refraction characteristics of the fast and slow MHD waves.     The above examples are only a few where the influence of the magnetic field on the wave properties is demonstrated to be important. Magnetic field not only changes the acoustic excitation rate and produces new wave modes. It also modifies the wave propagation paths and the direction of the energy propagation, it produces wave refraction and changes the reflection characteristics at the near-surface layers. Magnetic field defines the wave propagation speeds and changes the acoustic cut-off frequency. It can also change the wave dissipation rates and provides an additional energy source. Magnetic field connects all the atmospheric layers facilitating the channeling of waves from the lower to the upper atmosphere. This makes the magnetic field an important ingredient the theories of the wave propagation in the atmospheres of the Sun and stars.    Apart from the problems set by the local helioseismology in magnetic regions and the wave heating theory, of pure physical interest is the interpretation of oscillations observed in different magnetic structures in terms of MHD waves. For example, the wave dynamics seen in high-resolution DOT movies of a sunspot region \\cite{Rouppe+etal2003} demonstrate that phenomena such as chromospheric umbral flashes and running penumbral waves are closely related. What type of waves are responsible for them? Solar small-scale and large-scale magnetic structures have distinct magnetic and thermal properties and support different wave types. The observed frequency spectrum of waves in these structures is not the same (see the introduction in \\cite{Khomenko+Collados2007}). Numerical simulations of waves in non-trivial magnetic structures (\\eg\\ \\cite{Bogdan+etal2003, Hasan+Ballegooijen2008, Hasan+etal2005, Khomenko+Collados2006, Khomenko+Collados2007, Rosenthal+etal2002}) have shown the complex pattern formed by waves of various types that can propagate simultaneously in various directions. Different wave modes can be detected in observations depending on the magnetic field configuration and the height where acoustic speed, $c_S$, and Alfv\\'en speed, $v_A$, are equal relative to the height of formation of the spectral lines used in observations.  During the last years we applied efforts to develop a numerical code aimed at calculating the non-linear wave propagation inside magnetic fields in 2 and 3 dimensions. Using this code we focused our analysis on several problems described above, namely: magneto-acoustic wave propagation and refraction in sunspots and flux tubes; channeling the five-minute photospheric oscillations into the solar outer atmosphere through small-scale magnetic flux tubes; influence of the magnetic field on local helioseismology measurements in active regions. The results of these studies are published in \\cite{Khomenko+etal2008, Khomenko+Collados2006, Khomenko+Collados2007, Khomenko+etal2008b}. In the rest of the paper we briefly summarize the results and conclusions of these works. In addition, the last section gives our recent contribution to the problem of the interpretation of observations of waves in magnetic roAp stars.  \\begin{figure*} \\includegraphics[width=12cm]{plen_talk_khomenko_fig1.eps} \\caption{Variations of the velocity, magnetic field and pressure at an elapsed time $t=100$ sec after the beginning of the simulations. In each panel, the horizontal axis is the radial distance $X$ from the sunspot axis and the vertical axis is height from the photospheric level. The black inclined lines are magnetic field lines. The two more horizontal lines are contours of constant $c_S^2/v_A^2$, the thick line corresponding to $v_A=c_S$ and the thin line to $c_S^2/v_A^2=0.1$. Top left: transversal variations of the magnetic field. Top right: relative pressure variations. Bottom panels: transversal and longitudinal variations of the velocity. \\label{fig:long}} \\end{figure*} ",
        "conclusions": ""
    },
    "0812/0812.1791_arXiv.txt": {
        "abstract": "We formulate the problem of the formation and subsequent evolution of fragments (or cores) in magnetically-supported, self-gravitating molecular clouds in two spatial dimensions. The six-fluid (neutrals, electrons, molecular and atomic ions, positively-charged, negatively-charged, and neutral grains) physical system is governed by the radiative, nonideal magnetohydrodynamic (RMHD) equations. The magnetic flux is not assumed to be frozen in any of the charged species. Its evolution is determined by a newly-derived generalized Ohm's law, which accounts for the contributions of both elastic and inelastic collisions to ambipolar diffusion and Ohmic dissipation. The species abundances are calculated using an extensive chemical-equilibrium network. Both MRN and uniform grain size distributions are considered. The thermal evolution of the protostellar core and its effect on the dynamics are followed by employing the grey flux-limited diffusion approximation. Realistic temperature-dependent grain opacities are used that account for a variety of grain compositions. We have augmented the publicly-available Zeus-MP code to take into consideration all these effects and have modified several of its algorithms to improve convergence, accuracy and efficiency. Results of magnetic star formation simulations that accurately track the evolution of a protostellar fragment from a density $\\simeq 10^{3}$ cm$^{-3}$ to a density $\\simeq 10^{15}$ cm$^{-3}$, while rigorously accounting for both nonideal MHD processes and radiative transfer, are presented in a separate paper. ",
        "introduction": "\\label{section:intro}  The formulation of a theory of star formation is a formidable task. It requires understanding of the nonlinear interactions among self-gravity, magnetic fields, rotation, chemistry (including grain effects), turbulence, and radiation. Stars form in fragments within interstellar molecular clouds, which have sizes ranging from 1 to 5 pc, masses from a few tens to $10^5$ M$_\\odot$, mean densities $\\simeq 10^3$ cm$^{-3}$, and temperatures $\\simeq 10$ K \\citep{myers85,heiles87}. Their spectral lines have Doppler-broadened linewidths that suggest supersonic (but subAlfv\\'{e}nic) internal motions. In the deep interiors of such clouds, high-energy cosmic rays ($>100$ MeV) maintain a degree of ionization $x_{\\rm i} \\lesssim 10^{-7}$, whereas ultraviolet (UV) ionization is responsible for a much greater degree of ionization $x_{\\rm i}\\gtrsim 10^{-5}$ in the outer envelopes.  \\subsection{Magnetic Fields and Star Formation}  The possible importance of magnetic fields to the support of interstellar clouds and to the regulation of star formation was first studied by \\citet{cf53}, \\citet{ms56}, and \\citet{mestel65} using the virial theorem. Similar investigations by \\citet{strittmatter66a,strittmatter66b} and \\citet{spitzer68} followed. \\citet{mestel66} calculated the magnetic forces on a spherically-symmetric, gravitationally-bound cloud. Self-consistent calculations by \\citet{mouschovias76a,mouschovias76b} produced exact equilibria of initially uniform, isothermal, magnetic clouds embedded in a hot and tenuous, electrically-conducting external medium. \\citet{ms76} used these equilibrium states to find the critical mass-to-flux ratio  $(M/\\Phi_{\\rm B})_{\\rm cr} = 1/(63G)^{1/2}$ that must be exceeded for collapse against the magnetic forces to set in. \\citet{sb80} performed numerical simulations of the collapse of a supercritical (as a whole) magnetic cloud. A picture of molecular clouds emerged in which magnetic fields play a central role in their support and evolution \\citep{mouschovias78}. To this date, magnetic braking remains the only mechanism that has been shown quantitatively to resolve the most significant dynamical problem of star formation, namely, the angular momentum problem (Mouschovias \\& Paleologou 1979, 1980; see also summary below).  Subsequent observations lent credence to this picture by revealing the importance of magnetic fields through both dust polarization measurements and Zeeman observations. Polarization studies have exhibited large-scale ordered magnetic fields connecting protostellar cores to their surrounding envelopes \\citep{vct81,hvsssgs87,nghpd89,nddgh97,lcgr01,lgc03,cnwtk04,afg08}, often with an hourglass morphology \\citep{schleuning98,hddsv99,gcr99,svhdndd00,lcgr02,cc06,grm06}, as predicted by the theoretical calculations \\citep{mouschovias76b,fm93}. A large body of Zeeman observations \\citep{ck83,kc86,tck86,ctk87,gchmt89,ctghkm93,cmtc94,crmt96,tcghkm96, crtg99,ctlpk99,crutcher99,hc05,ccw05} revealed magnetic fields in the range $\\simeq \\, 10 - 200 \\, \\mu$G in molecular clouds, from small isolated ones to massive star-forming ones. These values are more than sufficient to establish the importance of magnetic fields in molecular cloud dynamics.  It was recognized early on (e.g., see Babcock \\& Cowling 1953, p. 373) that the magnetic flux of an interstellar blob of mass comparable to a stellar mass is typically several orders of magnitude greater than that of magnetic young stars. This is the so-called ``magnetic flux problem\" of star formation. It lies in the fact that substantial flux loss must take place at some stage during star formation. Ambipolar diffusion (the relative motion between plasma and neutrals) was first proposed by \\citet{ms56} as a means by which an interstellar cloud as a whole would reduce its magnetic flux and thereby collapse. \\citet{pm65} undertook a detailed calculation of the collapse of such (spherical) cloud. \\citet{spitzer68} calculated the ambipolar-diffusion timescale by assuming that the magnetic force on the ions is balanced by the (self-)gravitational force on the neutrals. \\citet{nakano79} followed the quasistatic contraction of a cloud due to ambipolar diffusion using a sequence of Mouschovias' (1976b) equilibrium states, each one of which had a smaller magnetic flux than the previous one.  A new solution for ambipolar diffusion by \\citet{mouschovias79} showed that the essence of this process is a redistribution of mass in the central flux tubes of a molecular cloud, rather than a loss of magnetic flux by the cloud as a whole. He found the ambipolar-diffusion timescale to be typically three orders of magnitude smaller in the interior of a cloud than in the outermost envelope, where there is a much better coupling between neutral particles and the magnetic field because of the much greater degree of ionization. This suggested naturally a self-initiated fragmentation of (or core formation in) molecular clouds on the ambipolar-diffusion timescale \\begin{equation*} \\tau_{\\rm AD}=1.8\\times 10^6\\,\\left(\\frac{x_{\\rm i}}{10^{-7}}\\right)\\;{\\rm yr}\\, \\end{equation*} (Mouschovias 1979, eq. [37]). The inefficiency of star formation was thereby attributed to the self-initiated formation and contraction of molecular cloud fragments (or cores) due to ambipolar diffusion in otherwise magnetically supported clouds \\citep{mouschovias76b,mouschovias77,mouschovias78,mouschovias79}. The {\\em central} mass-to-flux ratio eventually exceeds its critical value for collapse, \\begin{equation*} \\left(\\frac{dm}{d\\Phi_{\\rm B}}\\right)_{\\rm c,cr} = \\frac{3}{2}\\left(\\frac{M}{\\Phi_{\\rm B}}\\right)_{\\rm cr}\\,, \\end{equation*} (see Mouschovias 1976a, eq. [44]), and dynamic contraction ensues. Detailed numerical calculations in slab \\citep{pm83,mpf85}, cylindrical \\citep{mm91,mm92a,mm92b}, and axisymmetric geometries \\citep{fm92,fm93,cm93,cm94,cm95,bm94,bm95a,bm95b,ck98,dm01,tm05a,tm05b,tm07a,tm07b,tm07c} transformed this scenario of star formation into a theory with predictive power.  \\subsection{Rotation}  During the early, isothermal phase of star formation, a cloud (or a fragment) must also lose a large fraction of its angular momentum (e.g., see Spitzer 1968, p. 231). Observations show that molecular clouds and embedded fragments (or cores) rarely exhibit rotation significantly greater than that of the background medium \\citep{ga85}. If angular momentum were conserved from the initial galactic rotation (i.e., starting from angular velocity $\\Omega_0\\simeq 10^{-15}$ s$^{-1}$ at the mean density of the interstellar medium $\\simeq 1$ cm$^{-3}$), centrifugal forces would not allow even the formation of interstellar clouds (Mouschovias 1991, \\S~2). Fragmentation does not alter this conclusion (Mouschovias 1977, \\S~1). This is referred to as the ``angular momentum problem\" of star formation. As far as clouds and their cores are concerned, the angular momentum problem has been shown to be resolved by magnetic braking (i.e., the transport of angular momentum from a fragment to its surrounding medium through the propagation of torsional Alfv\\'{e}n waves along magnetic field lines connecting the fragment to the cloud envelope) analytically by \\citet{mp79,mp80} and numerically by \\citet{bm94,bm95a,bm95b} and \\citet{ml08}. Hence the centrifugal forces resulting from the cloud's or core's rotation have a negligible effect on the evolution of the contracting core, at least up to central densities of $\\approx 10^{14}$ cm$^{-3}$ (see the last paragraph of Tassis \\& Mouschovias 2007b).  \\subsection{Grain Effects}  Interstellar grains comprise about 1\\% of the mass in the interstellar medium \\citep{spitzer78}. \\citet{baker79} and \\citet{elmegreen79} suggested that charged grains may couple to the magnetic field and thereby play a role in ambipolar diffusion and star formation. \\citet{elmegreen79} and \\citet{nu80} compared and ambipolar-diffusion timescale and the free-fall timescale and concluded that ambipolar diffusion occurs over too long a timescale (roughly 10 times greater than free-fall) to be a relevant process in star formation. Refinements by the same authors \\citep{elmegreen86,un90,nnu91} led to similar conclusions. Through detailed numerical simulations of core formation and evolution including the effects of (negative and neutral) dust grains, \\citet{cm93,cm94} found that grains lengthen the timescale for the formation of a core because of grain-neutral collisions, but cautioned that the ambipolar-diffusion timescale should not be compared to the free-fall timescale in determining its relevance in magnetically-supported clouds, as originally pointed out by \\citet{mouschovias77}, because molecular clouds are not free-falling. Velocities characteristic of such collapse have not been observed. \\citet{cm95} extended these calculations by including UV ionization and a variety of atomic metal ions (C$^+$, S$^+$, Si$^+$, Mg$^+$, Na$^+$, Fe$^+$). Attention was also paid to the complementary effect of protostellar evolution on the microscopic physics and chemistry \\citep{cm96,cm98}.  \\subsection{MHD Waves and/or Turbulence}  The extent to which turbulence (or waves) may or may not affect the evolution of a protostellar fragment has been a topic of debate for several decades, receiving increased attention in recent years \\citep{mlk04,mtk06}. A consensus has yet to be reached concerning even what causes and maintains the observed broad linewidths long thought to be indicative of supersonic turbulence \\citep{ze74,am75,larson81,zj83,mouschovias87,mp95,mg99,es04,mtk06}, although \\citet{mp95} and \\citet{mtk06} showed quantitatively that observations contradict the key assumption of turbulence simulations, namely, that molecular clouds are magnetically supercritical by a factor $4 - 10$. Despite a lack of agreement on the origin of the linewidths, analytical \\citep{mouschovias91} and numerical calculations (Eng 2002; Eng \\& Mouschovias 2009, in preparation) have demonstrated that turbulence (or waves) plays an insignificant role in the star formation process once dynamical contraction of a fragment (or core) ensues. Observations showing narrowing and eventual thermalization of linewidths in protostellar cores \\citep{bpddc81,mb83,mlb83,bapabwt00} are in agreement with this conclusion and with an earlier version of it \\citep{mouschovias87}.  \\subsection{Radiative Transfer}  During the early phases of star formation the energy produced by compressional heating is radiated away by the dust grains in the infrared. At higher densities ($\\simeq 10^{11}$ cm$^{-3}$), the core traps and retains part of this heat and its temperature begins to rise. The evolution of the temperature in this nonisothermal regime may be approximated (but substantially overestimated) by using an adiabatic equation of state \\citep{boss81,dm01,tm07a,tm07b,tm07c}. More realistic equations of state have also been employed by, for example, \\citet{bate98}. To accurately model the nonisothermal phase of protostellar contraction, however, one needs to include a proper treatment of radiative transfer. Early efforts to include radiative transfer in (nonmagnetic) star formation calculations were confined to the use of the diffusion approximation \\citep{bodenheimer68,larson69,larson72,bb75,bb76,tscharnuter75,yk77}. While the diffusion approximation is strictly applicable only to optically thick regions, its ease of implementation and relatively low computational cost make it an attractive choice. The Eddington approximation offers a slight improvement in that it retains some of the rigor of using moments of the radiative transfer equation, while making the simplifying assumption that the radiation field is everywhere isotropic. Its use in numerical calculations of (nonmagnetic) star formation has been documented in \\citet{tscharnuter78}, \\citet{tw79}, \\citet{wn80a,wn80b}, \\citet{boss84,boss86,boss88}, and \\citet{bm95}. By implicitly assuming that photons always travel a distance comparable to their mean-free path (even if this distance exceeds the free-flight distance $c\\Delta t$, where $\\Delta t$ is the computational timestep), the Eddington approximation gives unphysical behavior in optically thin regions, in which the mean-free-path is huge. The result is a signal speed unbound by the speed of light, i.e., it violates causality (e.g., see \\S~97 of Mihalas \\& Mihalas 1984).  Increasing the accuracy and realism of a radiative transfer algorithm often requires making limiting assumptions about the hydrodynamics in order to make the problem tractable (e.g., Yorke 1980; Masunaga et al. 1998). A full frequency- and angle-dependent treatment of the radiation is nearly always confined to postprocessing the results of a hydrodynamic calculation \\citep{yorke77,ys81,as85,as86} or a grey (i.e., independent of frequency) radiation hydrodynamic calculation \\citep{by90,byrt90}. By contrast, the flux-limited diffusion (FLD) approximation \\citep{lp81} is a propitious compromise that retains some of the advantages of the diffusion and Eddington approximations, while preserving causality and coupling self-consistently to the hydrodynamic equations.  \\subsection{Outline}  In this paper we formulate the problem of the formation and evolution of protostellar fragments (or cores) in magnetically-supported, self-gravitating molecular clouds, including the effects of both ambipolar diffusion and Ohmic dissipation (which becomes important at high densities), grain chemistry and dynamics, and radiation. Using the results of Eng (2002) and Eng \\& Mouschovias (2009, in preparation), and \\citet{bm94}, we may safely ignore the effects of turbulence and rotation, respectively, on the evolution of the protostellar core for the densities considered here. The physical and chemical properties of the model cloud are summarized in Section \\ref{section:modelcloud}. The radiation magnetohydrodynamic (RMHD) equations governing the evolution of the model cloud are presented and discussed in Section \\ref{section:equations}. In Section \\ref{section:chem} we present the chemical model used in the calculations. The physics of magnetic diffusion (ambipolar and ohmic) is handled by using a generalized Ohm's law, which is derived in Section \\ref{section:fluxloss}. We treat the radiative transfer using the grey FLD approximation, with realistic grain opacities accounting for a variety of grain compositions (\\S~\\ref{section:radtrans}). The numerical method of solution is discussed in Section \\ref{section:zeusmp}. Finally, we give the simplified set of equations and a brief summary in Section \\ref{section:summary}. Details, mostly mathematical, are left for the Appendix. Results are presented in a separate paper (Kunz \\& Mouschovias 2009, in preparation). ",
        "conclusions": "\\label{section:summary}  We have formulated the problem of the formation and evolution of fragments (or cores) in magnetically-supported, self-gravitating molecular clouds in two spatial dimensions. The evolution is governed by the six-fluid RMHD equations. The magnetic flux is not assumed to be frozen in any of the charged species. Its evolution is determined by a newly-derived generalized Ohm's law, which accounts for the contributions of both elastic and inelastic collisions to ambipolar diffusion and Ohmic dissipation. The species abundances (electrons, atomic and molecular ions, positively-charged grains, negatively-charged grains, and neutral grains) are calculated using an extensive equilibrium chemical network. Both MRN and uniform grain size distributions are considered. The thermal evolution of the protostellar core and its effect on the dynamics are followed by employing the grey FLD approximation. Realistic temperature-dependent grain opacities are used that account for a variety of grain compositions. We have augmented the publicly-available Zeus-MP code to take into consideration all these effects and have modified several of its algorithms to increase its accuracy and efficiency.  We summarize here for convenience the simplified evolutionary equations discussed above and used in our modified version of the Zeus-MP code: \\begin{subequations} \\begin{equation} \\D{t}{\\dr{n}} + \\del\\bcdot(\\dr{n}\\vv{n}) = 0\\,, \\end{equation} \\begin{equation} \\D{t}{\\dr{g}} + \\del\\bcdot(\\dr{g}\\vv{n}) = -\\del\\bcdot\\bigl(\\eta_{\\rm cont,H}\\bb{j}\\btimes\\eb\\bigr)\\,, \\end{equation} \\begin{equation} \\D{t}{(\\dr{n}\\vv{n})} + \\del\\bcdot(\\dr{n}\\vv{n}\\vv{n}) = -\\del P_{\\rm n} -\\rho\\del\\psi + \\frac{1}{4\\pi}(\\del\\btimes\\bb{B})\\btimes\\bb{B} - \\lambda_{\\rm FLD}\\del\\mc{E}\\,, \\end{equation} \\begin{equation} \\D{t}{\\bb{B}} = \\del\\btimes\\left(\\vv{n}\\btimes\\bb{B} - \\frac{c^2\\eta_\\perp}{4\\pi}\\del\\btimes\\bb{B}\\right)\\,, \\end{equation} \\begin{equation} \\del^2\\psi = 4\\pi G\\dr{n}\\,, \\end{equation} \\begin{equation} \\D{t}{\\inteng} + \\del\\bcdot(\\inteng\\vv{n}) = -P_{\\rm n}\\del\\bcdot\\vv{n} - 4\\pi\\kappap\\mc{B} + c\\kappap\\mc{E}\\,, \\end{equation} \\begin{equation} \\D{t}{\\mc{E}} + \\del\\bcdot(\\mc{E}\\vv{n}) = \\del\\bcdot\\left(\\frac{c\\lambda_{\\rm FLD}}{\\chir}\\del\\mc{E}\\right) - \\del\\vv{n}\\,\\bb{:}\\,\\msb{P} + 4\\pi\\kappap\\mc{B} - c\\kappap\\mc{E}\\,. \\end{equation} \\end{subequations} These equations are considered together with the relations $P_{\\rm n}=(\\gamma-1)\\inteng$ and $\\msb{P}=\\msb{f}\\mc{E}$. Results will be presented in a forthcoming paper (Kunz \\& Mouschovias 2009, in preparation)."
    },
    "0812/0812.3277_arXiv.txt": {
        "abstract": "Since the suggestion of relativistic shocks as the origin of gamma-ray bursts (GRBs) in early 90's, the mathematical formulation of this process has stayed at phenomenological level. One of the reasons for the slow development of theoretical works has been the simple power-law behaviour of the afterglows hours or days after the prompt gamma-ray emission. It was believed that they could be explained with these formulations. Nowadays with the launch of the \\swift satellite and implementation of robotic ground follow-ups, gamma-ray bursts and their afterglow can be observed in multi-wavelength from a few tens of seconds after trigger onward. These observations have leaded to the discovery of features unexplainable by the simple formulation of the shocks and emission processes used up to now. Some of these features can be inherent to the nature and activities of the GRBs central engines which are not yet well understood. On the other hand {\\bf\\it devil is in details} and others may be explained with a more detailed formulation of these phenomena and without adhoc addition of new processes. Such a formulation is the goal of this work. We present a consistent formulation of the kinematic and dynamics of the collision between two spherical relativistic shells, their energy dissipation, and their coalescence. It can be applied to both internal and external shocks. Notably, we propose two phenomenological models for the evolution of the emitting region during the collision. One of these models is more suitable for the prompt/internal shocks and late external shocks, and the other for the afterglow/external collisions as well as the onset of internal shocks. We calculate a number of observables such as flux, lag between energy bands, and hardness ratios. One of our aims has been a formulation enough complex to include the essential processes, but enough simple such that the data can be directly compared with the theory to extract the value and evolution of physical quantities. To accomplish this goal, we also suggest a procedure for extracting parameters of the model from data. In a following paper we numerically calculate the evolution of some simulated models and compare their features with the properties of the observed gamma-ray bursts. ",
        "introduction": "\\label{sec:intro} The \\swift~\\citep{swift} observations of more than 200 Gamma-Ray Bursts (GRBs) and their follow-ups have been a revolution in our knowledge and understanding of these elusive phenomena. The rapid slew of the \\swift X-ray and UV/optical telescopes - respectively XRT~\\citep{xrt} and UVOT~\\citep{uvot} - as well as ground based robotic telescopes have permitted to observe GRBs and their afterglow in multi-wavelength from few tens of seconds after the prompt or precursor gamma-ray emission is detected by BAT~\\citep{bat}, up to days after trigger. These observations show that the emission can be essentially divided to three regimes: 1) The prompt gamma-ray emission which can be very short, few tenth of milliseconds, or long, up to few hundred of seconds; 2) A tail emission in X-ray which is observed for more than $90\\%$ of bursts. For some bursts this tail is also detected as a soft faint continuum in gamma-ray. In about $40\\%$ of bursts this early emission has been detected in optical and infrared too. In this regime for many bursts flares have been observed mainly in X-ray. Sometimes the counterpart of flares have been also observed in gamma-ray and/or optical/IR. In many bursts the early steep slope of the X-ray emission at the beginning of this regime becomes much shallower and somehow harder at the end; 3) The late emission can be considered as the epoch after the break of shallow regime in which the emission is usually a continuum and no or little flaring activity is observed (but there are exceptions such as GRB 070110~\\citep{grb070110} and GRB 081028~\\citep{grb081028} which had bright late flares). The duration and relative fluxes of these regimes can vary significantly between GRBs.  In one hand, it seems that the idea of synchrotron emission from accelerated particles in a relativistic shock as the origin of the prompt emission~\\citep{intext,intext1} is essentially correct. On the other hand, the early observations of what is usually called {\\it the afterglow} - the emission in lower energy bands usually observed from $\\lesssim 100$~sec after trigger onward - have been the source of surprises and raised a number of questions about many issues: the activity~\\citep{longact} and the nature of the engine~\\citep{progenitor,progenitor1,progenitorshort}, the concept of prompt/internal-afterglow/external shocks~\\citep{promptext}, the efficiency of energy transfer from the outflow -{\\it the fireball} - to synchrotron radiation~\\citep{enereff}, the collimation and jet break~\\citep{jetbreakex}, the behaviour of X-ray and optical light curves~\\citep{optag}, etc. Many of predictions such as the existence of a significant high latitude emission with a strict relation between the light curve time evolution slope and the spectrum index, and an achromatic jet break have not been observed. Moreover, the origin of unexpected behaviours such as a very steep decline in low energy bands after the prompt~\\citep{tailhighlat} and a very shallow regime which lasts for thousands of seconds are not well understood. Other unexpected observations are the existence of the chromatic multiple breaks in the X-ray light curves, flares in X-ray and optical bands hundred of seconds after the prompt even in some short bursts (ex: GRB 060313~\\citep{grb060313,grb060313-1}, GRB 070724A~\\citep{grb070724a}, a tail emission in short bursts (ex: GRB070714B~\\citep{grb070714b,grb070714b1}, GRB080426~\\citep{grb080426}, GRB 080503~\\citep{grb080503}), and very short, hard, and high amplitude spikes in long bursts that could lead to the classification of the burst as short if the instrument was not enough sensitive to detect the rest of the prompt emission (ex: GRB060614~\\citep{grb060614,swiftgrb060614}, GRB061006~\\citep{grb061006}). This makes the classification of bursts as short and long much more ambiguous~\\citep{grbnewclass}.  One conclusion that has been made from these observations is that the central engine can be active for up to thousands of seconds after the prompt emission~\\citep{longact}. But the nature of the fireball and its source of energy is not yet well understood, and we can not yet verify this interpretation or relate it to any specific process in the engine. It seems however that whatever the origin of the fireball, it must be baryon dominated otherwise it could not make long term effects correlated to the prompt emission. In this case, the internal and external shock models as the origin of the prompt and afterglow are good candidates. Nonetheless, the lack of a simple explanation for the observed complexities has encouraged authors to consider other possibilities, for instance associating both the prompt gamma-ray and the afterglows to external shocks and fast variations to abrupt density variation of the surrounding material~\\citep{promptext}. However, it has been shown that in such models it is not possible to explain the fast variations of the prompt even in presence of a bubbly environment or pulse-like density change~\\citep{promptext,promptext1}.  Here we suggest that at least some of the features of early afterglows can be related to a complex shock physics and/or features in the fireball/jet. In fact, simulations of the acceleration of electrons and positrons by the first and second Fermi processes show that the evolution of electric and magnetic fields as well as the energy distribution of accelerated particles are quite complex~\\citep{fermiacc,fermiacc1,fermiacc2}. Plasma instabilities lead to the formation of coherent electric and magnetic fields and acceleration of particles~\\citep{weibel,weibel1,weibeltemp,instabparal}. Their time evolution in relativistic shocks can significantly affect the behaviour of the prompt and the afterglow of GRBs. If the number density of particles in the ejecta is significant and the shock is collisional, the state of matter in the jet can be also an important factor in determining the behaviour of the fields, and thereby the synchrotron emission by accelerated electrons and positrons. Many aspects of these processes are not well understood, however realistic interpretations of observations should consider these complexities at least phenomenologically. For instance, the simple distributions such as a power-law distribution for Lorentz factor of electrons, or a constant magnetic field for the whole duration of prompt and afterglow can be quite unrealistic. Ideally, these quantities should come from the simulation of Fermi processes and plasma instabilities such as Weibel instability~\\citep{weibel,weibel1} that produces the coherent transverse magnetic field. However, these phenomena are complex and their simulations are very time and CPU consuming. For these reasons they can not yet explore the parameter space of the phenomena and are mostly useful for demonstrating the concepts and how they work. Therefore we are obliged to use simple analytical approximations for quantities related to the physics of relativistic shocks. In this situation a compromise between complex non-analytical expressions and too simplistic and too simplified but unrealistic analytical behaviour of the physical quantities can be the consideration of intervals in which a simple analytical function can be a good approximation. Then, by adding together these intervals - regimes - one can reconstruct the entire evolution of a burst and its afterglow.  Even with a simplified presentation of the physical processes one would not be able to explain GRB data without a model including both microphysics and dynamics of the fireball. The majority of works on the modelling of shocks and synchrotron emissions either deal with the emission~\\citep{emission,emission0,emission1,revshock,revshocksimul} or with the kinematics of the shock~\\citep{kinematic,kinematic1,kinematic2}, or both but in a phenomenological way~\\citep{kinematic3}. Few works~\\citep{shocksynch,shocksynch1} have tried to include both these aspects in a consistent model, but either they have not been very successful - their predictions specially for quantities such as lags in different bands were far from observed values and additional parametrization was necessary - or the formulation is too abstract to be compared directly with data~\\citep{grbboltz,grbboltz1}.  With these issues in mind, in Sec.\\ref{sec:model} we present a simplified shock model that includes both the kinematics of the ejecta and the dynamics of the synchrotron emission. The microphysics is included by the means of a simple parametrization. We calculate a number of observables such as flux, hardness ratios, and lags between different energy bands. In this paper and paper II in which we simulate part of prompt and afterglow regimes of GRBs in some time intervals, we show that this model can explain many aspects of bursts as long as we divide the data to separate regimes. The reason is that the simple parametrization of microphysics in this model can be valid at most in a limited time interval. Each regime should be separately compared to analytical and numerical results for extracting the parameters. The results will show how parameters that are considered as constant in this model evolve during the lifetime of the burst. This is the best we can do until a better understanding of relativistic shock models and Fermi processes become available. If the model and the estimation of its parameters for each regime is sufficiently correct, adding them together should give us an overall consistent picture of characteristics of the burst, its afterglow, and its surrounding material. Apriori this knowledge should help to better understand the engine activities and eventually its nature and classification.  The model presented here depends on a large number of parameters and we need an extraction procedure permitting to extract as much as possible information about the physical properties of the shock from the available data. In Sec.\\ref{sec:extract} we explain how in the frame work of this model one can extract various quantities from data. Evidently the success of the modelling strongly depends on availability and quality of simultaneous multi-wavelength observations.  The \\swift observations show that during the first few hundred seconds after the trigger there is usually a very close relation between the prompt gamma-ray emission and the emission in lower energy bands~\\citep{earlyemi, earlyemi1}, therefore, most probably they have a common origin, presumably internal shocks. However, historically and even in the present literature (and sometimes believes) any emission after the prompt gamma-ray is called {\\it the afterglow} - meaning due to a shock with ISM or surrounding material, presumably external shocks. Therefore, for clarity of context here we define {\\it the afterglow} as the emission in any energy band and at any time after the main prompt peaks regardless of its origin. If by {\\it afterglow} we mean the external shocks, this is mentioned explicitly in the text.  We finish this paper with some outlines and two appendices containing the details of calculation of the dynamics and flux for power-law distribution of electrons Lorentz factor. ",
        "conclusions": "We presented a formulation of the relativistic shocks and synchrotron emission that includes more details than the dominant term considered in the previous calculations. Although we consider spherical shells, most of our results are valid also for non-spherical collimated jets as long as the collimation angle due to the relativistic boost is smaller than collimation angle.  We showed that the lags between light curves at different energies exist in the dominant order and are not only due to the high latitude emissions which are negligible for ultra-relativistic ejecta.  The main reason for such behaviour is the evolution of electric and magnetic fields as well as the evolution of the emitting region which can be in addition energy dependent. This fact is more evident in the simulations presented in Paper II. Despite the absence of high latitude terms in our simulations, the presence of lags between the light curves of different energy bands is evident. For the same reasons the change in the slope of the light curves - what is called breaks - are also energy dependent. This explains chromatic breaks of the GRBs detected by \\swift.  The two phenomenological models we considered for the evolution of the active region are physically motivated, but do not have rigorous support from microphysics of the shock. Nonetheless, they can be easily replaced if future simulations of Fermi processes lead to a better estimation of the size of the region in which electric and magnetic fields are formed and particles are accelerated and dissipated. Presence of an external magnetic field in the environment of the candidates for central engine of GRBs is very plausible. The formulation present here does not include such a possibility, but an external magnetic field can be added to (\\ref{magener}). The modification of the evolution equation of $\\beta'$ and the flux is straightforward.  Many other details such as the effect of metalicity of the ejecta and surrounding material both on the low energy emission and absorption is not considered in this work. We have also neglected synchrotron self-abortion. It only affects the low energy bands, nonetheless in hard bursts even optical emission can be affected by self absorption. We leave the study of this issue, the effect of ionization on the emissions, and the thermalization of shocked material to future works.  \\appendix"
    },
    "0812/0812.3749_arXiv.txt": {
        "abstract": "{The properties of the early-type binary \\object{Cyg\\,OB2 \\#5} have been debated for many years and spectroscopic and photometric investigations yielded conflicting results.}{We have attempted to constrain the physical properties of the binary by collecting new optical and X-ray observations.}{The optical light curves obtained with narrow-band continuum and line-bearing filters are analysed and compared. Optical spectra are used to map the location of the He\\,{\\sc ii} $\\lambda$\\,4686 and H$\\alpha$ line-emission regions in velocity space. New XMM-Newton as well as archive X-ray spectra are analysed to search for variability and constrain the properties of the hot plasma in this system.}{We find that the orbital period of the system slowly changes though we are unable to discriminate between several possible explanations of this trend. The best fit solution of the continuum light curve reveals a contact configuration with the secondary star being significantly brighter and hotter on its leading side facing the primary. The mean temperature of the secondary star turns out to be only slightly lower than that of the primary, whilst the bolometric luminosity ratio is found to be 3.1. The solution of the light curve yields a distance of $925 \\pm 25$\\,pc much lower than the usually assumed distance of the Cyg\\,OB2 association. Whilst we confirm the existence of episodes of higher X-ray fluxes, the data reveal no phase-locked modulation with the 6.6\\,day period of the eclipsing binary nor any clear relation between the X-ray flux and the 6.7\\,yr radio cycle.}{The bright region of the secondary star is probably heated by energy transfer in a common envelope in this contact binary system as well as by the collision with the primary's wind. The existence of a common photosphere probably also explains the odd mass-luminosity relation of the stars in this system. Most of the X-ray, non-thermal radio, and possibly $\\gamma$-ray emission of Cyg\\,OB2 \\#5 is likely to arise from the interaction of the combined wind of the eclipsing binary with at least one additional star of this multiple system.} ",
        "introduction": "} The eclipsing nature of Cyg\\,OB2 \\#5 (= V\\,729\\,Cyg = BD\\,$+40^{\\circ}$\\,4220) was first reported by Miczaika (\\cite{Miczaika}). Photometric data were subsequently published by Hall (\\cite{Hall}), who gathered 143 $UBV$ observations. Hall (\\cite{Hall}) noted the existence of intrinsic variability as well as a slight asymmetry in the light curve which made the system appear brighter during the maximum that followed the secondary eclipse. The light curve obtained by Hall (\\cite{Hall}) revealed unequal eclipse depths, thus suggesting that the primary should have a larger surface brightness than the secondary.  From the analysis of a set of optical spectra, Bohannan \\& Conti (\\cite{BC}) concluded that the two components of Cyg\\,OB2 \\#5 should be of equal brightness, whilst they simultaneously inferred a mass ratio (primary/secondary) of $q = 4.3$. Therefore, the secondary star would appear to be far too luminous for its mass, and these authors suggested accordingly that the secondary might be an O-star on its way to becoming a Wolf-Rayet object. A re-analysis of the same spectra by Massey \\& Conti (\\cite{MC}) yielded a somewhat lower mass ratio of 3.3, but did not alter the basic conclusions of the Bohannan \\& Conti (\\cite{BC}) paper.  Leung \\& Schneider (\\cite{LS}) analysed Hall's (\\cite{Hall}) $UBV$ light curve adopting the mass ratio proposed by Bohannan \\& Conti (\\cite{BC}). Their description of the system was that of an overcontact binary with the primary accounting for almost 90\\,\\% of the total light, unlike the almost equal luminosities inferred from the strengths of the He absorption lines (Bohannan \\& Conti \\cite{BC}). Whilst the luminosity difference of 2.1\\,mag inferred by Leung \\& Schneider (\\cite{LS}) is in qualitative agreement with a large mass ratio, it is clearly at odds with the strong spectral signature of the secondary in the spectrum of the binary.  Vreux (\\cite{JMV}) studied the variations in the H$\\alpha$ emission line profile and suggested that the primary is currently losing material towards the secondary. This was interpreted as further evidence that the secondary star could be a Wolf-Rayet type object with an unusual spectrum produced by the accretion of the material.  Rauw et al.\\ (\\cite{Rauw}) analysed an extensive set of optical spectra and presented a revised orbital solution along with a discussion about the phase-locked spectroscopic variability. The new orbital solution yielded a mass ratio of $q = 3.31 \\pm 0.25$ and revealed a large difference in the systemic velocities of the primary and the secondary (the latter having a more negative $\\gamma$). Based on the relative strengths of classification absorption lines, Rauw et al.\\ (\\cite{Rauw}) assigned an O6.5-7 spectral type to the primary and Ofpe/WN9 to the secondary. They also compared the strengths of the primary and secondary absorption lines to infer a visual brightness ratio of $1.4 \\pm 0.6$. The complex behaviour of the absorption lines (some of which go into emission during part of the orbit) prevents however a more accurate determination of this ratio. The radial velocities of the emission lines imply that the emissions most likely originate in material that is moving from the primary towards the secondary. Rauw et al.\\ (\\cite{Rauw}) also reported on variations in the equivalent width of the He\\,{\\sc ii} $\\lambda$\\,4686 emission line, which suggest an occultation effect of the emission zone in addition to the modulation of the continuum light curve due to the photometric eclipses. The variations in the line EW suggest that the emission arises from a region close to the secondary star, slightly preceding this star on its orbit. One of the main motivations of the present photometric campaign was to study this behaviour using a set of specially designed narrow-band photometric filters.  Another goal of this paper is to study the variability of the X-ray emission of Cyg\\,OB2 \\#5. X-ray emission from this star was first discovered with the {\\it EINSTEIN} satellite in December 1978 during a calibration observation of Cyg\\,X-3 (Harnden et al.\\ \\cite{Harnden}). Since then, our target has been observed with many X-ray satellites, including {\\it ROSAT} (Waldron et al.\\ \\cite{Waldron}) and {\\it ASCA} (Kitamoto \\& Mukai \\cite{KM}). We conducted an observing campaign with {\\it XMM-Newton} to improve our knowledge of the X-ray emission of Cyg\\,OB2 \\#5.  Finally, we emphasize that important information about this enigmatic system can be derived from the radio domain. Indeed, Miralles et al.\\ (\\cite{Miralles}) noted that the radio emission of Cyg\\,OB2 \\#5 alternates between a low state with a spectral index that is roughly consistent with thermal wind emission and a high state with a much flatter (partially non-thermal) spectral index on a $\\sim 7$\\,yr period. Abbott et al.\\ (\\cite{Abbott}) and Miralles et al.\\ (\\cite{Miralles}) reported on the existence of a radio companion 0.8\\,arcsec to the NE of the main radio source (which is associated with the eclipsing binary). Subsequent observations by Contreras et al.\\ (\\cite{Contreras}) revealed that this radio companion has an elongated shape and lies in-between the short-period binary and a third star\\footnote{The latter star is 3 -- 4\\,mag fainter than the binary system, and its colours and magnitude suggest a B0-B2\\,V spectral type. The angular separation between this astrometric companion and the close binary system suggests an orbital period of several thousand years.}, which was first reported by Herbig (\\cite{Herbig}). Contreras et al.\\ (\\cite{Contreras}) suggested accordingly that the radio companion corresponds to the wind interaction zone between the binary system and the tertiary component. Recently, Kennedy et al.\\ (\\cite{Kennedy}) re-analysed all VLA observations of Cyg\\,OB2 \\#5 and showed that the primary radio source, associated with the eclipsing binary varies on a period of $6.7 \\pm 0.2$\\,yr whilst the secondary source remains constant. The variations in the main radio component can be represented by a simple model in which a fourth (unresolved) object orbits the eclipsing binary in an eccentric orbit and the varying radio emission results from a variable non-thermal emission produced in the wind interaction between the short period binary and the fourth star (Kennedy et al.\\ \\cite{Kennedy}). ",
        "conclusions": "} Using optical narrow-band photometry and spectroscopy, we have shown that Cyg\\,OB2 \\#5 is in a contact configuration, which has allowed us to solve one of the oldest mysteries associated with this system, i.e.\\ the apparent discrepancy between its mass ratio and its luminosity ratio. The binary appears to be in a rather `normal' stage of (likely case-A) mass transfer in binary evolution.  However, whilst we have made progress in understanding the system, there remain a number of major issues. The most important ones are: \\begin{itemize} \\item[$\\bullet$] New theoretical models need to be designed to simulate accurately the structure of a common envelope in an early-type, contact binary system. \\item[$\\bullet$] The origin of the period variations needs to be investigated by long-term photometric monitoring of the system. \\item[$\\bullet$] The origin of the X-ray emission in this system remains uncertain as is the cause of the variations in the radio emission (Kennedy et al.\\ \\cite{Kennedy}). The most likely explanations involve the existence of a fourth component in the system. However, to date no direct evidence that such a component exists has been obtained and again long term monitoring of the system is required. For instance, one needs to acquire optical spectroscopy to determine orbital solutions that sample the 6.7\\,yr cycle. \\item[$\\bullet$] And finally, the distance of Cyg\\,OB2 needs to be checked by studying more eclipsing binary systems in this important cluster. \\end{itemize} \\acknowledgement{We are grateful to the referee, Dr.\\ A.F.J.\\ Moffat for helpful comments, to Dr.\\ H.\\ Sana who took part in the 1998 OHP photometric observing campaign as well as to Dr.\\ S.\\ Dougherty and M.\\ Kennedy for sharing their results on the radio data of Cyg\\,OB2 \\#5 with us before publication. The Li\\`ege group acknowledges financial support from the FRS/FNRS (Belgium), as well as through the XMM and INTEGRAL PRODEX contract (Belspo). The travels to OHP were supported by the `Communaut\\'e Fran\\c caise' (Belgium). PE acknowledges support through CONACyT grant 67041.}"
    },
    "0812/0812.0326_arXiv.txt": {
        "abstract": "We present spectral energy distribution modelling of 6875 stars in $\\omega$ Centauri, obtaining stellar luminosities and temperatures by fitting literature photometry to state-of-the-art {\\sc marcs} stellar models. By comparison to four different sets of isochrones, we provide a new distance estimate to the cluster of $4850 \\pm 200$ (random error) $\\pm 120$ (systematic error) pc, a reddening of $E(B-V) = 0.08 \\pm 0.02$ (random) $\\pm 0.02$ (systematic) mag and a \\emph{differential} reddening of $\\Delta E(B-V) < 0.02$ mag for an age of 12 Gyr. Several new post-early-AGB candidates are also found. Infra-red excesses of stars were used to measure total mass-loss rates for individual stars down to $\\sim 7 \\times 10^{-8}$ M$_\\odot$ yr$^{-1}$. We find a total dust mass-loss rate from the cluster of $1.3 \\pm ^{0.8}_{0.5} \\times 10^{-9}$ M$_\\odot$ yr$^{-1}$, with the total gas mass-loss rate being $> 1.2 \\pm ^{0.6}_{0.5} \\times 10^{-6}$ M$_\\odot$ yr$^{-1}$. Half of the cluster's dust production and 30\\% of its gas production comes from the two most extreme stars -- V6 and V42 -- for which we present new Gemini/T-ReCS mid-infrared spectroscopy, possibly showing that V42 has carbon-rich dust. The cluster's dust temperatures are found to be typically $\\gtsim$ 550 K. Mass loss apparently does not vary significantly with metallicity within the cluster, but shows some correlation with barium enhancement, which appears to occur in cooler stars, and especially on the anomalous RGB. Limits to outflow velocities, dust-to-gas ratios for the dusty objects and the possibility of short-timescale mass-loss variability are also discussed in the context of mass loss from low-metallicity stars. The ubiquity of dust around stars near the RGB-tip suggests significant dusty mass loss on the RGB; we estimate that typically 0.20--0.25 M$_\\odot$ of mass loss occurs on the RGB. From observational limits on intra-cluster material, we suggest the dust is being cleared on a timescale of $\\ltsim 10^5$ years. ",
        "introduction": "\\subsection{Mass loss and globular clusters}  Galactic globular clusters (GCs) are unique stellar laboratories, containing roughly co-eval populations of stars at known distances, covering the range [Fe/H] $\\approx$ --2.3 to solar \\citep{Harris96}. They are ideal locations in which to probe the later stages of development of low-mass ($\\sim$0.8 M$_\\odot$) stars, and their large ($\\approx 10^{5}-10^{6}$) populations mean they can harbour objects in very short phases of evolution.  Stellar mass loss, and its implications for both stellar evolution and the cluster's fate are of particular importance. All stars over $\\sim$0.8 M$_\\odot$ are thought to lose $\\gtsim$30\\% of their mass on the Red and Asymptotic Giant Branches (RGB/AGB) (e.g.\\ \\citealt{Rood73,LDZ94,CDA08}). Warm giants may drive an outflow of $\\sim 10^{-9}$ M$_\\odot$ yr$^{-1}$ via magnetic or hydrodynamic processes, (e.g.~\\citealt{DHA84,MCP06}), while the coolest, most luminous giants are thought to be able to sustain a wind of at least $\\sim 10^{-6}$ M$_\\odot$ yr$^{-1}$ through a combination of pulsations and radiation pressure on circumstellar dust (e.g.~\\citealt{GW71,BW91}). Mid-infrared (IR) spectra can be used to identify the species, temperature, mass and nature of these circumstellar dust grains and improve estimates of the total mass-loss rate from the star \\citep{LPH+06,vLMO+06}.  Mass loss has important consequences for the remainder of the star's evolution and the chemical enrichment of interstellar material. Significant mass loss can decrease the number of thermal pulses such a star undergoes, thus inhibiting the growth of the stellar core, and reducing dredge-up, which affects the chemistry and characteristics of the post-AGB and planetary nebula phase \\citep{VW93}. Mass-loss evolutionary processes also determine whether the star remains an oxygen-rich AGB star or becomes a carbon star \\citep{vLZW+98,FCL+98}. Ultimately, mass loss will limit the white dwarf mass. Many aspects of the mass-loss process remain poorly understood. There is no clear consensus even on the primary driving mechanism behind mass loss, likely pulsation or continuum-driving of dust. By examining mass loss along the giant branches of populations of different metallicities, we can gain insight into the poorly-determined relationships between metallicity and the mass-loss rate, the wind speed and gas-to-dust ratio of the stars; and the correlation between mass loss and gravity (e.g.~\\citealt{MvL07}). The clearing of lost mass from within the cluster will also liberate mass from the cluster and could exacerbate its evaporation. Study of the clearing can also allow us to infer conditions in the cluster's local environment \\citep{OKYM07}.  The \\emph{Spitzer Space Telescope} \\citep{WRL+04,GRW+07} has, for the first time, allowed a complete census of dust production within GCs, due to its unprecedented sensitivity and angular resolution. A comprehensive analysis of mass loss from stars in 47 Tuc was recently performed by \\cite{ORFF+07} (hereafter O+07). In this study, mid-IR flux excesses from \\emph{Spitzer} data were used to construct the mass-loss rate from individual stars in the cluster, attaining an empirical relation for mass loss from metal-poor giant stars. This relation suggests that significant dust production occurs along a considerable part of the giant branch. We compare our results to this analysis in Section \\ref{MdotRGBSect}.    \\subsection{Omega Centauri}  On the basis of our own \\emph{Spitzer} data, published in \\cite{BMvL+08} (hereafter B+08), and new mid-IR spectra of the two stars (V6 \\& V42) with the strongest IR excess, we here present an estimate of the mass-loss evolution in the most massive Galactic globular cluster: $\\omega$ Centauri.  $\\omega$ Cen is unique in the wealth of information available -- it is comparatively nearby at a relatively well-known distance of $\\sim$5 kpc \\citep{Harris96,vLlPB+00,vdVvdBVdZ06,dPPS+06}; its high radial velocity of $v_{\\rm LSR} = +232$ km s$^{-1}$ \\citep{MM86,vdVvdBVdZ06} allows easy spectroscopic membership and radial velocity determination (\\citealt{vLvLS+07}; hereafter vL+07), and it also has independent membership determinations from proper motion measurements (\\citealt{vLlPB+00}; hereafter vL+00). The importance of this is highlighted in B+08, which also contains an assessment of the effects of stellar blending on our longer-wavelength (lower-resolution) data. The cluster contains a statistically-significant sample of stars in most advanced stages of evolution (vL+07). Coupling this with a large existing photometric dataset, it is possible to positively identify cluster members and probe the stellar population well down the giant branch at all wavelengths relevant to spectral energy distribution (SED) modelling.  The bulk of the stars in the cluster have a relatively low metallicity -- [Fe/H] $\\approx$ --1.7 to --1.6 \\citep{Norris96,SSC+00}. However, a helium-enriched, metal-intermediate population at [Fe/H] $\\approx$ --1.2 \\citep{Norris96,Norris04} also exists; along with a metal-rich component, the `anomalous' RGB (RGB-a), with metallicities up to [Fe/H] $\\approx$ --0.7, which together comprise perhaps 10\\% of the cluster \\citep{LJS+99,PFB+00,PPH+02}. Variations in surface abundances are also present, including oxygen-rich, M-type stars with titanium oxides; stars enhanced in CH and/or CN; and genuine carbon stars with molecular carbon (C$_2$; vL+07).  The origin of these sub-populations is undecided \\citep{SPF+05,SdCNC06a,VPK+07}, though there is evidence that $\\omega$ Cen may be the remnant nucleus of a tidally disrupted dwarf spheroidal galaxy (dSph) \\citep{ZKD+88,Freeman93}. Understanding $\\omega$ Cen may therefore assist our understanding of mass loss and chemical enrichment in other nucleated dSphs, as well as in the earlier, more metal-poor Universe.    \\subsection{Individual stars}  Within this paper, we also present the first mid-infrared spectra of the two brightest, most IR-excessive stars in the cluster, V6 (LEID 33062, ROA 162) and V42 (LEID 44262, ROA 90). \\cite{GF73,GF77} measured the IR colours of V6, confirming it as a very bright IR source and showing it to be variable at near-IR wavelengths, with a possible $L$-band excess. Its period is uncertain, being listed as 73.513 days in \\cite{SawyerHogg73} and 100--120 days in \\cite{DFLE72}. Classified as an M4--5 emission line variable, it has a possible radial velocity variation of up to 40 km s$^{-1}$ with a mean velocity of +213.6$\\pm$3.7 km s$^{-1}$ (Dickens et al.~1972; \\citealt{Webbink81}; quoted uncertainty $\\sim$13 km s$^{-1}$). Its temperature has been estimated at 3300--3600 K with log($g$) of 0.0 (assuming $M = 0.8$ M$_\\odot$; \\citealt{PCM+80,Frogel83}). It contains relatively large amounts of H$_2$O compared to the other cluster long-period variables (LPVs) and is CN and NH enhanced \\citep{CB86}. It is also known to be a TiO variable and shows variable hydrogen emission (\\citealt{LloydEvans83b,LloydEvans83c,LloydEvans86}). V6 appears to belong to the metal-intermediate sub-population, with [Fe/H] $\\sim$ --1.19 \\citep{ZW84,NdC95b,VWS02}.  \\cite{Feast65} implies V42 should be classified as a semi-regular variable of type SRd, but its spectrum is somewhat later than the F--K-type this implies. An M1--2.5 emission line variable, it may also exhibit radial velocity variations between about +253 and +272 km s$^{-1}$ (Dickens et al.~1972), though this variation may be largely attributable to insufficient signal-to-noise. V42 may represent a star bridging the gap between SRd stars and emission-line LPVs. As an apparent fundamental-mode pulsator, its 148.64$\\pm$0.03 day period (vL+00) is the longest in $\\omega$ Cen and among the longest in globular cluster variables as a group \\citep{Clement97}. Dickens et al.\\ also report that the TiO bands weaken and hydrogen emission lines are present only near optical photometric maximum (see also \\citealt{LloydEvans83c,LloydEvans83d}), suggesting an additional opacity source. It is also CN and NH enhanced \\citep{CB86}. \\cite{CF83} estimate log($g$) = 0.5 and $T$ = 3950 K, whereas \\cite{MW85} calculate a much lower temperature of 2818 K, based on IR data, and show V42 undergoes substantial variability -- some $\\sim$60\\% of its mean luminosity -- even in the $L$-band. No reliable information on the star's metallicity is available, although vL+07 place it at [Fe/H] $\\sim -1.25$; a high-resolution optical spectrum shows significant H-$\\alpha$ emission, attributable to shocks propagating in the stellar wind (\\citealt{MvL07}, hereafter MvL07).  The cluster also contains a number of other interesting objects. Five carbon stars have now been reported (vL+07). Several post-AGB stars are also known. Most famously, Fehrenbach's Star (LEID 16018, HD 116745) appears to have already undergone thorough mixing and mass loss \\citep{FD62,DP73,GW92}. Another, V1 (LEID 32029), is an irregularly-pulsating star with multiple periods that is thought to have undergone gas-dust separation, but not thought to have undergone third dredge-up (surface enrichment during thermal pulses) and may thus be a post-early-AGB star (\\citealt{Gonzalez94,MHLN98}; \\citealt{TKD+06,TKD+07}; vL+07). V29 (LEID 43105), V43 (LEID 39156), V48 (LEID 46162) and V92 (LEID 26026) have also been suggested to be post-AGB stars \\citep{Gonzalez94,GW94}.  The remainder of the paper is organised as follows: Section 2 describes the SED input data and models; Section 3 details the corrections to the photometry and the process of creation of the SEDs; Section 4 compares our observed temperatures and luminosities with those predicted by a variety of stellar evolution models, providing a new estimate of the distance and reddening to the cluster; Section 5 presents new mid-infrared spectra of V6 and V42, and discusses subsequent estimation of mass loss from individual stars; Section 6 discusses the implications of the dataset, including calculating the total mass-loss rate of the cluster, correlations with various stellar parameters, and the evolutionary nature of mass loss; finally, Section 7 presents our conclusions. ",
        "conclusions": "We have here presented stellar parameters derived from spectral energy distribution fitting to stars over two orders of magnitudes in luminosity down the RGB and AGB of $\\omega$ Centauri, creating a physical HRD of luminosity versus temperature, outlining the RGB, HB and AGB in great detail and accuracy, and identifying several new post-AGB star candidates. From isochrone and spectral energy distribution fitting, we estimate the following parameters for the cluster: \\begin{itemize} \\item Distance: $d = 4850 \\pm 200$ (statistical) $\\pm 120$ (systematic) pc; \\item Reddening: $E(B-V) = 0.08 \\pm 0.02$ mag (statistical) $\\pm 0.02$ mag (systematic) $\\pm 0.02$ mag (differential); \\item Total dust mass loss: $\\dot{M}_{\\rm dust} = 9 \\pm ^{6} _{4} \\times 10^{-10}$ M$_\\odot$ yr$^{-1}$; \\item Total gas mass loss: $\\dot{M}_{\\rm gas} \\sim 1.2 \\times 10^{-5}$ M$_\\odot$ yr$^{-1}$ ($\\sim 2/3$ from dusty winds, $\\sim 1/3$ from chromospheric mass-loss); \\item Timescale to clear the intra-cluster medium from cluster: $\\ltsim 10^5$ yr; \\end{itemize} under the following assumptions: \\begin{itemize} \\item gas-to-dust ratio, $\\psi \\propto 10^{-{\\rm [Fe/H]}}$; \\item wind velocity, $v \\propto L^{1/4} \\psi^{-1/2}$; \\item absence of substantial mid-IR variability; \\item constant dust chemistry along the giant branch; \\item inner dust envelope temperature of 1000 K; \\item our low wind velocities ($\\sim 1$ km s$^{-1}$) are accurate and still produce an accurate mass-loss rate. \\end{itemize} We derive inner dust envelope temperatures for a handful of stars (including V6 and V42), suggesting they are typically $\\gtsim$ 600 K.  We show that V6 and V42 contribute $\\sim$25\\% of the cluster's dust production. V42 may be variable at mid-infared wavelengths on timescales of as little as a few weeks, suggesting equally variable dust production, meaning a single estimate of mass-loss rate and wind conditions may not accurately reflect the long-term conditions in this and similar stars. V42, an oxygen-rich star, may also be producing carbon-rich dust.   Dusty mass loss appears to start suddenly at a threshold of $\\sim$4400 K, or $\\sim$ 1000 L$_\\odot$, for a significant fraction of stars. There appear to be additional star-to-star differences, suggesting that other factors also influence the mass-loss rate. Empirical literature relations do not accurately reflect the dust production rates we measure in $\\omega$ Cen.  By comparing our stellar parameters with optical line strengths, we have deduced that the Ba-rich stars in $\\omega$ Cen belong to the more metal-rich populations and not to the AGB. These stars are characterised by dustier winds.  Mass loss along the RGB appears quite efficient with typically 0.20--0.25 M$_\\odot$ lost on the RGB and $\\gtsim$0.05 M$_\\odot$ on the AGB, possibly leading to large numbers of AGB-\\emph{manqu\\'e} and post-early-AGB stars. Dust production in $\\omega$ Cen appears not confined to the AGB, with most dust producing stars in the cluster near the RGB-tip.  \\ \\\\  \\noindent {\\bf Acknowledgements}  \\noindent We thank Aaron Dotter, Noriyuki Matsunaga, Greg Sloan, and the many others who have given help in this project for their words of wisdom. I.M.\\ acknowledges a PPARC/STFC studentship award. L.D. acknowledges financial support from the Fund for Scientific Research -- Flanders (Belgium). M.L.B., C.E.W., and R.D.G.\\ are supported in part by NASA through \\emph{Spitzer} contracts 1276760, 1256406, and 1215746 issued by JPL/Caltech to the University of Minnesota. A.K.D.\\ acknowledges research support from \\emph{Spitzer} contract 1279224. Finally, we thank the referee for his very helpful advice. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. This manuscript is also based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Science and Technologies Facilities Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), CNPq (Brazil), and CONICET (Argentina)."
    },
    "0812/0812.2609_arXiv.txt": {
        "abstract": "A new model for source counts from 8-1100 $\\mu$m is presented, which agrees well with source-count data and the observed background spectrum.  The model is similar to that of Rowan-Robinson (2001), but with different evolution for each of the four assumed infrared template types.   The evolution is modified in two ways; (i) the exponential factor is modified so that it tends to a constant value at late times, (ii) the power-law factor is modified so that it tends to zero at redshift $z_f$, rather than 0 as assumed by Rowan-Robinson (2001). I find strong evidence from the 850 and 1100 $\\mu$m counts, and from the infrared background, that $z_f$ = 4-5, with some preference for a value at the low end of the range, implying that star-forming galaxies at z $>$ 5 are not significant infrared emitters, presumably due to a low opacity in dust at these early epochs.  The model involves zero or even negative evolution for starbursts and AGN at low redshifts ($<$0.2), suggesting that the era of major mergers and strong galaxy-galaxy interactions is over. ",
        "introduction": "Source-counts at infrared and submillimetre wavelengths, combined with the spectrum of the infrared background, give us important constraints on the star-formation history of the universe. First indications of strong evolution in the properties of star-forming galaxies came from IRAS 60 $\\mu$m counts (Hacking and Houck 1987, Lonsdale et al 1990, Rowan-Robinson et al 1991, Franceschini et al 1991, 1994, Pearson and Rowan-Robinson 1996)  ISO gave us deep counts at 15 $\\mu$m providing strong evidence for evolution (Oliver et al 1997,  Rowan-Robinson et al 1997,  Guiderdoni et al 1998, Aussel et al 1999, Elbaz et al 1999, 2002, Serjeant et al 2000, Gruppioni et al 2002, Lagache et al 2003) and useful counts at 90 and 175 $\\mu$m (Kawara et al 1998, Dole et al 2001, Efstathiou et al 2000a, Heraudeau et al 2004).  850 $\\mu$m counts with SCUBA on JCMT (Smail et al 1997, Hughes et al 1998, Eales et al 1999, 2000, Barger et al 1999, Blain et al 1999, Fox et al 2002, Scott et al 2002, Smail et al 2002, Cowie et al 2002, Webb et al 2003, Borys et al 2004, Coppin et al 2006) have given important insight into the role of cool dust and also strong constraints on the high-redshift evolution of hyperluminous infrared galaxies. Evidence for luminosity evolution of submillimetre galaxies was given by Ivison et al (2002). 1200 $\\mu$m counts with MAMBO have been reported by Greve et al (2004) and modelled in terms of a strongly evolving population. Recently the AZTEC collaboration have reported differential counts at 1100 $\\mu$m (Perera et al 2008, Austermann et al 2008), which provide even stronger constraints on high redshift evolution.  The detection of the infrared background with COBE (Puget et al 1996, Fixsen et al 1998, Hauser et al 1998, Lagache et al 1999) provided an important constraint on evolutionary models (Guiderdoni et al 1997, 1998, Franceschini et al 1998, 2001, Dwek et al 1998, Blain et al 1999, Gispert et al 2000, Dole et al 2001, Rowan-Robinson 2001,  Chary and Elbaz 2001, Elbaz et al 2002,  King and RR 2003, Balland et al 2003, Xu et al 2003, Lagache et al 2003).  Dole et al (2006) used stacking analysis on deep {\\it Spitzer} counts at 24, 90 and 160 $\\mu$m to estimate the integrated background radiation from sources at these wavelengths.  With {\\it Spitzer} we also have deep counts at 8, 24, 70, 160 $\\mu$m (Fazio et al 2004, Chary et al 2004, Marleau et al 2004, Papovich et al 2004, Dole et al 2004, Le Floch et al 2005, Frayer et al 2006, Shupe et al 2008), with 24 $\\mu$m providing an especially complete picture of the contribution of individual sources to the  infrared background.  While some pre-{\\it Spitzer} models were quite successful at 70 and 160 $\\mu$m, none captured full details of the 24 $\\mu$m counts.  Lagache et al (2004) provided early revised models in the light of the Spitzer data.  Rowan-Robinson (2001) modeled infrared source-counts and background in terms of four types of infrared galaxy: quiescent galaxies in which the infrared radiation (infrared 'cirrus') is emission from interstellar dust illuminated by the general stellar radiation field, starbursts for which M82 is the prototype, more extreme, higher optical depth starbursts with Arp 220 as prototype, and AGN dust tori.  The same evolution history, intended to reflect the global star-formation history, was used for each galaxy type.  This model was consistent with counts, luminosity functions and colour-luminosity relations from IRAS, counts from ISO and  available integral 850 $\\mu$m counts.  However the advent of deep source-count data from {\\it Spitzer} showed that this model, along with other pre-{\\it Spitzer} models, failed, especially at 24 $\\mu$m.  In this paper I show how the Rowan-Robinson (2001) model has to be modified to achieve consistency with {\\it Spitzer} and other modern data and give predictions for the {\\it Herschel} and {\\it Planck} wave-bands. Some preliminary results were given in Shupe et al (2008).  A cosmological model with $\\Lambda$ = 0.7, $h_0$=0.72 has been used throughout. ",
        "conclusions": "I have presented a new model for source counts from 8-1100 $\\mu$m, which agrees well with source-count data and the observed background spectrum.  The model is similar to that of Rowan-Robinson (2001), but with different evolution for each of the four assumed infrared template types.   The evolution is modified in two ways; (i) the exponential factor is modified so that it tends to a constant value at late times, (ii) the power-law factor is modified so that it tends to zero at redshift $z_f$, rather than 0 as assumed by Rowan-Robinson (2001). I find strong evidence from the 850 and 1100 $\\mu$m counts and from the infrared background that $z_f$ = 4-5, with some preference for a value at the low end of the range, implying that star-forming galaxies at z $>$ 5 are not significant infrared emitters, presumably due to a low opacity in dust at these early epochs.  The model involves zero or even negative evolution for starbursts and AGN at low redshifts ($<$0.2), suggesting that the era of major mergers and strong galaxy-galaxy interactions is over.  {\\it Herschel} and {\\it Planck} submillimetre counts will provide much stronger tests of these models."
    },
    "0812/0812.3219_arXiv.txt": {
        "abstract": "{Within the hierarchical framework for galaxy formation, merging and tidal interactions are expected to shape large galaxies to this day. While major mergers are quite rare at present, minor mergers and satellite disruptions -  which result in stellar streams -  should be common, and are indeed seen in both the Milky Way and the Andromeda Galaxy. As a pilot study, we have carried out ultra-deep, wide-field imaging of some spiral galaxies in the Local Volume, which has revealed external views of such stellar tidal streams at unprecedented  detail, with data taken at small robotic telescopes (0.1-0.5-meter) that provide exquisite surface brightness sensitivity. The goal of this project is to undertake the first systematic and comprehensive imaging survey of stellar tidal streams, from a sample of $\\sim$ 50 nearby Milky-Way-like spiral galaxies within 15 Mpc, that features a surface brightness sensitivity of $\\sim$ 30 mag/arcsec$^{2}$. The survey will result in estimates of the incidence, size/geometry and stellar luminosity/mass distribution of such streams. This will not only put our Milky Way and M31 in context but, for the first time, also provide an extensive statistical basis for comparison with  state-of-the-art, self-consistent cosmological simulations of this phenomenon.} ",
        "introduction": "\\label{sec:1}  Within the hierarchical framework for galaxy formation (e.g., White \\& Rees 1978), the stellar bodies of galaxies are expected to form and evolve through dark-matter-driven mass in-fall and successive coalescence of smaller, distinct sub-units that span a wide mass range. Mergers of initially bound sub-halos (which we refer to as {\\it satellites}; they consist of dark matter, gas, and in most cases stars) are effected by dynamic friction, either through gradual orbital decay or by a single encounter (depending on the initial orbit), its eccentricity and the satellite-to-main-galaxy mass ratio. It is likely that a satellite becomes disrupted by the tidal forces of the larger companion before its orbit spirals all the way to the center. If such a tidal disruption is complete, and no bound satellite is left, dynamical friction ceases to act. If the disruption is only partial at this epoch, the surviving satellite fragment displays extensive tidal tails, leading and trailing its current position in the galactic halo.  While in $\\Lambda$-Cold Dark Matter ($\\Lambda$CDM) the interaction rate is expected to drop to the present-day epoch, such disruption of satellites should still occur around normal spiral galaxies. The  fossil records of these merger events may be detected nowadays in the form of distinct coherent stellar structures in the outer regions of massive systems. The most spectacular cases of tidal debris are long, dynamically cold stellar streams from a disrupted dwarf satellite, which have wrapped around the host galaxy's disk and roughly trace the orbit of the progenitor satellite. The now well-studied Sagittarius tidal stream surrounding the Milky Way (Majewski et al. 2003) and the giant stream in Andromeda galaxy (Ibata et al. 2007) are archetypes of these satellite galaxy merger 'fossils'  in the Local Group. They provide sound qualitative support for the scenario that tidally disrupted dwarf galaxies are important contributors to stellar halo formation in the Local Group spirals.   State-of-the-art, high-resolution numerical simulations of galaxy formation, built within the $\\Lambda$CMD context (e.g. Moore et al, 1999; Springel et al. 2008), can guide the quest for observational signatures of such star-streams (e.g. Bullock \\& Johnston 2005; Johnston et al. 2008). Recent  simulations have demonstrated that the characteristics of substructure currently visible in the stellar halos are sensitive to the last (0-8 Gyr ago) merger histories of galaxies, a timescale that corresponds to the last few to tens of percent of mass accretion for a spiral galaxy like the Milky Way. While stellar streams in the Milky Way and Andromeda can be studied in detail, comparison with cosmological models is limited by 'cosmic variance'. However, the current models imply that a survey of 50-100 parent galaxies reaching to a surface brightness of 30 mag arcsec$^{-2}$ would reveal many tens of tidal features, perhaps nearly one detectable stream per galaxy (Johnston et al. 2008). However, a specific comparison of these simulations with observations is missing because no suitable data sets exist. Such a comparison, which could quantify the present sub-halo merger rate, is not only important as a test of $\\Lambda$CDM models, but also as a more direct probe of how resilient disks are to minor mergers.  \\begin{figure}[t] \\begin{center} \\epsfig{figure=Slide1.ps, height=9cm, width=12cm, angle=0} \\caption{ \\footnotesize{{\\it(left)} Deep image of the stellar tidal stream around NGC 5907  obtained with the 0.5-meter Black Bird Observatory (BBO) telescope  (Martinez-Delgado et al 2008a). A  N-body model of this structure is shown in Fig. 2; {\\it (right, top)} A low-galactic latitude stellar tidal stream of NGC 4013, discovered by our team from deep images taken with the BBO telescope; {\\it (right, bottom)} Deep images taking with a FSQ-106ED telescope of only 10cm aperture allowed the discovery of a  giant tidal stream in the halo of the spiral galaxy Messier 63 (Martinez-Delgado et al. 2009, in preparation). A colour inset of the disk of each galaxy has been inserted with reference purpose.}} \\label{fig_test} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0812/0812.1500_arXiv.txt": {
        "abstract": "We present optical, X-ray, high energy ($\\lessapprox 30$ GeV) and very high energy ($\\gtrapprox 100$ GeV; VHE) observations of the high-frequency peaked blazar Mrk 421 taken between 2008 May 24 and June 23. A high energy $\\gamma$-ray signal was detected by AGILE with $\\sqrt{TS}=4.5$ on June 9--15, with $F(E>100~\\mathrm{MeV})= 42^{+14}_{-12}\\times 10^{-8}$ photons cm$^{-2}$ s$^{-1}$. This flaring state is brighter than the average flux observed by EGRET by a factor of $\\sim$3, but still consistent with the highest EGRET flux. In hard X-rays (20-60 keV) SuperAGILE resolved a 5-day flare (June 9-15) peaking at $\\sim$ 55 mCrab. SuperAGILE, RXTE/ASM and Swift/BAT data show a correlated flaring structure between soft and hard X-rays. Hints of the same flaring behavior are also detected in the simultaneous optical data provided by the GASP-WEBT. A Swift/XRT observation near the flaring maximum revealed the highest 2-10 keV flux ever observed from this source, of $2.6$ $\\times 10^{-9}$ erg cm$^{-2}$ s$^{-1}$ (i.e. $> 100$ mCrab). A peak synchrotron energy of $\\sim$3 keV was derived, higher than typical values of $\\sim$0.5-1 keV. VHE observations with MAGIC and VERITAS on June 6-8 show the flux peaking in a bright state, well correlated with the X-rays. This extraordinary set of simultaneous data, covering a twelve-decade spectral range, allowed for a deep analysis of the spectral energy distribution as well as of correlated light curves. The $\\gamma$-ray flare can be interpreted within the framework of the synchrotron self-Compton model in terms of a rapid acceleration of leptons in the jet. ",
        "introduction": "Mrk 421 is a nearby blazar ($z=0.031$) and one of the brightest BL Lac objects given its distance of 134.1 Mpc ($H_0=71$ km s$^{-1}$  Mpc$^{-1}$, $\\Omega_{m}=0.27, \\Omega_{\\lambda}=0.73$) . It was observed in $\\gamma$-rays by EGRET (Lin et al. 1992) and it was the first extragalactic object detected at $E> 500$ GeV \\cite{punch92}. It belongs to the class of High-energy peaked BL Lac objects (HBLs) (Padovani \\& Giommi 1995), i. e. radio-loud active galactic nuclei with high radio and optical polarization. Its spectral energy distribution (SED) is double-humped with a first peak usually in the soft to medium X-ray range, and a second one at GeV-TeV energies (Sambruna et al. 1996; Fossati et al. 1998). The first hump is commonly interpreted as due to synchrotron radiation from high-energy electrons in a relativistic jet, while the origin of the second peak is still uncertain. In the leptonic scenario it is interpreted as inverse Compton (IC) scattering of the synchrotron (Synchrotron self-Compton, SSC) or external photons (External Compton, EC)  by the same population of relativistic electrons. The observed correlated variability between X-rays and TeV $\\gamma$-rays (Maraschi et al. 1999; Fossati et al. 2008, Wagner 2008) is well explained in the SSC framework (Ghisellini et al. 1998), whereas the EC scenario is unlikely to apply in HBLs, due to the low density of ambient photons. Alternatively, hadronic models invoke proton-initiated cascades and/or proton-synchrotron emission (Aharonian 2000, M\\\"{u}cke et al. 2003).   Leptonic and hadronic scenarios for HBLs predict different properties of the $\\gamma$-ray emission in relation to emissions in other energy bands. Specifically, the hadronic models (as opposed to the SSC ones) predict a flatter slope of the $\\sim 100$ MeV IC emission than that of the synchrotron emission in the optical-UV energy bands. $\\gamma$-ray observations of flaring BL Lacs and simultaneous multiwavelength data are thus the keys to investigating these two scenarios.  A hard X-ray flare of Mrk 421 was detected by SuperAGILE on 2008 June 10 (Costa et al. 2008). This detection was later followed by a detection in $\\gamma$-rays (Pittori et al. 2008) by the AGILE/GRID (Gamma Ray Imaging Detector) and prompted a ToO observation by Swift/XRT, complementing the ongoing multifrequency observing campaign of Mrk 421 with WEBT (optical), MAGIC and VERITAS (TeV). We report on the observations and the analysis of these data, complemented by the publicly-available data from RossiXTE/ASM (2-12 keV) and Swift/BAT (15-50 keV), and discuss the spectral energy distribution of the source during this bright $\\gamma$-ray flare. ",
        "conclusions": "Mrk 421 showed a very interesting broad-band activity during the first half of 2008 June as derived from AGILE data combined with those of WEBT, Swift, MAGIC and VERITAS. Using our multi-frequency data we were able to derive time-resolved SEDs (Fig. 2). We distinguish two time periods: {\\it period 1}: 2008 June 6, with the inclusion of optical, X-ray (XTE and BAT) and TeV data (VERITAS); and {\\it period 2}: 2008 June 9--15, including optical, UV, X-rays (XRT and SuperAGILE) and gamma-ray data (AGILE). The source shows a very interesting time-variable broad band emission that appears to be in overall agreement with an SSC model. The optical, soft and hard X-ray bands strongly constrain the SED around the synchrotron peak, and its daily variability reveals the physical processes of Mrk~421. Possible correlated variability is shown in Fig.~1 between the optical (an overall decreasing trend with superimposed spikes of emission), the X-rays (several emission peaks lasting few days), and the high-energy parts of the spectrum. Based on the physical constraints obtained for the synchrotron peak, we can model both the HE and VHE $\\gamma$-ray emission. The data collection and broad-band SED extends over 12 decades in energy. Taking advantage of the overlapping MAGIC and VERITAS observations, we present a combined VHE light curve using the current generation of Northern imaging air Cherenkov systems.        We first model the synchrotron peak of emission using the period 1 optical, soft and hard X-ray data. The short time-variability (Fig.~1) constrains the size of the emitting region to $R<cT\\delta\\sim 5\\times 10^{16}(\\delta/20)$ cm. Hence,  we consider a one-zone SSC model (Tavecchio et al. 1998) based on a blob of comoving size $R=4\\times 10^{16}$ cm, with a relativistic Doppler factor $\\delta=20$  and characterized by non-thermal comoving electron energy distribution function described by a double power law: \\begin{equation} n_{e}(\\gamma)=\\frac{K\\gamma_{b}^{-1}}{(\\gamma /\\gamma_{b})^{p_1}+(\\gamma /\\gamma_{b})^{p_2}} \\end{equation} where the comoving Lorentz factor ($\\gamma$) varies in the range  $\\gamma_{min}=4 \\times 10^{3}<\\gamma<\\gamma_{max}=1.3\\times 10^{6}$, the normalization (density) constant  $K=4\\times 10^{-4}$ cm$^{-3}$, and the break energy $\\gamma_{b}= 3.6\\times 10^{5} $ and with $p_1=2.22$, $p_2=4.5$, the low-energy and high-energy power-law indices, respectively (see Table 1).  With these parameters we found that the data for period 1  are best fitted with a comoving magnetic field $B=0.1$ G.  Variability may be caused by several factors; we consider two cases: (A) hardening/softening of the electron energy distribution function caused by particle acceleration processes; (B) increase/decrease of the comoving particle density, as a consequence of additional particle injection/loss by shock processes.  We expect TeV variability to be comparable with the X-ray one if case (A) applies: this is because the emission is in the Klein-Nishina regime. Alternately, for case (B) we expect the TeV relative variability ($\\Delta F/F$) to be a factor of 2 greater than that of the X-ray flux variability.  Our AGILE, MAGIC and VERITAS data appear to support case (A). We compare the SEDs for period 1 and period 2, to better assess the spectral evolution. In Fig.~2 we show our optimized modelling of the time-resolved synchrotron peak and consequent SSC high-energy emission for the period 1 as well as for 2008 June 12--13 of the period 2. In the last case, the adopted model parameters are $p_1=2.1$, $p_2=5$, $\\gamma_{b}=4.2\\times 10^{5}$, $K=6\\times 10^{-4}$ cm$^{-3}$. Our model predicts an even larger TeV flux for period 2 (no TeV observations exist, however) than detected in period 1.    A detailed discussion of the complex optical vs. X-ray variability of Mrk~421 as shown in Fig.~1 will be presented elsewhere. We notice here a few remarkable points. The optical light curve shows variations of the order of $10\\%$ on a time scale $\\sim$few days, superimposed on a long decay during the entire period. Individual soft and hard X-ray peaks result in increased fluxes by a factor of $\\sim$2.5 and $\\sim$5, respectively: no long term decay appears. This different behaviour of the X-ray radiation and the bulk of the optical emission may interestingly suggest a more complex scenarios than A) and B) ones: optical and X radiation comes from two different jet regions, each one characterized by its own variability. A possible scenarios is one in which the inner jet region would produce the X-rays and it would be at least partially transparent to the optical radiation. In contrast, the outer region can only produce lower-frequency emission. The signature of the X-ray events visible in the optical light curve would come from the inner region and would be diluted by the optical radiation emitted from the outer region (see Villata \\& Raiteri 1999 for the case of Mrk 501; Villata et al. 2004).   Interestingly, the 2-10 keV flux measured by XRT on June 12 -- 13,  $\\sim 2.6\\times 10^{-9}$ erg cm$^{-2}$ s$^{-1}$, is higher than all previous observations ($< 2 \\times 10^{-9}$ erg cm$^{-2}$ s$^{-1}$; Fossati et al. 2008; Lichti et al. 2008). A joint analysis of the XRT and SuperAGILE data, covering the range from 0.7 to 60 keV, provides a best-fit spectral model consistent with a log-parabolic shape, with parameters implying a peak energy $\\sim$ 3 keV, in good agreement with the steeper positive correlation between the peak energy and the maximum of the SED found by Tramacere et al. 2007 (see their Fig.~3), although our value of the peak energy shows a significant shift with respect to typical values of 0.5-1 keV for this source.   We conclude our analysis of the broad-band variable emission from Mrk~421 by emphasizing that our multi-telescope/instrument data show a very interesting variability that provides support for an SSC model of the source. The $\\gamma$-ray emission detected by AGILE during period 2 and the TeV emission detected during period 1 can be successfully modelled from the characteristics of the corresponding synchrotron peaks.   \\begin{figure}[hbt] \\begin{center} \\epsscale{1.3} \\plotone{f2.eps}  \\caption{SEDs of Mrk 421 obtained by combining the GASP-WEBT, SWIFT/UVOT, RossiXTE/ASM, XRT, SuperAGILE, BAT, GRID and VERITAS data in period 1 and period 2 (red empty cirles and black filled circles, respectively). Both  are one-zone SSC models (red dashed line for period 1 and black solid line for period 2).} \\end{center} \\end{figure}"
    },
    "0812/0812.4780_arXiv.txt": {
        "abstract": "We present a systematic temporal and spectral study of all \\swift-XRT observations of GRB afterglows discovered between 2005 January and 2007 December.  After constructing and fitting all light curves and spectra to power-law models, we classify the components of each afterglow in terms of the canonical X-ray afterglow and test them against the closure relations of the forward shock models for a variety of parameter combinations.  The closure relations are used to identify potential jet breaks with characteristics including the uniform jet model with and without lateral spreading and energy injection, and a power-law structured jet model, all with a range of parameters. With this technique, we survey the X-ray afterglows with strong evidence for jet breaks ($\\sim 12\\%$ of our sample), and reveal cases of potential jet breaks that do not appear plainly from the light curve alone (another $\\sim30\\%$), leading to insight into the missing jet break problem.  Those X-ray light curves that do not show breaks or have breaks that are not consistent with one of the jet models are explored to place limits on the times of unseen jet breaks.  The distribution of jet break times ranges from a few hours to a few weeks with a median of $\\sim 1$ day, similar to what was found pre-\\swiftnsp. On average \\swift GRBs have lower isotropic equivalent $\\gamma$-ray energies, which in turn results in lower collimation corrected $\\gamma$-ray energies than those of pre-\\swift GRBs.  Finally, we explore the implications for GRB jet geometry and energetics. ",
        "introduction": "One of the most surprising puzzles to emerge from the \\swift \\markcite{gehrels04}({Gehrels} {et~al.} 2004) mission and its dynamic study of Gamma-ray Bursts (GRBs) is the lack of expected jet breaks in X-ray afterglow emission.  It is vital to the entire study of GRBs to understand jet geometry, because of the inferred effects on the total output energy, GRB rate, afterglow structure, interactions with environment, and jet physical mechanisms. Pre-\\swift optical observations showed tens of cases of steepening in the light curves several days after the GRB triggers \\markcite{frail01,bloom01,zeh06}({Frail} {et~al.} 2001; {Bloom}, {Frail}, \\& {Sari} 2001; {Zeh}, {Klose}, \\& {Kann} 2006).  This steepening was interpreted as evidence for the collimation of the burst ejecta with physical half-angle $\\theta_j$.  The ejecta moves at relativistic velocities with a bulk Lorentz factor, $\\Gamma$, and the radiation is relativistically beamed into an angle $\\theta=1/\\Gamma$. As the ejecta sweeps up surrounding material, the fireball decelerates, with the beaming angle eventually exceeding the physical collimation angle, causing a sudden increase in the rate of decay of the flux (i.e. the jet break).  At the same time, sideways expansion of the ejecta with relativistic speeds also causes a sudden flux decrease \\markcite{sari99,rhoads99,zhang04}({Sari}, {Piran}, \\& {Halpern} 1999; {Rhoads} 1999; {Zhang} \\& {M{\\'e}sz{\\'a}ros} 2004). Most likely both of these effects contribute to the jet breaks, and therefore, both models must be considered. Breaks are expected to be achromatic based on the assumption that the afterglow emission regions and mechanisms are the same for various spectral regimes, and should therefore only reflect ejecta geometry. Achromaticity has indeed been confirmed in the optical/near-infrared bands in pre-\\swift GRBs \\markcite{kulkarni99,harrison01,klose04}({Kulkarni} {et~al.} 1999; {Harrison} {et~al.} 2001; {Klose} {et~al.} 2004).  In the \\swift era, it is X-ray afterglow light curves that provide the most homogeneous data set to study GRB afterglows.  With the rapid GRB triggers provided by the \\swift-BAT, and the autonomous prompt \\swift-XRT observations that frequently begin within 1-2 minutes of the trigger, X-ray afterglows have gone from sparsely sampled single power-laws (Beppo-SAX era, \\markcite{depasquale06}{de Pasquale} {et~al.} 2006) to a rich database of light curves with widely varying properties and durations.  A common canonical shape of the \\swift-XRT X-ray light curves emerged (\\markcite{nousek06,zhang06}{Nousek} {et~al.} (2006); {Zhang} {et~al.} (2006), Figure \\ref{fig:canon}) with five components (an initial steep decay, a shallow-decay ``plateau'' phase, a normal decay, a jet-like decay component, and flares) that could be used to explain the overall structure of the afterglows.  While elements of this canonical picture are seen in most X-ray afterglows, few afterglows contain all 5 components. Surprisingly, fully ``canonical'' jet breaks are rarely observed in the XRT light curves \\markcite{burrows07,liang08,evans09}({Burrows} \\& {Racusin} 2007; {Liang} {et~al.} 2008; {Evans} {et~al.} 2008).  In this work, we assume that all X-ray afterglows have inherently similar shapes (the canonical shape) with deviations in behavior due to environment, electron spectral shape, spectral regime, presence of energy injection, and jet properties.  We also explore observational biases such as late beginning and early ending of observation, flares, and observing gaps that lead to missing portions of individual light curves which can produce ambiguities in the identification of segments.  We assume that the light curve segments following the initial steep decay (Figure \\ref{fig:canon}, segments II-IV) are due to external forward shocks. Although we recognise that several alternative models have recently been proposed to explain the origin of the X-ray afterglows (e.g. \\markcite{genet07,shao07,ghisellini07,liang07,uhm07,liang08,panaitescu08,kumar08}{Genet}, {Daigne}, \\& {Mochkovitch} (2007); {Shao} \\& {Dai} (2007); {Ghisellini} {et~al.} (2007a); {Liang}, {Zhang}, \\& {Zhang} (2007); {Uhm} \\& {Beloborodov} (2007); {Liang} {et~al.} (2008); {Panaitescu} (2008); {Kumar}, {Narayan}, \\& {Johnson} (2008)), as discussed below (see also \\markcite{liang07,liang08}{Liang} {et~al.} 2007, 2008), the X-ray data can be generally interpreted within the framework of the forward shock, so that invoking non-forward-shock models is not absolutely demanded by the data. We also focus solely on the X-ray behavior of the afterglow light curves, while acknowledging that chromatic behavior inferred from optical observations provides important clues into jet properties \\markcite{liang08}({Liang} {et~al.} 2008).  Based on the optical afterglow observations in the pre-\\swift era \\markcite{frail01,bloom03}({Frail} {et~al.} 2001; {Bloom}, {Frail}, \\& {Kulkarni} 2003), we expected to find jet breaks occurring within several days after the bursts, with the light curves breaking to decay slopes of $\\sim 2.2$ \\markcite{sari99}({Sari} {et~al.} 1999). Several recent studies \\markcite{burrows07,liang08,kocevski07,willingale07,evans09}({Burrows} \\& {Racusin} 2007; {Liang} {et~al.} 2008; {Kocevski} \\& {Butler} 2008; {Willingale} {et~al.} 2007; {Evans} {et~al.} 2008) have searched for jet breaks in the XRT data and agree that there is a substantial deficit relative to pre-{\\it Swift} expectations.  \\markcite{panaitescu07}{Panaitescu} (2007) suggests additional potential jet breaks in the sample, but is very broad in his jet break definition, attributing even breaks with shallower decays occurring after plateaus to jet breaks without discussion of the more global context of the light curves.  This method does not follow the framework of the canonical picture observed in many afterglows, though it suggests that some jet breaks are buried in the existing measured breaks, which we explore in more detail.  Previous studies of large afterglow samples have applied only the simplest afterglow models, rather than the detailed interpretations needed to explain GRBs in individual cases.  For example, the end of the plateau phase is often attributed to a cessation of energy injection without considering the possibility that it might represent a jet break during energy injection (see also \\markcite{depasquale08}{de Pasquale} {et~al.} 2008). Other model variations including a flat electron spectrum, different progenitor environments, and jet geometry and dynamics would slightly alter the properties of the canonical behavior. Therefore, in this study we perform a more generalized characterization of all XRT afterglows, reexamining a wide variety of closure relations to evaluate which segments of each light curve are consistent with each family of closure relations, whether there are any jet breaks that have been previously misinterpreted, and what limits can be placed on jet break statistics and energetics.  There are multiple reasons that we may not be detecting jet breaks in XRT data. We ask the following questions to explore this problem:  Are jet breaks subtle and buried within our observing errors (see also \\markcite{curran08}{Curran}, {van der Horst}, \\&  {Wijers} 2008)?  Do the jet breaks occur after the XRT observations end?  Do observational biases that cause us to miss parts of the light curves result in ambiguous classifications?  Could some of the breaks at the end of the plateau phases actually be jet breaks that are masked by continuing energy injection?  Are those GRBs for which jet breaks are not detected somehow intrinsically or observationally different than those for which they are detected?  The goal of this study is to attempt to answer these questions.  We describe the data sample selection, temporal analysis, and spectral analysis in \\S\\ref{sec:data}, the closure relations in \\S\\ref{sec:cr}, the results in \\S\\ref{sec:results}, discussion and implications in \\S\\ref{sec:disc}, and conclude in \\S\\ref{sec:conc}. Throughout this paper, we adopted the convention $F \\propto t^{-\\alpha}\\nu^{-\\beta}$ where $\\alpha$ is the temporal index and $\\beta$ is the spectral index, $F$ is the energy flux (with cgs units of $erg\\ cm^{-2}\\ s^{-1}$), and we use cosmological parameters $H_0 = 70~{\\mathrm{km~s^{-1}~Mpc^{-1}}}$, $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "\\label{sec:conc}  Pre-\\swift expectations for GRB X-ray afterglows have been substantially revised with the great wealth of XRT observations. While we try to categorize and classify their properties, there is still a wide range of unexplained diversity in GRB afterglow properties.  Within the limits of theoretical expectations and observational biases, we have attempted to survey the properties of the X-ray afterglows.  In agreement with some previous studies (\\markcite{burrows07}{Burrows} \\& {Racusin} 2007, \\markcite{liang08}{Liang} {et~al.} 2008, \\markcite{willingale07}{Willingale} {et~al.} 2007, \\markcite{kocevski07}{Kocevski} \\& {Butler} 2008, \\markcite{evans09}{Evans} {et~al.} 2008), we find only a small fraction ($\\sim 12\\%$) of our total sample has a late-time break that is clearly a jet break justified by the closure relations.  We find an additional $\\sim 30\\%$ with observational biases that make segments IV non-distinct but with a strong case for post-jet break temporal and spectral properties.  Some of the bursts in our sample remain ambiguous in the jet break designation.  Despite not being able to make absolute claims about these specific bursts, we demonstrate that there are jet breaks hidden within the data and observational biases.  This suggests that there are $\\sim 20\\%$ more jet breaks in the XRT afterglows than previous studies have revealed and at least $40\\%$ of the missing jet breaks can be attributed to observational biases.  Some of our light curves that require energy injection to continue post-jet break may have been previously misidentified as the end of the energy injection phase.  \\markcite{evans09}{Evans} {et~al.} (2008) also explores the canonical X-ray afterglow form using the XRT sample, where they find that less than half of all light curves behave canonically, and one quarter are ``oddballs''.  Many of these ``oddballs'' can be explained by the scenarios we use to describe the ambiguous cases discussed in this paper.  While their approach is somewhat different, their conclusions are similar to ours.  Our study requires post-jet break energy injection to explain 4 cases of the Prominent jet break sample and 11 others in the Hidden jet break category.  This modification to the canonical X-ray afterglow form alters expectations from simply studying the light curve alone, and adds to the theory needed to explain the diversity of observed properties.  These explanations do not solve all of the remaining problems related to jet breaks.  Several afterglow light curves, particularly the ones who can only be fit by a single power-law, persist with a constant slope prior to and beyond the times for which we would expect a jet break.  These bursts require either an exceptionally early jet break (sometimes before $100$ sec) or an exceptionally late jet break, requiring a large jet opening angle and an enormous ($\\gg 10^{51}$ erg) collimation corrected energy output.  \\swift GRBs are on average at higher redshifts, smaller jet opening angles, lower isotropic equivalent energies, and lower collimated $\\gamma$-ray energies compared to GRBs observed prior to {\\it Swift}. Some of these effects can be attributed to the lower energy coverage and superior sensitivity of the \\swift-BAT.  However, the consequences of these observational biases towards selecting different sorts of bursts is unexpected.  One of the fundamental predictions of the jet break models used in this work is achromatic behavior in a single spectral component.  The jet break should be a purely geometrical effect and should therefore not be limited to the X-ray afterglows.  Other components of the afterglow geometry may be more closely tied to emission segments and mechanisms making direct afterglow comparison difficult.  Modeling of the complete spectral energy distribution would be necessary to understand how the spectral breaks might influence the chromatic light curve behavior.  This would be further complicated by uncertainties in the optical extinction and X-ray absorption.  These detailed spectral studies, which are beyond the scope of this work, would provide additional information if simultaneous optical, infrared, and radio observations were available to narrow down the closure relations models and better constrain the physical models. \\markcite{liang08}{Liang} {et~al.} (2008) have already made some  progress on exploring multi-wavelength approach to this problem, specifically evaluating cases of chromatic versus achromatic breaks."
    },
    "0812/0812.3958_arXiv.txt": {
        "abstract": "The existence of large-scale dynamos in rigidly rotating turbulent convection without shear is studied using three-dimensional numerical simulations of penetrative rotating compressible convection. We demonstrate that rotating convection in a Cartesian domain can drive a large-scale dynamo even in the absence of shear. The large-scale field contains a significant fraction of the total field in the saturated state. The simulation results are compared with one-dimensional mean-field dynamo models where turbulent transport coefficients, as determined using the test field method, are used. The reason for the absence of large-scale dynamo action in earlier studies is shown to be due to the rotation being too slow: whereas the $\\alpha$-effect can change sign, its magnitude stays approximately constant as a function of rotation, and the turbulent diffusivity decreases monotonically with increasing rotation. Only when rotation is rapid enough a large-scale dynamo can be excited. The one-dimensional mean-field model with dynamo coefficients from the test field results predicts reasonably well the dynamo excitation in the direct simulations. This result further validates the test field procedure and reinforces the interpretation that the observed dynamo is driven by a turbulent $\\alpha$-effect. This result demonstrates the existence of an $\\alpha$-effect and an $\\alpha^2$-dynamo with natural forcing. ",
        "introduction": "Convective instability drives turbulence in the outer layers of late-type stars, such as the Sun. The large-scale magnetic fields of these stars are thought to arise from the interaction of turbulent convection and the overall rotation of the object. In mean-field dynamo theory (e.g.\\ Moffatt 1978; Parker 1979; Krause \\& R\\\"adler 1980; R\\\"udiger \\& Hollerbach 2004), this process relies on large-scale differential rotation ($\\Omega$-effect) producing toroidal field by shearing and the $\\alpha$-effect which regenerates the poloidal field. In simple situations, the $\\alpha$-effect is related to the kinetic helicity of the flow (e.g.\\ Steenbeck \\& Krause 1969). Dynamos where differential rotation is important, e.g.\\ the solar dynamo, are therefore often called $\\alpha\\Omega$-dynamos.  The $\\alpha\\Omega$ dynamo process was first invoked by Parker (1955) to explain solar magnetism. Mean-field models have been used extensively ever since to study various aspects of dynamos. Although very useful in their own right, these models often rely on ill-known parameterizations of turbulence, such as the $\\alpha$-effect and turbulent diffusivity. Only during recent years numerical simulations have reached a level of sophistication where large-scale $\\alpha\\Omega$-dynamos have been obtained self-consistently within the framework of local Cartesian simulations. These simulations operate in highly simplified situations and the turbulence is driven by external forcing (e.g.\\ Brandenburg et al.\\ 2001; Brandenburg \\& K\\\"apyl\\\"a 2007; K\\\"apyl\\\"a \\& Brandenburg 2009). In rapidly rotating stars differential rotation is likely to play only a minor role (e.g.\\ Hall 1991) and there the $\\alpha$-effect generates also the toroidal field. These systems are called $\\alpha^2$-dynamos. Again, such solutions have been found from direct simulations of forced turbulence (Brandenburg 2001; Mitra et al.\\ 2009b,c).  On the other hand, numerical simulations of magnetoconvection have been around for at least two decades, but somewhat surprisingly large-scale dynamos were not found until quite recently (Jones \\& Roberts 2000; Rotvig \\& Jones 2002; Browning et al.\\ 2006; Brown et al.\\ 2007; K\\\"apyl\\\"a et al.\\ 2008, hereafter Paper I; Hughes \\& Proctor 2009). The first two studies are related to the geodynamo and therefore convection is rotationally dominated. The next two use a global models with an imposed tachocline (Browning et al.\\ 2006) or rapid rotation (Brown et al.\\ 2007), respectively. The last two are local models with imposed shear flows, reminiscent of the shear dynamos reported from non-helical turbulence (Yousef et al.\\ 2008a,b; Brandenburg et al.\\ 2008). The origin of large-scale fields in the non-helically forced simulations cannot be due to the mean-field $\\alpha$-effect. However, the incoherent $\\alpha$-shear dynamo (e.g.\\ Vishniac \\& Brandenburg 1997; Proctor 2007) or the mean-field shear-current effect (Rogachevskii \\& Kleeorin 2003, 2004; Kleeorin \\& Rogachevskii 2008) could drive these dynamos. In the shearing local convection simulations the same two effects have been suggested mainly due to the fact that rotating convection alone has so far been found unable to generate large-scale magnetic fields (e.g.\\ Nordlund et al.\\ 1992; Brandenburg et al.\\ 1996; Cattaneo \\& Hughes 2006; Tobias et al.\\ 2008), although it should produce a mean $\\alpha$-effect according to theory. This drawback has prompted speculations that the $\\alpha$-effect in its mean-field incarnation simply does not work (Hughes \\& Proctor 2009).  A conclusive proof of the existence of an $\\alpha^2$-dynamo should be the demonstration of large-scale dynamo action in a simulation without shear. An additional property of such a setup is that the Vishniac \\& Cho (2001) flux of magnetic helicity is absent. This might lead to slow saturation of the large-scale magnetic field, unless there are other ways of shedding small-scale magnetic helicity.  The most natural way of explaining the large-scale magnetic fields seen in the rapidly rotating global convection simulations (Brown et al.\\ 2007) would be in terms of a mean-field $\\alpha^2$- or $\\alpha\\Omega$-dynamo. In the global models large-scale shear flows can also be generated, so it is not straightforward to determine what type of dynamo is operating there. In the present study, we use local simulations to demonstrate that large-scale dynamos can indeed be excited in rapidly rotating convection, i.e.\\ in the absence of shear, as long as rotation is rapid enough. We determine the turbulent transport coefficients from the simulations using the test field procedure (Schrinner et al.\\ 2005, 2007) and use them in a one-dimensional dynamo model to test their consistency with the direct simulations. A more thorough study of the turbulent transport coefficients from convection using the test field method is presented elsewhere (K\\\"apyl\\\"a et al.\\ 2009, hereafter Paper II). ",
        "conclusions": "\\label{sec:conclusions} We use numerical simulations to demonstrate that rigidly rotating convection can generate a large-scale dynamo. The importance of this result lies in the fact that, according to mean-field theory, rotating turbulent convection leads to an $\\alpha$-effect, which should result in a large-scale dynamo. So far this has not been seen in simulations without shear, leading to doubts of the applicability of the mean-field approach. The present results show that the lack of large-scale dynamos in rotating convection previously reported is probably due to too slow rotation. We find that an appreciable large-scale field is generated in the direct simulations if $\\Co\\ga4$.  With the smallest system size explored here, the scale separation is small even in the most rapidly rotating case (Run~A6), i.e.\\ $k_{\\rm max}/k_1\\approx2$. In this case the large-scale magnetic field is concentrated in the $k=0$ mode. Doubling the domain size increases the scale separation to $k_{\\rm max}/k_1\\approx5$ (Run~B6). Here the $k=0$ contribution is smaller and most of the magnetic energy is found for $k/k_1=1$ and 2 modes. Doubling the size of the box once more, $k_{\\rm max}/k_1\\approx10$. The magnetic energy peaks at the largest scales, although the $k=0$ mode is significantly weaker than in the cases with smaller system sizes. This suggests that for the present parameters there is a maximum size for the large-scale magnetic field structures which is larger than $2d$ and smaller than $4d$.  We compute the turbulent transport coefficients using the test field method and find that the magnitude of the $\\alpha$-effect remains relatively unaffected when rotation increases. However, at the same time the turbulent diffusion is severely quenched by the rotation. When these coefficients are used in a one-dimensional mean-field dynamo model, the two most rapidly rotating runs exhibit a growing dynamo in accordance with the direct simulations. Although the mean-field model is too simple to fully describe the field in the direct simulations, our results seem to validate the test field method further and lend support to our interpretation that the large-scale magnetic fields observed in the simulations are due to a turbulent $\\alpha$-effect.  On a more general note, the present results also represent a nice demonstration of the usefulness of mean-field theory as a predictive tool. In fact, as explained in \\S~\\ref{sec:magnetic}, our work was motivated by the earlier findings of Paper~II that the $\\alpha$-effect increases and the turbulent diffusivity decreases as the rotation rate is increased. This did already suggest the existence of an $\\alpha^2$-dynamo for sufficiently rapid rotation -- a suggestion that we have now been able to confirm in this paper.  It should be noted that the magnetic field structure of mean-field dynamos depends crucially on the geometry of the domain and the nature of the boundary conditions. Therefore, our results are not directly relevant to astrophysical bodies, because their geometry is not Cartesian nor the boundaries periodic. Our models can also not be thought of as local re\\-pre\\-sen\\-ta\\-tions of a star, although it is possible to get some idea about the dependence of the transport coefficients on latitude (Ossendrijver et al.\\ 2002; K\\\"apyl\\\"a et al.\\ 2006; Paper II). Furthermore, the values of the magnetic Reynolds number are obviously not in the astrophysically relevant regime. Nevertheless, our simulations provide the first concrete evidence of dynamo action from rotating convection without the additional help of shear. This is an important point, because it proves for the first time that an $\\alpha^2$ dynamo from rotating convection exists and that it is strong enough to produce large-scale fields. This was thought impossible until now (Cattaneo \\& Hughes 2006; Hughes \\& Cattaneo 2008; Hughes \\& Proctor 2009)."
    },
    "0812/0812.3169_arXiv.txt": {
        "abstract": "The reflex motion of a star induced by a planetary companion is  too small to detect by photographic astrometry. The apparent discovery in the 1960s of  planetary systems around certain nearby stars, in particular Barnard's star, turned out to be spurious. Conventional stellar radial velocities determined from photographic spectra at that time were also too inaccurate  to detect the expected reflex velocity changes. In the late 1970s and early 1980s, the introduction of solid-state,  signal-generating detectors and absorption cells to impose wavelength fiducials directly on the starlight, reduced radial velocity errors to the point where such a search became feasible. Beginning in 1980, our team from UBC introduced an absorption cell of hydrogen fluoride gas in front of the CFHT  coud\\'{e} spectrograph and, for 12 years,  monitored the radial velocities of some 29  solar-type stars. Since it was assumed that extra-solar planets would most likely resemble Jupiter in mass and orbit, we were awarded only three or four two-night observing runs each year. Our survey highlighted  three potential planet hosting stars,  $\\gamma$ Cep (K1 IV), $\\beta$ Gem (K0 III), and $\\epsilon$ Eri (K2 V). The putative planets all resembled Jovian systems with periods and masses of:  2.5 yr and  1.4 $M_{J}$, 1.6 yr and 2.6 $M_{J}$, and 6.9 yr and 0.9 $M_{J}$, respectively. All three were subsequently confirmed from more extensive data by the Texas group led by Cochran and Hatzes who also derived the currently accepted orbital elements.  None of these three systems is simple. All 5  giant stars and the supergiant in our survey proved to be intrinsic velocity variables. When we first drew attention to a possible planetary companion to  $\\gamma$ Cep in 1988 it was classified as a giant,  and there was the possibility that its radial velocity variations and those of $\\beta$ Gem (K0 III) were intrinsic to the stars. A further complication for $\\gamma$ Cep was the presence of an unseen secondary star in an orbit with a period initially estimated at some 30 yr. The implication was that the planetary orbit might not be stable, and a Jovian planet surviving so close to a giant then seemed improbable. Later observations by others showed the stellar binary period was closer to 67 yr, the primary was only a sub-giant and a weak, apparently synchronous chromospheric variation  disappeared. Chromospheric activity was considered important because $\\kappa^{1}$ Cet, one of our program stars, showed a significant correlation of its radial velocity curve with chromospheric activity.  $\\epsilon$ Eri is a young, magnetically active star with spots making it a  noisy target for radial velocities. While the signature of  a highly elliptical orbit ($e=0.6$)  has persisted for more than 3 planetary orbits, some feel that even more extensive coverage is needed to confirm the identification  despite an apparent complementary astrometric acceleration detected with the Hubble Space Telescope.  We confined our initial analyses of the program stars to looking for circular orbits. In retrospect, it appears that some 10\\% of our sample did in fact have Jovian planetary companions in orbits with periods of years. ",
        "introduction": "\\label{}  It is quite hard nowadays to realise the atmosphere of skepticism and indifference in the 1980s to proposed  searches for extra-solar planets. Some people felt that such an undertaking was not even a legitimate part of astronomy. It was against such a background that we began our precise radial velocity (PRV) survey of certain bright solar-type stars in 1980 at the Canada France Hawaii 3.6-m Telescope (CFHT). Up to that time, the only serious searches for extra-solar planets had been astrometric studies that looked for long term periodic displacements of stars on the sky caused by unseen companions. Unfortunately, the tiny angular displacements were much smaller than the proper motions and parallaxes for nearby stars and beyond the capability of ground-based astrometric telescopes using photographic emulsions.  For a circular planetary orbit, the semi-angular reflex displacement, $\\theta$, of the parent star on the sky caused by a companion planet is: \\begin{equation} \\theta \\approx \\frac{R_{pl}M_{pl}}{~D~M_{st}} ~~~~ {\\rm milli\\mbox{-}arc sec}, \\end{equation} where $R_{pl}$ is the distance of the planet from its parent star in AU, $M_{pl}$ and $M_{st}$ are the planet and stellar masses in Jovian and solar masses, respectively. $D$ is the distance of the star from the Earth in parsecs. $\\theta$ increases linearly with $R_{pl}$ (and as the $\\frac{2}{3}$ power of the period). For the Sun-Jupiter system seen from a distance of 5 pc, $\\theta$ would be $\\sim$1 milli-arcsec and take nearly 12 years to complete a cycle. Considering that the average ground-based photographic image was spread over $>$1000 milli-arcsec, the challenge of detecting such a miniscule displacement over several decades is obvious.  Nonetheless, tentative extra-solar planetary discoveries were claimed for several nearby stars from astrometry with the Sproul Observatory 24 inch refractor, but the signals were close to the errors of measurement. The best known identifications were for Barnard's Star \\citep{Kam69}. The orbital periods in all cases were close to multiples of 8 years, and \\citet{Her73} demonstrated that the epochs of the `discontinuities' in the Sproul data coincided with times when the telescope primary lens had been removed for cleaning.  The amplitude of the periodic reflex variation of a planet-hosting star's radial velocity complements the astrometric variation. Again, for circular orbits, the amplitude of the stellar radial velocity change, $\\delta RV$, would be: \\begin{equation} \\delta RV \\approx \\frac{30M_{pl}~{\\rm sin}i}{(M_{st}~R_{pl})^{1/2}}~~~~ {\\rm km~~ s^{-1}}, \\end{equation} where $i$ is the inclination of the orbital axis to the line of sight and the units are the same as those in equation 1. The radial velocity excursion is independent of the distance from the Earth and increases inversely as $\\surd R_{pl}$. As a star, $\\delta RV$ for the Sun as caused by Jupiter would be $\\sim\\pm$13 m s$^{-1}$ over 12 years when viewed at $i$=0. Unlike astrometry, radial velocities favor the detection of tightly bound planets.  In the early 1970s, quality radial velocities determined from photographic plates had typical errors of $\\sim$1 km s$^{-1}$. The large error was in part due to the photographic process but, more seriously, to the effect of guiding errors at the spectrograph slit which cause the stellar spectrum to shift relative to the bracketting comparison arc lines. Obviously, an improvement of more than an order of magnitude coupled with a long term stability of years, was necessary  if radial velocities were to be useful in looking for extra-solar Jovian analogs.  The above comments also reflect the belief at the time that planetary systems beyond the solar system would most likely be similar to our own, a belief strongly bolstered by theoretical simulations even down to the Titus Bode relation \\citep{Isa77}! No one predicted the populations of Jovian planets in few-day, or highly elliptical, orbits which eventually turned up. As a result, early search strategies concentrated on long-term monitoring with observations spaced out over months and years. ",
        "conclusions": "By 1993, when our PRV program ended, the HF technique had passed its prime. Our observing list was necessarily small because the 100 \\AA~ coverage was an order of magnitude less than with I$_{2}$ vapor and  CCD arrays coupled to an  \\'{e}chelle spectral format \\citep[e.g.][]{Mar92,Wit06}.  Indeed,  \\citet{May95} had even demonstrated that with sufficient care one could perhaps dispense with an absorption cell altogether.  The handling and preparation of the HF system was also dangerous. On the other hand, the HF system had been straight-forward to implement. The natural width and spacing of the HF lines had avoided the additional layer of processing necessary  with the sharp overlapping  I$_{2}$ lines. The photometric fidelity of the diode arrays was ideal, and we already had a `small-shift' progam in place as a basis for the analysis. Our early success in reducing PRV errors  had, we liked to think, encouraged other groups to press on.  As with any new level of precision, our PRV turned up unexpected results such as the variability of the yellow giants, the correlated variations of  the PRV and chromospheric activity in $\\kappa^{1}$ Cet, and a paucity of close brown dwarf companions. These finds complicated interpretation of our results for the program stars. Above all,  a general skepticism still greeted any announcement of an exo-planet discovery. Nonetheless, three stars that we flagged as possible planet hosts, $\\gamma$ Cep, $\\beta$ Gem, and $\\epsilon$ Eri  were later confirmed by the Texas group, albeit not everyone is yet convinced about  $\\epsilon$ Eri b despite the complementary astrometric acceleration detected by \\citet{Ben06}.  In simple terms, roughly some 10\\% of our program stars turned out to have Jovian planets with orbital periods of years. This roughly matches the statistics from the current huge sample of stars being monitored by others  \\citep[see for example][]{Mar00}, although each of our three candidates was unusual.  Since our stars were only observed a few times per year, we had little chance of detecting closely bound systems like 51 Peg, but the results for  $\\kappa^{1}$ Cet are tantalising in that respect.  The success of our extra-solar planet search depended on many people, but particularly on Bruce Campbell, Stephenson Yang, Alan Irwin and Ron Johnson who built successive low-noise versions of the Reticon camera systems. It is a great pleasure to be able to acknowledge my debt to them here."
    },
    "0812/0812.2997_arXiv.txt": {
        "abstract": "We present a detailed analysis of 256 radio sources from our deep (flux density limit of $42~\\mu$Jy at the field centre at 1.4 GHz) Chandra Deep Field South 1.4 and 5 GHz VLA survey. The radio population is studied by using a wealth of multi-wavelength information in the radio, optical, and X-ray bands. The availability of redshifts for $\\sim 80\\%$ of the sources in our complete sample allows us to derive reliable luminosity estimates for the majority of the objects. X-ray data, including upper limits, for all our sources turn out to be a key factor in establishing the nature of faint radio sources. Due to the faint optical levels probed by this study, we have uncovered a population of distant Active Galactic Nuclei (AGN) systematically missing from many previous studies of sub-millijansky radio source identifications. We find that, while the well-known flattening of the radio number counts below 1 mJy is mostly due to star forming galaxies, these sources and AGN make up an approximately equal fraction of the sub-millijansky sky, contrary to some previous results. The AGN include radio galaxies, mostly of the low-power, Fanaroff-Riley I type, and a significant radio-quiet component, which amounts to approximately one fifth of the total sample. The ratio of radio to optical luminosity depends more on radio luminosity, rather than being due to optical absorption. ",
        "introduction": "The extragalactic radio source population ranges from normal galaxies with luminosities $\\sim 10^{20}$ W Hz$^{-1}$ to galaxies whose radio emission is as much as $10^4 -10^7$ times greater owing to regions of massive star formation or to an active galactic nucleus (AGN). The population of radio sources in the sky with flux densities $> 1$ mJy is dominated by AGN driven emission generated from the gravitational potential associated with a supermassive black-hole in the nucleus. For these sources, the observed radio emission includes the classical extended jet and double lobe radio sources as well as compact radio components more directly associated with the energy generation and collimation near the central engine. Below 1 mJy there is an increasing contribution to the radio source population from synchrotron emission resulting from relativistic plasma ejected from supernovae associated with massive star formation in galaxies or groups of galaxies, often associated with mergers or interactions \\citep[e.g.,][]{win95,ric98,fom02}. However, the mix of star forming galaxies (SFG) and AGN and their dependence on epoch is not well determined.  For example, \\cite{gru03}, by using 15-$\\mu$m Infrared Space Observatory (ISO) and 1.4 GHz data have suggested that starburst galaxies start to become a significant fraction of the radio source population for $S_{\\rm r} \\la 0.5 - 0.8$ mJy, exceeding $\\sim 60\\%$ only at $S_{\\rm r} < 50~ \\mu$Jy. On the other hand, \\cite{hu05}, by fitting radio number counts of the Hubble Deep Field South and other surveys with evolutionary starburst models, have estimated that the number of SFG exceeds $50\\%$ of the total already at $\\sim 0.25$ mJy.  The most commonly accepted paradigm has been that the sub-mJy population is largely made up of SFG. A (small but growing) number of researchers, however \\citep[e.g.,][]{gru99,cil03,sey08,smo08}, have pointed out that radio sources (radio galaxies and radio-loud AGN) could also make a significant contribution to the sub-mJy population. It has been further suggested \\citep{jar04,sim06} that radio-quiet AGN (see below) might become relevant at sub-mJy levels.  The Chandra Deep Field South (CDFS) area is part of the Great Observatories Origins Deep Survey (GOODS) and as such is one of the most intensely studied region of the sky. High sensitivity X-ray observations are available from Chandra \\citep{gia02,leh05,luo08}. The GOODS and Galaxy Evolution from Morphology and Spectral energy distributions (GEMS) multiband imaging programs using the HST Advanced Camera for Surveys (ACS) give sensitive high resolution optical images \\citep{gia04,rix04}. Ground based imaging and spectroscopy are available from the ESO 2.2m and 8m telescopes \\citep[see references in][Paper II]{mai08}.  We observed the CDFS with the NRAO Very Large Array (VLA) for 50 hours at 1.4 GHz mostly in the BnA configurations in October 1999 and February 2001, and for 32 hours at 5 GHz mostly in the C and CnB configurations between June and October 2001 \\citep[details are given in][Paper I]{kel08}. The effective angular resolution was $3.5''$ and the minimum root-mean-square (rms) noise was as low as $8.5~\\mu$Jy per beam at 1.4 GHz and $7~\\mu$Jy per beam at 5 GHz.  A total of 266 radio sources were catalogued at 1.4 GHz, 198 of which are in a complete sample with signal-to-noise ratio (SNR) greater than 5 and located within 15$^{\\prime}$ of the field center. The corresponding flux density ranges from 42 $\\mu$Jy at the field center to 125 $\\mu$Jy near the field edge. These deep radio observations complement the larger area, but less sensitive, lower resolution observations of the CDFS discussed by \\cite{afo06} and by \\cite{nor06}.  In addition to our deep radio data, our set of ancillary data is quite unique and allows us to shed new light on the nature of the sub-mJy radio source population. These data include reliable optical/near-IR identifications for $94\\%$ of the radio sources, optical morphological classification for $61\\%$ of the sample, redshift information for $73\\%$ of the objects, and X-ray detections for $33\\%$ of the sources, with upper limits for all but three of the others. This is the largest and most complete sample of $\\mu$Jy sources in terms of redshift information.  The purpose of this paper is to study the composition of the radio-detected VLA-CDFS sample in terms of source properties by using all available information.  More details on the VLA-CDFS can be found in Paper I, which addresses the radio data, in Paper II, which discusses the optical and near IR counterparts to the observed radio sources, and in \\cite{toz08} (Paper III), which studies the X-ray properties.  The sample radio-optical data are described in \\S~2, while \\S~3 deals with the X-ray data. \\S~4 analyzes the overall properties of our sample and derives the radio number counts for various classes together with some model predictions, while \\S~5 discusses our results. Finally, \\S~6 summarizes our conclusions. Throughout this paper spectral indices are defined by $S_{\\nu} \\propto \\nu^{-\\alpha}$ and the values $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M} = 0.3$, and $\\Omega_{\\rm \\Lambda} = 0.7$ have been used. This paper supersedes the preliminary results presented by \\cite{pad07b}. ",
        "conclusions": "We have used a deep radio sample, which includes 256 objects down to a 1.4 GHz flux density of $42~\\mu$Jy selected in the Chandra Deep Field South area, to study the nature of sub-mJy sources. Our unique set of ancillary data, which includes reliable optical/near-IR identifications, optical morphological classification, redshift information, and X-ray detections or upper limits for a large fraction of our sources, has allowed us to shed new light on this long standing astrophysical problem. By analyzing the ratio of radio to optical luminosity and the radio and X-ray powers of the sources with morphological and redshift information, we have selected candidate star-forming galaxies and AGN from a complete sample of 194 objects. As the first two parameters by themselves are not very good discriminants between star-forming galaxies and AGN, optical morphology and especially X-ray data turned out to be vital in establishing the nature of faint radio sources. We have also complemented our data analysis with model calculations based on brighter samples and counts in other bands. Our main results can be summarized as follows:  \\begin{enumerate} \\item The well-known flattening of the radio number counts below $\\approx 1$ mJy is mostly due to star-forming galaxies, which are missing above $\\sim 2$ mJy but become the dominant population below $\\approx 0.1$ mJy. \\item AGN exhibit the opposite behavior, as their counts drop at lower flux densities, going from 100\\% of the total at $\\sim 10$ mJy down to $< 50\\%$ at the survey limit. This is driven by the fall of radio-loud sources, as radio-quiet objects, which make up $\\sim 20\\%$ of sub-mJy sources, display relatively flat counts. \\item Radio-quiet AGN make up about half of all AGN. Their counts appear to be a scaled down version, by a factor $\\approx 3 - 4$, of those of star-forming galaxies, and are very different from those of the radio-loud AGN population. This should provide a clue to the origin of their radio emission. \\item Star-forming galaxies make up $\\la 60\\%$ of sub-mJy sources down to the flux limit of this survey. This has to be regarded as a robust upper limit, as whenever we had to make some assumptions, we chose to maximize their numbers. This result is at variance with the many papers, which over the years have suggested a much larger dominance of these sources. On the other hand, our results are in broad agreement with a number of recent papers, which found a significant AGN component down to $\\approx 50~\\mu$Jy. \\item The results of our model calculations agree quite well with the observed number counts and provide supporting evidence for the scenario described above. Moreover, they imply that sub-mJy radio-loud AGN are dominated by low-power, Fanaroff-Riley type I radio galaxies, as their high-power counterparts and radio-loud quasars are expected to disappear below $\\sim 0.5 - 1$ mJy. \\item We find a correlation between X-ray and radio power for star-forming galaxies, in agreement with most previous studies. \\item Despite some previous claims, we find that the ratio of radio to optical luminosity depends more on radio luminosity, rather than being due to optical absorption, and is therefore not related to a high degree of extinction. \\end{enumerate}  Considering the apparent emerging population of low luminosity AGN at microjansky levels, care is needed when interpreting radio source counts in terms of the evolution of the star formation rate in the Universe.  We plan to expand on this work by using our deeper ($7~\\mu$Jy per beam over the whole Extended CDFS region) radio observations \\citep{mil08} and the recently released 2 Msec Chandra data \\citep{luo08}."
    },
    "0812/0812.2029_arXiv.txt": {
        "abstract": "We study the stellar populations of a sample of 14 elliptical galaxies in the Virgo cluster. Using spectra with high signal-to-noise ratio (S/N$\\simgt 100$\\AA\\ $^{-1}$) we propose an alternative approach to the standard side-band method to measure equivalent widths (EWs).  Our \\emph{Boosted Median Continuum} is shown to map the EWs more robustly than the side-band method, minimising the effect from neighbouring absorption lines and reducing the uncertainty at a given signal to noise ratio. Our newly defined line strengths are more sucessful at disentangling the age-metallicity degeneracy. We concentrate on Balmer lines (H$\\beta$,H$\\gamma$,H$\\delta$), the G band and the combination [MgFe] as the main age and metallicity indicators. We go beyond the standard comparison of the observations with simple stellar populations (SSP) and consider four different models to describe the star formation histories, either with a continuous star formation rate or with a mixture of two different SSPs. These models improve the estimates of the more physically meaningful mass-weighted ages. Composite models are found to give more consistent fits among individual line strengths and agree with an independent estimate using the spectral energy distribution.  A combination of age and metallicity-sensitive spectral features allows us to constrain the average age and metallicity. For a Virgo sample of elliptical galaxies our age and metallicity estimates correlate well with stellar mass or velocity dispersion, with a significant threshold around $5\\times 10^{10} M_\\odot$ above which galaxies are uniformly old and metal rich. This threshold is reminiscent of the one found by Kauffmann et al. in the general population of SDSS galaxies at a stellar mass $3\\times 10^{10}M_\\odot$. In a more speculative way, our models suggest that it is formation \\emph{epoch} and not formation timescale what drives the Mass-Age relationship of elliptical galaxies. ",
        "introduction": "Unveiling the star formation histories of elliptical galaxies is key to our understanding of galaxy formation. Being able to resolve their seemingly homogenous distribution is hampered by the fact that their light is dominated by old, i.e. low-mass stars, which do not evolve significantly even over cosmological times. Furthermore, the presence of small amounts of young stars as recently discovered in NUV studies \\citep{fs00,yi05,kav07} reveals a complex history of star formation that requires proper estimates of mass-weighted ages, in contrast with the luminosity-weighted ages that simple stellar populations (SSPs) can only achieve. The majority of papers dealing with age estimates of the stellar populations of elliptical galaxies rely on such SSPs \\citep[see e.g. ][]{kd98,sct00,tm05}, and it is only recently that special emphasis has been made on the need to go beyond simple populations \\citep{fyi04,serra07,idi07}  Dating the (old) stellar populations of elliptical galaxies has been fraught with difficulties, the most prominent being the age-metallicity degeneracy, whereby the photo-spectroscopic properties of a galaxy of a given age and metallicity can be replicated by a younger or older galaxy at a suitably higher or lower metallicity, respectively. To some extent, this problem has been overcome by the measurement of pairs of absorption line indices \\citep{wo94}, one whose change in equivalent width (EW) is dominated by the average metallicity of the population and the other dominated by the average age of the population. Typical metal-sensitive line strengths are the Mg feature at 5170\\AA\\ , the iron lines around 5300\\AA , or a combination such as [MgFe] \\citep{gon93}. Balmer lines are more age-sensitive and are often combined with metal-sensitive lines to break the degeneracy. However, measurements of EWs of Balmer lines can be affected by the age-metallicity degeneracy because of the presence of nearby absorption lines. Such is the case of H$\\gamma$, with the prominent G band at 4300\\AA\\ or the CN bands in the vicinity of the H$\\delta$ line. This paper is partly motivated by the need to define a method to estimate EWs that minimise the sensitivity of metallicity on Balmer lines by a proper estimate of the continuum.  Even at relatively younger ages (a few Gyr) where such uncertainties are reduced, considerable degeneracies still remain.  The confirmation of recent star formation (RSF) occuring in early type galaxies \\citep{yi05,kav07} has raised the problem that the existence of a young population can considerably affect the parameters derived through SSP analysis. \\citet{sct00} and later \\citet{serra07} showed that even relatively small mass fractions ($\\sim$1\\%) of young stars can distort age and metallicity estimates and in moderate cases ($\\sim$10\\%) completely overshadow the older population.  In addition, the age and mass fraction of any younger sub-population will also be degenerate, with larger mass fraction of relatively older sub populations having the same effect as smaller fractions of younger ones.  \\citet{schiavon04} noticed that using different Balmer lines (H$\\beta$, H$\\gamma$ and H$\\delta$) to estimate the age gives slightly different results, which was suggested to show that the galaxy had undergone recent star formation.  Contrary to this, \\citet{tm04} find that H$\\gamma$ and H$\\delta$ equivalent widths are more affected by a non-solar $\\alpha$/Fe ratio on higher order Balmer lines.  This effect is due to the increase of metal lines at bluer wavelengths, thereby distorting the continuum as measured by a side-band method.  This was expanded by \\cite{serra07} who remodelled synthetic 2-burst models using H$\\beta$ and H$\\gamma_A$ and achieved a result consistent with both papers, concluding that a mistmatch between the three Balmer line estimates could possibly reveal underlying younger populations.  Moving forward with such analysis -- beyond simple populations and luminosity-weighted parameters -- requires improvements on both the H$\\gamma$ and H$\\delta$ measurements. Balmer line equivalent widths suffer from the effects of the metallicity due to the presence of such lines in the spectral region used to determine the continuum \\citep{wo97,tm04,pro07}. Although considerable work has already been done in this area \\citep[see e.g. ][]{rose94,jw95,va99,yam06}, we test a completely different approach.  In this paper we present a comprehensive analysis of the stellar populations of 14 Virgo cluster ellipticals using several models to describe the star formation history, exploring the discrepancies found between simple and composite stellar populations. In an attempt to combat the problems discussed above, we introduce a new method for the measurement of equivalent widths using a high percentile running ``median'' to describe the continuum.  The properties of this new method are exploited by using various age- and metallicity-sensitive spectral features.  \\begin{figure*} \\begin{minipage}{18cm} \\begin{center} \\includegraphics[width=3.3in]{BFPS_f1a.eps} \\includegraphics[width=3.3in]{BFPS_f1b.eps} \\caption{Dependence of the age-metallicity gradient -- ($\\Delta$ Age/Age)(/$\\Delta$ Z/Z) -- on the choice of BMC parameters, namely the level at which the 'boosted median' is taken (\\emph{left}) or the size of the kernel (\\emph{right}). On the left panel a number of kernel sizes is shown as labelled: $\\Delta\\lambda$=50\\AA\\ (dashed); 100\\AA\\ (solid) and 200\\AA\\ (dotted). On the right panel a number of confidence levels are considered: 50\\% (dotted); 70\\% (long dashed); 90\\% (solid) and 95\\% (short dashed).  The grey horizontal line in both panels is the estimate from the standard side-band method (SB), and the grey vertical shaded area marks our choice of 'boosted median' parameters.} \\label{fig:DADZ} \\end{center} \\end{minipage} \\end{figure*} ",
        "conclusions": "We have revisited the superb, high signal-to-noise spectra of 14 elliptical galaxies in the Virgo cluster presented by \\citet{yam06,yam08}. Our main goal was two-fold: a) to give an optimal but versatile definition of the equivalent width of spectral features that would minimise the contribution from nearby lines, b) to explore the possibility of discriminating between the standard treatment of galaxy spectra either as simple stellar populations or composite models with a distribution of ages and metallicities.  The former was achieved by defining a ``boosted median'' pseudo-continuum. This method is very easy to implement on any spectral data and it improves on the traditional side-band methods by reducing the final uncertainty of the EWs for the same spectra and reducing the age-metallicity degeneracy of age- and metal-sensitive lines by reducing the effect of unwanted, nearby spectral features. The method only requires two numbers to determine the pseudo-continuum (i.e. confidence level of the boosted median and size of the kernel over which the analysis is performed). We propose 90\\% and 100\\AA\\ for these two parameters when dealing with stellar populations of galaxies at moderate resolution (R$\\sim 1000-2000$).  The second goal -- going beyond SSPs -- is approached by comparing SSPs and three more sets of models which assume some distribution of ages and metallicities. We find simple populations fail to give a consistent marginalised distribution of ages when fitting independently different age-sensitive lines. We propose either a 2-Burst model or a $\\tau$ star formatiom history with a phenomenological prescription for chemical enrichment (CSP models). They give similar results, with a clear trend between average age and either stellar mass or velocity dispersion, as shown in figure~\\ref{fig:AgeMet}.  As suggested above when discussing figure~\\ref{fig:AgeHist}, if younger (older) galaxies are better fit by 2-Burst (CXP) models, one would expect that young stars in elliptical galaxies appear in a random way, and not as a time-coherent, smooth distribution. This result supports minor merging -- possibly involving metal-poor sub-systems -- as the main mechanism to generate recent star formation in early-type galaxies \\citep{kav09}. This process will be readily detected in low-mass galaxies, whereas a similar amount of young stars in a massive galaxy will be harder to detect.  Doubtlessly average ages and metallicities are the observables best constrained by these models. Parameters like formation timescale and formation epoch can only be considered ``next-order'' observables, which will be fraught with larger uncertainties.  This is largely due to the massive degeneracies which are present in such parameters as mentioned in the introduction. Its is well known that the mass and age of a subpopulation are degenerate in that additional mass fraction can compensate for an older age or vice versa \\citep{leo96}. Another considerable degeneracy also exists between the formation redshift and the star formation timescale for both the EXP and CXP models \\citep[see e.g. ][]{gobat08}.  While our indices cover a large wavelength range we find that these degeneracies are still significant. Figure~\\ref{fig:Degen} shows the degeneracy between parameters for a typical case using 2-Burst models (\\emph{left}) and CXP models (\\emph{right}). We take the probability weighted values across the entire parameter space as a more robust measure of these parameters.  Although the derived parameters are still subject to considerable uncertainty, as expressed by the error bars shown in figure~\\ref{fig:zF}, we note that our grids sample the parameter space quite finely meaning our recovered likelihood distribution should be a reasonable approximation to the complete distribution.  We proceed in a more speculative way, taking the continuous CXP models at face value and putting some trust on the predictions of their formation epochs and timescales. Figure~\\ref{fig:zF} reveals an intriguing trend that suggests it is not just formation timescale but also formation epoch which drives the mass-age relationship in early-type galaxies. The solid lines in the top/bottom right panels gives a simple linear fit to the data, namely: \\be z_F \\sim  1.8 + 3.9\\log\\sigma_{100} \\label{eq:zF} \\ee \\be \\tau ({\\rm Gyr}) \\sim  0.3 - 1.14\\log\\sigma_{100} \\label{eq:Tau} \\ee  \\noindent with $\\sigma_{100}$ given in units of 100~km/s. Figure~\\ref{fig:SFH} shows our proposed model for the star formation history of Virgo cluster elliptical galaxies. We model the formation histories as skewed Gaussians whose spread maps the star formation timescale, and we include a correlation between formation epoch and velocity dispersion, as labelled. Notice that the \\emph{observed} low-mass galaxies (black in the bottom panel) are best fit by late formation epochs, although this fact does not prevent smaller early-type galaxies in the past to have formed at earlier epochs (represented by the grey Gaussians). Our sample is too small and too local to explore this issue further.  One could argue that the lack of a correlation between formation timescale ($\\tau_1$; \\emph{top}) and mass would contradict the observed relationship between mass and abundance ratios \\citep[see e.g. ][]{tm05}. One possibility would involve minor merging of dwarf galaxies. Those mergers will have a more prominent effect on the spectra of low-mass galaxies. Dwarf galaxies have a very extended period of star formation and eject a big fraction of their metals, resulting in metal-poor gas with solar or sub-solar abundance ratios that will reduce the observed (luminosity-weighted) [$\\alpha$/Fe] in low-mass galaxies. Alternatively, notice that the formation epoch ($z_F$; \\emph{bottom}) correlates quite strongly with velocity dispersion, implying that the structures that form the bulk of the stars in low-mass galaxies are more likely to be contaminated by the ejecta from type~Ia supernova, thereby reducing the abundance ratios to reproduce the observed correlation between [$\\alpha$/Fe] and mass. A more detailed analysis of the abundance ratios -- although beyond the scope of this paper -- is under way to explore this very interesting possibility.   \\begin{figure*} \\begin{minipage}{18cm} \\includegraphics[width=3.2in]{BFPS_f15a.eps} \\includegraphics[width=3.2in]{BFPS_f15b.eps} \\caption{Probability distributions of NGC 4489 (left) and NGC 4467 (right) for 2-Burst and CXP models, respectively.  \\emph{Left:} The contour plot shows the degeneracy between the age (t$_{\\rm Y}$) and mass fraction (f$_{\\rm Y}$) of the younger component. Contours are shown at the 75, 90 and 95\\% confidence level.  The solid circle identifies the best fit and quoted error bars.  The inset shows the probability distributions of the average age and metallicity, both of which are well constrained.  \\emph{Right Panel} The degeneracy between the formation redshift (z$_{\\rm F}$) and star formation timescale ($\\tau_1$) is shown in the contour plot (75,90 and 95\\% confidence levels).The probability distribution of the average age and metallicity (inset) show that their values are well constrained.} \\label{fig:Degen} \\end{minipage} \\end{figure*}"
    },
    "0812/0812.0506_arXiv.txt": {
        "abstract": "We examine the variability in the intrinsic absorption in the Seyfert 1 galaxy \\mrk\\ using three epochs of observations from the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} and two epochs of observations with the Space Telescope Imaging Spectrograph on the {\\it Hubble Space Telescope}. Rather than finding simple photoionization responses of the absorbing gas to changes in the underlying continuum, the observed changes in the absorption profiles can be understood more clearly if the effective covering fraction of the gas in all emission components, continuum and broad and intermediate velocity width emission lines, is accounted for. While we do not uniquely solve for all of these separate covering fractions and the ionic column densities using the spectral data, we examine the parameter space using previously well-constrained solutions for continuum and single emission component covering fractions.  Assuming full coverage of the continuum, we find that of the two velocity components of the Mrk~279 absorption most likely associated with its outflow, one likely has zero coverage of the intermediate line region while the other does not.  For each component, however, the broad line region is more fully covered than the intermediate line region. Changes in the \\ion{O}{6} column densities are unconstrained due to saturation, but we show that small changes in the nonsaturated \\ion{C}{4} and \\ion{N}{5} column densities are consistent with the outflow gas having zero or partial covering of the intermediate line region and an ionization parameter changing from $\\sim$0.01 to $\\sim$0.1 from 2002 to 2003 as the UV continuum flux increased by a factor of $\\sim$8. The  absence of a change in the \\ion{C}{3} absorbing column density is attributed to this species arising outside the Mrk279 outflow. ",
        "introduction": "Mass outflows from low-redshift active galactic nuclei (AGNs) are inferred from blueshifted, variable X-ray and UV absorption which often shows evidence for partial coverage of the contintuum source in nearly one half of all nearby Seyfert galaxies (see review in Crenshaw, Kraemer, \\& George 2003 and recent results in Dunn et al.\\ 2007). In addition to possibly playing a strong role in the accretion process (Blandford \\& Begelman 1999, 2004), AGN outflows likely have substantial impacts on their environments, providing feedback energy to regulate the formation of galaxies and clusters of galaxies (Granato et al.\\ 2004, Scannapieco \\& Oh 2004, Scannapieco, Silk, \\& Bouwens 2007); to set the luminosity-temperature relation in galaxy clusters (Cavaliere et al.\\ 2002); to distribute metals in galaxy clusters (Moll et al.\\ 2007) and into the intergalactic medium (Adelberger et al.\\ 2003).  Theoretical models of AGN outflows, from accretion disk winds (K\\\"{o}nigl \\& Kartje 1994, Murray et al.\\ 1995, Proga 2000, 2003, Proga, Stone, \\& Kallman 2000, Proga \\& Kallman 2004), or from  ablation from the obscuring torus (Krolik \\& Kriss 1995, 2001) place the absorbing material at different distances from the central engine and with different orientations with respect to the broad line region (BLR). A large sample of high resolution X-ray and UV spectra of ouflows in nearby Seyferts from {\\it Chandra}, \\xmm, the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} and from the Space Telescope Imaging Spectrograph (STIS) onboard the {\\it Hubble Space Telescope} has been building in the literature over the last decade, providing much insight into the physical conditions in these ouflows. However, a broad understanding of these outflows has yet to emerge, and several outstanding questions remain, namely: What is the relationship between the X-ray- and UV- absorbing gas?; What is the location and structure of the absorbing material?; Is the observed variability in some absorption profiles a response to flux variations in the central engine (Krolik \\& Kriss 1995, Shields \\& Hamann 1997, Crenshaw et al.\\ 2000, Kraemer et al.\\ 2002) or to bulk motions of the absorbing gas (Crenshaw \\& Kraemer 1999, Kraemer et al.\\ 2001, Gabel et al.\\ 2003b)?  In general, it is difficult to associate a particular UV absorption component with the X-ray absorber unambiguously, in part due to relatively low spectral resolution in X-ray data and in part because the ionization modeling is sensitive to the shape of the UV-to-X-ray continuum (Kaspi et al.\\ 2001). However, with  the exception of campaigns on Mrk~279, using \\fuse, \\stis, and \\chand\\ with both the {\\it High Energy Transmission Grating (HETG)} (Scott et al.\\ 2004, Gabel et al.\\ 2005a, S04 and G05 hereafter) and the {\\it Low Energy Transmission Grating} (Costantini et al.\\ 2007), on NGC~7469 (Scott et al.\\ 2005) using \\fuse, \\stis, and \\chand/{\\it HETG}, the long-term observing campaign on NGC~3783, using \\fuse, \\stis, and \\chand/{\\it HETG} (Kaspi et al.\\ 2002, Netzer et al.\\ 2003, Gabel et al.\\ 2003a, 2005b), and the campaign on NGC~4151 using \\fuse, \\stis, and \\chand/{\\it HETG} (Kraemer et al.\\ 2005, Kraemer et al. 2006, Crenshaw \\& Kraemer 2007), UV and X-ray observations are generally obtained at different times.  Several observations of the absorption profiles  taken over a period of time are particularly useful for addressing the last two questions mentioned above, as they provide contraints on the response time of the gas to variability in the central engine and thus limits on the distance of the absorber from the continuum-emitting region. Mrk~279 ($z_{em}=0.0305$) has been the subject of an intensive campaign of simultaneous observations in the UV and X-ray (S04, Kaastra et al.\\ 2004, Arav et al.\\ 2005, 2007, G05, Costantini et al.\\ 2007). In this paper, we use several epochs of UV spectral obsevations of Mrk~279 over 1999 to 2003 from \\fuse\\ and \\stis\\ presented in  S04 and G05 to investigate the variability of the column density and/or covering fraction of the outflow over time. We choose to concentrate on the UV data because the 2002 epoch of \\chand\\ observations coincided with a low flux state.  The grating data therefore  have signal-to-noise (S/N) too low to perform detailed analyses (S04).  The spectral coverage of \\fuse\\ provides information on \\ion{H}{1} absorption in Ly~$\\beta$ and higher order Lyman transitions, and \\ion{C}{3}~\\lam977, while the \\stis\\ observations cover Ly$\\alpha$ and the \\ion{N}{5}, \\ion{Si}{4}, and \\ion{C}{4} doublets. The different epochs permit us to investigate the effects of UV continuum variability on the intrinsic absorption. The analysis presented here will also take into account changes in absorption profiles that may be a result of differing contributions of the various emission sources, i.e. continuum versus broad lines, in different observation epochs. ",
        "conclusions": "\\label{sec-disc} The results for \\ion{O}{6}, \\ion{N}{5}, and \\ion{C}{4} are shown in Figures~\\ref{fig:o6land}-\\ref{fig:c4land}.  In each lower left panel, we show the blue and red components of the doublet profiles from the 2002 data. In the other three quadrants, we show the blue component of the 2002 profile in black.  Superimposed on this are the attempts to reproduce this profile from the  parameter constraints from the high signal-to-noise 2003 data described above, under variations in both the column densities and the covering fractions. The column density changes from 2002 to 2003 increase clockwise from top left to bottom right, with one panel for each ion showing how the profiles would change if {\\it no} change in the ion column density occured between 2002 and 2003.  Figure~\\ref{fig:o6land} shows that because the \\ion{O}{6} profile is saturated, the column density is unconstrained. For Component 2, regardless of the column density the only value of $C^{\\rm ILR}$ consistent with the profile is zero, while for Component 4, $C^{\\rm ILR}=0.5 C^{\\rm BLR}$.  The photoionization model predicts an increase of about a factor of four in the \\ion{O}{6} column density from 2002 to 2003, given the change in the continuum level over that period.  For \\ion{N}{5}, shown in Figure~\\ref{fig:n5land}, the best fit solution for both Components 2 and 4 is a column density decrease of $\\sim$30\\% from 2002 to 2003 and $C^{\\rm ILR}=0$, although $C^{\\rm ILR}=0.5C^{\\rm BLR}$ also provides an acceptable fit for Component 4. Our pure photoionization model would predict a decrease of about one half that amount in the \\ion{N}{5} column density between 2002 and 2003.  For \\ion{C}{4}, Figure~\\ref{fig:c4land} shows that a decrease in the column density from 2002 to 2003 is required to match the profiles in both components, regardless of the value of $C^{\\rm ILR}$.  A change in the column density by larger than a factor of $\\sim$2, however, is not consistent with the observed absorption depths. Our photoionization modeling indicates that we would expect a larger decrease in the \\ion{C}{4} column density from photoionization alone, a factor of nearly six. This covering fraction is relatively poorly constrained because the relative contributions of the BLR and ILR change very little between the 2002 and 2003 epochs, as noted above. Of our models, the best fits to the 2002 spectrum come from $N$(\\ion{C}{4})$_{2002}$=$2 \\times N$(\\ion{C}{4})$_{2003}$ and $C^{\\rm ILR}=0$ for Component 2 and $C^{\\rm ILR}=C^{\\rm BLR}$ for Component 4.   This ILR covering fraction solution for Component 4, however, is not significantly different from $C^{\\rm ILR}=0.5C^{\\rm BLR}$.  In the formal solution for these covering fractions we assumed that the effective line covering fractions for \\ion{O}{6}, \\ion{N}{5}, and \\ion{C}{4} are all the same in the 2003 epoch but that the resulting individual BLR and ILR covering fractions are not necessarily equal, and that those individual BLR and ILR covering fractions did not change between 2002 and 2003. However, from the success of the G05 global CNO covering fraction fits, the covering fractions for each of these highly ionized species are likely to be comparable. Thus, we summarize the results detailed above in this context and find that (1) the BLR covering fraction of \\ion{C}{4}, \\ion{N}{5}, and \\ion{O}{6} in the cores of Components 2 and 4 are $\\sim$0.8-0.9; (2) the  ILR covering fractions are zero for Component 2 and $\\sim$0.4 for Component 4; and (3) the column density in \\ion{O}{6} decreased by an indetermiate factor while the column density in \\ion{N}{5} decreased by one third, and that of \\ion{C}{4} decreased by a factor of $\\sim$2 between the two epochs. Should we choose to constrain the individual broad and intermediate emission line covering fractions of \\ion{C}{4}, \\ion{N}{5}, and \\ion{O}{6} to be equal (like G05, but unlike S04), we could use this set of three doublets to solve for six variables, the continuum, broad line, and emission line covering fractions and the column densities of \\ion{O}{6}, \\ion{N}{5}, and \\ion{C}{4} as a function of velocity across the absorption profiles. We leave this for future work.  The \\ion{C}{3} profile presents somewhat of a puzzle in light of the results we have outlined thusfar.  We expect the column density of this species to be quite sensitive to the factor of $\\sim$8 increase in the ionizing flux between the 2002 and 2003 observations. Figure~\\ref{fig:ratio} shows that for $U_{2002}=0.01$, the \\ion{C}{3} column density should decrease by a factor of about 25-30 from 2002 to 2003. In Figure~\\ref{fig:c3}, we demonstrate that we see no significant change in the \\ion{C}{3} profile in either Component 4 or in Component 2.  One way to resolve this problem is to invoke a long photoionization response time for the gas. To test this possibility, we follow an equilibrium photoionization model of the gas over time.  We began with a model that matches the well-constrained 2003 column densities and dropped the photoionizing flux by a factor of 8 to evolve the absorber back in time to the 2002 epoch. Following the trend of ionic abundances, we find that the \\ion{C}{3} column density should decrease by a factor of two in about one hour for $n_{\\rm e} = 3 \\times 10^4$ cm$^{-3}$, in about one day for $n_{\\rm e} = 1500$ cm$^{-3}$ and in about one year for $n_{\\rm e} = 5$ cm$^{-3}$. Therefore, for a very small change in the \\ion{C}{3} column density, a small total density is required if the absorbing gas is in equilibrium. This density, along with $\\log U = -1$ and the luminosity of Mrk~279, require that the gas lie at $\\sim$2 kpc from the ionizing source if it is part of the Mrk~279 outflow. This would certainly present us with a conundrum given the nonunity covering fractions we have consistently found for the emission lines in absorption Components 2 and 4.  The stable \\ion{C}{3} profiles are not likely solely the result of problems with the normalization of the spectra.  We confirm this by comparing the profiles of two nearby ISM lines in the three \\fuse\\ epochs. In Figure~\\ref{fig:ism} we show two prominent but unsaturated ISM lines, \\ion{S}{3}~\\lam1012 and \\ion{Ar}{1}~\\lam1048, in all three \\fuse\\ observation epochs. We do see some normalization problems, particularly in the  \\ion{Ar}{1} region, likely attributable mainly to the low S/N in the 2002 data. The absorption depth in the 2002 \\ion{Ar}{1} profile is higher than in the other epochs when it is not expected to vary at all. This may be a contributor to the apparent lack of variability in the \\ion{C}{3} profiles. We note, however, that the deepest absorption in the \\ion{C}{3} profile peaks at -540 \\kms, -460 \\kms, and at -385 \\kms. As discussed by S04, the low ionization absorption at these velocities may not be part of the Mrk~279 outflow but may in fact be attributable to either the host galaxy of the AGN or to a companion galaxy, MCG +12-13-024.  Photoionization models of quasar line emission indicate that intermediate line region lies at approximately 1 pc from the central ionizing source, a factor of $\\sim$10 further than than the broad line region (Brotherton et al.\\ 1994). However, the UV and X-ray observations of the Seyfert galaxy NGC~4151 place the D+E subcomponents of its outflow absorption in an accretion disk wind which is the ILR and which lies at only 0.1 pc from the continuum source (Kraemer et al.\\ 2006, Crenshaw \\& Kraemer 2007, see also Corbin 1997). For Mrk~279, we have found ILR covering fractions that are consistently smaller than those of the BLR and continuum, and in some cases $C^{\\rm ILR}=0$.  This would place the absorbing gas within or interior to the ILR.  Assuming that the width of the intermediate lines, $\\sim$2900~\\kms, is produced gravitationally, a black hole mass of $10^{7.54}$ M$_{\\sun}$ (Peterson et al.\\ 2004) gives a radius of $\\sim$0.02 pc for the ILR, an estimate that is only slightly larger than measured constraints on the size of the BLR, $\\sim$0.01 pc (Maoz et al.\\ 1990). We note, however, that this estimate of the ILR size is influenced by the relatively broad ILR emission features, certainly as compared with NGC~4151."
    },
    "0812/0812.4218.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {To understand the origin of stellar activity in pre-main sequence Herbig Ae/Be stars and to get a deeper insight in the interior of these enigmatic stars, the pulsational instability strip of Palla and Marconi is investigated. In this article we present a first discovery of non radial pulsations in the Herbig Ae spectroscopic binary star RS Cha.} % aims heading (mandatory) {The goal of the present work is to detect for the first time directly by spectrographic means non-radial pulsations in a Herbig Ae star and to identify the largest amplitude pulsation modes.} % methods heading (mandatory) {The spectroscopic binary Herbig Ae star RS Cha was monitored in quasi-continuous observations during 14 observing nights (Jan 2006) at the 1m Mt John (New Zealand) telescope with the Hercules high-resolution echelle spectrograph. The cumulated exposure time on the star was 44 hrs, corresponding to 255 individual high-resolution echelle spectra with $R = 45000$. Least square deconvolved spectra (LSD) were obtained for each spectrum representing the effective photospheric absorption profile modified by pulsations. Difference spectra were calculated by subtracting  rotationally broadened artificial profiles; these residual spectra were analysed and non-radial pulsations were detected. A subsequent analysis with two complementary methods, namely Fourier Parameter Fit (FPF) and Fourier 2D (F2D) has been performed and first constraints on the pulsation modes have been derived. } % results heading (mandatory) {For the very first time we discovered by direct observational means using high resolution echelle spectroscopy non radial oscillations in a Herbig Ae star. In fact, both components of the spectroscopic binary are Herbig Ae stars and both show NRPs. The FPF method identified 2 modes for the primary component with (degree $\\ell$, azimuthal order $m$) couples ordered by decreasing probability:   $f_1$ = 21.11 \\cd with ($\\ell$,$m$) = (11,11), (11,9) or (10,6) and  $f_2$ = 30.38 \\cd with ($\\ell$,$m$) = (10,6) or (9,5). The F2D analysis indicates for $f_1$ a degree $\\ell$ =  8-10. For the secondary component, the FPF method identified 3 modes with ($\\ell$,$m$) ordered by decreasing probability:   $f_1$ = 12.81 \\cd with ($\\ell$,$m$) = (2,1) or (2,2),  $f_{\\mathrm 2b}$ = 19.11 \\cd with ($\\ell$,$m$) = (13,5) or (10,5) and  $f_3$ = 24.56 \\cd with ($\\ell$,$m$) = (6,3) or (6,5). The F2D analysis indicates for $f_1$ a degree $\\ell$ =   2 or 3, but  proposes a contradictory identification of $f_2$ as a radial pulsation ($\\ell$ = 0). } % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Asteroseismology represents a modern tool for studying the stellar interiors of the enigmatic group of pre-main sequence (PMS) stars of intermediate mass ($\\sim$ 2-8 M$_{\\odot}$),  so called Herbig Ae/Be stars (\\cite{Herbig};  \\cite{Strom}; \\cite{fimu}; \\cite{fija}). Since their first systematic classification in 1960 (\\cite{Herbig}) this group of young stars has been extensively studied, but one of the major questions remains unanswered: how can the intense stellar activity and strong stellar winds (\\cite{praderie};  \\cite{catala}; \\cite{caku}; \\cite{tba6} and  \\cite{tba8}), as well as the many highly variable emission lines observed in the spectra of these stars be explained?  Magnetism is frequently invoked as being responsible for active stellar phenomena. The position of the Herbig stars in the HR diagram indicates that they are in the radiative phase of their contraction towards the main sequence (\\cite{iben}; \\cite{gilliland}), and should in principle not possess any outer convective zone. In its absence, the classical magnetic dynamo mechanism could not be responsible for the observed active phenomena. Moreover, recent spectropolarimetric observations of Herbig stars indicate the presence of significant large-scale magnetic fields only in a small fraction of them (\\cite{tba16}, \\cite{tba18} and \\cite{alecian2008}). This fraction, close to 10\\%, is consistent with a magnetic field distribution similar to that of the main sequence stars. In that scenario, magnetic Herbig Ae stars would be the progenitors of main sequence Ap/Bp stars. These recent results are in agreement with the primordial fossil field hypothesis. Unless magnetic fields  in Herbig Ae/Be stars are of very complex nature and can therefore not be detected by spectropolarimetric means, in opposition to clearly organized dipolar fields, the pure magnetic origin of activity in these stars is more than doubtful.  For many years non-classically generated magnetic fields have been used as a potential explanation for the observed active phenomena. Even if the presence of complex fields can not be ruled out it is of major importance to investigate other possible external or internal origins of this tremendous amount of dissipated energy, as witnessed by chromospheres and coronae, but also variable spectral lines, winds and bipolar jets.  The only way of studying in detail the internal stellar structure is the analysis and the modelling of stellar pulsations, if observed. As an example, PMS stars gain their energy from gravitational contraction and therefore differ signi\u00decantly from post-main-sequence stars having already processed nuclear material - stellar pulsations are sensitive to these differences of internal stellar structure (see e.g. \\cite{suran}). As of today,  the internal stellar structure of PMS stars is not yet well constrained. Since few years the existence of pulsating intermediate mass PMS stars is known (\\cite{breger1972}; \\cite{kuma}; \\cite{donati}). This observational result motivated Marconi \\& Palla (\\cite{mapa}) to investigate the  pulsation characteristics of HR\\,5999 theoretically, which enabled them to predict the existence of a pre-main-sequence instability strip, which is being crossed by most of the intermediate mass PMS objects for a significant fraction of their evolution to the main sequence. This strip covers approximately the same area in the HR diagram as the $\\delta$ Scuti variables. Zwintz (\\cite{zwintz08}) compared, based on photometry, the observational instability regions for pulsating pre-main sequence and classical $\\delta$ Scuti stars and concluded that the hot and cool boundaries of both HR diagram instability regions seem to coincide. This preliminary result deserves further study by aim of full asteroseismological approach based on spectroscopy.  As of today, more than 30 intermediate-mass PMS stars have revealed to be pulsating at time-scales typical of $\\delta$ Scuti stars (see e.g. \\cite{kuma}; \\cite{kuca}; \\cite{donati};  \\cite{tba15};  \\cite{marconi2002};  \\cite{rima}; \\cite{zwintz}; \\cite{ca} and references therein).  RS Chamaeleontis is a bright spectroscopic eclipsing binary star. Both components are Herbig Ae PMS stars of similar mass (close to 1.9 M$_{\\odot}$).  Recently the age of RS\\,Cha has been determined to 6$^{+2}_{-1}$\\,Myr  (\\cite{luhman}), which verifies it's PMS nature. \\cite{andersen} already reported small amplitude radial velocity variations on top of the binary radial velocity curve for both components of RS Cha, suggesting the possible presence of stellar pulsations. Photometric observations by \\cite{mcinally} revealed short-term variations in at least one of the two components, possibly linked to stellar pulsations. Very recently, Alecian et al. (\\cite{alecian2005}) reported radial velocity variations  in the residual velocity frame (cleaned for orbital velocity) with amplitudes up to a few \\kms\\ and periods of the order of 1h, indicative of $\\delta$ Scuti type pulsations.  The aim of our study of the two components of RS Cha is to provide a first set of asteroseismic constraints for forthcoming non-radial pulsation models by determining unambiguously a higher number of periodicities and identifying, in a second step, the corresponding pulsation modes with their respective degree $\\ell$ and azimuthal number $m$.  To achieve this goal,  we decided to perform high resolution spectroscopic observations on a large time basis and with optimized time coverage.  Section \\ref{previous} reviews previous related work, Sect. \\ref{observations}  describes the observations and data reduction, Sect. \\ref{orbit}  summarizes results of the orbit determination, Sect. \\ref{detection} reveals the detection of non-radial pulsations in both components of RS Cha, Sect. \\ref{pulsaprim}  and \\ref{pulsasecond}  present frequency analysis and moment identification in the primary and secondary component, respectively. A discussion is proposed and a conclusion is drawn in Section \\ref{dissconc}.    %__________________________________________________________________ ",
        "conclusions": "\\label{dissconc}  For the first time we detected by direct spectroscopic means non-radial pulsations in both components of a binary Herbig Ae star. We identified  with a high level of confidence in the profile variations of the primary  and secondary component two, respectively three pulsation frequencies. A first identification by two complementary methods revealed very different main pulsation modes for the two components: while the dominant mode of the primary component seems to be a high degree pro-grade mode with $\\ell$ = 10 or 11, the dominant mode of the secondary component is identified as a low degree mode with $\\ell$ = 0,1 or 2. Two different methods have been applied to perform a mode identification of both components. The Fourier 2D method as presented by \\cite{kennelly} provides direct insights on the spatial structure of the oscillation. However, the derived apparent $\\left | m \\right |$ value seems to be linked to $\\ell$ in the way described here-above, but more detailed simulations for different configurations have to be carried out in order to understand potential biases in the degree $\\ell$ determination. We therefore privilege the $\\ell$ determination as provided by the FPF method. Table \\ref{mainresults} summarizes the main results for the dominant pulsation modes of each component. As can be seen, both methods yield very different results on the determination of the degree $\\ell$ of $f_{\\mathrm 2b}$ of the secondary component, perhaps due to an insensibility of the F2D method to tesseral mode signatures.  \\begin{table} \\caption[]{Frequency search and mode identification. Summary of the main results for the dominant frequencies of each component. The ($\\ell$,m) couples indicated for the FPF method correspond to the best identification, in decreasing order. For each of the frequencies detected by the FPF method, we identified an $\\ell$ value as detected by the F2D method. \\\\} \\begin{tabular}{cccc} \\hline\\hline \\multicolumn{4}{c}{\\bf Primary component} \\\\ &&&\\\\ & d$^{-1}$ & Identification FPF  & Identification F2D \\\\\\hline &                 &     best ($\\ell$,m)    &   \\\\ ~~~~~ $f_1$ ~~~~~~ &   21.11 & (11,11), (11,9), (10,6) &  $\\ell$ = 8-10\\\\ $f_2$  & 30.38  &   (10,6), (9,5), (11,7) & \\\\\\hline &&&\\\\\\hline\\hline \\multicolumn{4}{c}{\\bf Secondary component} \\\\ &&&\\\\ & d$^{-1}$ & Identification FPF  & Identification F2D \\\\\\hline &                 &     best ($\\ell$,m)    &   \\\\ $f_1$  &  12.81  &  (2,1), (2,2) &  $\\ell$ =  2 or 3\\\\ $f_{\\mathrm 2b}$  &  19.11  & (13,5), (10,5) &  $\\ell$ = 0\\\\ $f_3$  &  24.56  & (6,3), (6,5) & \\\\   \\hline \\hline \\end{tabular} \\label{mainresults} \\end{table}  The precise redetermination of the binary orbit reveals a slight gradient change of the observed - calculated phase shift curve, which might be the signature of changes in the orbital period of the order of  $\\Delta P$/P = (7.5 $\\pm$ 0.7) 10$^{-6}$. This tendency, if confirmed, might be due to the presence of a third body in the RS\\,Cha system or due to intrinsic orbital period changes, which then need to be explained.  Next steps require a precise redetermination of the fundamental parameters and a search for the intrinsic rotation periods of both components. The first frequencies are now known and their associated modes are constrained by the present work. In subsequent work, a numerical model will be developed in order to constrain the internal structure of both components."
    },
    "0812/0812.4524_arXiv.txt": {
        "abstract": "The two-component, core-crust, model of a neutron star with homogenous internal and dipolar external magnetic field is studied responding to quake-induced perturbation by substantially nodeless differentially rotational Alfv\\'en oscillations of the perfectly conducting crustal matter about axis of fossil magnetic field frozen in the immobile core. The energy variational method of the magneto-solid-mechanical theory of a viscoelastic perfectly conducting medium pervaded by magnetic field is utilized to compute the frequency and lifetime of nodeless torsional vibrations of crustal solid-state plasma about the dipole magnetic-moment axis of the star. It is found that obtained two-parametric spectral formula for the frequency of this toroidal Alfv\\'en mode provides fairly accurate account of rapid oscillations of the X-ray flux during the flare of SGR 1806-20 and SGR 1900+14, supporting the investigated conjecture that these quasi-periodic oscillations owe its origin to axisymmetric torsional oscillations predominately driven by Lorentz force of magnetic field stresses in the finite-depth crustal region of the above magnetars. ",
        "introduction": "The quasi-periodic oscillations (QPOs) recently detected in the X-ray flux of flaring SGR 1806-20 and SGR 1900+14 and interpreted as a manifest of quake-induced torsional vibrations of underlying magnetar (Israel et al 2005, Watts \\& Strohmayer 2006) are among the topical themes in current development of the astroseismology of this subclass of neutron stars (e.g. Glampedakis et al 2007; Samuelssen \\& Andersson 2007; Levin 2007, van Hoven \\& Levin 2008, Lee 2007, 2008; Sotani et al 2007, 2008; Vavoulidis et al 2008) whose bursting and quiescent electromagnetic activity is dominated, as is commonly believed, by frozen-in fossil magnetic field of extraordinary intensity (e.g., Woods \\& Thompson   2006, Mereghetti 2008). Motivated by this interest, in (Bastrukov et al 2007, 2008, 2009) several possible scenarios of post-quake relaxation of magnetar has been investigated with emphasis on regime of nodeless torsional vibrations about the dipole magnetic-moment axis of these seismically active neutron stars. In particular, in (Bastrukov et al 2007b, 2008), a case of the elastic-force-driven nodeless shear oscillations, both torsional --  $_0t_\\ell$ and spheroidal -- $_0s_\\ell$, locked in the finite-depth crust $\\Delta R$ has been studied in some details. It was found that obtained two-parametric spectral equations for the frequency of pure elastic spheroidal and torsional modes as a function of multipole degree $\\ell$ of nodeless shear oscillations of the form (Bastrukov et al 2008) \\begin{eqnarray} \\label{e1.1} && \\omega(_0s_\\ell)=\\omega_s(\\ell,\\omega_e,\\lambda),\\quad \\omega(_0t_\\ell)=\\omega_t(\\ell,\\omega_e,\\lambda),\\quad\\quad \\ell\\geq 1,\\\\ \\label{e1.2} && \\omega_e=\\frac{c_t}{R},\\quad c_t=\\sqrt{\\frac{\\mu}{\\rho}},\\quad \\lambda=1-\\frac{\\Delta R}{R},\\quad \\Delta R=R-R_c \\end{eqnarray} provide proper account of the low-frequency QPOs frequencies, supporting $\\ell$-pole identification proposed in works (Watts \\& Strohmayer 2007, Samuelssen \\& Andersson 2007) and it was shown that the unidentified before lowest overtones of QPOs in light-curve of flaring SGR 1806-20 with frequencies 18 and 26 Hz can be interpreted as manifest of dipole, $\\ell=1$, spheroidal and torsion nodeless elastic shear vibrations, exhibiting features generic to Goldstone soft modes.  On the other hand, one can perhaps cast doubt on such interpretation of QPOs as being produced by seismic vibrations driven by solely elastic restoring force, because it rests on a poorly justifiable assumption about dynamically passive role of ultra-strong magnetic field defining only direction of axis about which the stellar matter undergo torsional nodeless oscillations and remaining unaltered in the process of post-quake relaxation of magnetar. Instead, it seems more plausible to expect that explosive liberation of magnetic field energy resulting in the crust fracture by relieved magnetic field stresses -- magnetic breaking associated with disruption of the magnetic field lines -- must be followed, after the magnetic field lines reconnection, by the Lorentz-force-driven Alfv\\'en differentially rotational oscillations of the perfectly conducting star matter about axis of frozen-in homogeneous magnetic field, either in the entire volume or in the peripheral finite-depth crust. This line of argument has been taken into account in the magneto-solid-mechanical model of global torsional nodeless vibrations about the dipole magnetic moment axis of neutron star driven by the joint action of Hooke's force of elastic shear stresses and Lorentz force of magnetic field stresses (Bastrukov et al 2009). The computed in this latter work the two-parametric frequency spectrum of the form \\begin{eqnarray} \\label{e1.3} && \\omega(_0t_\\ell)=\\omega_t(\\ell,\\omega_e,\\beta),\\quad\\quad \\beta=\\frac{\\omega_A^2}{\\omega_e^2}=\\frac{B^2}{4\\pi\\mu},\\quad\\quad \\ell\\geq 2,\\\\ \\label{e1.4} && \\omega_e^2=\\frac{c_t^2}{R^2},\\quad c_t^2=\\frac{\\mu}{\\rho},\\quad\\omega_A^2=\\frac{v_A^2}{R^2},\\quad v^2_A=\\frac{B^2}{4\\pi\\rho} \\end{eqnarray} was found to be consistent with data on QPOs with frequencies $\\nu=\\omega/2\\pi$ from the range $30\\leq \\nu \\leq 200$ Hz interpreted as being produced by these global seismic vibrations of multipole degree $\\ell$ from interval $2\\leq \\ell\\leq 20$, but unable to explain the above mentioned lowest overtones in the QPOs spectrum of SGR 1806-20. Also, from the viewpoint of both outlined models the nature of restoring force responsible for high-frequency QPOs with $\\nu=626$ and  $\\nu=1837$ Hz remains uncertain.    In this paper, continuing investigations reported in (Bastrukov et al 2007, 2008, 2009), we focus on different, to the above mentioned scenarios, case of post-quake vibrational relaxation of a neutron star. Specifically, we consider torsional nodeless Alfv\\'en oscillations of the crustal solid-state plasma, that is, seismic vibrations restored by solely Lorentz force of magnetic field stresses and excited in the peripheral finite-depth seismogenic layer of neutron star model with homogeneous internal and dipolar external magnetic field. In section 2, a brief outline is given of equations of magneto-solid-mechanics appropriate for the perfectly conducting viscoelastic continuous medium pervaded by magnetic field. In section 3, the spectral formulae for the frequency and lifetime of this toroidal Alfv\\'en vibration mode are obtained followed by discussion of inferences that can be made from considered scenario of seismic vibrations regarding electrodynamic properties of the neutron star matter with homogeneous configuration of the frozen-in magnetic field. The relevance of considered asteroseismic model to the detected QPOs in the X-ray flux during the flare of SGR 1806-20 and SGR 1900+14 is assessed in section 4. The newly obtained results are briefly summarized in section 5. ",
        "conclusions": "The neutron stars, both pulsars and magnetars, are seismically active compact objects of the end point of the evolutionary stellar track. They are, as is widely believed, to be formed in the magnetic-flux-conserving supernova collapse of the cores of massive stars coming into existence, as sources of pulsed electromagnetic emission, with a fairly complex internal constitution (e.g. Weber 1999, Bisnovatyi-Kogan 2002, Lattimer \\& Prakash 2007). The past two decades have seen increasing recognition that valuable tool of looking into the neutron star interior provides asteroseimology. The explosive interest to current development of this domain of pulsar astrophysics has been stimulated by the above mentioned discovery of quasi-periodic oscillations during the giant X-ray flares of SGR 1806-20 and SGR 1900+14, which has also prompted extensive investigations of radiative processes caused by ultra-strong magnetic fields of magnetars (e.g. Heyl \\& Hernquist 2005, Harding \\& Lai 2006, Zhang et al 2007, Timokhin et al 2008, Bo Ma et al 2008, Rea et al 2008) and processes of their cooling (Yakovlev et al 2008).  Many agree today that detected QPOs can be explained on the basis of asteroseismic models of a magnetar undergoing quake-induced torsional oscillations about its the dipole magnetic-moment axis. Adhering to this attitude and taking into account that the main energy source of bursting high-energy emission of magnetars is the immense store of the magnetic field energy, we have set up and investigated one of such models interpreting the above QPOs in terms of axisymmetric Alfv\\'en nodeless oscillations of the crustal solid-state plasma about axis of homogeneous internal and dipolar external magnetic field frozen in the motionless core. As is demonstrated in Figs. 7 and 8, the obtained two-parametric spectral equation of this model provides proper account of general trends in data on QPOs frequencies as well as quite reasonable their $\\ell$-pole identification. The main argument in favor of such interpretation is fairly accurate description of both the low-frequency (18 and 26 Hz) and high-frequency (627 and 1837 Hz) QPOs in X-ray flaring SGR 1806-20, which is presented in Fig.8, supporting the underlying this model conjecture about dominate role of Lorentz force of magnetic field stresses in maintaining the nodeless torsional oscillations of perfectly conducting solid-state plasma entrapped in the peripheral finite-depth seismogenic zone of magnetar.    The physically meaningful framework for discussion of considered scenario of post-quake relaxation of magnetar provides the two-component, core-crust, model of quaking paramagnetic neutron star (Bastrukov et al 2002a, 2002b, 2003) which attributes the core matter incapability of supporting Alfv\\'en differentially rotational shear oscillations to poor electrical conductivity of central region of the star, considering the magnetar core as a motionless spherical bar magnet composed of degenerate Fermi-matter of non-relativistic neutrons with spin magnetic moments polarized along the uniform fossil magnetic field frozen in the star on the stage of catastrophic contraction of its massive progenitor. Together with this, one cannot of course state that such an image of the quaking magnetar interior is unique, that is, that the above QPOs cannot be explained by alternative conjectures about electrodynamic properties of neutron star material, restoring forces and volumetric character of torsional oscillations (e.g. Bastrukov et al 2007b 2008, 2009). One more of such uninvestigated alternatives is the regime of nodeless global differentially rotational vibrations of neutron star with non-homogenous internal and dipolar external magnetic field presuming, thus, that stellar material in the total volume of magnetar possesses properties of perfect conductor. We shall return to such a model of quaking neutron star in a forthcoming paper."
    },
    "0812/0812.1906_arXiv.txt": {
        "abstract": "Cosmic rays accelerated by a shock form a streaming distribution of outgoing particles in the foreshock region. If the ambient fields are negligible compared to the shock and cosmic ray energetics, a stronger magnetic field can be generated in the shock upstream via the streaming (Weibel-type) instability. Here we develop a self-similar model of the foreshock region and calculate its structure, e.g., the magnetic field strength, its coherence scale, etc., as a function of the distance from the shock. Our model indicates that the entire foreshock region of thickness $\\sim R/(2\\Gamma_{\\rm sh}^2)$, being comparable to the shock radius in the late afterglow phase when $\\Gamma_{\\rm sh}\\sim1$, can be populated with large-scale and rather strong magnetic fields (of sub-gauss strengths with the coherence length of order $10^{17}\\ {\\rm cm}$) compared to the typical interstellar medium magnetic fields. The presence of such fields in the foreshock region is important for high efficiency of Fermi acceleration at the shock. Radiation from accelerated electrons in the foreshock fields can constitute a separate emission region radiating in the UV/optical through radio band, depending on time and shock parameters. We also speculate that these fields being eventually transported into the shock downstream can greatly increase radiative efficiency of a gamma-ray burst afterglow shock. ",
        "introduction": "Do gamma-ray bursts (GRBs) accelerate cosmic rays (CRs)? There are arguments in favor of such an idea \\citep{DA04}. It is supposed that CRs are accelerated via the Fermi mechanism in which a particle crosses the shock many times and gradually gains its energy. For an ultra-relativistic shock, a steady-state universal power-law energy distribution of particles shall form \\citep{Kirk+00,Acht+01}. The shock shall accelerate all particle species, but the electrons being much lighter than protons and ions will also loose their energy via synchrotron cooling. This radiation is thought to be observed as the delayed afterglow emission of GRB sources. A problem immediately arises here from a simple estimate: if the X-ray afterglow observed on a day timescale after the prompt burst is indeed the synchrotron radiation from the shock-accelerated electrons, then the pre-shock medium has to be highly magnetized with the fields of milligauss strengths \\citep{LW06}. Thus, either the magnetic field is somehow generated in the shock upstream, or the conventional paradigm of the GRB afterglow needs revision.  In this paper, we present a self-similar model of the large region in front of a relativistic shock -- the foreshock. This region is populated with the shock-accelerated particles, which stream away from the shock into the collisionless ambient medium and generate magnetic field via a streaming (Weibel-type) instability. The model predicts the generation of strong, sub-gauss, magnetic fields in the entire foreshock whose thickness is $\\sim R/(2\\Gamma_{\\rm sh}^2)$ and is comparable to the shock radius, $\\sim10^{17}-10^{18}\\ {\\rm cm}$, in the afterglow phase a day or more after the explosion, when the shock is weakly relativistic or non-relativistic. The fields are sustained against dissipation by the anisotropy of newly accelerated particles. Moreover, these fields are relatively large-scale, with the coherence length being as large as a fraction of the foreshock size, $\\sim10^{16}\\ {\\rm cm}$, which makes them effectively decoupled from dissipation. We speculate, that these mesoscale magnetic fields being ultimately advected into the shock downstream can significantly increase the radiative efficiency of GRB afterglows and, perhaps, explain the origin of the magnetic field in an external shock of a GRB. We remark here, however, that our study is analytical and cannot account for a number of nonlinear feedback effects of the generated fields and pre-conditioned external medium onto the shock structure and particle acceleration. Kinetic and hybrid computer modeling is essential for better and more accurate understanding of the foreshock structure. ",
        "conclusions": "Here we presented a model of a self-similar foreshock produced by protons scattered by a relativistic shock into the unmagnetized or weakly magnetized external medium. It is immediately applicable to the external shock producing a GRB afterglow. The model predicts that a large region in front of the shock, of thickness of order the shock radius, shall be filled with relatively strong and large-scale magnetic fields. The upstream magnetic field strength and correlation length depend on the distance from the shock and the power-law index of accelerated protons (cosmic rays) and are given by Eqs. (\\ref{lambda-x}), (\\ref{B-x}), (\\ref{x0B0}) in the shock co-moving frame. The overall energetics of the field is dominated by large distances; hence the average foreshock B-field is of sub-milligauss strength and is increasing toward the shock while its typical coherence length is of order of few percent of the shock radius and is decreasing toward the shock.  The result is interesting, especially in the light of observational constraints on the particle acceleration in GRB afterglows. It has been shown that a few-milligauss magnetic fields are needed in front of the shock in order to efficiently Fermi-accelerate the electrons to the energies required to produce the observed X-ray emission \\citep{LW06}. Our model provides a possible and rather natural mechanism for generation of such fields in the far upstream medium. We also make a prediction that the shock-accelerated electrons will be radiating in the foreshock fields at a characteristic synchrotron frequency given by Eqs. (\\ref{nu-obs}). For reference, $\\nu_m\\sim 10\\ {\\rm THz}$ at about a day after the burst. It is possible that nonlinear effects omitted in this analysis (see discussion below) may limit the field strength to lower values compared to the present analysis, so the synchrotron peak can move to cm/mm-wave band. We can speculate that the emission from the foreshock can form an emission region separate from the afterglow shock and show up in the X-ray/optical band in the early afterglow phase and in the radio at the very late times.  The presented results can have interesting implications for the radiative efficiency of external shocks. Unlike internal shocks, where electron cooling is extremely fast and a thin shock layer of thickness of a hundred ion inertial lengths may be enough to produce the observed prompt emission \\citep{MS09}, the Weibel shock model \\citep{ML99} seems to face the efficiency problem when it is applied to an external shock. In such a shock, magnetic fields shall occupy a much larger region, perhaps the entire downstream region, in order for the shock to produce enough photons that will be observed as a delayed afterglow emission. This is not very likely (though not proved to be impossible, yet) provided that the small-scale fields generated {\\em at} the shock can be subject to rapid dissipation. However, dissipation shall be of much lesser importance for the foreshock fields, which have a (much) larger coherence length, Eq. (\\ref{lambda-x}). In the steady state, the fields generated in the upstream are advected to the shock and their strength is maintained against dissipation by the anisotropy of the continuously ``refreshed'' CR distribution. Hence, one can expect that the magnetic field near the shock will have the spectrum given in Eq. (\\ref{Bspec}). Once these fields pass though the shock into the downstream, they are enhanced by the shock compression and begin to decay. Commonly, dissipation is proportional to the inverse gradient scale squared $\\propto\\partial_x^2\\propto\\lambda^{-2}$, so that the skin-length-scale fields may eventually disappear. However, the fields above a certain coherence length, $\\lambda_{\\rm diss}$, shall survive and fill up the post-shock medium. The mechanism of dissipation is not yet understood in detail, so it is premature to make any quantitative conclusions about $\\lambda_{\\rm diss}$, but it will certainly be much smaller than $\\lambda(x_{\\rm max})\\lesssim R$, see Eq. (\\ref{lambda-xmax}). Since $B_\\lambda$ is a weak function of $\\lambda$ one can speculate that relatively strong fields, perhaps of order tens or hundreds milligauss, to occupy the post-shock medium.  We want to note that a number of simplifying assumptions has been made in our analysis. In particular, {\\em nonlinear feedback} effects of the upstream magnetic field on the particle distribution, on the shock structure and on Fermi acceleration were omitted. The inclusion of these effects is hardly possible in any analytical model. We also assumed that a steady state exists for the shock-foreshock system at hand. Apparently, it is not at all clear whether the steady state is at all possible or the system exhibits an intermittent behavior. One can envision a scenario in which the CRs overproduce upstream magnetic fields leading to enhanced particle scattering and the overall preheating of the ambient medium, which, in turn, can cause the shock to weaken,  disappear and then re-appear in a different place further upstream. Presently available 2D PIC simulations of an electron-position shock do show the upstream field amplification and no steady state has been achieved: both upstream and downstream fields continue to grow for the duration of the simulations \\citep{KKSW08}. We argue that extensive PIC or/and hybrid simulations of a shock are imperative for further study.  Finally, we mention that our model compliments other models of the magnetic field generation. It is reasonable to expect that the field can be amplified by vortical motions produced by the Richtmeyer-Meshkov instability, if the ambient medium is clumpy or if the shock velocity is not perfectly uniform \\citep{GM07,SG07,Milos+08}. On the other hand, if the ambient magnetic fields are strong enough, the fields can be generated via non-resonant \\citep{BellLucek01,Bell04,PLM08} or resonant \\citep{DM06,DM07,Z03,V+06,E+07} mechanisms, or both \\citep{Kirk+08}.    \\ack  The authors thank colleagues at IKI and RRC ``KI\" for discussions. This work has been supported by NSF grant AST-0708213, NASA ATFP grant NNX-08AL39G, Swift Guest Investigator grant NNX-07AJ50G and DOE grant  DE-FG02-07ER54940."
    },
    "0812/0812.0288_arXiv.txt": {
        "abstract": "A new multi-dimensional Hierarchical Structure Finder (HSF) to study the phase-space structure of dark matter in $N$-body cosmological simulations is presented. The algorithm depends mainly on two parameters, which control the level of connectivity of the detected structures and their significance compared to Poisson noise.  By working in 6D phase-space, where contrasts are much more pronounced than in 3D position space, our HSF algorithm is capable of detecting subhaloes including their tidal tails, and can recognise other phase-space structures such as pure streams and candidate caustics.  If an additional unbinding criterion is added, the algorithm can be used as a self-consistent halo and subhalo finder. As a test, we apply it to a large halo of the Millennium Simulation, where $19\\%$ of the halo mass are found to belong to bound substructures, which is more than what is detected with conventional 3D substructure finders, and an additional $23-36\\%$  of the total mass belongs to unbound HSF structures.  The distribution of identified phase-space density peaks is clearly bimodal: high peaks are dominated by the bound structures and show a small spread in their height distribution; low peaks belong mostly to tidal streams, as expected. However, the projected (3D) density distribution of the structures shows that some of the streams can have comparable density to the bound structures in position space.  In order to better understand what HSF provides, we examine the time evolution of structures, based on the merger tree history. Given the resolution limit of the Millennium Simulation, bound structures typically make only up to 6 orbits inside the main halo. The number of orbits scales approximately linearly with the redshift corresponding to the moment of merging of the structures with the halo.  At fixed redshift, the larger the initial mass of the structure which enters the main halo, the faster it loses mass. The difference in the mass loss rate between the largest structures and the smallest ones can reach up to $20\\%$. Still, HSF can identify at the present time at least 80\\% of the original content of structures with a redshift of infall as high as $z \\leq 0.3$, which illustrates the significant power of this tool to perform dynamical analyses in phase-space. ",
        "introduction": "\\label{sec:introduction}  When \\cite{Zwicky1933} studied galaxy velocities in clusters, he was the first to notice that there should be about one order of magnitude more matter in the universe than the observed amount of baryonic matter to explain the proper motions of galaxies through gravitational forces. Dark Matter (DM) was introduced to overcome this problem. Later on, the existence of a dark matter component was confirmed by the analysis of galaxy rotation curves \\citep{Rubin1970}. Recent studies of gravitational lensing \\citep[e.g.][]{lensing} and, more generally, multiwavelength observations in e.g. the COSMOS project \\citep{COSMOS} provide additional proofs for the existence of DM. Other constraints on the non-baryonic nature of DM were also set by the analysis of the Cosmic Microwave Background \\citep[e.g.][]{WMAP5}.  For the last three decades, the DM paradigm has been studied extensively in the context of cosmological $N$-body simulations. The comparison of structures formed in such simulations to observed ones excluded some of the theoretical models, such as Hot Dark Matter and led to the nowadays commonly accepted $\\Lambda$-Cold Dark Matter ($\\Lambda$CDM) model. In the $\\Lambda$CDM model, dark matter is collisionless, with a very small velocity dispersion at high redshift; structures are built in a hierarchical, bottom-up process, where small structures arise first, seeded from initial fluctuations, and then merge together to build up larger and larger structures, designated commonly as {\\it haloes}.  Inside the gravitational wells of these dark matter haloes, baryonic matter forms galaxies \\citep{White1978}.  Recently, the efforts to finally identify the physical nature of dark matter particles, either directly through detecting them in ground-based dark matter particle detectors or indirectly by observing their annihilation radiation, have intensified.  At the same time, the ever increasing resolution of N-body simulations \\citep[e.g.][]{Springel2008} puts new levels of demand on the field of theoretical study of non-linear halos. The careful analysis of cosmological structures moves from the study of spherically averaged three-dimensional density profiles \\citep{NFW} to the study of the full six dimensional phase-space. For example, this concerns the analyses of the properties of caustics described analytically in e.g. Bertschinger's secondary infall model \\citep{Bertschinger} and recently reviewed in the context of numerical simulations \\citep{Roya2008,White2008}. Investigations of full phase-space structures include accurate simulations of two dimensional phase-space \\citep{Alard2005,Colombi2007} and analyses relying on a metric approach to six dimensional phase-space in $N$-body simulations \\citep{Vogelsberger2008,White2008}.  A particularly important step in understanding dark matter clustering lies in an analysis of the bound structures found in $N$-body simulations. This is at present usually carried out with structure finders such as {\\small SUBFIND} \\citep{Springel2001}, {\\small ADAPTAHOP} \\citep{Aubert2004} or {\\small PSB} \\citep{Kim2006}. Following this path, we present a new multi-dimensional Hierarchical Structure Finder which complements all the above numerical methods with an effective and robust analysis of phase-space structures in full 6D space.  The paper is organised as follows. First, we review current structure finders in Section~\\ref{sec:structure}.  We then present our new multi-dimensional Hierarchical Structure Finder (HSF) in Section~\\ref{sec:hsf}. In Section~\\ref{sec:test}, we use our algorithm to detect and analyse phase-space structures of a large halo taken from the Millennium Simulation. We investigate the space of parameters on which our HSF algorithm depends and try to find the best choice of the parameters according to the application under consideration.  We also introduce the simulation merger tree to follow the evolution of structures in phase-space.  This allows us to analyse in detail a few representative cases.  This is followed by a quantitative analysis of HSF structures in the space and time domain. We also discuss the bimodal nature of the substructure population, in terms of bound structures versus tidal tails and tidal streams.  Finally, in Section~\\ref{sec:discussion} we give a summary and present our conclusions. ",
        "conclusions": "\\label{sec:discussion}  We introduced a new universal multi-dimensional hierarchical structure finder (HSF) which was employed here to study dark matter structures in six dimensional phase-space.  The algorithm used, for each particle, the phase-space density and local neighbourhood estimated with the SPH method with adaptive metric implemented in the EnBiD package \\citep{Sharma2005}. To detect structures, HSF builds on the SUBFIND and ADAPTAHOP algorithms, with the introduction of a new simple but robust cut or grow criterion depending on a single {\\rm connectivity parameter} $\\alpha$.  The main steps of the algorithm are as follows: (i) local phase-space density maxima are detected and structures around them are grown by following local gradients up to saddle points; (ii) at each saddle point level, the density $f$ of each structure is compared to Poisson noise: the structure is kept if $f$ is $\\beta$ times more significant than the Poisson r.m.s. noise level. At the same time its mass and the mass of its partner (connected to it through the saddle point) are measured and the cut or grow criterion is applied: if one structure is $\\alpha$ times smaller than its neighbour, then all particles below the saddle point are attached to the neighbour. When the two structures have comparable mass within a factor $\\alpha$, they are both set to grow down as before.  This criterion allows us to better trace substructures in phase-space, with a good control of the effect of Poisson noise, which is very important in this rather sparsely sampled space.  We demonstrated the potential of HSF on a large FOF dark matter halo taken from the Millennium Simulation. Our tests show that $\\beta$ and especially $\\alpha$ are important control parameters. To better study the smallest possible structures, $\\beta$ should be set close to $0$. The smaller $\\alpha$, the more subtle the structures found by the algorithm. In our analysis, we give preference to $\\alpha=0.2$, which provides a good balance between finding the finest possible substructures and not overgrowing them. This value of $\\alpha$ is particularly appropriate when an additional binding step is performed. In contrast, an analysis of tidal tails is best carried out with small $\\alpha$, around $0.01-0.001$, which separates structures into smaller pieces. It is possible to use it in combination with $\\beta=4-10$, which tends to reconnect the structures together in a consistent way, to reconstruct tidal tails rather well. A more advanced method of reconnecting phase-space structures, by using the topology of the hierarchical tree created by the HSF algorithm is under investigation.  We used the Millennium Simulation merger tree \\citep{Springel2005b} to compare the HSF phase-space structures found at the present time with the same structures traced back to the time just before they enter the main halo.  While the best three dimensional algorithms used presently, such as SUBFIND, manage to find only the main part of bound structures, HSF is capable of finding more extended bound components along with their tidal tails.  There is much more information about structure evolution still stored in phase-space than in 3D and this information can be potentially fully recovered from the data by a six dimensional algorithm such as HSF.  The main results of our analysis in time and space domain are the following: \\begin{itemize}  \\item HSF structures contain on average $80-100\\%$ of the mass inside the initial structures up to a redshift of merging $z=0.3-0.4$. This value drops down to $50\\%$ for $z=1$. On the other hand, bound HSF structures contain on average $80-100\\%$ of the mass inside the initial halos only up to $z=0.09$ and $50\\%$ up to $z=0.6$. This shift in the mass loss is caused by the existence of tidal tails, which are joined to HSF structures, but do not belong to their bound part. In other words we can say that HSF is able to reconstruct in most cases the full dynamical structures which enter the halo at redshifts as high as $z=0.3-0.4$.  \\item The distribution function of the phase-space density maxima $f_{\\rm max}$ of HSF structures clearly shows a bimodality. We can explain it by partitioning the structures into two distinct groups. In the first group, corresponding to the high phase-space density peak regime, $f_{\\rm max}\\approx0.2 \\funit$, with a small spread around that value, there are mostly bound structures. In the second group, corresponding to a $3$ orders of magnitude smaller phase-space density regime, $f_{\\rm max}\\approx3.3\\times10^{-5} \\funit$, and a larger spread around this value, there are all unbound structures i.e.~tidal tails, streams and possibly some caustics. In terms of ``peakness'', $c_f=f_{\\rm max}/f_{\\rm min}$, where $f_{\\rm min}$ is minimum value of $f$ in each substructure, this translates into $c_f=1.1\\times10^5,8.7$ for bound and unbound structures, respectively.  \\item We noticed, similarly as \\cite{Vass2008}, that the distribution function of the values of $f$ around each dark matter particle is close to lognormal for the smooth component of the halo and the unbound part of the substructures (tidal streams).  \\item We found that there is more mass in bound HSF substructures than in SUBFIND ones. Figure~\\ref{mass_profile} shows the cumulative mass of substructures divided by the total halo mass as a function of substructure mass.  Around $18.5\\%$ of halo mass is stored in bound HSF structures, and this quantity almost does not depend on $\\alpha$. In comparison, about $12.4\\%$ of the mass is stored in SUBFIND bound structures. The additional mass in HSF bound structures comes mainly from the fact that subhaloes are better defined in phase-space and are more extended. However, the set of identified bound substructures is nearly identical in both methods, and hence the cumulative abundance of substructures as a function of maximum circular velocity is the same as well.  \\begin{figure}% \\includegraphics[width=8.5cm,height=8.5cm]{mass_profile5.eps} \\caption{Cumulative mass in substructures found by SUBFIND, bound HSF method, and HSF with different choices of the connectivity parameter $\\alpha$.} \\label{mass_profile} \\end{figure}  We note that in our test halo, $41.2\\%$ of the mass belongs to substructures for $\\alpha=0.2$, $50.4\\% $ for $\\alpha=0.01$, and $55.2\\%$ for $\\alpha=0.001$. When we substract from these numbers the contribution of bound structures, we find that $22.9\\%-36.4\\%$ of halo mass is stored in unbound structures.  This should be taken into account when analytical models of halos with substructures are proposed.  \\item While we would need a larger statistical sample of halos to perform robust measurements, we noticed that at fixed redshift of merging with the main halo, small structures tend to lose less mass than larger ones, in agreement with expectations based on the higher concentration of smaller halos. Also, we found a strong correlation between mass loss and the number of orbits a substructure can make inside the main halo. \\end{itemize}   When we observe our own Galaxy, we do not have access to different ``snapshots'' anymore, in stark difference with the world of simulations. Instead, we have to be content with the data at the present time. However, because we now know that our phase-space structure finder can identify dynamical structures that were bound before tidal disruption, it can provide totally new insights about the past dynamical history of our Galaxy. Within the hierarchical framework, we expect that our Galaxy should be made through the merging of more than about 100 smaller subcomponents. Comparing structures in observational data and simulations can be one of the best tests for the theory of hierarchical galaxy formation, and provide important constraints on cosmological models such as $\\Lambda$CDM.  Up to now, we studied only the evolution of dark matter, but we can also similarly study the evolution of baryons in gas and stars. Galaxies are observed in many different ways, ranging from star distributions, velocity and chemical properties, to HI measurements, etc. Our phase-space structure finder with local metric fitting is in fact implemented in such a way that it can be used in any number of dimensions, where each dimension can have completely different physical units. So it is in principle straightforward to use it for studying galaxy structure evolution in multi-dimensional space with the appropriate probabilistic weightings to take into account the noise and holes (missing measurements) in the data hypercube. We think that such an approach can yield a deeper understanding of galaxy evolution, and looks especially promising in light of the upcoming GAIA mission \\citep{Gilmore1998} which plans to map the positions of around one billion stars in our Galaxy."
    },
    "0812/0812.2244_arXiv.txt": {
        "abstract": "We present a case that current observations may already indicate new gravitational physics on cosmological scales. The excess of power seen in the Lyman-$\\alpha$ forest and small-scale CMB experiments, the anomalously large bulk flows seen both in peculiar velocity surveys and in kinetic SZ, and the higher ISW cross-correlation all indicate that structure may be more evolved than expected from $\\Lambda$CDM. We argue that these observations find a natural explanation in models with infinite-volume (or, at least, cosmological-size) extra dimensions, where the graviton is a resonance with a tiny width. The longitudinal mode of the graviton mediates an extra scalar force which speeds up structure formation at late times, thereby accounting for the above anomalies. The required graviton Compton wavelength is relatively small compared to the present Hubble radius, of order 300-600 Mpc. Moreover, with certain assumptions about the behavior of the longitudinal mode on super-Hubble scales, our modified gravity framework can also alleviate the tension with the low quadrupole and the peculiar vanishing of the CMB correlation function on large angular scales, seen both in COBE and WMAP. This relies on a novel mechanism that cancels a late-time ISW contribution against the primordial Sachs-Wolfe amplitude. ",
        "introduction": "\\label{intro}  In the ten years since the discovery of cosmic acceleration by the supernovae teams~\\cite{sn1a}, the $\\Lambda$-cold dark matter ($\\Lambda$CDM) model has emerged as the standard paradigm for cosmology. With a single parameter ($\\Lambda$), the model predicts expansion and growth histories consistent with observations. As near-future experiments will soon probe the large scale structure with unprecedented accuracy, it is worth asking at this juncture whether ten years from now the $\\Lambda$CDM model will be hailed as the ultimate cosmological theory, or whether it will have by then joined its Einstein de Sitter predecessor in the ranks of defunct cosmologies. A host of recent observations may have already revealed cracks in $\\Lambda$CDM's armor:  \\begin{enumerate}  \\item The cross-correlation between galaxy surveys and the cosmic microwave background (CMB), resulting from the Integrated Sachs-Wolfe (ISW) effect, has recently been measured at the $4\\sigma$ level using the Wilkinson Microwave Anisotropy Probe (WMAP) data and a combination of large-scale structure surveys~\\cite{seljak,Giannantonio:2008zi}. This provides non-trivial and independent evidence for cosmic acceleration. Remarkably, the combined amplitude is larger than the $\\Lambda$CDM prediction by $\\gsim 2\\sigma$'s: $2.23\\pm 0.60$~\\cite{seljak}, with 1 being the $\\Lambda$CDM expectation.  \\item Independent recent analyses suggest anomalously large bulk flows on large scales. Using a compilation of peculiar velocity surveys, Watkins {\\it et al.}~\\cite{hudson} obtained a bulk flow of $407\\pm 81$~km~s$^{-1}$ on 50$h^{-1}$~Mpc scales, in contrast with the $\\Lambda$CDM expectation of $\\sim 190$~km~s$^{-1}$ for the r.m.s. value of the bulk flow, which are inconsistent at the $2\\sigma$ level. Meanwhile, using the kinetic Sunyaev-Zeldovich effect to estimate peculiar velocities of galaxy clusters, Kashlinsky {\\it et al.}~\\cite{kash} found a coherent bulk motion of 600-1000~km~s$^{-1}$ out to $\\gsim 300h^{-1}$~Mpc.  \\item High resolution observations of the CMB at $\\ell \\sim 3000$ by the Cosmic Background Imager (CBI) indicate an excess power compared to the theoretical primary CMB power spectrum \\cite{Readhead:2004gy}. This excess could be explained by the thermal Sunyaev-Zel'dovich effect from galaxy clusters at $z\\sim 1$. However, the amplitude of matter power spectrum, necessary to explain CBI excess is about $35\\%$ (at $\\sim 2\\sigma$ level) larger than the current best-fit concordance $\\Lambda$CDM model \\cite{Reichardt:2008ay}.  \\item The Lyman-$\\alpha$ forest in the spectra of high redshift quasars provides us with the highest resolution precision measurement of the cosmological matter power spectrum. However, Lyman-$\\alpha$ forest constraints on the linear matter power spectrum, based on the SDSS quasar spectra at $z \\sim 3$ \\cite{McDonald:2004eu,McDonald:2004xn} is also $\\sim 35\\%$ (at $\\sim 2.5\\sigma$ level) larger than the $\\Lambda$CDM prediction (Fig. 1 in \\cite{Seljak:2006bg}).  \\item As first observed by the Cosmic Background Explorer (COBE)~\\cite{cobe} and later confirmed by WMAP~\\cite{wmap1}, the CMB temperature anisotropy shows a lack of correlation on large ($\\gsim 60\\,^{\\circ}$) angular separations. While related to the low quadrupole, the real-space anomaly is more robust statistically --- less than 0.03\\% of random realizations exhibit smaller power on those scales~\\cite{wmap1,huterer,others}. The statistical significance of this anomaly has only grown with subsequent WMAP data releases~\\cite{huterer2}.  \\end{enumerate}  Of course, these anomalies may eventually go away, as systematic effects are better understood. After all, these represent for the most part $\\sim 2\\sigma$ discrepancies. But it is striking that the first four observations on this list of misfits all point towards the same physics: that structure is more evolved on large scales than predicted by $\\Lambda$CDM. This suggests that gravity may be stronger at late times and on large scales.  In this paper we show that this is precisely what happens in infrared-modified gravity theories, inspired by brane-world constructions with infinite-volume extra dimensions.  Because the extra dimensions are infinite in extent (at least cosmologically large), the 4d graviton is no longer massless but is instead broadened into a resonance --- a continuum of massive states --- with a tiny width $r_c^{-1}$. At first sight this may seem at odds with our earlier conclusion, since a massive/resonance graviton implies weaker gravity. While this is certainly true on the largest distances, on intermediate (albeit cosmologically-relevant) scales the longitudinal (or helicity-0) mode mediates an extra scalar force which enhances gravitational attraction by order unity.  This extra force is of course at the origin of the well-known van Dam-Veltman-Zakharov (vDVZ) discontinuity~\\cite{vDVZ}. As conjectured by Vainshtein~\\cite{vainshtein}, however, non-linear interactions can suppress the effects of the longitudinal mode near astrophysical sources. This screening is also at play cosmologically, suppressing the extra force at early times and on small scales. Thus gravity becomes stronger only at late times and on sufficiently large scales. This suppression at high density or curvature is qualitatively similar to the chameleon mechanism \\cite{cham1,cham2,cham3}, except that the Vainshtein effect arises from derivative interactions as opposed to a scalar potential.  The theories of interest are higher-dimensional generalizations of the Dvali-Gabadadze-Porrati (DGP) model~\\cite{DGP}, in which our visible universe is confined to a 3-brane. The confrontation of its cosmological predictions with observations has been the subject of much literature since its advent~\\cite{lue,mustafa,koyama,hu,hu2,wiley0,amin,song,wiley}. Instead of having one extra dimension, here the bulk consists of 6 or more space-time dimensions.  Extending the DGP scenario to higher dimensions has proven difficult historically. Early studies demonstrated that naive extensions are plagued by ghost-like instabilities, even around flat space~\\cite{sergei,gigashif}. Recently, however, it was shown that such instabilities are absent if our 3-brane lies within a succession of higher-dimensional branes, each with their own induced gravity term, and embedded in one another in a flat bulk space-time~\\cite{oriol,us}. We refer to this framework as Cascading Gravity~\\cite{claudiareview}. In the simplest codimension-2 case, for instance, our 3-brane is embedded in a 4-brane within a 6-dimensional bulk. A similar cascading behavior of the gravitational force law was also obtained in a different codimension-2 framework~\\cite{nemanjacharting}.   Due to the higher-dimensional nature of these constructions, extracting cosmological predictions presents a daunting technical challenge. Even deriving the modified Friedmann equation is highly non-trivial. Here we circumvent these difficulties by developing a phenomenological description for the background cosmology, while using the parametrized post-Friedmann (PPF)~\\cite{ppf} framework to encode modifications to Einstein gravity into the evolution of perturbations. This approach has been shown to accurately reproduce the cosmological behavior of the standard DGP model. Inspired by these results, we can infer the parametric dependence of the required PPF input functions and parameters for Cascading Gravity models. This allows us to calculate various cosmological observables, all within a 3+1-dimensional framework.  One of our key observations is that in 6 or more dimensions the background evolution becomes nearly indistinguishable from $\\Lambda$CDM. This allows us to consider relatively small graviton Compton wavelengths, {\\it e.g.} $r_c\\sim  300-600$~Mpc, without violating constraints on the expansion history. Compared to the usual assumption $r_c\\sim H_0^{-1}$ in DGP, this expands considerably the duration and range of scales over which enhanced gravity is effective.  As mentioned earlier, the culprit for most of the interesting phenomenology in these models is the longitudinal or helicity-0 mode of the massive graviton. This degree of freedom is ultimately responsible for addressing all of the aforementioned anomalies:  \\begin{itemize}  \\item The longitudinal degree of freedom modifies the time-evolution of the lensing potential, which results in a stronger ISW cross-correlation compared to $\\Lambda$CDM.  \\item In the non-relativistic limit, the helicity-0 mode can be thought of as mediating an extra scalar force, which enhances gravitational attraction on intermediate scales. This speed-up in structure formation at late times results in larger bulk flows, more clusters at $z\\sim 1$, and enhanced power on Lyman-$\\alpha$ scales.  \\item As a consequence of this mode, our modified perturbation equations exhibit an effective anisotropic stress component. With certain assumptions about the form of this component on super-Hubble scales, we find that the late-time ISW effect can destructively interfere with the primordial Sachs-Wolfe amplitude, resulting in a power deficit in the CMB on large scales.  \\end{itemize}  Looking forward, our modified gravity models make several key predictions that will be tested by near-future experiments. For instance, because photons do not couple directly to the longitudinal mode, the weak lensing potential is considerably less affected than the Newtonian potential. Thus a distinguishing prediction is a discrepancy in the value of $\\sigma_8$ estimated from the matter power spectrum and that inferred from either CMB or weak lensing observations. For $r_c\\sim  300-600$~Mpc, the difference is a 20-30\\% effect, depending on the redshift of the observation.   The paper is organized as follows. In Sec.~\\ref{motiv}, we describe the building blocks of higher-dimensional generalizations to DGP, following~\\cite{giaalpha,degrav}. The result is a two-paramater family of massive or resonance gravity theories, with $\\alpha$ specifying the form factor, and $r_c$ the Compton wavelength for the graviton. Motivated by the DGP Friedmann equation, we propose in Sec.~\\ref{backcosmo} a generalized Friedmann equation for higher-dimensional theories. We then turn to cosmological perturbations in Sec.~\\ref{cosmopert}, and motivate our fiducial input parameters for PPF. Using these, we derive the resulting CMB power spectrum and angular correlation function in Sec.~\\ref{cancel}. In Sec.~\\ref{strucform}, we study the predictions for structure formation, including the weak lensing power spectrum (Sec.~\\ref{WL}), the Integrated Sachs-Wolfe cross-correlation (Sec.~\\ref{ISW}), large scale bulk flows (Sec.~\\ref{BF}), the Lyman-$\\alpha$ forest (Sec.~\\ref{lalpha}), and secondary CMB anisotropies from thermal Sunayev-Zeldovich (SZ) effect (Sec.~\\ref{CBISZ}). We conclude with a summary of our results in Sec.~\\ref{conc}. ",
        "conclusions": "\\label{conc}  Arguably the most pressing question in cosmology is whether the observed cosmic acceleration is due to a breakdown of Einstein gravity or to conventional vacuum energy. Fortunately, the $\\Lambda$CDM model is highly predictive --- a single dark energy parameter $\\Lambda$ suffices to predict the entire expansion and growth histories. Near future probes of the large scale structure will subject this standard model to increasingly stringent tests of its predictions.  In this paper we have presented a case that current observations may already be pointing towards new gravitational physics on cosmological scales. The excess of power seen in the Lyman-$\\alpha$ forest and small-scale CMB experiments, the larger bulk flows seen both in peculiar velocity surveys and in kinetic SZ, and the higher ISW cross-correlation all present potential challenges for $\\Lambda$CDM. And although their individual statistical significance is only at the $\\sim 2\\sigma$ level, it is certainly intriguing that these observations suggest a common explanation: that structure is more evolved than expected from $\\Lambda$CDM. On the other hand, weak lensing measurements give a value for $\\sigma_8$ consistent with the CMB.  We have argued that these observations find a natural explanation in models where the graviton is a resonance with a tiny width $r_c\\sim 300-600$~Mpc. The longitudinal mode $\\pi$ of the graviton can be thought of as mediating an extra scalar force which, thanks to the Vainshtein effect, only kicks in at late times and on sufficiently large scales. This speeds up structure formation, thereby accounting for the above anomalies. Meanwhile, since photons do not couple directly to $\\pi$, weak lensing is much less affected, which explains its consistency with the CMB.  Our approach has been motivated by higher-dimensional generalizations of the DGP scenario, such as Cascading Gravity models. A crucial advantage of considering more extra dimensions is that the modifications to the Friedmann equation become indistinguishable from a cosmological constant. This makes it possible to consider much smaller values of the graviton Compton wavelength than usually allowed by constraints on the expansion history. And this in turn amplifies the impact of the longitudinal mode on structure formation.  The obvious danger with enhancing structure growth is that one risks boosting the ISW contribution to the CMB low-$\\ell$ power. In fact, observations show the opposite trend, starting with the well-known low quadrupole. A more statistically significant anomaly is the peculiar fact that the angular correlation function vanishes on $\\gsim 60\\,^{\\circ}$ scales, as seen both in COBE and WMAP.  We have shown that our modified gravity models can account for these large scale CMB anomalies. With certain assumptions about the super-horizon behavior of $\\pi$, we have presented a novel mechanism in which a late-time ISW contribution cancels against the primordial Sachs-Wolfe amplitude. This ISW contribution is different from the usual one due to $\\Lambda$, but arises because the dynamics of $\\pi$ undergo a transition when $H\\sim r_c^{-1}$. Our mechanism relies on conservation of $\\zeta$ and is analogous to the jump in the Newtonian potential incurred during the radiation- to matter-dominated era in standard cosmology.  Many issues remain to be addressed:  \\begin{itemize}  \\item To compare our predictions with galaxy surveys, such as SDSS, we need a handle on the halo bias in our model. As argued in~\\cite{Hui:2007zh}, the bias is generically expected to be scale dependent in modified gravity theories. This issue is currently under study using N-body simulations~\\cite{mark}.  \\item A more thorough comparison with data requires a full likelihood analysis, including the PPF parameters used here.  \\item The effects of modified gravity is suppressed by the Vainshtein effect at large densities. How does this affect the non-linear collapse process, and the halo mass function in our model?  \\item While the phenomenological approach adopted here constitutes an important and instructive first step, it is imperative to justify the choice of PPF input functions through full-fledged brane-world calculations, for instance in Cascading Gravity models. In particular, the suppression of the large-angle CMB power relies on certain assumptions about the behavior of perturbations on super-Hubble scales, as mentioned above, and it will be essential to check whether this behavior can be explicitly realized in higher-dimensional constructions. If not, are there other consistent theories of modified gravity that can reproduce the phenomenology presented here?  \\end{itemize}"
    },
    "0812/0812.2843_arXiv.txt": {
        "abstract": "In models of coupled dark energy, in which a dark energy scalar field couples to other matter components, it is natural to expect a coupling to the inflaton as well. We explore the consequences of such a coupling in the context of single field slow-roll inflation. Assuming an exponential potential for the quintessence field we show that the coupling to the inflaton causes the quintessence field to be attracted towards the minimum of the effective potential. If the coupling is large enough, the field is heavy and is located at the minimum. We show how this affects the expansion rate and the slow-roll of the inflaton field, and therefore the primordial perturbations generated during inflation. We further show that the coupling has an important impact on the processes of reheating and preheating. ",
        "introduction": "\\label{sec:intro} The nature and origins of dark energy, the energy component which is responsible for the observed accelerated expansion of the universe, remain a mystery. Although the observations can be accounted for by the cosmological constant, scalar fields \\cite {Wetterich:1987fm,Ratra:1987rm} and modified gravity theories \\cite{Carroll:2004de} have also been suggested (for reviews and references see e.g. \\cite{Copeland:2006wr, Martin:2008qp, Nojiri:2006ri}). One of the major aims of modern cosmology is to determine the properties of dark energy, many of the parameters of which (equation of state, matter couplings etc.) can be constrained by considering data from supernovae at high redshifts, observations of anisotropies in the cosmic microwave background radiation (CMB) and the large scale structures (LSS) in the universe. At even higher redshifts and in the early universe, constraints from varying constants and big bang nucleosynthesis (BBN) give further constraints on the dark energy evolution \\cite{Kneller:2002zh, Dent:2008vd}. Some models of dark energy can even be tested in the laboratory \\cite{Brax:2007ak,Brax:2007vm,Brax:2007hi,Stojkovic:2007dw,Greenwood:2008qp}. In this paper, we study the impact of couplings of a quintessence-like scalar field to the inflaton field, the scalar field responsible for an accelerated expansion in the very early universe. In this important epoch, the seeds for the structures we observe in the universe were created. It is usually assumed that dark energy is not important during inflation. In the case in which dark energy and the inflaton field are not coupled, the vacuum expectation value (VEV) of the dark energy field is driven by quantum fluctuations to large field values, but otherwise there are no consequences for the inflationary dynamics (\\cite{Malquarti:2002bh}; see also \\cite{Martin:2004ba}). We show that this is not necessarily the case if the inflaton field couples to the dark energy field. In models such as coupled quintessence (\\cite{Wetterich:1994bg,Amendola:1999er}; see \\cite{Martin:2008qp} for a recent review) or quintessence models with a growing matter component \\cite{Amendola:2007yx}, dark energy couples to at least one species, which is  thought to be the decay product of the inflaton field. Therefore,  it is natural in these types of models to consider a coupling between dark energy and the inflaton field  as well. We will show in this paper that for large enough couplings (to be specified below) one can expect modifications to the predictions of the spectral index, its running and the tensor-to-scalar ratio. We find also that  the details of the physics of reheating and preheating are affected by the presence of a coupling between dark energy and the inflaton. The paper is organized as follows: in Section \\ref{sec:2} we describe our model  and study the inflationary epoch and discuss the effect of the quintessence field on the primordial perturbations. In Section \\ref{sec:3} we discuss the consequences for reheating and preheating. We conclude in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have described a model of coupled quintessence in which the scalar field responsible for dark energy is coupled to the inflaton field. To be specific, we have used the exponential potential in eqn. (\\ref{eq:pot}) and assumed that the coupling between the fields is of the form $A(Q)=\\exp(\\beta Q/M_{Pl})$ with $\\beta$ positive, so that the effective potential has a minimum. The large effective mass of the $Q$-field drives it into this minimum for a large range of $\\lambda$ and $\\beta$ values, and in Section \\ref{sec:21} we used this to evaluate the slow-roll parameters. We found that the expressions for $\\eta$ and $\\epsilon$ are modified by factors of $A^{-2}(1+4\\beta/\\lambda)^{-1}$ and $A^{-2}(1+4\\beta/\\lambda)^{-2}$ respectively, meaning that the end of inflation is delayed in this model, so that the inflaton field has a smaller field amplitude during its oscillations.  We further discussed the modification to the cosmological perturbations. The equations for the spectral index and its running (in terms of the new slow-roll parameters) vary with the ratio $\\beta/\\lambda$. Normalising the power spectrum to COBE, we found that the mass of the inflaton field is required to be smaller than in the standard case, whilst the effect of the presence of a dark energy coupling is to increase the value of $n_s$, for a given $\\lambda$. The parameters $\\beta$ and $\\lambda$ can be constrained by observations if one assumes that the coupling between the inflaton and dark energy has the same equal magnitude as the coupling between dark matter (or neutrinos) and dark energy.  We found that the presence of the $Q$ field during reheating leads to an unnaturally large amount of Hubble damping in the equation for the inflaton oscillations during reheating. We showed that if reheating is to be successful, $\\lambda$ must be large. The dark energy field rolls faster down its potential, which in turn reduces the Hubble damping and allows the radiation produced during reheating to become the dominant component of the universe. We also considered how the mechanism of parametric resonance during preheating might work in this model. We found that parametric resonance could still produce large amounts of radiation as the smaller value of $m$ required to match the COBE normalisation decreases the lower bound on the amplitude of the inflaton oscillations required for resonance. However, if $\\lambda$ is small, the energy density of the inflaton quickly becomes less than that of the $Q$-field so a large value of $\\lambda$ is still required to reheat the universe in our model. The fact that a large $\\lambda$ is required for successful reheating has interesting consequences. For example, it is not compatible with the condition $\\lambda<0.95$ stated in \\cite{Bean:2008ac}. Such models would be only consistent if the quintessence field does not couple to the inflaton field. Alternatively, the dark energy field could couple to a subdominant species, such as neutrinos, as e.g. in \\cite{Amendola:2007yx}.  \\vspace{1cm}  \\noindent{\\bf Acknowledgements:} JW is supported by EPSRC, CvdB and LHMH are supported by STFC. CvdB thanks the Institut de Physique, Saclay and Christof Wetterich and the Institut f\\\"ur Theoretische Physik of the University of Heidelberg for hospitality while parts of this work have been completed."
    },
    "0812/0812.2302_arXiv.txt": {
        "abstract": "If a pulsar orbits a supermassive black hole, the timing of pulses that pass close to the hole will show a variety of strong field effects.  To compute the intensity and timing of pulses that have passed close to a nonrotating black hole we introduce here a simple formalism based on two ``universal functions,'' one for the bending of photon trajectories and the other for the photon travel time on these trajectories. We apply this simple formalism to the case of a pulsar in circular orbit that beams its pulses into the orbital plane. In addition to the ``primary'' pulses that reach the receiver by a more-or-less direct path, we find that there are secondary and higher order pulses. These are usually much dimmer than the primary pulses, but they can  be of comparable or even greater intensity if they are emitted when pulsar is on the side of the hole furthest from the receiver. We show  that there is a phase relationship of the primary and secondary pulses that is a probe of the strongly curved spacetime geometry. Analogs of these phenomena are expected in more general configurations, in which a pulsar in orbit around a hole emits pulses that are not confined to the orbital plane. ",
        "introduction": "\\label{sec:intro}  Pulsars are rotating neutron stars that emit beams of radiation that can be detected on Earth.  As the star rotates, this beam sweeps past the Earth, producing regular pulses.  The timing of these pulses is tied to the large inertial moment of a compact body, making it an extremely stable clock: pulsars have been found whose pulse arrival times fluctuate by less than 200 nano-seconds~\\citep{2008ApJ...679..675V}.  Pulsar timing is thus an excellent probe of delicate phenomena, including gravity, in the vicinity of the pulsar.  It has been used to detect the tug of planets orbiting a pulsar, to observe the gradual loss of orbital energy to gravitational waves in binary systems containing a pulsar, and has been proposed as a method of detecting cosmological gravitational waves \\citep{EW75,Sazhin78,Detweiler79}. In this paper we explore how pulsar timing might probe the gravitational environment of pulsars orbiting supermassive black holes.  Supermassive black holes are now thought to be ubiquitous in the Universe, residing in the cores of most large galaxies, including our own.  They have masses in the range of millions to billions of Solar masses: our own Galaxy's black hole has an estimated mass of $4\\times10^6M_\\odot$~\\citep{2005ApJ...620..744G,2005ApJ...628..246E,2008arXiv0808.2870G}. The best estimates of its mass come from observing the orbits of O-type stars in the Galactic nucleus.  These observations have also revealed that the Galactic nucleus is home to a significant population of young massive stars, contrary to earlier expectations \\citep{2006JPhCS..54..279L}.  Perhaps the best model for this population has it forming \\textit{in situ} from the dense molecular hydrogen disk surrounding the black hole, with an initial mass function that is strongly tilted towards massive stars \\citep{2005MNRAS.364L..23N,2007MNRAS.374..515L,2007ApJ...669.1024M}. The Galactic nucleus is thus a likely environment for neutron stars to form, some of which may be pulsars orbiting deep within the potential well of the supermassive black hole.  Although no such pulsars have been detected, they may be discovered by future searches for periodic signals with corrections for high accelerations (due to gravity) and/or dispersion (due to plasma in the Galactic core).  The systems of interest here will consist of neutron stars at distances down to within several million kilometers (a few Schwarzschild radii) of the central black hole.  Even at these close separations, it would take many years for the neutron star orbit to decay due to gravitational radiation emission, making it possible to conduct lengthy timing observations of a pulsar deep within the strong field of a black hole.  (By contrast, a pulsar orbiting within a thousand Schwarzschild radii of a $10M_\\odot$ black hole would decay within a year.)  In this paper we consider a pulsar to be in orbit around a supermassive black hole, and we ask what effects the strong field of the hole would have on pulsar observations.  Such effects would depend on the details of the hole/pulsar system, and there are many details: the pulsar orbital elements, the alignment of the pulsar spin axis and the orbital plane, the pulsar spin rate, the angle between the pulsar spin axis and the pulse emission direction, and the black hole spin. (The black hole mass can be treated as a scaling parameter.)  An important step in understanding the details of strong field effects on pulsar obsrvations is to understand them in simple cases. This is what we do here, in two stages. First, we consider pulses emitted from a pulsar in the neighborhood of a Schwarzschild (nonrotating) black hole, with no restrictions on the pulsar orbit, spin, or emission direction. We point out that in the case of a spherically symmetric hole it is not necessary to do extensive computing of null geodesics. Rather, it is necessary only to compute two functions of the emission direction, one representing the bending of the light path, and the second the time delay along that path. These two ``universal functions'' of emission direction are parameterized only by the distance of the pulsar from the hole at the emission event.  The second stage in our analysis is to apply these universal functions to a particularly simple astrophysical scenario, that of a pulsar emitting in the orbital plane as it travels in a circular orbit around a Schwarzschild hole. It turns out that even in this simplest possible case, there is a rich set of interesting, and potentially observable phenomena. Analogs of these phenomena could occur also in binary pulsar systems and in binary neutron star/black hole systems of comparable mass.    The paper is organized as follows. In Sec.~\\ref{sec:unversal} we define the two universal functions, and we present quantitative results for them. Next, in Sec.~\\ref{sec:TOAs}a we show how we can use these functions to find any timing effect in the Schwarzschild geometry. We then, in Sec.~\\ref{sec:TOAs}b, apply this method explicitly to pulses beamed in or near the orbital plane.  These results are then applied, in Sec.~\\ref{sec:examples}, to the simple case of a circular orbit and equatorial beaming. In particular, we show effects that would appear to a distant receiver of the pulse trains.  A summary and a consideration of future applications of this method are given in Sec.~\\ref{sec:conc}. To allow the main ideas of the paper to be as clear as possible, several sets of details have been relegated to appendices.  Appendix A gives the details of finding the direction of emission as a function of time for a pulsar in a relativistic circular orbit.  Appendix B gives the details of the computation of the universal functions.  The relativistic calculations presented use the notational conventions of the text by Misner et al. \\citep{MTW}; in particular we use $c=G=1$. ",
        "conclusions": "\\label{sec:conc}    We have introduced the problem of black hole effects on pulses from an orbiting pulsar. We have shown that in the case of a spherically symmetric (i.e., nonrotating) hole, the use of two ``universal functions'' removes the need for extensive computation of null geodesics. The universal function approach to computation has been used for a first exploration of the phenomena that might be encountered with pulses from a pulsar-hole system.  This first exploration investigates pulse emission in the orbital plane, for a pulsar in circular orbit around a nonrotating hole.  Even for this highly simplified configuration we have found a number of interesting phenomena: (i)~``Primary'' pulses (pulses that travel from the pulsar to the receiver with relatively little gravitational bending) are continually accompanied by higher order (secondary, tertiary,...) pulses emitted during earlier pulsar orbits. (ii)~As would be expected, the primary pulses do not have a constant period, but rather have a period that is modulated by orbital motion. (iii)~The period of primary and of higher order pulses is not precisely the same, thus the higher order pulses arrive with a phase shift, relative to the primary pulses, that varies from one pulse to the next. (iv)~Strong field effects dominate the pulsar observations when the pulsar is on the far side of the hole, the side opposite that of the receiver. In this case the emitted pulse can reach the receiver only by a highly bent path. (v)~During the epoch of emission from the far side, the sequence of primary pulses and the sequence of secondary pulses, exchange roles. One consequence is that there is no clear distinction of primary and secondary for pulses emitted when the pulsar is close to the middle of its passage through the far side of the  hole. (vi)~For emission during most of the pulsar orbit, the secondary pulses have much lower intensity than the primary pulses. As the pulsar moves toward the dark side, however, the secondary pulses increase in intensity and the primary pulses decrease.  (vii)~Among the pulses emitted from the dark side are pulses that are {\\em amplified} by strong field effects analogous to gravitational lensing.  The simple configuration studied would be expected to exaggerate strong field effects when compared to a more realistic configuration in which pulse emission is directed well out of the orbital plane. For example, if the pulsar spin axis is perpendicular to the orbital plane (as in our simple configuration), and emission is not close to the orbital plane (contrary to our simple configuration), the pulse trajectory will never cross the orbital plane, and hence never come sufficiently near the hole for strong field effects to be significant. More interesting is the case in which the spin axis and beaming details are such that the beam does cross (or pass close to) the orbital plane.  Particularly noteworthy would be a configuration in which the receiver receives no primary pulses, and receives pulses only with the aid of strong field bending. It is also of note that some of the phenomena we have described, especially the existence of higher order pulses, can also occur in principle in binary pulsar systems, and pulsar-hole binaries of comparable mass.  Work is underway on investigating configuratins with out of orbit beaming.  An exploration of the large parameter space will be feasible with the efficiency provided by the universal function approach, so we are, at least at first, restriciting attention to nonrotating supermassive holes."
    },
    "0812/0812.2905_arXiv.txt": {
        "abstract": "Recently Eisenstein and collaborators introduced a method to `reconstruct' the linear power spectrum from a non-linearly evolved galaxy distribution in order to improve precision in measurements of baryon acoustic oscillations. We reformulate this method within the Lagrangian picture of structure formation, to better understand what such a method does, and what the resulting power spectra are. We show that reconstruction does {\\it not\\/} reproduce the linear density field, at second order. We however show that it does reduce the damping of the oscillations due to non-linear structure formation, explaining the improvements seen in simulations. Our results suggest that the reconstructed power spectrum is potentially better modeled as the sum of three different power spectra, each dominating over different wavelength ranges and with different non-linear damping terms. Finally, we also show that reconstruction reduces the mode-coupling term in the power spectrum, explaining why mis-calibrations of the acoustic scale are reduced when one considers the reconstructed power spectrum. ",
        "introduction": "The baryon acoustic oscillation (BAO) method \\cite{EisReview} is an integral part of current and next-generation dark energy experiments. Oscillations in the baryon-photon fluid, frozen into the matter distribution at decoupling, provide a standard ruler to constrain the expansion of the Universe. These sound waves imprint an almost harmonic series of peaks in the power spectrum $P(k)$, corresponding to a feature in the correlation function $\\xi(r)$ at $\\sim$100 Mpc, with width $\\sim 10$\\% due to Silk damping (see \\cite{Eis98,MeiWhiPea99} for a detailed description of the physics, and \\cite{ESW} for a comparison of Fourier and configuration space pictures). While the early Universe physics is linear and well understood, the low redshift observations are complicated by the non-linear evolution of matter (not to mention galaxy bias and redshift space distortions \\cite{SSS}, but we will defer these to future work) which erases the oscillations on small scales and shifts the peaks \\cite{ESW,CroSco,Mat08a,Seo08} \\begin{equation} P_{\\rm obs}(k) = e^{-k^2\\Sigma^2/2} P_{\\rm lin}(k) + P_{\\rm mc}(k)  + \\cdots \\label{eq:processed} \\end{equation} by coupling individual $k$-modes which are at early times independent. The exponential damping of the linear power spectrum (or equivalently the smoothing of the correlation function) reduces the contrast of the feature and thereby the precision with which the size of ruler may be measured. Neglect or incorrect modeling of the ``mode-coupling'' term $P_{\\rm mc}$ may bias the resulting distance measurements.  In \\cite{ESW} it was pointed out that much of the modification to the power spectrum comes from large-scale modes, bulk flows and super-cluster formation, in principle enabling their effects to be corrected. Eisenstein et al~\\cite{ES3} introduced a method for removing the non-linear degradation of the acoustic signature, sharpening the feature in configuration space or restoring/correcting the higher $k$ oscillations in Fourier space. Given the ambitious nature of future experiments, there has been considerable interest \\cite{ES3,Huff,Wagner} in ``reconstruction'' schemes which remove the effects of non-linearities, reducing the damping and mode coupling terms above.  Since the method proposed in \\cite{ES3} is an inherently non-linear mapping of the observed density field, it is difficult to intuitively understand. It is however easily formulated within the Lagrangian picture of structure formation, where the fundamental quantity is the displacement of particles from their initial positions (contrasted with the Eulerian picture where one tracks the evolution of the density field at a fixed location).  Motivated by recent developments in Lagrangian perturbation theory (LPT) \\cite{Mat08a,Mat08b}, we discuss reconstruction within the context of LPT, both to elucidate how it works and to expose possible shortcomings. Although we use the method of \\cite{ES3} for specificity, the lessons learned have broader validity.  We proceed as follows : \\S\\ref{sec:recon} introduces the essential aspects of both LPT as well as reconstruction. We then compute the reconstructed density field to second order, and demonstrate that there are corrections to the linear density at this order. \\S\\ref{sec:power} then explains why the BAO feature is enhanced in the reconstructed power spectrum. We conclude in \\S\\ref{sec:discuss}, highlighting potential avenues for improvements. ",
        "conclusions": "It is now generally understood (\\cite{ESW, CroSco, Mat08a} and this work) that the dominant effect of the non-linear evolution of matter perturbations on the baryon oscillations is to damp the higher harmonics, $P_{\\rm obs}(k) = \\exp(-k^2 \\Sigma^2/2) P_{\\rm lin}(k) + \\cdots$, or equivalently, smooth the feature in the correlation function. Eisenstein et al~\\cite{ES3} proposed a ``reconstruction'' method, demonstrated on simulations, that undoes this non-linear smoothing and appears to restore the linear power spectrum. Motivated by recent progress in Lagrangian perturbation theory \\cite{Mat08a, Mat08b}, we revisit this algorithm in order to better understand why it works as well as its shortcomings. Our principal conclusions are \\begin{itemize} \\item[(i)] The field generated by the reconstruction process is {\\it not\\/} the linear density field at second order. Note that this is a general statement, independent of assumptions about the smoothing of the initial density field. \\item[(ii)] Reconstruction does reduce the damping of the oscillations, by about a factor of 2 when the input density field is smoothed on the non-linear scale. \\item[(iii)] Reconstruction also reduces the mode coupling terms which introduce an out of phase component of the oscillations or shift the peak. \\item[(iv)] The reconstructed power spectrum is the sum of three power spectra (the auto-power spectra of the displaced and shifted fields, and their cross-spectrum), each of which have different damping terms (Eq.~\\ref{eq:damptransform}). An appropriate model for the reconstructed power spectrum should take this into account, instead of modeling it as a single damping scale. \\item[(v)] When the smoothing scale is close to the non-linear scale, the correlation between the shifted and displaced fields plays a crucial role. \\end{itemize}  Our results suggest a number of natural extensions. The effects of bias and redshift space distortions have been incorporated into the Lagrangian formalism \\cite{Mat08a, Mat08b}, and could therefore be folded in to the LPT formulation of reconstruction. We have observed that the reconstructed density field is not the linear density field; an interesting possibility is to explore whether higher order reconstruction schemes actually yield dividends. Even within the context of the existing reconstruction schemes, it is possible that a different weighting of the three power spectra may yield improved accuracy in measuring the distance scale. We leave these avenues open for future investigation."
    },
    "0812/0812.4418_arXiv.txt": {
        "abstract": "}[2]{{\\vspace{1mm}\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {\\bf Key words:~~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}  \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\AN}{Astron.~Nachr.} \\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 }} {High-frequency twin peak quasiperiodic oscillations (QPOs) are observed in four microquasars, \\ie Ga\\-lactic black hole binary systems, with frequency ratio very close to 3\\,:\\,2. In the microquasar GRS 1915+105, the structure of QPOs exhibits additional frequencies, and more than two frequencies are observed in the Galaxy nuclei Sgr\\,A$^*$, or in some extragalactic sources (NGC 4051, MCG-6-30-15 and NGC 5408 X-1). The observed QPOs can be explained by a variety of the orbital resonance model versions assuming resonance of oscillations with the Keplerian frequency $\\nu_{\\mathrm{K}}$ or the vertical epicyclic frequency $\\nu_{\\theta}$, and the radial epicyclic frequency $\\nu_{\\mathrm{r}}$, or some combinations of these frequencies. Generally, different resonances could arise at different radii of an accretion disc. However, we have shown that for special values of dimensionless black hole spin $a$ strong resonant phenomena could occur when different resonances can be excited at the same radius, as cooperative phenomena between the resonances may work in such situations. The special values of $a$ are determined for triple frequency ratio sets $\\nu_{\\mathrm{K}} : \\nu_{\\theta} : \\nu_{\\mathrm{r}} = s:t:u$ with $s,t,u$ being small integers. The most promising example of such a special situation arises for black holes with extraordinary resonant spin $a=0.983$ at the radius $r = 2.395\\,M$, where $\\nu_{\\mathrm{K}}$\\,:\\,$\\nu_{\\theta}$\\,:\\,$\\nu_{\\mathrm{r}} = $ 3\\,:\\,2\\,:\\,1. We also predict that when combinations of the orbital frequencies are allowed, QPOs with four frequency ratio set 4\\,:\\,3\\,:\\,2\\,:\\,1 could be observed in the field of black holes with $a = 0.866,\\,0.882$ and $0.962$. Assuming the extraordinary resonant spin $a=0.983$ in Sgr\\,A$^*$, its QPOs with observed frequency ratio $\\approx$ 3\\,:\\,2\\,:\\,1 imply the black hole mass in the interval $4.3 \\times 10^6 ~\\mathrm{M}_{\\odot} < M < 5.4 \\times 10^6~\\mathrm{M}_{\\odot}$, in agreement with estimates given by other, independent, observations.}{Accretion, accretion disks -- Black hole physics -- X-rays: general} ",
        "introduction": "Quasiperiodic oscillations (\\emph{QPOs}) of X-ray brightness had been observed in many Galactic low-mass X-ray binaries containing neutron~stars (see, \\eg van der Klis 2000, 2006, Barret \\etal 2005, Belloni \\etal 2005, 2007) or black holes (see, \\eg McClintock and Remillard 2004, Remillard 2005, Remillard and McClintock 2006). Some of the QPOs are in the kHz range and often come in pairs $(\\nu_{\\rm upp}, \\nu_{\\rm down})$ of {\\it twin peaks} (often called \\emph{double peaks}) in the Fourier power spectra. Since the peaks of high frequencies are close to the orbital frequency of the marginally stable circular orbit representing the inner edge of Keplerian discs orbiting black holes (or neutron~stars), the strong gra\\-vi\\-ty effects have to be relevant in explaining high-frequency QPOs (Abramowicz \\etal 2004).  The twin peak QPOs were observed in four microquasars, namely GRO 1655-40, XTE 1550-564, H 1743-322, GRS 1915+105 (T{\\\"{o}}r{\\\"{o}}k \\etal 2005). In all of the four cases, the frequency ratio of the twin peaks is very close to 3\\,:\\,2 and the orbital resonance model assuming non-linear resonances between oscillations in Keplerian $\\nu_{\\mathrm{K}}$ and epicyclic frequencies $\\nu_{\\theta}$ (vertical), or $\\nu_{\\mathrm{r}}$ (radial), or their combinations, seems to be the potential explanation of the microquasar kHz QPOs. The resonances of oscillations with the epicyclic frequencies were firstly mentioned in Aliev and Galtsov (1981). In the context of QPOs observed in the neutron star and black hole systems, the orbital resonance model was introduced and developed by Klu{\\'z}niak and Abramowicz (2000), see Abramowicz \\etal (2004).  According to the resonance hypothesis, the two modes in resonance should have eigenfrequencies $\\nu_{\\rm r}$ (equal to the radial epicyclic frequency) and $\\nu_{\\rm v}$ (equal to the vertical epicyclic frequency $\\nu_{\\rm{\\theta}}$ or to the Keplerian frequency $\\nu_{\\rm K}$). While models based on the \\emph{parametric resonance} identify the two observed frequencies of the twin peaks ($\\nu_\\mathrm{upp}$, $\\nu_\\mathrm{down}$) directly with the eigenfrequencies of a resonance, models based on the \\emph{forced resonance} allow to observe combinational (beat) frequencies of the modes, or oscillations with the combinational frequencies entering the resonance (Landau and Lifshitz 1976). Both parametric and forced resonance models make clear and precise predictions about the values of observed frequencies in connection with spin and mass of the observed object, at least in the case of black holes (T{\\\"{o}}r{\\\"{o}}k \\etal 2005).  In some sources, more than two high-frequency peaks are observed. The microquasar GRS 1915+105 reveals high-frequency QPOs appearing at four frequencies with the lower and upper pairs in the ratio close to 3\\,:\\,2 (Remillard and McClintock 2006) and even a fifth frequency was reported (Belloni \\etal 2001). An additional sixth frequency was mentioned (Strohmayer 2001), although not confirmed. In Sgr\\,A$^*$, three frequencies were reported with ratio close to 3\\,:\\,2\\,:\\,1 (Aschenbach 2004, Aschenbach \\etal 2004, T{\\\"{o}}r{\\\"{o}}k 2005). In the galactic nuclei MCG-6-30-15 and NGC 4051, two pairs of QPOs were reported with the ratios close to 3\\,:\\,2 and 2\\,:\\,1, respectively (Lachowicz \\etal 2006). In the source NGC 5408 X-1 oscillations with three frequencies of ratio close to 6\\,:\\,4\\,:\\,3 were observed (Strohmayer \\etal 2007).  In the microquasar GRS1915+105, an almost extreme Kerr black hole with $a \\approx 1$ is expected (McClintock \\etal 2006), and all the five (six) frequencies of QPOs can be explained in the framework of the extended resonance model with the hump-induced oscillations, predicting the black hole spin $a=0.9998$ and its mass $M=14.8~\\mathrm{M}_{\\odot}$ (Stuchl{\\'{\\i}}k \\etal 2006, 2007d). In the extended resonance model, forced resonances of the epicyclic oscillations with an additional oscillation induced by the humpy orbital velocity profile (related to the physically privileged locally non-rotating frames) that occurs in Keplerian discs orbiting Kerr black holes with $a > 0.9953$ are assumed (Aschenbach 2004, 2006, Stuchl{\\'{\\i}}k \\etal 2004, 2005, 2007c). In the ``humpy'' extended resonance model, all the oscillations in resonance can occur at, and be related to the exclusively defined, ``humpy radius'' with extremal orbital velocity gradient within the humpy profile (Stuchl{\\'{\\i}}k \\etal 2007c). However, this model can be relevant only for near-extreme Kerr black holes with spin $a > 0.9953$ for Keplerian discs, and with even higher spin $a > 0.9998$ for marginally stable thick accretion discs (Stuchl{\\'{\\i}}k \\etal 2004, 2005).  In order to explain complex frequency structures observed in some black hole systems, we expect that more than one resonance occur in a Keplerian disc, and different versions of the orbital resonance model could be acting. Of course, one version of the orbital resonance model appearing at different resonant points, at different radii of the disc, is also possible (Stuchl{\\'{\\i}}k and T{\\\"{o}}r{\\\"{o}}k 2005). Note that the so called total precession resonance model, assuming resonance of oscillations with the Keplerian frequency $\\nu_{\\mathrm{K}}$ and the total precession frequency $\\nu_{\\mathrm{T}}=\\nu_{\\theta}-\\nu_{\\mathrm{r}}$, can explain quite well the multi-resonant phenomena observed in the neutron star atoll source 4U 1636-53 (Stuchl{\\'{\\i}}k \\etal 2007e) and in some other atoll sources (Bakala \\etal 2008, in preparation).  Generally, two resonances could occur at different radii, and for special values of the dimensionless spin $a$, the bottom, top, or mixed frequencies coincide, reducing two frequency pairs into a triple frequency set (Stuchl{\\'{\\i}}k \\etal 2007b). Physically more interesting seems to be the case, when a triple frequency set $\\nu_{\\mathrm{K}}$, $\\nu_{\\theta}$, $\\nu_{\\mathrm{r}}$ with rational ratio arises at a given radius. The resonant phenomena can be then expected stronger than in the case of internally independent resonances occurring at two different radii. Clearly, while sharing a given radius, the resonances could be causally related and could cooperate efficiently (Landau and Lifshitz 1976). That is the reason why we focus here attention on the possibility of causally related resonances appearing at a fixed radius. Such a situation is possible only in the field of Kerr black holes with special values of the dimensionless spin $a$, when the value of $a$ is related to the corresponding triple frequency ratio set, and concrete versions of the resonance model that are realized. Since the strength of the resonance and the resonant frequency width decrease rapidly with the order of the resonance $n+m$ (see Landau and Lifshitz 1976), in the following we restrict ourselves to the cases with $n$, $m \\leq 5$. ",
        "conclusions": "We present the possibilities for strong resonant phenomena arising in a shared radius in the field of black holes with spin appropriately tuned. We assume that the strong resonances could occur and influence each other causally, if $\\nu_{\\mathrm{K}}:\\nu_{\\theta}:\\nu_{\\mathrm{r}}$ is in ratio of small integers $s:t:u$ with $s \\leq 5$, when the order of the resonances could be low enough in order to enable strong resonance with relatively wide resonance frequency width (Landau and Lifshitz 1976). The strong resonances with lowest values of the frequency ratio occur for the extraordinary resonant spin $a_{3:2:1} = 0.983$.  Due to the allowed frequency detuning in the resonant phenomena (with maximal relative frequency detuning limited to few percent), the strong resonant phenomena could appear in black holes with spin values concentrated around the special resonant spin values. For a given relative frequency detuning $\\delta$, the spin scatter differs for different special resonant spin values, being dependent on the position of the resonant radius (and the gradient of the frequency profiles at this radius), see Table \\ref{rozp-tab}. In the most interesting case of the extraordinary resonant spin $a_{3:2:1}=0.983$, the spin scatter $\\approx \\pm 0.003$ for detuning $\\delta = 0.01$ gives probably a realistic interval of black hole spins where the strong resonant phenomena could be observable. Of course, the value of $\\delta$ (and the spin value scatter) depend on physical conditions and resonant phenomena in concrete sources and have to be discussed carefully for each individual source.  There is an indication that the QPOs data observed in the Galactic centre source Sgr\\,A$^*$ imply the black hole spin close to the extraordinary resonant value of $a \\approx 0.983$, when $\\nu_{\\mathrm{K}}:\\nu_{\\theta}:\\nu_{\\mathrm{r}} = $ 3\\,:\\,2\\,:\\,1 holds at the radius $x_{3:2:1} = 2.395$ and the black hole mass can be estimated to be in the interval of $(4.8 \\pm 0.5)\\times 10^6 ~\\mathrm{M}_{\\odot}$. Therefore, it is interesting to check concordance of the QPOs observation induced data with a variety of the other observations of the Sgr\\,A$^*$. It should be stressed that more precise measurement of the QPOs frequencies will enable more precise determination of the black hole mass. Moreover, we expect NGC 5408 X-1 to be another candidate for a black hole admitting strong resonant phenomena because of the observed frequency set (Strohmayer \\etal 2007).  It is important from the principal reason to note that if a black hole with appropriately tuned spin $a$ admits strong resonant phenomena with different resonances sharing the same radius of the accretion disc, we could expect observable QPOs at resonant frequencies even in situations when the internal disc conditions are not convenient for starting up resonant phenomena, as observed, \\eg in Sgr\\,A$^*$ source (Genzel \\etal 2003).  On the other hand, the four (or five) QPOs frequencies observed in the microquasar GRS 1915+105 can not be explained by the strong resonances model because of specific distribution of the observed frequency ratios (two different pairs of frequencies with 3\\,:\\,2 ratio). In this case, the extended resonant model including the so called humpy oscillations induced by the ``humpy'' velocity profile of the accretion disc, as measured in the privileged family of locally non-rotating frames, can explain all the observed frequencies, if the black hole parameters are fixed at the spin $a \\approx 0.9998$ and the mass $M \\approx 14.4~\\mathrm{M}_{\\odot}$ (Stuchl{\\'{\\i}}k \\etal 2007c), in agreement with the spectral fit of the spin $a \\approx 1$ (McClintock \\etal 2006).  The conditions for strong resonant phenomena could be realized only for black holes with high values of dimensionless spin ($a \\geq 0.75$). Therefore, the idea of strong resonant phenomena probably could not be extended to the neutron star systems, where we expect spin $a < 0.5$, at least if the Hartle--Thorne metric parameters corresponding to the spin $a$ and the quadrupole momentum $q$ are close to the quasi-Kerr values with $q\\sim j^2$. Note that for the neutron star systems probably another version of the multiresonant idea could be realized in close agreement with observational data, where the so called total precession resonance model with resonant frequencies $\\nu_{\\mathrm{K}}$ and $\\nu_{\\theta}-\\nu_{\\mathrm{r}}$ (related to the relativistic precession model, Stella and Vietri 1998) is involved and realized at different radii, as shown in Stuchl{\\'{\\i}}k \\etal (2007e).  In the field of Kerr black holes with $a \\approx 0.866$, the four frequency set ratio 4\\,:\\,3\\,:\\,2\\,:\\,1 could be observed, if the resonant conditions allow combinational frequencies to be observable. For the spin $a \\approx 0.962$ or $a \\approx 0.882$, even the five frequency set 5\\,:\\,4\\,:\\,3\\,:\\,2\\,:\\,1 could be observable in some special circumstances. Clearly, it is not necessary that all the resonances are realized simultaneously and that the full five frequency set is observed at the same time.  Undoubtly, it is worth to search for such special cases of frequency sets in observational data as these could lead to precise determination of the black hole spin enabling a deeper understanding of a variety of astrophysical phenomena involved in determining behaviour of accretion discs in strong gravitational fields. We should note that when  simple combinational frequencies not exceeding the Keplerian frequency are allowed, the lowest triple frequency ratio set 3\\,:\\,2\\,:\\,1 can be realized for black holes with the extraordinary resonant spin $a \\approx 0.983$, but also for the black holes with three other values of the spin $a = 0.866, 0.882, 0.962$, if the uppermost frequencies are not observed for some reasons.  Finally, it has to be stressed that observation of some characteristic frequency sets with low integer ratio could be a signature of a specific value of the black hole spin $a$ enabling strong resonances at a fixed radius. However, generally, there exist few values of the spin $a$ and the corresponding shared resonance radius allowed for a given frequency ratio set (Stuchl{\\'{\\i}}k \\etal 2007b). Therefore, detailed analysis of the resonance phenomena, including the resonance strength and resonant frequency width, has to be considered in a concrete candidate system, and it has to be further confronted with the spin estimates coming from spectral analysis of the black hole system as given, \\eg in McClintock \\etal (2006) and Middleton \\etal (2006) for GRS 1915+105, and Shafee \\etal (2006) for GRO J1655-40, the line profiles (Fabian and Miniutti 2005, Dov\\v{c}iak \\etal 2004, Zakharov 2003, Zakharov and Repin 2006), and the orbital periastron precession of some stars moving in the region of Sgr\\,A$^*$ (Kraniotis 2005, 2007), in order to establish the black hole spin. Very promising from this point of view seem to be studies of the energy dependencies of high-frequency QPO determining the QPO spectra at the QPO radii ({\\.Z}ycki \\etal 2007) as they could be credited by the strong resonance conditions giving precisely the black hole spin and the radius where the observed QPOs are generated.  \\Acknow{This work was supported by the Czech grant MSM 4781305903.}"
    },
    "0812/0812.3148_arXiv.txt": {
        "abstract": "Our heuristic understanding of the abundance of dark matter halos centers around the concept of a density threshold, or ``barrier'', for gravitational collapse.  If one adopts the {\\it ansatz} that regions of the linearly evolved density field smoothed on mass scale $M$ with an overdensity that exceeds the barrier will undergo gravitational collapse into halos of mass $M$, the corresponding abundance of such halos can be estimated simply as a fraction of the mass density satisfying the collapse criterion divided by the mass $M$. The key ingredient of this {\\it ansatz} is therefore the functional form of the collapse barrier as a function of mass $M$ or, equivalently, of the variance $\\sigma^2(M)$. Several such barriers based on the spherical, Zel'dovich, and ellipsoidal collapse models have been extensively discussed. Using large-scale cosmological simulations, we show that the relation between the linear overdensity and the mass variance for regions that collapse to form halos by the present epoch resembles expectations from dynamical models of ellipsoidal collapse. However, we also show that using such a collapse barrier with the excursion set {\\it ansatz} predicts a halo mass function inconsistent with that measured directly in cosmological simulations. This inconsistency demonstrates a failure of the excursion set {\\it ansatz} as a physical model for halo collapse. We discuss implications of our results for understanding the collapse epoch for halos as a function of mass, and avenues for improving consistency between analytical models for the collapse epoch and the results of cosmological simulations. ",
        "introduction": "\\label{section:introduction}  A central concept in the modern theory of galaxy formation is the connection between characteristics of the linear density field and the abundance and properties of virialized dark matter halos in the contemporary universe. The power spectrum of density perturbations seeded by inflation \\citep[e.g.,][]{guth1982a,bardeen1983a,starobinsky1983a} and the cosmological transfer function \\citep[e.g.,][]{peebles1982a,bardeen1986a,eisenstein1998a} determine the character of the subsequent nonlinear growth of structure through gravitational clustering \\citep{white1978a}. Growing perturbations in the initial density field serve as the sites of galaxy formation \\citep[e.g.,][]{peebles1965a,sachs1967a,white1978a}.  In the context of a cold dark matter cosmology, these processes give rise to the characteristic mass scale of observed galaxies \\citep{rees1977a,blumenthal1984a}.  Cosmological observations have both motivated and verified this picture, most recently with measurements of galaxy clustering \\citep[e.g.,][]{percival2007a}, the linear power spectrum of cosmological structures \\citep[e.g,][]{mcdonald2006a}, and high-precision measurements of the cosmological microwave background radiation \\citep[e.g.,][]{dunkley2008a}.   Methods for calculating the abundances of nonlinear, collapsed structures have been developed to link the growth of density perturbations with the observed number densities of galaxy- and cluster-scale objects.  Dynamical models for the collapse of individual dense patches into virialized structures, such as the spherical collapse \\citep[][see Appendix \\ref{section:spherical_collapse}]{gunn1972a} and ellipsoidal collapse \\citep[][see Appendix \\ref{section:ellipsoidal_collapse}]{eisenstein1995a,bond1996a} models, provide physically-motivated methods for estimating the necessary, linearly-extrapolated overdensity (the ``collapse barrier'') for a region to break from the cosmic expansion, condense, and form a high-density, virialized structure (i.e., a dark matter halo).  When combined with the statistics of the initial density field, the collapse barrier can thereby be utilized to estimate the abundance of dark matter halos as a function of mass and redshift.  The purpose of this paper is to re-examine the connection between the collapse barrier and halo abundance, and test the common assumptions and methodologies used to calculate the mass function of dark matter halos from a dynamical model for their collapse~\\citep[e.g., the ``excursion set'' formalism,][]{bond1991a}.   \\cite{press1974a} first used the spherical collapse model to calculate the abundance of galaxies. They assumed the probability distribution function $\\dd P(\\deltar)/\\dd \\deltar$ (PDF) of the smoothed overdensity field $\\deltar = (\\rhom-\\rhobarm)/\\rhobarm$, where the mean matter density is $\\rhobarm$ and the density is averaged over a region of typical size $R$ containing mass $M \\propto \\rhobarm R^3$, was a Gaussian with a scale-dependent variance $\\sigma^{2}(M)$. They integrated this Gaussian PDF above the typical collapse overdensity $\\deltac$ (i.e., $P(\\deltam > \\deltac) \\propto \\erfc[-\\deltac/\\sqrt{2}\\sigma(M)]$) and differentiated with respect to mass $M$ (i.e., $\\dd P/\\dd M \\propto \\exp[-\\deltac^{2}/2\\sigma^{2}(M)]\\times\\dd \\sigma^{-1}/\\dd M$) to arrive at the fraction of all mass contained in objects of mass $M$. The \\citet{press1974a} calculation accounts for only half of the total universal mass density in bound objects because it does not address underdense regions contained within still larger regions for which the threshold $\\deltam > \\deltac$ is satisfied.  To remedy this shortcoming, \\citet{press1974a} multiplied their final answer by a factor of two to account for all mass with little justification.    \\citet{bond1991a} studied the properties of sets of regions above the threshold, the ``excursion sets'' of the density field. Using the excursion set formalism, \\cite{bond1991a} demonstrated that by filtering the initial overdensity field on a variety of mass scales, the results of \\cite{press1974a} could be derived in a manner that accounts for patches of low density embedded in large, high-density regions collapsing on larger scales (the ``cloud-in-cloud'' problem). Moreover, \\cite{bond1991a} and \\cite{lacey1993a} showed that the excursion set formalism provided a means to compute other halo properties such as their mass acquisition histories. The excursion set theory of halo abundance was later extended by \\cite{mo1996a} to describe their spatial clustering through the bias parameter $b^{2}\\equiv\\xi_{hh}/\\xi_{m}$ relating the halo ($\\xi_{hh}$) and mass ($\\xi_{m}$) correlation functions \\citep[a computation that recovers the ``peak-background split'' result developed in][]{kaiser1984a,efstathiou1988a,cole_kaiser1989}. The details of the excursion set theory are collected in the recent review by \\citet{zentner2007a}.   Numerical simulations of cosmological structure formation, which were developed concurrently with the analytical collapse calculations, demonstrated that nonlinear gravitational collapse produces halo mass functions that are inconsistent with the predictions of \\cite{press1974a}. Evidence for this disagreement, as well as corresponding discord in the spatial clustering of halos, developed over twenty years \\citep[e.g.,][]{efstathiou1988a,white1993a,lacey1994a,eke1996a,gross1998a,tormen1998a,jing1998a,lee1999a,jing1999a,porciani1999a,governato1999a}. The inability of the simple, spherical collapse model set within the excursion set formalism to describe the results of cosmological simulations and the need for robust predictions of the abundance of halos for comparisons with observations led to accurate formulae for mass functions determined by numerical fits to the results of N-body simulations \\citep[e.g.,][for more recent results see \\citealt{tinker2008a}]{sheth1999a,jenkins2001a,warren2006a}. These fitting formulae were not the results of specific dynamical models for the collapse of dark matter halos.  Rather, they were developed through a practical approach of trying to reproduce accurately the results of cosmological simulations.   Additional dynamical models for the growth of structure were developed in an attempt to better reproduce simulated mass functions while retaining a comparably simple, physically-motivated framework. \\cite{lee1998a} presented results for the halo mass function motivated by the \\cite{zeldovich1970a} pancake collapse model, which they extended to account for the collapse of halos along each of their principal axes. \\citet[][SMT01]{sheth2001a} argued that the functional form of the \\citet[][ST99]{sheth1999a} mass function can be motivated by the ellipsoidal collapse model of \\cite{bond1996a}. For a given ellipticity $e$ and prolaticity $p$ of the shear field about an overdensity $\\delta$, the \\cite{bond1996a} model provides a method for estimating the linearly-extrapolated overdensity at collapse (the ``ellipsoidal collapse barrier''). SMT01 built on a calculation by \\cite{doroshkevich1970a} to determine a probability distribution of shear ellipticities and prolaticities, which was then used to find the most-probable values for $e$ and $p$ as a function of $\\sigma(M)$. Combining these results, SMT01 found an effective ellipsoidal collapse barrier as a function of the ``peak height'', $\\nu_{c}=\\delta_{c}/\\sigma(M)$, alone.   With this new, mass-dependent collapse barrier SMT01 computed a halo mass function using the excursion set prescription of \\citet{bond1991a} and showed that the predicted functional form was close to that measured in cosmological simulations by ST99. The ST99 mass function has therefore become associated with the ellipsoidal collapse model.  Note, however, that the ellipsoidal collapse model predicts that the collapse barrier converges to the spherical collapse barrier $\\deltac=1.69$ (for $\\Omega_{m}=1$), in the high-mass limit, because the rarest peaks have a preferentially spherical shape. SMT01, on the other hand, found that the collapse barrier had to be lowered in the high-mass limit to $\\sqrt{a_{\\mathrm{ST}}} \\deltac \\sim 0.84 \\deltac$, in order to reproduce the mass function measured in cosmological simulations \\citep[see also][]{sheth2002a}. Although they argued that the lower value can be motivated by the mass definition of the Friends-of-Friends algorithm which they used to identify halos in simulations, this rescaling is not well justified.  Indeed, it disagrees with the fact that the collapse of the highest mass halos should be described well by the spherical collapse model. We argue below the lowering of the collapse barrier is instead required by the internal inconsistencies of the excursion set {\\it ansatz} and explicitly demonstrate that such inconsistencies persist for a wide variety of halo mass definitions, including the Friends-of-Friends and spherical overdensity criteria (see Appendix \\ref{section:mass_definitions} for a detailed discussion).   In the context of the excursion set formalism, the adoption of a shape for the collapse barrier (along with choice of prescription with which to smooth the density field) effectively determines the mass function.  In this paper we test the excursion set {\\it ansatz} by measuring the effective collapse barrier for halos formed in cosmological simulations and comparing excursion set predictions for the halo abundance given such a barrier to the halo mass functions measured in the simulations.   The organization of this paper is as follows. In \\S \\ref{section:theory}, we review the theoretical background of the excursion set formalism, common dynamical models, and collapse barriers from the literature. In \\S \\ref{section:simulations}, we measure the linear overdensity of collapsed regions in cosmological simulations to study the consistency between the simulated collapse of halos and the predictions of dynamical models. We then compare the abundance of dark matter halos predicted from the excursion set formalism and dynamical models with the simulated halo mass function in \\S \\ref{section:fcds}. We discuss our results in \\S \\ref{section:discussion} and summarize our results and conclusions in \\S \\ref{section:summary}. The paper contains four appendices: a review of the spherical collapse model (Appendix \\ref{section:spherical_collapse}), a review of the \\cite{bond1996a} ellipsoidal collapse model (Appendix \\ref{section:ellipsoidal_collapse}), a summary of the \\cite{zhang2006a} analytical method for calculating the excursion set mass function (Appendix \\ref{section:zh06}), and a study of the connection between the excursion set {\\it ansatz } and the halo mass definition (Appendix \\ref{section:mass_definitions}). Unless noted otherwise, throughout the paper we assume a flat $\\Lambda$-Cold Dark Matter ($\\LCDM$) cosmology.  Cosmological parameters in our calculations and simulations are close to the values suggested by observations: $\\Omega_{m}\\approx 0.3$, $\\Omega_{\\Lambda}\\approx0.7$, and $H_{0}\\approx70~\\km~\\s^{-1}~\\Mpc^{-1}$. ",
        "conclusions": "\\label{section:discussion}   In the excursion set theory, the abundance, formation time, and bias of dark matter halos are assumed to be directly linked to the initial linear overdensity field extrapolated to a given epoch using linear growth rates. This assumption significantly simplifies modeling, but clearly needs to be tested against direct numerical simulations. Our results show that the linear overdensities around the Lagrangian positions of the centers-of-mass of collapsed halos in simulations behave in a manner analogous to the collapse overdensity barrier predicted by SMT01 using an ellipsoidal collapse model. The similar behavior suggests that this model captures the main physics behind the nonlinear gravitational collapse around peaks in the initial density field.  At the same time, the failure of the excursion set {\\it ansatz} to predict correct abundance of collapsed halos with the ellipsoidal collapse barrier demonstrates that the {\\it ansatz} is flawed. Several aspects of the {\\it ansatz} may be responsible for its failure to accurately describe simulation results.  For example, assumption that each mass element can be assigned to a collapsed halo using only the local overdensity independently of its environment is definitely problematic. The threshold overdensity for collapse (the barrier) in either the spherical or ellipsoidal collapse models is predicted for the volumes centered on a density peak, and not for a random mass element in the field.  While the actual collapse will occur around the density peaks, additional mass near a peak may collapse onto it even though the extended region may not satisfy the local collapse condition. Our study highlights the failure of the excursion set {\\it ansatz} on a statistical basis for the entire halo population, but a physical model for the collapse and formation of dark matter halos must also succeed object-by-object \\citep[see, e.g.,][]{katz1993a}. The methods for calculating the abundance of halos that treat individual peaks in the density field \\citep[][Dalal et al., in prep.]{bond1996a,monaco2002a} may therefore afford a better way to predict the formation, masses, and abundances of collapsed objects.  The shape of the effective collapse barrier is determined both by the distribution of overdensities in a smoothed Gaussian field (e.g., with ellipticities and prolaticities, or other properties, sufficient for collapse) and by the dependence of the collapse condition on the overdensity (and/or other properties). The halo formation time can be defined relative to the mean assembly history \\citep[e.g.,][]{wechsler2002a}, but the specific choice of definition can influence the relation between formation time and halo mass \\citep[for a recent discussion, see][]{li2008a}. In terms of the linear overdensity and the effective barrier $B$, one definition of halo collapse is simply that of Eq.~(\\ref{equation:collapse_barrier}).  Of course, determining the effective barrier $B$ or measuring the collapse epoch $z_{c}$ complicates the matter. The definition of collapse can be connected to the physical properties of halos through dynamical models, but identifying evolutionary phases in such simple models with the actual nonlinear growth of dark matter halos may be incorrect or inaccurate. For instance, the ellipsoidal collapse model of \\cite{bond1996a} associates halo formation with the collapse of the longest ellipsoid axis, while freezing the collapse of the shorter axis at a particular point to prevent a density singularity. Halo virialization, however, may be associated with the collapse of another axis or conditional properties of shape of the density or shear fields \\citep[e.g.,][]{lee1998a,monaco2002a}. These issues warrant a more careful examination that we defer for future work now that we have a statement of the problems at hand.   Dark matter halos reside in special locations of the density field, and have different clustering properties than bulk matter \\citep[e.g.,][]{kaiser1984a,efstathiou1988a,cole1989a}. Models of halo bias are tightly connected with the effective collapse barrier. \\cite{mo1996a} presented the idea that the biasing of halos is connected to the rate at which density trajectories cross two separate barriers, and used this concept to calculate the Lagrangian bias of halos in the Press-Schechter model. SMT01 adapted these ideas to calculate the halo bias implied by their modified ellipsoidal collapse model.  In principle, the same collapse barrier should predict consistently both the bias and mass function of halos.  However, while the ST99 mass function models the halo abundance reasonably well, \\cite{seljak2004a} found that the SMT01 bias model does not work well at low masses.  For massive halos at high $\\nuc$, \\cite{cohn2008a} found that the \\cite{mo1996a} bias scaling works better than the SMT01 model \\citep[however, see][]{hu2003a}. Interestingly, \\cite{reed2008a} found that early-forming halos in their simulations were well-modeled by SMT01 at peak heights as large as $\\nuc\\sim4$ (at small $\\nuc$, the large-scale bias they measure drops below the SMT01 model, in a manner similar to that found by \\citealt{seljak2004a}).  We note that if the effective collapse barrier converges to the spherical collapse overdensity $\\deltac$ for large mass or highly-biased halos, then one might expect the halo bias to mimic the \\cite{mo1996a} scaling at large $\\nuc$.  We speculate that the intriguing deviations between the models for halo bias and abundance of halos are connected to the current discrepancy between the effective collapse barrier and the abundance of halos in the excursion set {\\it ansatz} that we discuss in this paper.  Lastly, the effective collapse barrier may be connected with the assembly bias phenomenon where dark matter halo clustering correlates with formation time.  \\cite{gao2005a} found that low-mass halos ($M<\\Mstar$) that formed early were much more strongly clustered than late-forming halos of the same mass \\citep[see also][]{sheth2004a,harker2006a}. \\cite{wechsler2006a} demonstrated that the sense of assembly bias reverses at high-mass ($M>\\Mstar$) such that late-forming halos were more strongly clustered than early-forming halos at fixed mass, and that these correlations are reflected in the relative bias of halos with different concentrations at fixed mass \\cite[see also][]{wetzel2007a,jing2007a}.   \\citet{zentner2007a} in his review showed that window functions that are local in real-space rather than Fourier space naturally result in early-forming halos that reside in underdense regions, as reported for high-mass halos ($M>\\Mstar$) in the simulations. The standard implementation of the excursion set formalism assumes that the process of nonlinear collapse can be encapsulated into the assignment of the collapse barrier.  For high-mass halos, tidal influences and nonlinear interactions with nearby objects should be minimal because halos with masses significantly larger than $\\Mstar$ form from nearly spherical peaks and usually dominate their local environments.  As a result, the excursion set assumption should be most valid for high-mass ($M \\gg \\Mstar$) halos and leads to a natural picture where early-forming, high-mass halos become less strongly clustered than their late-forming counterparts.   The reversal of the environment-dependent halo formation trend at low-mass may owe to the truncation of small halo growth by nearby structure as suggested by \\citet{wang2007a}. \\cite{dalal2008a} greatly extended this work, validating the high-mass trend in a set of scale-free numerical simulations and showing that environmental influences on halo bias, concentration, and formation time at fixed mass could be accounted for by considering the ``peak curvature'', $\\dd \\delta_{\\Rw}/\\dd \\sigma$, in addition to peak height.  The peak curvature serves as some proxy for environment as peaks with greater curvature lie in relatively underdense environments. \\cite{dalal2008a} also used a toy model to demonstrate that at small masses environmental effects, such as those suggested by \\cite{wang2007a}, can truncate halo growth and drive early-forming halos to become less anti-biased as they are advected by the larger-scale matter field.   Our results relate to these findings by providing a new outlook on the connection between halo abundance, the effective collapse barrier, and the excursion set formalism.  If the excursion set formalism fails to account properly for the abundance of halos using an appropriate form for the collapse barrier, then it may also fail to describe reliably the connection between formation time, mass, and properties of the density field.  The effective collapse barrier may well be a function of additional parameters beyond the local density \\citep[e.g.,][]{chiueh2001a,sandvik2007a}, may incorporate information about the larger-scale field \\citep[e.g.,][]{zentner2007a,desjacques2008a,dalal2008a}, and may require additional parameter dependencies that to account for the nonlinear collapse of overdensities \\citep[e.g.,][]{wang2007a,dalal2008a}.   At present, the excursion set theory provides the main framework used to develop heuristic understanding of simulation results and to formulate fits to simulation results for halo abundance, clustering, and other halo properties.  The method succeeds in a gross sense.  Excursion set theory identifies the fundamental scale in the problem, the mass where $\\sigma(M) \\sim \\deltac$. Below this characteristic mass the halo abundance per unit mass has a simple, power-law form, while  above this mass the halo abundance drops rapidly. However, a precise understanding of halo abundance and clustering beyond the gross accuracy provided by the excursion set {\\it ansatz} is now necessary. Contemporary and forthcoming efforts to use measurements of the abundance and clustering of galaxy clusters to constrain cosmological parameters, as well as comprehensive statistical studies of galaxy formation and evolution, only highlight the need for a sound understanding of halo abundance and assembly, both globally and as a function of environment. We need to explore amendments and alternatives to the excursion set model to make progress in our understanding of halo formation.     In this paper we have presented tests of the excursion set {\\it ansatz} against cosmological simulations.  Using a subset of cosmological simulations from the \\cite{tinker2008a} study of the halo mass function, we identify the locations in the linear overdensity field that later collapse to form dark matter halos. We demonstrate that the dependence of the linear overdensity of these regions on mass or smoothing scale $\\sigma(M)$ resembles predictions of the ellipsoidal collapse model. While the effective collapse barrier of simulated halos behaves analogously to the simple ellipsoidal collapse barrier, the simulated halo mass function is inconsistent with what the excursion set {\\it ansatz} predicts for such a barrier. This inconsistency implies that the excursion set {\\it ansatz} is not valid and cannot be used reliably to predict halo abundance or bias.   The modified collapse barrier of \\cite{sheth2001a} differs significantly from the physical behavior calculated by the ellipsoidal collapse model, which for example predicts convergence to the spherical barrier ($\\delta\\to\\deltac$) for the rarest peaks ($\\sigma\\to0$). In view of this, the interpretation of the \\cite{sheth1999a} mass function as a prediction of the ellipsoidal collapse model for the abundance of dark matter halos is not correct.  The impressive statistics of ever-larger dark matter simulations will likely continue to uncover increasingly subtle variations on the classical picture of dark matter halo formation.  In this work, we identify a striking inconsistency between the effective collapse barrier of simulated halos and excursion set formalism predictions for their abundance. Our results also demonstrate that there is still much to learn and understand about the conditions for and the process of halo collapse, which warrants further studies that critically revisit these issues using modern large cosmological simulations."
    },
    "0812/0812.2187_arXiv.txt": {
        "abstract": "The fate of massive cold clumps, their internal structure and collapse need to be characterised to understand the initial conditions for the formation of high-mass stars, stellar systems, and the origin of associations and clusters. We explore the onset of star formation in the 75\\,M$_\\sun$ SMM1 clump in the region ISOSS J18364-0221 using infrared and (sub-)millimetre observations including interferometry. This contracting clump has fragmented into two compact cores SMM1 North and South of 0.05\\,pc radius, having masses of 15 and 10\\,M$_\\sun$, and luminosities of 20\\,L$_\\sun$ and 180\\,L$_\\sun$. SMM1 South harbours a source traced at 24 and 70\\,$\\micron$, drives an energetic molecular outflow, and appears supersonically turbulent at the core centre. SMM1 North has no infrared counterparts and shows lower levels of turbulence, but also drives an outflow. Both outflows appear collimated and parsec-scale near-infrared features probably trace the outflow-powering jets. We derived mass outflow rates of at least $4\\times10^{-5}$\\,M$_\\sun$\\,yr$^{-1}$ and outflow timescales of less than 10$^4$\\,yr. Our HCN(1-0) modelling for SMM1 South yielded an infall velocity of 0.14\\,km\\,s$^{-1}$ and an estimated mass infall rate of $3\\times10^{-5}$\\,M$_\\sun$\\,yr$^{-1}$. Both cores may harbour seeds of intermediate- or high-mass stars. We compare the derived core properties with recent simulations of massive core collapse. They are consistent with the very early stages dominated by accretion luminosity. ",
        "introduction": "\\subsection{Star formation in massive cold clumps}  The early phases of star formation occur in the densest and coldest regions within molecular clouds, i.e.\\ dense cold cores. Detailed studies of nearby star-forming molecular clouds have led to a conception of the evolution of individual low-mass stars from prestellar cores to pre-main-sequence stars \\citep[e.g.][]{1994ApJ...420..837A,2007prpl.conf...17D,2007prpl.conf...33W}. The formation of high-mass stars (M $\\gtrsim 10$\\,M$_\\sun$) is less understood, in particular its earliest stages and the link to the origin of associations and clusters \\citep{2007ARA&A..45..481Z,2008ASPC..387.....B}. Previous studies, mostly aimed towards luminous infrared sources, have provided evidence for collimated outflows \\citep[recently e.g.][and ref.\\ therein]{2007A&A...470..269Z,2005ccsf.conf..105B} and large-scale rotating structures around forming high-mass objects \\citep[][and ref.\\ therein]{2005IAUS..227..135Z}. This indicates that they also grow via accretion through a disk and that the high radiation pressure could be released in outflow cavities \\citep{2005ApJ...618L..33K}. In the low-mass regime, accretion rates are found to be of the order 10$^{-6}$\\,M$_\\sun$\\,yr$^{-1}$ \\citep{1997ApJ...485..703W}. Much higher values of 10$^{-3}$\\,M$_\\sun$\\,yr$^{-1}$ are expected during the build-up of high-mass stars \\citep{2003ApJ...585..850M,2007ApJ...660..479B} and corresponding increased outflow energetics have been observed \\citep[e.g.\\ see compilation by][and ref.\\ therein]{2004A&A...426..503W}. Recently, simulations of protostars accreting at such high rates predicted a bloating by several tens of solar radii or more and the accompanying high luminosity at low surface temperature before they reach the main sequence \\citep{2008ASPC..387..189Y,2008ASPC..387..255H}. This gives further indication that the formation process differs from a scaled-up version of the low-mass case.  It is presumed that the onset of high-mass star formation occurs in cores (diameters of 0.1\\,pc and less) within massive cold clumps (diameters of around 0.5\\,pc). Surveys in the far-infrared and at millimetre wavelengths are well suited to search for such objects, but so far the earliest, i.e.\\ prestellar, stages have eluded discovery. \\citet{2007A&A...476.1243M} derive statistical lifetimes of less than 10$^3$\\,yr for the former, and of 10$^4$\\,yr for massive dense cores that are forming stars but exhibit low infrared emission. Interferometric observations provide the essential spatial resolution to study the substructure of clumps, and to infer the properties of embedded cores further requires a good coverage of their spectral energy distribution. In this work we present such a detailed study of one clump. In particular, we investigate the amount of fragments in the clump and their properties, the arising outflow activity, and the indications for collapse. The clump is located in the region \\object{ISOSS J18364-0221} that has been identified using the ISOPHOT Serendipity Survey (ISOSS, \\citet{1996A&A...315L..64L,2007A&A...466.1205S}) at 170\\,$\\micron$ to establish far-infrared colour temperatures. \\citet{2003PhDT.........3K} and \\citet{2004BaltA..13..407K} selected sources containing a large fraction of cold dust, and subsequent studies revealed a population of cold clumps \\citep{2003A&A...398.1007K,2008A&A...485..753H} and a collapsing massive core \\citep{2007A&A...474..883B}.   \\subsection{The ISOSS J18364-0221 star-forming region and the SMM1 clump}  The star-forming region \\object{ISOSS J18364-0221} (R.A.\\,$18^h36^m24.7^s$, Decl.\\,-02\\arcdeg21\\arcmin49\\arcsec [J2000]) was initially studied by \\citet{2006ApJ...637..380B} and is located at a distance of about 2.2\\,kpc. Derived from the near-infrared extinction over a $15\\arcmin\\times15\\arcmin$ field ($\\sim$\\,100\\,pc$^2$), the cloud complex associated with the star-forming region comprises a mass M $\\approx 3200$\\,M$_\\sun$. In the same manner, a lower mass limit M $\\ge 460$\\,M$_\\sun$ was found for the central region ($\\sim$\\,1\\,pc$^2$), while the far-infrared and submillimetre measurements give a luminosity L $\\approx 800$\\,L$_\\sun$, an average dust temperature of about 15\\,K and a mass of $900^{+450}_{-330}$\\,M$_\\sun$. This region contains two clumps detected in the submillimetre continuum named SMM1 and SMM2 that are separated by about $1.5\\arcmin$ along the east-west direction. The SMM1 clump is subject of this publication. Its effective radius is about 0.2\\,pc, and from the thermal emission a characteristic dust temperature of $16.5^{+6}_{-3}$\\,K and a clump mass of M$_{\\rm SMM1} = 75\\pm30$\\,M$_\\sun$ were found. The molecular line observations reported in \\citet{2006ApJ...637..380B} show red-shifted self-absorption that is interpreted as signature of collapse motions. Furthermore, at least one outflow is present. The outflow properties resulting from the single-dish observations, i.e.\\ the outflow mass of about 18\\,M$_\\sun$ and the mass outflow rate of 10$^{-3}$\\,M$_\\sun$\\,yr$^{-1}$, are comparable to values derived for outflows from presumed high-mass star precursors. These results render the SMM1 clump a promising object to study the early phases of collapse and fragmentation occurring in massive cold cores.  The detailed study of this source we present here made use of recently collected infrared and (sub-)millimetre data described in the next section. They reveal two star-forming cores and molecular outflows (Section~3). In Section~4, we discuss the core properties in the context of early-stage star formation and present a comparison of the observed HCN spectra with simple radiative transfer models. Finally, the results are summarised in Section~5. ",
        "conclusions": "We have investigated the onset of star formation in the massive cold clump \\object{ISOSS J18364-0221} SMM1. The derived results regarding the clump fragmentation, the core and outflow properties, and the collapse indications are summarises as follows: \\begin{enumerate} \\item Two compact, embedded cores are the only objects we identified. Thus, this clump shows little fragmentation, but seems to rather produce a low number of compact objects. \\item The cores dubbed SMM1 North and South are of about 10000 and 9000\\,AU radius. The dust temperatures are 15 and 22\\,K and they have masses of about 15 and 10\\,M$_\\sun$, respectively. SMM1 South harbours an infrared source not detected at 8, but at 24 and 70\\,$\\micron$. The luminosity of SMM1 South is about 180\\,L$_\\sun$ and a molecular outflow is detected in its vicinity. SMM1 North lacks infrared emission, but a second molecular outflow is emerging from this core. The luminosity of SMM1 North is approximately 20\\,L$_\\sun$. Both outflows appear collimated and we derived lower outflow mass limits of about 0.2 and 0.3\\,M$_\\sun$ of presumably entrained gas, and mass outflow rates of about $4\\times10^{-5}$\\,M$_\\sun$\\,yr$^{-1}$. In Fig.~\\ref{fig:schem} we show a schematic illustration of these results. \\item We detected filaments of shock-excited H$_2$ emission at 2.122\\,$\\micron$ towards the lobes of the SMM1 South molecular outflow, and also at distances up to 4.6\\,pc from SMM1. Each filament is roughly aligned with one of the outflow axes. They are possibly tracing the outflow-generating jets, and the dynamical timescales we derived are consistent with those of the molecular outflows considering the uncertain jet velocities. For both molecular outflows, ages of less than $10^4$\\,yr were found. \\item SMM1 South is bright in HCN(1-0) and our modelling results support that the core is collapsing. We got an infall velocity of 0.14\\,km\\,s$^{-1}$. This results in a mass infall rate estimate of $3\\times10^{-5}$\\,M$_\\sun$\\,yr$^{-1}$. For both cores the spectra can be interpreted with turbulent internal motions, and in the case of SMM1 South the innermost part appears highly supersonic. \\end{enumerate}  Both SMM1 North and South are more massive than typical low-mass dense cores and embedded in the presumably contracting 75\\,M$_\\sun$ SMM1 clump. This suggests that the two cores harbour protostellar seeds that may become intermediate- to high-mass stars. The presence of a 24\\,$\\micron$ source, a very turbulent central region and the jet features at large distances to the core support that SMM1 South is more evolved than SMM1 North. We interpret it as having an embedded young protostar that constitutes the driving source of the outflow. The outflows as well as the infall indicate that it is accreting. The associated outflow indicates that SMM1 North is also a star-forming core, and the outflow energetics are similar. However, the forming object remains undetected, and the relatively low level of turbulence implies that it has only marginally affected the surrounding core. The presumed difference in evolution of the forming objects contrast with the similar sizes, densities and outflow properties of the cores. These findings suggest a rapid evolution of the luminosity as predicted for high accretion rates \\citep{Yorke08}.  In comparison to the core collapse simulations of \\citet{2007ApJ...656..959K} we find that the luminosity of SMM1 South would be reached at a stage where the embedded protostellar mass is around 0.3\\,M$_\\sun$. Accretion onto the protostar is the dominant luminosity source, and they find an accretion rate on the order of 10$^{-4}$\\,M$_\\sun$\\,yr$^{-1}$ from directly after the formation of the protostar until this stage. The resulting timescale of about 3000\\,yr is consistent with the outflow timescales. Also the mass outflow rate as well as the infall rate estimate may be consistent with the high accretion rate. SMM1 North is less luminous, and hence the protostellar mass would be even lower, but may still be consistent with a high accretion rate in their model. However, the formation of outflows is not incorporated therein.  \\begin{figure} \\epsscale{1.15} \\plotone{f7.eps} \\caption{Schematic illustration of the identified cores, the corresponding outflows, and their properties. The 3.4\\,mm contours are plotted above the 24\\,$\\micron$ emission (cf.\\ Fig.~\\ref{fig:4map}). } \\label{fig:schem} \\end{figure}"
    },
    "0812/0812.2464_arXiv.txt": {
        "abstract": "Oe stars are a subset of the O-type stars that exhibit emission lines from a circumstellar disk. The recent detection of magnetic fields in some O-type stars suggests a possible explanation for the stability of disk-like structures around Oe stars. According to this hypothesis, the wind of the star is channeled by a dipolar magnetic field producing a disc in the magnetic equatorial plane. As a test of this model, we have obtained spectropolarimetric observations of the hottest Galactic Oe star HD 155806. Here we discuss the results and implications of those observations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.5066_arXiv.txt": {
        "abstract": "Rotating black holes in the brany universe of the Randall--Sundrum type with infinite additional dimension are described by the Kerr geometry with a tidal charge $b$ representing the interaction of the brany black hole and the bulk spacetime. For $b<0$ rotating black holes with dimensionless spin $a>1$ are allowed. We investigate the role of the tidal charge in the orbital resonance model of quasiperiodic oscillations (QPOs) in black hole systems. The orbital Keplerian frequency $\\nu_{\\mathrm{K}}$ and the radial and vertical epicyclic frequencies $\\nu_{\\mathrm{r}}$, $\\nu_{\\theta}$ of the equatorial, quasicircular geodetical motion are given. Their radial profiles related to Keplerian accretion discs are discussed, assuming the inner edge of the disc located at the innermost stable circular geodesic. For completeness, naked singularity spacetimes are considered too. The resonant conditions are given in three astrophysically relevant situations: for direct (parametric) resonances of the oscillations with the radial and vertical epicyclic frequencies, for the relativistic precession model, and for some trapped oscillations of the warped discs, with resonant combinational frequencies involving the Keplerian and radial epicyclic frequencies. It is shown, how the tidal charge could influence matching of the observational data indicating the $3\\!:\\!2$ frequency ratio observed in GRS 1915+105 microquasar with prediction of the orbital resonance model; limits on allowed range of the black hole parameters $a$ and $b$ are established. The ``magic'' dimensionless black hole spin enabling presence of strong resonant phenomena at the radius, where $\\nu_{\\mathrm{K}}\\!:\\!\\nu_{\\theta}\\!:\\!\\nu_{\\mathrm{r}} = 3\\!:\\!2\\!:\\!1$, is determined in dependence on the tidal charge. Such strong resonances could be relevant even in sources with highly scattered resonant frequencies, as those expected in Sgr\\,A$^*$. The specific values of the spin and tidal charge are given also for existence of specific radius where $\\nu_{\\mathrm{K}}\\!:\\!\\nu_{\\theta}\\!:\\!\\nu_{\\mathrm{r}} = s\\!:\\!t\\!:\\!u$ with $5\\geq s>t>u$ being small natural numbers. It is shown that for some ratios such situation is impossible in the field of black holes. We can conclude that analysing the microquasars high-frequency QPOs in the framework of orbital resonance models, we can put relevant limits on the tidal charge of brany Kerr black holes. ",
        "introduction": "In recent years, one of the most promising approaches to the higher-dimensional gravity theories seems to be the string theory and M-theory describing gravity as a truly higher-dimensional interaction becoming effectively 4D at low enough energies and these theories inspired braneworld models, where the observable universe is a 3-brane (domain wall) to which the standard-model (non-gravitational) matter fields are confined, while gravity field enters the extra spatial dimensions, the size of which may be much larger than the Planck length scale $l_\\mathrm{P}\\sim 10^{-33}~\\mathrm{cm}$ \\cite{Ark-Dim-Dva:1998:}. The braneworld models could therefore provide an elegant solution to the hierarchy problem of the electroweak and quantum gravity scales, as these scales become to be of the same order ($\\sim\\mathrm{TeV}$) due to large scale extra dimensions \\cite{Ark-Dim-Dva:1998:}. Therefore, future collider experiments can test the braneworld models quite well, including the hypothetical mini black hole production on the TeV-energy scales \\cite{Emp-Mas-Rat:2002:,Dim-Lan:2001:}. On the other hand, the braneworld models could be observationally tested since they influence astrophysically important properties of black holes.  Gravity can be localized near the brane at low energies even with a non-compact, infinite size extra dimension with the warped spacetime satisfying the 5D Einstein equations with negative cosmological constant as shown by Randall and Sundrum \\cite{Ran-Sun:1999:}. Then an arbitrary energy-momentum tensor could be allowed on the brane \\cite{Shi-Mae-Sas:2000:}.  The Randall--Sundrum model gives the 4D Einstein gravity in low energy limit, and the conventional potential of weak, Newtonian gravity appears on the 3-brane with high accuracy. Significant deviations from the Einstein gravity occur at very high energies, e.g., in the very early universe, and in vicinity of compact objects (see, e.g., \\cite{Dad-etal:2000:,Maa:2004:,Ger-Maa:2001:,Ali-Gum:2005:}). Gravitational collapse of matter trapped on the brane results in black holes mainly localized on the brane, but their horizon could be extended into the extra dimension. The high-energy effects produced by the gravitational collapse are disconnected from the outside space by the horizon, but they could have a signature on the brane, influencing properties of black holes \\cite{Maa:2004:}. There are high-energy effects of local character influencing pressure in collapsing matter, and also non-local corrections of ``back-reaction'' character arising from the influence of the Weyl curvature of the bulk space on the brane -- the matter on the brane induces Weyl curvature in the bulk which makes influence on the structures on the brane due to the bulk graviton stresses \\cite{Maa:2004:}. The combination of high-energy (local) and bulk stress (non-local) effects alters significantly the matching problem on the brane, as compared to the 4D Einstein gravity; for spherical objects, matching no longer leads to a Schwarzschild exterior in general \\cite{Dad-etal:2000:,Ger-Maa:2001:}. Moreover, the Weyl stresses induced by bulk gravitons imply that the matching conditions do not have unique solution on the brane; in fact, knowledge of the 5D Weyl tensor is needed as a minimum condition for uniqueness \\cite{Ger-Maa:2001:}.\\footnote{At present, no exact 5D solution in the braneworld model is known.}  There are two kinds of black hole solutions in the Randall--Sundrum braneworld model with infinite extension of the extra dimension. One kind of these solutions looks like black string from the viewpoint of an observer in the bulk, while being described by the Schwarzschild metric for matter trapped on the brane \\cite{Cha-Haw-Rea:2000:}. The generalizations to rotating black string solution \\cite{Mod-Pan-Sen:2002:} and solutions with dilatonic field \\cite{Noj-etal:2000:} were also found. However, the black string solutions have curvature singularities at infinite extension along the extra dimension \\cite{Ali-Gum:2004:}. There is a proposal that the black hole strings could evolve to localized black cigar solutions due to the classical instability near the anti-de~Sitter horizon \\cite{Cha-Haw-Rea:2000:,Gre-Laf:1993:}, but it is not resolved at present \\cite{Hor-Mae:2001:,Ali-Gum:2005:}.  Second kind of these solutions representing a promising way of generating exact localized solutions in the Randall--Sundrum brane\\-world models was initiated by Maartens and his coworkers \\cite{Maa:2004:,Ger-Maa:2001:,Dad-etal:2000:}. Assuming spherically symmetric metric induced on the 3-brane, the effective gravitational field equations on the brane could be solved, giving Reissner--Nordstr\\\"{o}m static black hole solutions endowed with a ``tidal'' charge parameter $b$ instead of the standard electric charge parameter $Q^2$ \\cite{Dad-etal:2000:}. The tidal charge reflects the effects of the Weyl curvature of the bulk space, i.e., from the 5D graviton stresses with bulk graviton tidal effect giving the name of the charge \\cite{Dad-etal:2000:,Maa:2004:}. Note that the tidal charge can be both positive and negative, and there are some indications that the negative tidal charge should properly represent the ``back-reaction'' effects of the bulk space Weyl tensor on the brane \\cite{Maa:2004:,Dad-etal:2000:,Sas-Shi-Mae:2000:}.  The exact stationary and axisymmetric solutions describing rotating black holes localized in the Randall--Sundrum braneworld were derived in \\cite{Ali-Gum:2005:}. They are described by the metric tensor of the Kerr--Newman form with a tidal charge representing the 5D correction term generated by the 5D Weyl tensor stresses. The tidal charge has an ``electric'' character again and arises due to the 5D gravitational coupling between the brane and the bulk, reflected on the brane through the ``electric'' part of the bulk Weyl tensor \\cite{Ali-Gum:2005:}, in close analogy with the spherically symmetric case \\cite{Dad-etal:2000:}.  When both the tidal and electric charge are present in a brany black hole, its character is much more complex and usual Kerr--Newman form of the metric tensor is allowed only in the approximate case of small values of the rotation parameter $a$. For linear approximation in $a$, the metric arrives at the usual Boyer--Lindquist form describing ``tidally'' charged and slowly rotating brany black holes \\cite{Ali-Gum:2005:}. For large rotational parameters, when the linear approximation is no longer valid, additional off-diagonal metric components $g_{r\\phi}$, $g_{rt}$ are relevant along with the standard $g_{\\phi t}$ component due to the combined effects of the local bulk on the brane and the dragging effect of rotation, which through the ``squared'' energy momentum tensor on the brane distort the event horizon that becomes a stack of non-uniformly rotating null circles having different radii at fixed $\\theta$ while going from the equatorial plane to the poles \\cite{Ali-Gum:2005:}. In the absence of rotation, the metric tensor reduces to the Reissner--Nordstr\\\"{o}m form with correction terms of local and non-local origin \\cite{Cha-etal:2001:}.  Here we restrict attention to the Kerr--Newman type of solutions describing the brany rotating black holes with no electric charge, when the results obtained in analysing the behaviour of test particles and photons or test fields around the Kerr--Newman black holes could be used assuming both positive and negative values of the brany tidal parameter $b$ (used instead of the charge parameter $Q^2$) \\cite{Mis-Tho-Whe:1973:Gra:}.  It is very important to test the role of the hypothetical tidal charge, implied by the theory of multidimensional black holes in the Randall--Sundrum braneworld with non-compactified additional space dimension, in astrophysical situations, namely in the accretion processes and related optical phenomena, including the oscillatory features observed in the black hole systems. There are two complementary reasons for such studies. First, the observational data from the black hole systems (both Galactic binary systems or Sgr\\,A$^*$ and active galactic nuclei) could restrict the allowed values of the tidal charge, giving a relevant information on the properties of the bulk spacetime and putting useful additional limits on the elementary particle physics. Second, the presence of the tidal charge could help much in detailed understanding of some possible discrepancies in the black hole parameter estimates coming from observational data that are obtained using different aspects of modelling accretion phenomena \\cite{Ter-etal:2007:spin-problem,McCli-Nar-Sha:2007:}.  In fact, the black hole parameter estimates come from a variety of astrophysical observations \\cite{Kli:2000:ARASTRA:,Kli:2004:,McCli-Rem:2004:CompactX-Sources:,Rem:2005:ASTRN:,Rem-McCli:2006:ARASTRA:,McCli-Nar-Sha:2007:}. The black hole spin estimates are commonly given by the optical methods, namely by X-ray line profiles \\cite{Laor:1991:ASTRJ2:,Kar-Vok-Pol:1992:ASTRJ2L:,Dov-Kar-Mar:2004:RAGtime4and5:CrossRef,Fab-Min:2005:XSpectraKerr:Book,Zak:2003:,Zak-Rep:2006:} and X-ray continuum spectra \\cite{McCli-etal:2006:astro-ph/0606076:,Mid-etal:2006:,Sha:2006::astro-ph/0508302}, and by quasiperiodic oscillations, the frequency of which enable, in principle, the most precise spin estimates, because of high precision of the frequency measurements.\\footnote{But see some problems connected with the wide variety of the resonance models \\cite{Ter-Abr-Klu:2005:ASTRA:QPOresmodel}.}  Therefore, we discuss here in detail the orbital resonance model of QPOs, which seems to be the most promising in explaining the observational data from four microquasars, namely GRO J1655-40, XTE 1550-564, H 1743-322, GRS 1915+105 \\cite{Ter-Abr-Klu:2005:ASTRA:QPOresmodel,Ter:2005:astro-ph0412500:,Tor:2005:ASTRN:,Stu-Sla-Ter:2007:ASTRA:} and in Sgr\\,A$^*$ \\cite{Asch:2004:ASTRA:,Asc-etal:2004:ASTRA:,Ter:2005:astro-ph0412500:}, or some extragalactic sources as NGC 5408 X-1 \\cite{Str:2007:astro-ph/0701390:}.  It is well known that in astrophysically relevant situations the electric charge of a black hole becomes zero or negligible on short time scales because of its neutralization by accreting preferentially oppositely charged particles from ionized matter of the accretion disc \\cite{Zel-Nov:1971:RelAstr:,Mis-Tho-Whe:1973:Gra:,Dam-etal:1978:}. Clearly, this statement remains true in the braneworld model, and that is the reason why it is enough to consider properties of brany Kerr black holes endowed with a tidal charge only. Of course, the tidal charge reflecting the non-local gravitational effects of the bulk space is non-negligible in general and it can have quite strong effect on the physical processes in vicinity of the black hole. Recently, some authors tested the tidal charge effects in the weak field limit for optical lensing \\cite{Kee-Pet:2006:brane-lensing:} and relativistic precession or time delay effect \\cite{Boh-etal:2008:SolarSys:}. Here we develop a framework of testing the tidal charge effects in the strong field near black holes with accretion discs giving rise to kHz QPOs.  In Section~\\ref{sec:tri}, the Kerr black holes with a tidal charge, introduced by Aliev and G{\\\"u}mr{\\\"u}k{\\c c}{\\\"u}o{\\u g}lu \\cite{Ali-Gum:2005:}, are described and their properties are briefly summarized. In Section~\\ref{sec:ctyri}, the Carter equations of motion are given, the equatorial circular geodesics, i.e., Keplerian circular orbits reflecting properties of thin accretion discs, are determined and properties of photon circular orbits and innermost stable orbits are discussed. In Section~\\ref{sec:pet}, the radial and vertical (latitudinal) epicyclic frequencies $\\nu_{\\mathrm{r}}$ and $\\nu_{\\theta}$, together with the Keplerian orbital frequency $\\nu_{\\mathrm{K}}$, are given. In Section~\\ref{sec:sest}, their properties are discussed, namely their radial profiles through the Keplerian accretion disc with its inner radius assumed to be located at the radius of the innermost stable circular geodesic, where the radial epicyclic frequency vanishes. In Section~\\ref{sec:sedm}, we shortly discuss the resonance conditions for the direct resonance of the both epicyclic frequencies ($\\nu_{\\theta}\\!:\\!\\nu_{\\mathrm{r}} =  3\\!:\\!2$) assumed to be in a parametric resonance \\cite{Ter-Abr-Klu:2005:ASTRA:QPOresmodel}, the relativistic precession model \\cite{Ste-Vie:1998:ASTRJ2L:} where we assume a resonance of oscillations with $\\nu_{\\mathrm{K}}\\!:\\!\\left(\\nu_{\\mathrm{K}}-\\nu_{\\mathrm{r}}\\right)= 3\\!:\\!2$, and the resonance of trapped oscillations assumed in warped disc as discussed by Kato \\cite{Kato:2007:PASJ:} [$(2\\nu_{\\mathrm{K}}-\\nu_{\\mathrm{r}})\\!:\\!2(\\nu_{\\mathrm{K}}-\\nu_{\\mathrm{r}}) = 3\\!:\\!2$]. In Section~\\ref{sec:osm} we determine the ``magic'' (dimensionless) spin of brany Kerr black holes in dependence on the (dimensionless) tidal charge, enabling presence of strong resonant phenomena because of the very special frequency ratio $\\nu_{\\mathrm{K}}\\!:\\!\\nu_{\\theta}\\!:\\!\\nu_{\\mathrm{r}} = 3\\!:\\!2\\!:\\!1$; possibility of other small integer ratios of the three frequencies at a shared radius is also discussed. Concluding remarks on the resonant phenomena in strong gravity of brany black holes are presented in Section~\\ref{sec:devet}. ",
        "conclusions": ""
    },
    "0812/0812.4898_arXiv.txt": {
        "abstract": "A method is presented for finding anisotropic distribution functions for stellar systems with known, spherically symmetric, densities, which depends only on the two classical integrals of the energy and the magnitude of the angular momentum. It requires the density to be expressed as a sum of products of functions of the potential and of the radial coordinate. The solution corresponding to this type of density is in turn a sum of products of functions of the energy and of the magnitude of the angular momentum. The products of the density and its radial and transverse velocity dispersions can be also expressed as a sum of products of functions of the potential and of the radial coordinate. Several examples are given, including some of new anisotropic distribution functions. This device can be extended further to the related problem of finding two-integral distribution functions for axisymmetric galaxies. ",
        "introduction": "\\label{intro} There is a long history of studying the structure of galaxies by construction of self-consistent distribution functions for a stellar system with a known gravitational potential.  The potential of the system determines the self-consistent mass density $\\rho$ of the system via Poisson's equation generated by the well-known Newtonian gravitational law, and also the structure of the stellar orbits according to Newton's equations of motion. The system is constructed from the building blocks of the orbits that can lie within the potential. This is called the ``from $\\rho$ to $f$'' approach for finding a self-consistent distribution function $f$ (Binney and Tremaine 1987, hereafter BT). The distribution function (hereafter DF) of the system represents how stars are distributed in the phase space of the system and the integration of the DF over velocity space yields the density. Therefore the problem of finding the DF is that of solving an integral equation; in the spherical case in particular this equation is of the first kind, in which the unknown DF $f$ only occurs inside the integral.  Some outstanding astronomers have contributed to this mathematical problem. Eddington (1916) showed that it can be solved for an isotropic DF that depends only on the energy in the spherical case, whose velocity dispersions in the radial and tangential directions must be equal, by first expressing the density as a function only of the potential, and then solving the Abel integral equation. It is also well known that the isotropic DF is unique for the potential of any spherical system.  Fricke's (1952) expansion method can be obviously generalized to the spherical case (e.g., Camm 1952; Bouvier 1962, 1963; Ossipkov 1979a; Kent and Gunn 1982; Dejonghe 1986, 1987; Dejonghe and Merritt 1988; Cuddeford 1991; Louis 1993), and similarly yields the result that DFs which are products of the two powers of the energy and the square of the angular momentum correspond to densities which are proportional to products of the potential and the spherical radial coordinate for spherical systems. Hence the DF for the system can be obtained by first expressing the density as a function of the potential and the spherical radial coordinate, and then expanding as a power series (Qian and Hunter 1995). Moreover, there may be an infinity of anisotropic DFs corresponding to any given mass density  in spherical stellar systems (Dejonghe 1987).  A class of typical anisotropic DFs that depend on the energy and the magnitude of the angular momentum,  whose velocity dispersions are anisotropic (but ellipsoidal), was independently found by Ossipkov (1979) and Merritt (1985), for a spherical density distribution. Such anisotropic DFs  of Ossipkov-Merritt type are in fact  analogues of Eddington's formula, mentioned above. In a number of papers, different integral transformation techniques can be used to obtain the solution of the spherical (or axisymmetric) problem (e.g. Lynden-Bell 1962; Hunter 1975; Kalnajs 1976; Dejonghe 1986; Qian and Hunter 1995) but there is the same difficulty of requiring not only the validity of these transformations of the density but also the complex analyticity of a density-related integral kernel to complex arguments. The contour integral method of Hunter and Qian (1993) can be also used to find anisotropic DFs for spherical systems (Qian and Hunter 1995) but it is valid for densities that are analytic, and whose singularities satisfy some conditions.  It is worth mentioning that Kuzmin and Veltmann (1967a, 1973) and Veltmann (1961, 1965, 1979, 1981) developed some more general classes when DF is a product of an unknown function over one argument on a known function of another.  The fundamental integral equations of the problem are given in {Sect.} \\ref{finteq}, and some new formulae for finding other anisotropic DFs for stellar systems with known spherically symmetric densities are presented in {Sect.} \\ref{anisdf}. These depend only on the two integrals of the energy and the magnitude of the angular momentum. This work constitutes the core of this paper. These formulae in fact come from an combination of the ideas of Eddington and Fricke (see above).  Of course, they can be also regarded as simply an extension of Eddington's formula. A type of anisotropic DF which is a sum of products of functions only of the energy and powers of the magnitude of the angular momentum is derived in {Sect.} \\ref{anisdf1} and another, which is a sum of products of functions only of a special variable and powers of the magnitude of the angular momentum, in {Sect.} \\ref{anisdf2}. More general formulae are given in the last part of {Sect.} \\ref{anisdf}. Various formulae of the velocity dispersions for such models of these DFs are also  shown in all the three parts of {Sect.} \\ref{anisdf}. Several examples are given in {Sect.} \\ref{appli}, including some of new anisotropic DFs.  Different anisotropic DFs for the Plummer model are first given in {Sect.} \\ref{apm}. Then anisotropic DFs of the H\\'{e}non isochrone model appear in {Sect.} \\ref{ahm} and those of the $\\gamma$-model are described in {Sect.} \\ref{agm}. Section {\\ref{con} is a summary and conclusion. ",
        "conclusions": "\\label{con} Although real galaxies are hardly spherically symmetric, it is  a very necessary and significant step to study the self-consistent anisotropic DFs for stellar systems with the known spherically symmetric density. This is not only because some important properties (e.g. surface density and cumulative mass) of the spherical model are similar to those of the models that can be generated from the spherically symmetric density by replacing the spherical radius $r$ by an axisymmetric or triaxial radius $m=\\sqrt{x^2+(y/q_2)^2+(z/q_3)^2},$  but also because both spherical and non-spherical models have the same typical behaviours of the dynamical quantities in the limits of small  and large radii. Such models are elliposoidal. Other classes of flattened models can be constructed by using the equipotential method (e.g., Kutusov and Ossipkov 1980; Jiang 2000).  Anisotropic DFs can be obtained for  stellar systems with known spherically symmetric density as a sum of products of functions only of the potential and a  special function (or power) only of the radial coordinate, i.e. these DFs are a sum of products of functions only of a special variable (or the energy) and a power only of the magnitude of the angular momentum. This comes from a combination of the ideas of Eddington and Fricke. It is an extension of Eddington's classical solution for the isotropic DF of a known spherical density. Like his method, it requires that the density be expressed as a special function of the potential, and now also of the radial coordinate. Also, part of it is a sum of real integrals of functions only of the potential. Most of these integrals can be calculated analytically in terms of elementary functions or generalized hypergeometric functions (e.g. Kuzmin and Veltmann 1967b).  These formulae concerning anisotropic DFs are obtained by use of the Abel integral equation, and they are suitable for all such densities that allow the real integrals of the potential to be valid. The anisotropic DFs of Ossipkov-Merritt type are their examples in special cases. After the application to Plummer's spherical model, it is found that the anisotropic DFs given by (\\ref{pdf}) are in fact the same as those given by Dejonghe (1986). Three different anisotropic DFs are given for H\\'{e}non's (1959a, b) isochrone model and they can be expressed in terms of elementary functions. Many anisotropic DFs of $\\gamma$-models, expressed in elementary functions or generalized hypergeometric functions, can also be obtained. One can further give formulae of the velocity dispersions for these anisotropic DFs.  These expressions can be extended further to axisymmetric systems. It is straightforward to find their analogues for axisymmetric systems and they can be used to obtain the even DF of Binney's (BT) logarithmic potential although Evans (1993) derived it using Lynden-Bell's (1962) method. For the well-known Lynden-Bell (1962) model, these analogues degenerate into the method of Fricke (1952).  \\begin{acknowledgement} The first author was supported in part by NSFC 10271121 and by SRF for ROCS, SEM. The numerical calculations were performed by use of the computer cluster at the Department of Mathematics of the University of Manchester. The first author would like to thank Dr.~David Moss for his helpful comments on this paper. The second author was supported by Leading Scientific School grant 1078.2003.02. The cooperation of authors was supported by joint grants of NSFC 10511120278/10611120371 and RFBR 04-02-39026. The two authors are very grateful to Professor Konstantin Kholshevnikov for his valuable discussions on this work. \\end{acknowledgement}"
    },
    "0812/0812.4851_arXiv.txt": {
        "abstract": " ",
        "introduction": "Recent observations by the PAMELA \\cite{Adriani:2008zr,Adriani:2008zq} and ATIC/PPB-BETS \\cite{chang:2008,Torii:2008xu} experiments have revealed the electron and positron excesses in the cosmic ray spectrum. These results indicate the presence of nearby sources of electron-positron pairs (less than $1$ kpc away). Possible candidates include astrophysical objects such as pulsars \\cite{Kawanaka:2009dk,Shen70,Buesching:2008hr,Hooper:2008kg, Yuksel:2008rf,chi:1996,zhang:2001,grimani:2007,Profumo:2008ms,Aha:95,Malyshev:2009tw,Grasso:2009ma}, supernova (SN) remnants, \\cite{Fujita:2009wk,ShenBerkey68,cowsik79,boulares89,erlykin02,pohl98,Kobayashi:2003kp, Shaviv:2009bu,Hu:2009bc,Blasi:2009hv,Blasi:2009bd,Mertsch:2009ph,Biermann:2009qi} or microquasars \\cite{Heinz:2002qj}, or dark matter annihilations/decays. \\cite{Asano:2006nr,ArkaniHamed:2008qn, Bergstrom:2008gr,Cirelli:2008jk,Chen:2008yi,Chen:2008qs,Chen:2008dh,Cholis:2008wq,Cirelli:2008pk,Hamaguchi:2008ta,Hisano:2008ah,Hisano:2004ds,Ishiwata:2008cv,Hall:2008qu,Zhang:2008tb,MarchRussell:2008tu,Hooper:2008kv} Instead, we might be observing the propagation effects \\cite{Delahaye:2007fr,Cowsik:2009ga,Katz:2009yd,Stawarz:2009ig,Schlickeiser:2009qq} or the proton contamination \\cite{Fazely:2009jb,Schubnell:2009gk,Israel:2009}.  The ATIC/PPB-BETS excess has a possible cutoff at $\\varepsilon_e \\sim 600$ GeV, which might fix the dark matter mass. From the astrophysical viewpoint, the cutoff implies a single or at least a few sources since many sources usually broaden the cutoff. The source age should be less than $10^{6{\\rm -}7}$ years because electrons lose energy through synchrotron and inverse Compton processes, suggesting the Galactic rate of $\\sim (10\\ {\\rm kpc}/1\\ {\\rm kpc})^2/10^{6{\\rm -}7}\\ {\\rm yr} \\sim 1/10^{4{\\rm -}5}$ yr, i.e., $\\sim 10^2$--$10^3$ times rarer than SNe. This ratio $\\sim 10^2$--$10^3$ is comparable to that of the energy density between cosmic ray nuclei and positrons. Therefore, the electron-positron source may also produce a huge energy $\\sim 10^{50}$ erg like an SN that releases $\\sim 10^{50}$ erg for providing cosmic ray nuclei.  In this paper, we propose a new possibility that a nearby ($d \\sim 1$ kpc) gamma-ray burst (GRB) or GRB-like pulsar/SN remnant/microquasar about $t_{\\rm age}\\sim 10^{5-6}$ years ago may be responsible for the PAMELA and ATIC/PPB-BETS excesses, and predict a sharp spectral cutoff that is similar to the dark matter predictions, in addition to a possible line. GRBs are the most luminous objects in the universe \\cite{Meszaros:2006rc,Zhang:2007nka}, brief ($\\sim 1$--$100$ sec) bursts of high-energy ($\\sim 0.1$--$1$ MeV) photons appearing at random in the sky about 1000 times per year (corresponding to the collimation-corrected local rate $\\sim 1/10^{5{\\rm -}6}$yr/galaxy \\cite{Guetta:2006gq}). Thanks to the discovery of afterglows (long-lasting counterparts in longer wavelengths) and SNe, it is widely accepted that (long) GRBs are associated with the deaths of massive stars. The central core of a massive star gravitationally collapses into a black hole or neutron star, which somehow launches a collimated outflow (jet) and produces GRBs and afterglows with a typical true energy of $\\sim 10^{51}$ erg. GRBs could emit a significant amount of energy that is comparable to the main MeV photons into GeV-TeV gamma-rays, as observed by the Fermi satellite \\cite{Abdo09}, and eV photons (so-called optical flashes) as in the famous naked-eye GRB 080319B \\cite{Racusin:2008pd}.  Very recently, the Fermi Large Area Telescope has measured the electron spectrum up to $\\sim 1$ TeV, which is very smooth $\\sim \\varepsilon_e^{-3}$ without any spectral peak as reported by ATIC/PPB-BETS. The HESS collaboration also provides the electron spectrum \\cite{Collaboration:2008aa,Aharonian:2009ah}, which is consistent with the Fermi result and appears to show a steepening above $\\sim 1$ TeV. The differences between ATIC/PPB-BETS and Fermi/HESS are still controversial because Fermi removes a significant fraction of electrons above $\\sim 300$ GeV to avoid the $\\sim 10^{3{\\rm -}4}$ times larger hadron contamination and reconstructs the real flux by the Monte Carlo simulations \\cite{Moiseev:2007js,Israel:2009}, while the statistical errors of ATIC are much worse than those of Fermi. Therefore, we discuss each case separately and show that a GRB/pulsar model with sightly different parameters may reproduce the Fermi/HESS smooth spectra as well.  In this paper, we try not to specify a source class to make the discussion as model-independently as possible. We consider a GRB-like astrophysical source, denoted as GRB/Pulsar for short, which produces electron-positron pairs from a compact region in a short timescale compared with its age. If the source has these properties, we can apply the results to pulsars, SN remnants, microquasars, GRBs, etc. The exception is \\S\\ref{sec:GRB}, which is devoted to GRB models. ",
        "conclusions": "\\label{sec:dis} We have proposed that a nearby gamma-ray burst (GRB) or GRB-like (old, single, and short-lived) pulsar/SN remnant/microquasar about $10^{5-6}$ years ago may be responsible for the excesses of cosmic ray positrons and electrons recently observed in the PAMELA, ATIC/PPB-BETS, and Fermi/HESS experiments. Such a scenario appears extreme but still consistent with the current observations. In particular, a GRB/pulsar model can reproduce the smooth Fermi/HESS spectra as well as the spiky ATIC/PPB-BETS spectra by slightly changing the parameters (see Fig.~\\ref{fig:flux}).  Although such a burstlike scenario was discussed previously, \\cite{Atoyan:1995,ShenBerkey68,Shen70} it is the first to argue the similarities (e.g., sharp cutoff) and differences (e.g., cutoff width) between the astrophysical and dark matter scenarios.\\footnote{ Our work was carried out independently of Profumo, \\cite{Profumo:2008ms} who also pointed out that an astrophysical source (pulsar) can produce a sharp spectral cutoff.} In particular, we propose a new method to discriminate models by using the cutoff width (see below). This new viewpoint arises from the confrontation between astrophysical and dark matter models as well as the developments of experiments with fine spectral resolutions, both of which did not exist previously. The GRB model is also a new one, and the markedly improved quality of current data would allow us to reconsider the problem of whether a GRB-like source can explain the data or not.  The spectral cutoff and line in Figs.~\\ref{fig:e+frac} and \\ref{fig:flux} should have a finite dispersion under realistic circumstances, in contrast with the dark matter origin \\cite{Chen:2008fx}. We may be able to discriminate models by observing the cutoff shape (or width) since the future CALET experiment has a resolution better than a few $\\%$ ($>100$ GeV) \\cite{torii:2006,torii:2008}.  Since Eq.~(\\ref{eq:emax}) yields ${\\Delta \\varepsilon_{\\rm cut}}/{\\varepsilon_{\\rm cut}} =-{\\Delta b}/{b}-{\\Delta t}/{t}$, the dispersion arises from (a) the fluctuation of the energy loss rate $\\Delta b$ due to the difference of starlight and magnetic fields by location and (b) the duration of the source $\\Delta t$. To estimate $\\Delta b$, we assume that the energy loss rate fluctuates by $\\delta b$ over the scale $d_{b}$. Cosmic rays travel a distance $c t_{\\rm age}$ and pass through $N_b \\sim c t_{\\rm age}/d_b$ patches, averaging the fluctuations as $\\Delta b \\sim \\delta b/\\sqrt{N_b}$. Then, we have \\beqa \\left(\\frac{\\Delta \\varepsilon_{\\rm cut}}{\\varepsilon_{\\rm cut}}\\right)_{\\Delta b} \\sim 6\\% \\ \\left(\\frac{\\delta b}{b}\\right) \\left(\\frac{d_b}{1\\ {\\rm kpc}}\\right) \\left(\\frac{t_{\\rm age}}{10^6\\ {\\rm yr}}\\right)^{-1/2}, \\label{eq:dcutdb} \\eeqa which may be detectable if the starlight and magnetic fields differ by $\\delta b/b \\sim 1$ over the disk thickness $d_b \\sim 1$ kpc [see also Ref. \\citen{Malyshev:2009tw}]. As for the duration effect, GRBs are too short, but a pulsar with magnetic field $B$ and initial rotation period $P_0$ has a spin-down duration $\\Delta t \\sim 3c^3 I/B^2 R_{*}^6 \\Omega_0^2 \\sim 6 \\times 10^3\\ {\\rm yr}$ $(B/10^{12}\\ {\\rm G})^{-2} (P_0/10\\ {\\rm ms})^2$, whereas a microquasar has an active time $\\Delta t \\sim 10^{50}\\ {\\rm erg}/L \\sim 10^{5}\\ {\\rm yr}$ $(L/10^{38}\\ {\\rm erg}\\ {\\rm s}^{-1})^{-1}$, yielding \\beqa \\left(\\frac{\\Delta \\varepsilon_{\\rm cut}} {\\varepsilon_{\\rm cut}}\\right)_{\\Delta t} \\sim 10\\% \\left(\\frac{\\Delta t}{10^{5}\\ {\\rm yr}}\\right) \\left(\\frac{t_{\\rm age}}{10^{6}\\ {\\rm yr}}\\right)^{-1}. \\eeqa [See Ref. \\citen{Kawanaka:2009dk} for a more detailed discussion.]  The line shape in Figs.~\\ref{fig:e+frac} and \\ref{fig:flux} is $\\propto (1-\\varepsilon_e/\\varepsilon_{\\rm cut})^{\\alpha-2}$ from Eqs.~(\\ref{eq:fsol}) and (\\ref{eq:emax}). For an energy resolution $\\Delta \\varepsilon_e/\\varepsilon_e=\\lambda$, the flux is enhanced at the line by a factor of $\\lambda^{\\alpha-2}/(\\alpha-1)$ ($\\sim 2.7$ for $\\lambda=2\\%$ and $\\alpha=1.8$). As $\\varepsilon_e<(b t +1/\\varepsilon_{\\max})^{-1} =(1/\\varepsilon_{\\rm cut}+1/\\varepsilon_{\\max})^{-1}$, there is no divergence in the line.  Note that there are still some uncertainties in the diffusion coefficients. The change in the diffusion coefficients can be adjusted by slightly changing the model parameters. The smaller $K$ makes the diffusion length $d_{\\rm diff}$ smaller, and the particle density inside that radius becomes higher, being proportional to $d_{\\rm diff}^{-3}$. For different $\\tilde{K}$ instead of $K$, we can apply our results by rescaling the distance and energy as $d\\rightarrow d\\sqrt{\\tilde{K}/K}$ and $E_{e}\\rightarrow E_{e}(\\tilde{K}/K)^{3/2}$, respectively. Here, the rescaling generally depends on the rigidity of cosmic rays.  Note also that the energy densities of radiation and magnetic fields differ by a factor of $\\sim 2$--$3$ within the diffusion region $\\sim 1$ kpc \\cite{Strong:1998fr}. These fluctuations do not lead to a large dispersion of the cutoff width as shown in Eq.~(\\ref{eq:dcutdb}), whereas the mean energy density (i.e., the cutoff energy itself) may have an uncertainty by a factor of $\\sim 2$--$3$. Above $\\sim 500$ GeV, the Klein-Nishina suppression of the inverse Compton cooling is also important, \\cite{Stawarz:2009ig,Schlickeiser:2009qq} since the optical and infrared photons dominate the radiation energy density, while the synchrotron cooling still operates. This effect would raise the cutoff energy by a factor of $\\sim 2$, but not affect the cutoff width considerably as long as the energy loss time is shorter for higher energy. The form of the source cutoff around $\\varepsilon_e \\siml \\varepsilon_{\\max}$ also affects the cutoff width of \\beqa \\left(\\frac{\\Delta \\varepsilon_{\\rm cut}}{\\varepsilon_{\\rm cut}}\\right)_{\\varepsilon_{\\max}} \\sim \\frac{1}{bt \\varepsilon_{\\max}} \\sim 20\\% \\left(\\frac{\\varepsilon_{\\rm cut}}{600\\ {\\rm GeV}}\\right) \\left(\\frac{\\varepsilon_{\\max}}{3\\ {\\rm TeV}}\\right)^{-1}. \\eeqa In addition, as we have shown in Figs.~\\ref{fig:e+frac} and \\ref{fig:flux}, a soft source spectrum $\\alpha>2$ leads to a round spectral shape below the cutoff, making it difficult to measure the cutoff width. Therefore, only under special conditions ($\\alpha \\siml 2$, $\\varepsilon_{\\max} \\simg 5 \\times \\varepsilon_{\\rm cut}$, $t_{\\rm age} \\simg 5 \\times \\Delta t$), a GRB-like source has a distinguishable feature of a sharp cutoff, while, once detected, it could pindown the GRB-like nature of the source.  Our GRB-like astrophysical scenario is compatible with the diffuse gamma-ray background observations because the electron plus positron flux from a GRB/pulsar is comparable at most to the local electron spectra (see Fig.~\\ref{fig:flux}) and the locally observed electron spectra predict a minor role of the inverse Compton emission far below the pion decay component as in the ``conventional model'' described in Ref. \\citen{Strong:2004de}. The constraints are less severe in the GRB case of small chance probability. The fluctuation of the diffuse gamma-ray background due to the nonuniform distribution of GRBs/pulsars could be an interesting future probe. Note that we can roughly estimate the diffuse gamma-ray flux from the electron flux as follows. Gamma-rays with $\\varepsilon_\\gamma \\sim 10$--$1000$ GeV are produced by electrons with $\\varepsilon_e \\sim 100$--$1000$ GeV via inverse Compton scattering of $\\sim 1$ eV photons, and the total gamma-ray energy created during the cooling time $\\sim 1/b \\varepsilon_e$ is comparable to the electron total energy. Since it takes only a time $\\sim 2[K(\\varepsilon_e)/b \\varepsilon_e]^{1/2}/c$ for gamma-rays to cross the diffusion length of electrons while electrons stay there for the cooling time $\\sim 1/b \\varepsilon_e$, the flux ratio is about the time ratio \\beqa \\frac{\\varepsilon_\\gamma^2 \\Phi_\\gamma(\\varepsilon_\\gamma)} {\\varepsilon_e^2 \\Phi_e(\\varepsilon_e)} \\sim 2 \\left[\\frac{K(\\varepsilon_e) b \\varepsilon_e}{c^2}\\right]^{1/2} \\sim 0.05 \\left(\\frac{\\varepsilon_e}{10^3\\ {\\rm GeV}}\\right)^{2/3}, \\eeqa which is below the pion emission in the GeV-TeV region.  Similar GRBs in our Galaxy may have been observed as mysterious TeV gamma-ray sources, the so-called TeV unidentified sources, which have no clear counterpart at other wavelengths \\cite{Ioka:2004kv,Ioka:2009dh} and/or the 511 keV electron-positron annihilation line from the Galactic bulge \\cite{Bertone:2004ek,Parizot:2004ph}. There is evidence that GRBs predominantly occur in galaxies with less metals than our own \\cite{Fruchter:2006py,Stanek:2006gc}. However, no metallicity-energy correlation is confirmed \\cite{Savaglio:2008fj}. There may also be a bias that a metal-rich region tends to have no optical afterglow due to dust absorption (i.e., dark GRB, which is about half of all the GRBs) and lacks host identifications and metal measurements.  In conclusion, the PAMELA, ATIC/PPB-BETS, and Fermi/HESS excesses can be reproduced well if a nearby GRB-like source, such as a pulsar, SN remnant, microquasar, or GRB, produces electron-positron pairs with energy $\\sim 10^{50}$ erg and a power-law spectral index $\\alpha \\sim 1.6$--$2.2$. The spectra have a sharp cutoff that is similar to the dark matter predictions. The cutoff energy marks the source age, whereas the cutoff width has information on the source duration, which may be resolved by future experiments like CALET. Whether the source is single or multiple may be probed with the anisotropy of electron and positron fluxes, whereas whether the source is leptonic or hadronic may be probed with the antiprotons \\cite{Fujita:2009wk} and secondary nuclei \\cite{Mertsch:2009ph}."
    },
    "0812/0812.3195_arXiv.txt": {
        "abstract": "GJ 436b is a Neptune-size planet with 23.2 Earth masses in an elliptical orbit of period 2.64 days and eccentricity 0.16.  With a  typical tidal dissipation factor ($Q' \\sim 10^6$) as that of a giant planet with convective envelope, its orbital circularization timescale under internal tidal dissipation is around $1$ Gyr, at least two times less than the stellar age ($>3$~Gyr). A plausible mechanism is that the eccentricity of GJ 436b is modulated by a planetary companion due to their mutual perturbation. Here we investigate this possibility from the dynamical viewpoint. A general method is given to predict the possible locations of the dynamically coupled companions, including in nearby/distance non-resonant or mean motion resonance orbits with the first planet. Applying the method to GJ 436 system,  we find it is very unlikely that the eccentricity of GJ 436b is maintained at the present location by a nearby/distance companion through secular perturbation or mean motion resonance.  In fact, in all these simulated cases, GJ 436b will undergo eccentricity damp and orbital decay, leaving the present location within the stellar age. However, these results do {\\em not} rule out the possible existence of  planet companions in nearby/distance orbits, although they are not able to maintain the eccentricity of GJ 436b. ",
        "introduction": "The discovery of the first extrasolar planet around a pulsar$^{[1]}$, which was quickly followed  by the detection of first Jupiter-like planet around the star 51 Peg$^{[2]}$, opened a new era  for planetary science. The planet around 51 Peg is known as a hot Jupiter: planets with masses comparable to Jupiter's, orbits typically within $0.1$ AU and surface temperatures $\\sim 1000$ K$^{[3]}$. To date, more than 310 planets have been discovered in about 250 planetary systems\\footnote{http://exoplanet.eu/}. With the improvement of detection precision and the use of various techniques, the minimum mass of discovered  planets around main sequence stars  is down to 3.3 Earth masses ($M_\\oplus$), which is MOA-2007-BLG-192-L b around a star with mass of $0.06M_\\odot$$^{[4]}$.  Among the planets discovered so far, GJ 436b is the only transiting Neptune-mass planet orbiting an M-type star.  It was first detected by  radial velocity techniques$^{[5]}$, with its orbital elements refined later by Maness et al.$^{[6]}$. The transiting signals of GJ 436b were first discovered by Gillon et al.$^{[7]}$ and followed by lots of work$^{[8-13]}$. These observations reveal GJ 436b with a mass of $23.17M_\\oplus$ and a radius of $4.22 R_\\oplus$(see Table 1) .  \\begin{table}[!htbp] \\begin{center} \\caption{The parameters of GJ 436 and its planet companion GJ 436.} \\begin{tabular}{lll} \\tableline\\tableline Parameter  & Value  & Ref. \\\\ \\tableline $M_*~(M_\\odot)$ & $0.452^{+0.014}_{-0.012}$ &  1 \\\\ $R_*~(R_\\odot)$ & $0.464^{+0.009}_{-0.011}$ & 1 \\\\ stellar age (Gyr) &    $6^{+4}_{-5}$ &  1 \\\\ \\tableline $M_p (M_\\oplus)$  & $23.17\\pm0.079$  & 1  \\\\ $R_p~(R_\\oplus)$ & $4.22^{+0.09}_{-0.10}$ & 1\\\\ $M_p sin i (M_J)$  & $ 0.0713 \\pm 0.006 $  & 2  \\\\ $a (AU)$ & $0.0285$ & 2  \\\\ $P (days)$  & $2.64385\\pm0.00009$ & 2  \\\\ $e$ &  $0.16\\pm0.019$ & 2 \\\\ \\tableline \\end{tabular} \\end{center} Refs.: 1. \\cite{tor07}, 2. \\cite{man07}. \\end{table}  One of the most interesting characteristics about  GJ 436b is its significant eccentricity (0.16) in an orbit (2.64 days) very close to the host star. Assuming a  tidal dissipation factor ($Q' \\sim 10^6$) as that of a gas giant planet, its orbit circularization timescale under internal tidal dissipation is around $1$ Gyr. On the other hand, the fiducial age   of the host star is $6^{+4}_{-5}$~Gyr$^{[11]}$,  and according to observation, GJ 436 has low rotation velocity and does not exhibit particular strong chromospheric activity nor  photometric variability$^{[5]}$, indicating an age $>3$ Gyr. As  the orbit is not circulated by planetary tide, either GJ 436b has $Q'>6\\times 10^6$, or there is a planetary companion which  induces a periodic modulation of its orbital eccentricity.  Considering that radial velocities of GJ 436 reveal a long-term trend, Maness et al. proposed the presence of a long-period ($\\sim 25$ yr) planet companion with mass $\\sim 0.27 M_J$ (Jupiter mass) in an eccentric orbit ($e\\sim 0.2$)$^{[6]}$. Recently, Ribas et al. suggested that the observed radial velocities of the system are consistent with an additional small, super-Earth planet in the outer 2:1 mean-motion resonance with GJ 436b$^{[14]}$. Such a possibility was also studied  from the dynamical viewpoint$^{[15]}$. More recent inspection of transit data implies that GJ 436b is perturbed by another planet with mass $\\le 12 M_\\oplus$ in a non-resonant orbit of 12days (0.08 AU)$^{[16]}$.  In this paper, we exam the possibility of a nearby or distant undiscovered  planet through dynamical considerations. The key issue here is to locate the undetected  companion that can both excite and maintain the eccentricity of GJ 436b. First we present the dynamical restrictions on the possible locations of the companion in section 2. Then in section 3 we show the analytical and numerical results on the eccentricity modulation by a companion in nearby or distant non-resonance orbits, followed by the investigation of the possible companion in resonant orbits in section 4. Conclusions are presented  in section 5.  Although the method is derived in GJ 436 system, it can be applied to any other systems in similar situations. ",
        "conclusions": "Perturbation from a companion $M_2$ seems to be the most plausible approach to maintain the  moderate eccentricity  of GJ 436b in a close-in orbit. In this paper, we study extensively all different locations of the planet companion in the following three  situations:  {\\em (i) In nearby orbits}. The eccentricity of GJ 436b can be excited to 0.16 with a broad range of planet mass (above few Earth-masses). An interesting case is the second planet can have a mass much smaller than $M_1$. However, since the orbital decay time is  short ($\\sim 20 $Gyr) for GJ 436b at the present location with the observed eccentricity, significant orbital decay $(\\sim 25\\%)$ is expect so that GJ 436b would be in a much closer orbit, and the eccentricity of GJ 436b would be damped within the stellar age. Thus the eccentricity of GJ 436b can  {\\em not} be maintained by a nearby companion. On the other hand, one can {\\em not} rule out the possibility that companions exist in nearby orbits of GJ 436b as long as they are in stable orbits, e.g., the possible existence of a planet with $\\le 12 M_\\oplus$ in an orbit of 12days from  transit data$^{[16]}$.  However, the companion, if exists, does not account for the eccentricity of GJ 436b.   {\\em (ii) Distance Orbits}. Distance companions on moderate eccentric orbits can {\\em not} excite and maintain  the significant eccentricity of GJ 436b. However, the presence of giant planets in extended orbits ($\\sim $ few AU) is also possible, although they can {\\em not} modulate the eccentricity of close-in planets. For example, there may exist a long-period ($\\sim 25$ yr) planet companion with mass $\\sim 0.27 M_J$ inferred from radial velocities of  GJ 436 $^{[6]}$.  {\\em (iii) Resonance Orbits}. Although a border range of companion mass can excite the eccentricity of $M_1$ to  a moderate value, similar to situation ({\\em i}), significant orbital decay would occur so that the two planets will leave 2:1 resonance soon, if GJ 436b has a normal dissipation factor ($Q'\\sim 10^6$). In this case, the presence of a companion in the outer 2:1 resonance  with GJ 436b is unlikely. However, if we have to confess that GJ 436b has a extremely high dissipation factor ($Q'> 6 \\times 10^6$), the existence of companions in resonance orbits can {\\em not} be ruled out in this extreme case, e.g., the planet ($\\sim 5 M_\\oplus$) in the outer  2:1 resonance place proposed by  Ribas et al.$^{[14]}$. However, even it exists, it does not account for the eccentricity of GJ 436b.  Based on the extensive investigations in this paper, we think $e_1=0.16$ of GJ 436b can not be maintained by a companion  in either nearby or distance orbits through secular perturbation or mean motion resonance. Thus the maintaining of its eccentricity remains a challenge problem, unless GJ 436b has a extremely high dissipation factor ($Q'> 6 \\times 10^6$)."
    },
    "0812/0812.4310_arXiv.txt": {
        "abstract": "% We describe a new image co-addition tool, AWAIC, to support the creation of a digital Image Atlas from the multiple frame exposures acquired with the Wide-field Infrared Survey Explorer (WISE). AWAIC includes preparatory steps such as frame background matching and outlier detection using robust frame-stack statistics. Frame co-addition is based on using the detector's Point Response Function (PRF) as an interpolation kernel. This kernel reduces the impact of prior-masked pixels; enables the creation of an optimal matched filtered product for point source detection; and most important, it allows for  resolution enhancement (HiRes) to yield a model of the sky that is consistent with the observations to within measurement error. The HiRes functionality allows for non-isoplanatic PRFs, prior noise-variance weighting, uncertainty estimation, and includes a ringing-suppression algorithm. AWAIC also supports the popular overlap-area weighted interpolation method, and is generic enough for use on any astronomical image data that supports the FITS and WCS standards. ",
        "introduction": "The goal of image co-addition is to optimally combine a set of (usually dithered) exposures to create an accurate representation of the sky, given that all instrumental signatures, glitches, and cosmic-rays have been properly removed. By ``optimally'', we mean a method which maximizes the signal-to-noise ratio (SNR) given prior knowledge of the statistical distribution of the input measurements.  The Wide-field Infrared Survey Explorer (WISE) mission will be generating over a million exposures (or frames) over the sky. WISE is a NASA Midex mission scheduled for launch in late 2009. It will survey the entire sky at 3.3, 4.7, 12, and 23$\\mu$m with sensitivities up to three orders of magnitude beyond those achieved with previous all-sky surveys. For details on the scientific goals, requirements, instrument and mission design, see Mainzer et al.\\ 2005. One of the primary products from WISE is a digital Image Atlas that combines the multiple 8.8 second, $47\\arcmin\\times47\\arcmin$ frame exposures within predefined tiles over the sky. To support this, we have developed a suite of software modules collectively referred to as AWAIC for execution in the automated processing pipeline at the Infrared Processing and Analysis Center. The modules are written in ANSI-compliant C and wrapped into a Perl script, and will be made portable in the near future.  Here we review AWAIC's co-addition algorithms, products, and extension to resolution enhancement (HiRes). It's important to note that HiRes is not in the WISE baseline plan. It was implemented primarily to support offline research. The statistical robustness and performance of algorithms will be addressed in more detail in future papers. ",
        "conclusions": "We have given a broad overview of the algorithms implemented in a new generic co-addition/HiRes'ing tool. The goal is to produce high fidelity science quality products with uncertainty estimates and metrics for validation thereof. The HiRes (MCM) algorithm contains considerable improvement over previous methods in that it includes {\\it a posteriori} uncertainty estimation, statistically motivated convergence criteria, a powerful diagnostic (the CFV) to locate inconsistencies in the input data and assess the overall quality of HiRes solutions, and the ability to handle non-isoplanatic PRFs. Algorithms will be discussed in more detail in future papers. Future work will explore methods to accelerate convergence in MCM, the ability to handle time-dependent PRFs (e.g., adapted to variable seeing), and an analysis of the completeness, reliability, and photometric accuracy of sources detected in HiRes'd images, especially in confused fields. More examples, analyses, user-interface details, and animations of MCM can be found at \\htmladdURL{http://web.ipac.caltech.edu/staff/fmasci/home/wise/awaic.html}"
    },
    "0812/0812.2911_arXiv.txt": {
        "abstract": "{Evidence has grown over the past few years that the neutral phase of blue compact dwarf (BCD) galaxies may be metal-deficient as compared to the ionized gas of their H\\2\\ regions. These results have strong implications for our understanding of the chemical evolution of galaxies, and it is essential to strengthen the method, as well as to find possible explanations. } {We present the analysis of the interstellar spectrum of Pox\\,36 with the \\textit{Far Ultraviolet Spectroscopic Explorer} (\\textit{FUSE}).  Pox\\,36 was selected because of the relatively low foreground gas content that makes it possible to detect absorption-lines weak enough that unseen components should not be saturated.} {Interstellar lines of H\\1, N\\1, O\\1, Si\\2, P\\2, Ar\\1, and Fe\\2\\ are detected. Column densities are derived directly from the observed line profiles except for H\\1, whose lines are contaminated by stellar absorption, thus needing the stellar continuum to be removed. We used the TLUSTY models to remove the stellar continuum and isolate the interstellar component. The best fit indicates that the dominant stellar population is B0. The observed far-UV flux agrees with an equivalent number of $\\sim300$ B0 stars. The fit of the interstellar H\\1\\ line gives a column density of $10^{20.3\\pm0.4}$\\,cm$^{-2}$. Chemical abundances were then computed from the column densities using the dominant ionization stage in the neutral gas. Our abundances are compared to those measured from emission-line spectra in the optical, probing the ionized gas of the H\\2\\ regions.} {Our results suggest that the neutral gas of Pox\\,36 is metal-deficient by a factor $\\sim7$ as compared to the ionized gas, and they agree with a metallicity of $\\approx1/35$\\,Z$_\\odot$. Elemental depletion is not problematic because of the low dust content along the selected lines of sight. In contrast, the ionized gas shows a clear depletion pattern, with iron being strongly depleted. } {The abundance discontinuity between the neutral and ionized phases implies that most of the metals released by consecutive star-formation episodes mixes with the H\\1\\ gas. The volume extent of the enrichment is so large that the metallicity of the neutral gas increases only slightly. The star-forming regions could be enriched only by a small fraction ($\\sim1\\%$), but it would greatly enhance its metallicity. Our results are compared to those of other BCDs. We confirm the overall underabundance of metals in their neutral gas, with perhaps only the lowest metallicity BCDs showing no discontinuity. } ",
        "introduction": "In the chemical downsizing scenario, in which massive galaxies are the first to form stars at a high rate (e.g., Cen \\&\\ Ostriker 1999), dwarf galaxies can remain chemically unevolved until later times. It is thus theoretically possible that some dwarf galaxies in the nearby Universe still contain pristine gas. The first observational constraint was provided by studies of chemical abundances in the ionized gas of blue compact dwarf galaxies (BCDs). These studies yield a fundamental result: a metallicity floor appears to exist, with no BCD having an ionized gas oxygen abundance lower than $12+\\log {\\rm (O/H)} \\sim 7.0$, or Z$_\\odot$/50\\footnote{We use hereafter the solar oxygen abundance $12+\\log {\\rm (O/H)} = 8.66$ of Asplund et al.\\ (2005)}. For more than three decades, I Zw 18 discovered by Zwicky (1966) and first studied by Sargent \\& Searle (1970) held the record as the most metal-deficient galaxy known, with an oxygen abundance $12+\\log {\\rm (O/H)} = 7.17\\pm0.01$ in its NW component and $7.22\\pm0.02$ in its SE component (Thuan \\& Izotov 2005). Only very recently has I\\,Zw\\,18 been displaced by two BCDs that are more metal-deficient, SBS\\,0335--052W with $12+\\log {\\rm (O/H)} = 7.12\\pm0.03$ (Izotov et al.\\ 2005) and DDO\\,68 with $12+\\log {\\rm (O/H)} = 7.14\\pm0.03$ (Izotov \\& Thuan 2007). The difficulty of finding more metal-poor objects was originally explained by the local mixing of the amount of metals produced and released in a single starburst episode (Kunth \\& Sargent 1986). This hypothesis has been since challenged by observations of uniform abundances in disconnected star-forming regions within a single dwarf galaxy (e.g., Skillman \\& Kennicut 1993; Kobulnicky \\& Skillman 1997; Noeske et al.\\ 2000).  Tenorio-Tagle (1996) proposed an attractive scenario in which supernovae (SNe) products are driven to high galactic latitudes in a hot phase ($\\sim10^6$\\,K) before they cool down, condense as molecular droplets, and rain back on the disk. The complete cycle (before mixing) would last several $\\sim100$\\,Myr. This scenario naturally explains the uniform abundances found in dwarf galaxies, while ruling out local enrichment within the star-forming regions. It must be noted that models of Rieschick \\& Hensler (2001; 2003), based on Tenorio-Tagle's scenario, predict that only $\\sim75$\\%\\ of the SNe products go through the outer galactic cycle as hot gas, while the remaining 25\\%\\ evolve in an inner local cycle and can enrich the star-forming regions in only 10\\,Myr. A few percent could eventually leave the gravitational field of the galaxy. On the other hand, {\\it Chandra} X-ray observations of SBS\\,0335--052W, SBS\\,0335--052E, and I\\,Zw\\,18 do not reveal any sign of a hot gas breaking out from the stellar body implying that the galactic fountain scenario may not be valid for BCDs (Thuan et al.\\ 2004). In fact, Recchi et al.\\ (2001; 2004) modeled star-formation in gas-rich dwarf galaxies and found that the majority of newly produced metals resides in a cold phase ($\\lesssim2\\times10^4$\\,K) after only a few tens of Myr, supporting the common instantaneous mixing hypothesis.  To understand the metal enrichment of galaxies, it is necessary to study all the gaseous phases involved in the gas mixing cycle. X-ray observations of starburst galaxies led to the conclusion that the hot gas is probably enriched by SNe products (e.g., Martin et al.\\ 2002). The following step in the mixing cycle is poorly constrained observationally. Do metals mix with the cold halo and thus enriching the H\\1\\ region? Do they eventually condense as molecular droplets and settle back onto the stellar body? The situation might be different in BCDs if galactic fountains do not form. However, it seems that winds from massive stars and SNe are still able to carve the ISM to form superbubbles in which metals are injected at high temperatures (Thuan et al.\\ 2004). Mixing could then occur at the interfaces between the hot and cold gas. Thus the study of the metal content in the neutral gas of dwarf galaxies provides us with essential constraints for their chemical evolution models.  The neutral phase usually consists of at least two distinct components (Field et al.\\ 1969), with a warm medium (WNM; $T_w\\sim10^4$\\,K, $n_w\\sim$ few tenths cm$^{-3}$) showing broad hyperfine 21\\,cm line emission, and a cold medium (CNM, $T_c\\sim100$\\,K, $n_c\\sim$ few tens cm$^{-3}$) showing narrow 21\\,cm emission and absorption. It must be noted that the distinction between CNM and WNM becomes indistinct in galaxies with low-pressure ISM, such as dwarfs (Lo et al.\\ 1993). Neutral metals can be observed through resonant lines in the far-ultraviolet (FUV). BCDs are ideal targets because they display large amounts of H\\1\\ gas (Thuan \\& Martin 1981), and because the massive stars provide a strong FUV continuum. Kunth et al.\\ (1994) used the \\textit{GHRS} instrument on board the \\textit{Hubble Space Telescope} to observe interstellar absorption lines along the lines of sight toward the many massive stars in IZw18. The authors detected the O\\1\\ line at 1302.2\\,\\AA\\ and, together with the neutral hydrogen content measured from the 21\\,cm line, estimated the metallicity of the neutral envelope to be at least 10 times lower than that of the ionized gas in the H\\2\\ regions. Later, Thuan \\& Izotov (1997) obtained a \\textit{GHRS} spectrum for SBS\\,0335--052E and found the metallicity of the neutral gas to be similar to that of the ionized gas, about 2\\% solar.  These results remained however inconclusive because of possible saturation effects of the O\\1\\ absorption line (see Pettini \\& Lipman 1995).  A new step forward was achieved with the launch of the \\textit{Far Ultraviolet Spectroscopic Explorer} (\\textit{FUSE}; Moos et al.\\ 2000), which allows the observation of absorption-lines of H\\1\\ together with many metallic species such as N\\1, O\\1, Si\\2, P\\2, Ar\\1, and Fe\\2. The spectral resolution of \\textit{FUSE} ($R\\sim20\\,000$) and its sensitivity made it possible to detect numerous absorption-lines, some of them apparently not saturated. The neutral gas chemical composition was derived in several BCDs, spanning a wide range in ionized gas metallicity, from 1/50 to 1/3 solar, with I\\,Zw\\,18 (Aloisi et al.\\ 2003; Lecavelier des Etangs et al.\\ 2004), SBS\\,0335--052 (Thuan et al.\\ 2005), I\\,Zw\\,36 (Lebouteiller et al.\\ 2004), Mark\\,59 (Thuan et al.\\ 2002), NGC\\,625 (Cannon et al.\\ 2005), and NGC\\,1705 (Heckman et al.\\ 2001). These investigations showed that the neutral gas of BCDs is not pristine, but that it has already been enriched with metals up to an amount of $\\gtrsim1/50$\\,Z$_\\odot$. The second most important result is that the metallicity of the neutral gas is generally lower than that of the ionized gas $-$ except for the two lowest-metallicity BCDs, I\\,Zw\\,18 and SBS\\,0335--052E $-$ implying that the neutral phase has been probably less processed.  The interpretation of these remarkable results has been limited so far by the lack of knowledge about the origin of the neutral gas in BCDs whose absorption is detected in the FUV. Direct comparison with the H\\1\\ gas seen \\textit{via} the 21\\,cm emission is hampered by a bias in selecting out relatively dust-free FUV lines of sight toward the massive stars. Even more importantly, FUV observations only probe gas in front of stellar clusters while radio observations in emission probe the whole system. Many uncertainties also remain concerning the method of deriving column densities from absorption-lines in complex systems. Most \\textit{FUSE} studies of BCDs assumed a single homogeneous line of sight with the exception of Mark\\,59 (Thuan et al.\\ 2002; see also Lebouteiller et al.\\ 2006 for the giant H\\2\\ region NGC\\,604). A more realistic approach would be to consider multiple lines of sight toward massive stars, each line of sight intersecting clouds with possibly different physical conditions (such as turbulent velocity $b$, radial velocity $v$) and chemical properties. If individual absorption components are unresolved, it is almost impossible to tell whether they are saturated, in which case the column density inferred from the global absorption line can be severely underestimated (see e.g., Lebouteiller et al.\\ 2006). Following the suggestion of Kunth \\&\\ Sargent (1986), Bowen et al.\\ (2005) got around the problem by deriving the abundances in the neutral gas of the dwarf spiral galaxy SBS\\,1543+593 along a quasar line of sight. The authors found that the abundances agree with those in the ionized gas of the brightest H\\2\\ region in the arms, contrasting with the \\textit{FUSE} results in BCDs. As far as chemical models are concerned, Recchi et al.\\ (2004) examined the case of I\\,Zw\\,18 for which they found that abundances in the H\\1\\ medium (defined by a temperature lower than $7\\,000$\\,K) are similar to those in the H\\2\\ gas. Hence, although some physics is still missing in the models (radiative transfer, clumpy ISM, dust phase), it seems that \\textit{FUSE} results still cannot be reproduced.  In this paper, we look anew at the problem of how the chemical compositions of the ionized and the neutral gas in BCDs compare by selecting an object with an extremely low H\\1\\ column density. This object has some absorption lines that are weak enough to be safely considered as unsaturated. The object we analyze is Pox\\,36 (also known as ESO\\,572--G\\,034, I\\,SZ\\,63, HIPASSJ1158-19b, or IRAS\\,11564--1844), an emission-line galaxy first studied by Rodgers et al.\\ (1978) and then rediscovered by Kunth et al.\\ (1981). It is characterized by strong optical emission-lines superimposed on a very blue continuum. Its spectral characteristics and morphological properties fit the definition of a BCD galaxy (see e.g., Thuan \\& Martin 1981, Kunth \\& \\\"Ostlin 2000). Its known observational properties are listed in Table\\,\\ref{tab:prop}. In particular, chemical abundances in the ionized gas of the H\\2\\ regions within Pox\\,36 have been derived by several independent studies (e.g., Kunth \\& Sargent 1983; Izotov \\& Thuan 2004). The ionized gas has an oxygen abundance $12+\\log {\\rm (O/H)} \\approx 8.05$ which corresponds to Z$_\\odot$/5. With a H\\1\\ velocity of 1114 km\\,s$^{-1}$ (Meyer et al.\\ 2004), a Hubble constant of 73 km\\,s$^{-1}$\\,Mpc$^{-1}$, its cosmology-corrected distance is 20\\,Mpc.  The {\\it FUSE} observations are presented in Sect.\\,\\ref{sec:observations} and an overview of the \\textit{FUSE} spectrum is given in Sect.\\,\\ref{sec:overview}. We infer the column densities of H\\1\\ and of the metals from the line profile in Sect.\\,\\ref{sec:fitting}. Chemical abundances are derived and discussed in Sect.\\,\\ref{sec:abundances}. Finally, we study the properties of the neutral gas in the current {\\it FUSE} BCD sample to discuss general trends in Sect.\\,\\ref{sec:global}. ",
        "conclusions": "We have analyzed the FUV \\textit{FUSE} absorption spectrum of the BCD Pox\\,36. The following results are found: \\begin{itemize} \\item  We detect the interstellar absorption lines of H\\1, N\\1, O\\1, Si\\2, P\\2, Ar\\1, and Fe\\2. Pox\\,36 appears to be have a particularly low foreground gas content, which allows us to measure column densities with minimal systematic errors due to hidden saturation effects. \\item Metal column densities have been derived from the line profiles. The interstellar H\\1\\ lines are strongly contaminated by stellar absorption. We have modeled the spectral continuum using TLUSTY models. The best model is constrained by the high-order Lyman line Ly$\\epsilon$ and gives a photospheric temperature of $\\approx30\\,000$\\,K (B0 type). After correcting Ly$\\beta$ for stellar absorption, we have used the interstellar damping wings to infer a H\\1\\ column density equal to $20.3\\pm0.4$. \\item The observed FUV continuum level is consistent with $\\sim300$ B0 stars. This represents only the unextincted stars. \\item Abundances have been computed from the column densities using the dominant ionization stage. The large [O/Fe] ratio in the ionized gas of Pox\\,36 is interpreted as caused by iron depletion onto dust grains. In contrast, the neutral gas does not show any depletion, which is attributed to the low dust content along the observed lines of sight. \\item We find that the neutral gas is metal-underabundant by a factor 5 to 8 as compared to the ionized gas of the H\\2\\ regions. This corresponds to a metallicity of $\\approx1/35$\\,Z$_\\odot$ in the neutral gas as compared to $\\approx1/5$\\,Z$_\\odot$ in the ionized gas and imply that the neutral gas is less chemically processed. \\item It is shown that the metal deficiency in the neutral gas can be achieved and maintained if most of the metals released by star-formation episodes mix with the H\\1\\ gas, increasing only slightly its metallicity, while a small fraction (1\\%) could mix locally and enrich greatly the ionized gas. The metals should come from previous bursts which occurred around the same location as the present one. \\item Based on the sample of the seven BCDs studied by FUSE, we have confirmed that there is a systematic underabundance by a factor $\\sim10$ of the neutral gas as compared to the ionized gas in the H\\2\\ regions. Besides, [Ar/H] and [N/H] seem to reach a minimum value for the lowest-metallicity objects ($\\sim1/8$\\,Z\\,$_\\odot$), which could be due to a default enrichment. \\item There is no convincing evidence that the neutral gas in BCDs contains pristine gas pockets, i.e., completely metal-free. However, it could be that the neutral gas of BCDs has not been affected yet by the star-formation episodes, and that its metallicity is that of the IGM. \\end{itemize}"
    },
    "0812/0812.3256_arXiv.txt": {
        "abstract": "Previous fits of sterile neutrino dark matter models to cosmological data assumed a peculiar production mechanism, which is not representative of the best-motivated particle physics models given current data on neutrino oscillations. These analyses ruled out sterile neutrino masses smaller than 8-10~keV. Here we focus on sterile neutrinos produced resonantly. We show that their cosmological signature can be approximated by that of mixed Cold plus Warm Dark Matter (CWDM). We use recent results on $\\Lambda$CWDM models to show that for each mass $\\ge 2$~keV, there exists at least one model of sterile neutrino accounting for the totality of dark matter, and consistent with Lyman-$\\alpha$ \u00a0and other cosmological data. Resonant production occurs in the framework of the $\\nu$MSM (the extension of the Standard Model with three right-handed neutrinos). The models we checked to be allowed correspond to parameter values consistent with neutrino oscillation data, baryogenesis and all other dark matter bounds. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.1253_arXiv.txt": {
        "abstract": "We have performed new Big Bang Nucleosynthesis calculations which employ arbitrarily-specified, time-dependent neutrino and antineutrino distribution functions for each of up to four neutrino flavors. We self-consistently couple these distributions to the thermodynamics, the expansion rate and scale factor-time/temperature relationship, as well as to all relevant weak, electromagnetic, and strong nuclear reaction processes in the early universe.  With this approach, we can treat any scenario in which neutrino or antineutrino spectral distortion might arise.  These scenarios might include, for example,  decaying particles, active-sterile neutrino oscillations, and active-active neutrino oscillations in the presence of significant lepton numbers.  Our calculations allow lepton numbers and sterile neutrinos to be constrained with observationally-determined primordial helium and deuterium abundances.  We have modified a standard BBN code to perform these calculations and have made it available to the community. ",
        "introduction": "There is a new paradigm in Big Bang Nucleosynthesis (BBN) studies which promises enhanced probes of the early universe and a window into new physics.  In the past, BBN predictions have been used to place constraints on the baryon number at three minutes after the Big Bang. This was done by comparing the observationally-inferred primordial light element abundances to abundances predicted by BBN calculations over a wide range of baryon-to-photon ratio values.  With the high precision results of the Wilkinson Microwave Anisotropy Probe (WMAP), however, the baryon-to-photon ratio, $\\eta$, is now independently determined -- at 300,000 years after the Big Bang -- from observations of the cosmic microwave background (CMB) relative acoustic peak amplitudes \\cite{WMAP, WMAP1, 3yrwmap}.  Currently, the WMAP Three Year Mean value for the baryon-to-photon ratio is $\\eta = \\left(6.11\\pm .22\\right) \\times 10^{-10} $.  Future missions ($e.g.$, Planck\\cite{bond}) promise considerably higher precision determinations of $\\eta$.  Since the baryon-to-photon ratio is known independently, and to excellent precision albeit at much later times,  BBN calculations can now be used to probe or constrain new physics or heretofore poorly determined parameters.  For example, we can use BBN predictions to constrain not only the lepton numbers but also the physics behind these lepton numbers .  The existence of a nonzero electron lepton number follows from charge neutrality and the observed proton content of the universe.  The contributions of neutrinos and antineutrinos to the electron, muon, and tau $(e, \\mu, \\tau)$ lepton numbers are not known, since we do not directly observe these relic particles.  The neutrino contribution to the lepton number for a given flavor,  $\\alpha={\\rm e},\\mu,\\tau$, is defined analogously to the baryon-to-photon ratio, $\\eta \\equiv (n_b-n_{\\bar b})/n_\\gamma$, as \\begin{equation} L_{\\nu_\\alpha} \\equiv {{n_{\\nu_\\alpha}-n_{\\bar\\nu_\\alpha}}\\over{n_\\gamma}}, \\label{lepton} \\end{equation} where $n_\\gamma = (2\\zeta(3)/\\pi^2) T^3_\\gamma$ is the proper photon number density at temperature $T_\\gamma$, and $n_{\\nu_\\alpha}$ and $n_{\\bar\\nu_\\alpha}$ are the neutrino and antineutrino number densities.  Observational bounds on the lepton numbers\\cite{abfw, kfs, Kneller:2001cd, abb, wong, dolgov, Simha:2008mt, Cuoco:2003cu, Serpico:2005bc} remain large compared to the values of these that could significantly affect BBN when there is new leptonic sector physics ($e.g.$, sterile neutrinos)\\cite{abfw}.  The neutrino lepton numbers influence BBN and the resulting primordial element abundances in a number of ways\\cite{wfh}.  The energy density in the neutrino sector contributes to the total energy density of the universe which determines the expansion rate.  The expansion rate is crucial to the outcome of BBN because it determines the weak freeze-out temperature which in turn effectively sets the neutron-to-proton ratio and, therefore, the primordial abundances of $^4$He and the other light elements.  Not only is the total number of neutrinos important to the outcome of BBN, but the neutrino distribution functions are key components of the phase space integrals in the weak reaction rates in BBN.  The weak reactions of greatest interest are those that inter-convert neutrons and protons: \\begin{equation} \\nu_e+n\\rightleftharpoons p+e^-, \\label{nuen} \\end{equation} \\begin{equation} \\bar\\nu_e+p\\rightleftharpoons n+e^+, \\label{nuebarp} \\end{equation} \\begin{equation} n \\rightleftharpoons p+e^-+\\bar\\nu_e. \\label{ndecay} \\end{equation} Since the rates for the weak reactions are strongly energy dependent, the energy distributions of the neutrinos and antineutrinos can figure prominently in both the forward and reverse rates in the processes in Eqs.~(\\ref{nuen}), (\\ref{nuebarp}), and (\\ref{ndecay}).  In standard BBN scenarios the neutrino distribution functions are assumed to be thermally-shaped Fermi-Dirac distributions.  However, it is possible that non-thermal neutrino distribution functions arise after the neutrinos decouple from the background plasma around $T \\approx 3\\,{\\rm MeV}$ and during times crucial to BBN.  There are many possible mechanisms that could alter the neutrino spectra. Altered neutrino energy spectra, in turn, could change the resulting primordial element abundances from what one would expect given a particular lepton number.  Neutrino energy spectrum-altering scenarios include, but are not limited to, active-active neutrino oscillations\\cite{dolgov, abb, abfw, wong}, active-sterile neutrino oscillations\\cite{abfw, kfs, sfka, cirelli, FV, fv97}, or particle decay into the neutrino sea\\cite{pastor}.  Moreover, active-sterile neutrino flavor mixing and other mechanisms for creating sterile neutrino dark matter before neutrino decoupling are a focus of current research\\cite{Dodelson:1993je, afp, Dolgov:2000ew, Shaposhnikov:2006xi, Kusenko:2006rh, Petraki:2007gq, Petraki:2008ef, Shi:1998fu, Chiu:1977ds, Boyanovsky:2007zz, Boyanovsky:2006it, Boyanovsky:2007ba, Abazajian:2002yz}, as is the constraint of these scenarios via x-ray observations and large-scale structure considerations\\cite{Abazajian:2001vt, Abazajian:2006yn, Boyarsky:2006fg, Boyarsky:2005us, Yuksel:2007xh, Watson:2006qb, Viel:2005qj, Abazajian:2005xn, Seljak:2006qw}. Though these models may not directly affect BBN through the spectral distortion of $\\nu_e$ and $\\bar\\nu_e$ energy distribution functions discussed here, they nevertheless may affect the overall values of lepton number, entropy, and energy density which are relevant to BBN.  In the end, the existence of sterile neutrino states changes the meaning and utility of lepton number\\cite{Foot:1995qk, Shi:1996ic}. To use BBN predictions to probe or constrain any such scenario requires an approach that self-consistently includes neutrino and antineutrino energy spectra of arbitrary shape.  We have performed detailed  calculations of primordial nucleosynthesis in which we include neutrino and antineutrino spectral distortion.  Our results are surprising.  We find that even modest distortions of the neutrino and/or antineutrino spectral shapes from Fermi-Dirac black body forms can result in significant modification of the net neutron-proton interconversion rates and, hence, alteration of the light element abundances.  To study the effects of neutrino spectral distortion, we have modified the original Kawano/Wagoner BBN code described in Ref. \\cite{skm} to calculate the primordial element abundances self-consistently with arbitrarily-specified non-thermal and/or time-dependent neutrino distribution functions.  This paper is  structured as follows: Section II describes the calculation of weak charge-changing reaction rates in the early universe and our prescription for employing non-thermal neutrino and antineutrino energy distribution functions; Section III discusses our new BBN code; Section IV will present example results for non-thermal neutrino distribution functions resulting from various physical scenarios; and Section V gives conclusions. ",
        "conclusions": "We have developed an approach to Big Bang Nucleosynthesis (BBN) calculations where we can treat arbitrarily-specified energy distributions for all neutrino types, including $\\nu_e$ and $\\bar\\nu_e$.  We can also allow these distribution functions to be altered dynamically and follow all nuclear and weak reactions self-consistently with these alterations. This new approach can extend the usefulness of BBN predictions for exploring and constraining new physics in the neutrino and weak interaction sectors.  Examples of such new physics include active-sterile neutrino mixing and particle decays that have neutrinos in the final state.  We have given an explicit example of the former scenario. In this example we have demonstrated how active-sterile neutrino oscillation physics can alter neutrino or antineutrino distribution functions on short time scales, alter the neutron-proton inter-conversions rates, and so modify BBN abundance yields over those of the standard scenario.  Our calculations hold out the promise that light element abundances could place the best constraints on primordial lepton numbers and active-sterile neutrino mixing parameters when the sterile neutrino mass is in the $\\sim 1\\,{\\rm eV}$ range. Present laboratory experiments, like mini-BooNE, are sensitive to neutrino flavor mixing in the active-sterile channel at the $\\sim 1,{\\rm eV}$ mass scale only when the appropriate effective $2\\times 2$ vacuum mixing angle satisfies $\\sin^22\\theta \\gg {10}^{-4}$. By contrast, in the presence of a net lepton number, BBN abundance yields might be significantly altered for  active-sterile neutrino mixing parameters for $\\sin^22\\theta > {10}^{-8}$. The greater reach in vacuum mixing angle afforded by BBN considerations stems from: (1) the long (gravitational) expansion time scale of the early universe which dictates the MSW resonance sweep rate and sets the minimum mixing angle required for adiabatic and efficient conversion of the active neutrinos into sterile species; and (2) the significant sensitivity of the neutron-proton weak inter-conversion rates to alterations of the neutrino or antineutrino energy distribution functions. Our new calculations allow us to follow simultaneously and self-consistently both of these effects along with all relevant weak, electromagnetic, and strong nuclear reaction rates.  This new approach is incorporated into an update of the Kawano/Wagoner BBN code -- which can now accommodate and integrate any arbitrary neutrino and/or antineutrino distribution function with any specified time dependence. We will soon make this code available to the community at bigbangonline.org."
    },
    "0812/0812.1409_arXiv.txt": {
        "abstract": "{ This paper explores single field inflation models with a constant, but arbitrary speed of sound $c_s$, obtained by deforming the kinetic energy terms to a Dirac-Born-Infeld form. Allowing $c_s<1$ provides a simple parameterization of non-gaussianity. The dependence of inflationary observables on the parameter $c_s$ is considered in the leading order slow roll approximation. The results show that in most cases the dependence is actually rather weak for the range of $c_s$ allowed by existing bounds on non-gaussianity. } ",
        "introduction": "Models of k-inflation \\cite{ArmendarizPicon:1999rj,Garriga:1999vw} have been the subject of much recent interest. These models modify standard single-field inflation by allowing a more general form of the kinetic energy terms. An important feature of k-inflation is the alteration of the speed of propagation of disturbances in the inflaton field -- the speed of sound  $c_s$. Models of inflation with canonical kinetic energy have $c_s=1$, but with a more general form of kinetic energy this is no longer the case -- the speed of sound i a field dependent quantity.  One particularly interesting example of this is  a model which replaces the canonical kinetic energy by the Dirac-Born-Infeld form \\cite{Silverstein:2003hf,Alishahiha:2004eh}, which has come to the fore in recent years in the context of D-brane inflation \\cite{Cline:2006hu}--\\cite{McAllister:2007bg}. It is however also quite natural to view models of this kind in the general context of k-inflation \\cite{Lidsey:2006ia,Spalinski:2007qy}. The DBI form of the kinetic energy terms involves a square root factor, $\\gamma>1$, reminiscent of the Lorentz factor of special relativity. Indeed, the square root is responsible for introducing a ``speed limit'' on the inflaton scalar. This facilitates a form of non-slow-roll inflation \\cite{Alishahiha:2004eh}, on which most effort has focused. For such trajectories of the inflaton the Lorentz factor $\\gamma$ is large. However, it is also interesting to consider what quantitative effect allowing $\\gamma>1$ has in the slow roll regime. In the case of DBI models the speed of sound is expressed in terms of $\\gamma$ as $c_s=1/\\gamma$.  A particularly simple situation arises when the $\\gamma$ factor is constant \\cite{Spalinski:2007un}, which means that the speed of sound is also constant, as in the canonical case, but no longer equal to the speed of light. This case can be considered as a leading approximation in an expansion of the field dependent speed of sound if it is assumed to vary slowly in the relevant region of field space. Some exactly solvable examples of this sort have recently been discussed in \\cite{Spalinski:2007un}--\\cite{Lorenz:2008et}.  From a phenomenological perspective the interest in k-inflationary models stems from the fact that they provide a fairly simple way of accounting for non-gaussianity \\cite{Bartolo:2004if} in the spectrum of density perturbations. So far no conclusive evidence for departures from gaussianity has been observed, but it is of great interest to see if, and what kind of, theoretical possibilities exist, given the high expectations that data which will become available in the coming years will allow for a significant tightening of existing bounds. The simplest measure of non-gaussianity is the parameter $f_{NL}$, which is defined in terms of 3-point functions in \\cite{Chen:2006nt,Creminelli:2006rz} (for example). For DBI models one has the simple relation \\cite{Chen:2006nt} $f_{NL}=35 (\\gamma^2 - 1)/108$, so the value of $\\gamma$ strongly affects the deviations from gaussianity. The current bounds on $f_{NL}$ imply roughly $\\gamma<35$. DBI models with constant speed of sound provide a very simple parameterization of non-gaussianity in models of single field inflation. Even if the observed non-gaussianity turns out to be large, the variation of the speed of sound with scale could be negligible.  This note reports the results of analyzing a series of frequently considered models of slow roll inflation and explores how sensitive their predictions are when one allows a small deviation from the canonical form of the kinetic energy, as measured by a constant $\\gamma>1$. The observables $n_s$ and $r$ have been computed for general models of k-inflation by Garriga and Mukhanov \\cite{Garriga:1999vw}, and they are easily adapted to the case of DBI kinetic energy \\cite{Chen:2006nt}. These formulae are evaluated at the time  when the present Hubble scale crossed the horizon during inflation. The corresponding number of e-folds is denoted by $N_\\star$\\footnote{As in a number of other studies, the number of e-folds is defined as decreasing to 0 at the end of inflation.}. This number is afflicted by some theoretical uncertainties \\cite{Kinney:2005in}, and the possibility of having $\\gamma>1$ adds another ones, as discussed in section \\rf{basics}. The remaining sections discuss a number of popular inflationary models case by case. In most cases, notably chaotic inflation, the results for the inflationary observables do not depend on $\\gamma$, or the dependence is very weak. However in some cases of modular inflation it is found that the tensor fraction $r$ effectively grows with $\\gamma$, so one can envisage that it might become observable due to this effect. ",
        "conclusions": "It is expected that soon the existing bounds on non-gaussianity will be significantly tightened or a measurement of it will be made. In view of this it is important to consider theoretical options which lead to non-gaussian perturbation spectra. One simple possibility was considered here. The conclusion is basically that as long as the speed of sound is constant, the basic predictions for the spectral tilt and tensor fraction are quite robust even if one allows $f_{NL}$ as large as existing limits permit. Some variation of the tensor fraction with $f_{NL}$ is possible however, and while generically $r$ goes down with rising non-gaussianity, in some examples (new inflation with $k>2$) it can increase somewhat.  \\begin{center} {\\bf Acknowledgements} \\end{center}  \\noindent I would like to thank Micha\\l\\ Spali\\'nski for helpful discussions  and pointing out the problem. This work has been supported by a grant PBZ/MNiSW/07/2006/37."
    },
    "0812/0812.1205_arXiv.txt": {
        "abstract": "It is shown that the recently detected acceleration of the universe can be understood by considering a modification of the teleparallel equivalent of General Relativity (TEGR), with no need of dark energy. The solution also exhibits phases dominated by matter and radiation as expected in the standard cosmological evolution. We perform a joint analysis with measurements of the most recent type Ia supernovae (SNe Ia), Baryon Acoustic Oscillation (BAO) peak and estimates of the CMB shift parameter data to constraint the only new parameter this theory has. ",
        "introduction": "The discovery of an unexpected diminution in the observed energy fluxes coming from type Ia supernovae \\cite{perlmutter-otros, kowalski} has opened one of the most puzzling and deepest problems in cosmology today. These observations have been interpreted as solid evidence for an accelerating universe dominated by something called dark energy. Although the cosmological constant seems to be the simplest explanation for the phenomenon, several dynamical scenarios have been tried out since 1998 (see e.g., \\cite{sahni, padma, frieman}). While some authors sustain the idea of the existence of a dark energy, others propose modifications of the Einstein-Hilbert Lagrangian known as $f(R)$ (\\cite{buch, staro, ker, barrow1, barrow2, carroll, starobinsky, hu, odintsov1} or \\cite{odintsov2, capo, sotiriou} for recent reviews) as a way to obtain a late accelerating expansion. A great difficulty these theories have, from the point of view of the metric formalism, is that the resulting field equations are 4th order equations, which in many cases makes these hard to analyze. Besides, the simplest cases of the kind $% f(R)=R-\\beta/R^n$ have shown difficulties with, weak field tests \\cite% {chiba, olmo}, gravitational instabilities \\cite{dolgov} and do not present a matter dominated era previous to the acceleration era \\cite{amendola1, amendola2}. Alternatively, the Palatini variational approach for such $f(R)$ theories leads to 2nd order field equations, and some authors have achieved to put observational constraints to these theories \\cite{amar, fay, santos}. However in many cases the equations are still hard to work with, as evidenced by the functional form of the modified Friedmann equation for a generic $f(R)$. Recently, models based on modified teleparallel gravity were presented as an alternative to inflationary models \\cite{franco,franco2}. In this paper we show a cosmological solution for the acceleration of the universe by means of a sort of theories of modified gravity, namely $f(L_T)$, based on a modification of the teleparallel equivalent of General Relativity (TEGR) Lagrangian \\cite{einstein, hayashi} where the torsion will be the responsible of the observed acceleration of the universe, and the field equations will always be 2nd order equations. ",
        "conclusions": "A theory $f(L_T)$ based on a modification of the teleparallel equivalent of General Relativity (TEGR) -where torsion is the geometric object describing gravity instead of curvature and its equations are always of 2nd order- is remarkably simpler than $f(R)$ theories. We have tested the theory $% f(L_T)=L_T-\\alpha(-L_T)^{-n}$ with the aim of reproducing the recently detected acceleration of the universe without resorting to dark energy. We have here performed observational viability tests for this theory by using the most recent SN Ia data, and combined them with the information coming from BAO peak and CMB shift parameter in order to find constraints in the $% n-\\Omega_m$ plane. At 68.3\\% c.l. we found that the values lie in the ranges $n\\in[-0.23, 0.03]$ and $\\Omega_m\\in[0.25, 0.29]$. The values for $\\Omega_m$ are consistent with recent estimations obtained by other authors (see, e.g., \\cite{spergel}). The model with the best-fit values minimizing the $\\chi^2$ that combines SNIa+BAO+CMB data ($n=-0.10$ and $\\Omega_m=0.27$) exhibits the last three phases of cosmological evolution: radiation era, matter era and late acceleration; this last stage having started at $z_{acc}\\simeq0.74$.  \\bigskip  \\bigskip"
    },
    "0812/0812.4420_arXiv.txt": {
        "abstract": "{Short supersoft X-ray source (SSS) states (durations $\\le $ 100 days) of classical novae (CNe) indicate massive white dwarfs that are candidate progenitors of supernovae type Ia.} {We carry out a dedicated optical and X-ray monitoring program of CNe in the bulge of \\m31.} {We discovered \\nova and determined its optical and X-ray light curve. We used the robotic Super-LOTIS telescope to obtain the optical data and \\xmm and \\chandra observations to discover an X-ray counterpart to that nova.} {Nova \\nova is a very fast CN, exhibiting a very short SSS state with an appearance time of 6--16 days after outburst and a turn-off time of 45--58 days after outburst.} {The optical and X-ray light curves of \\nova suggest a binary containing a white dwarf with $M_{WD} >$ 1.0 \\msun.} ",
        "introduction": "Classical novae (CNe) are thermonuclear explosions on the surface of white dwarfs (WDs) in cataclysmic binaries that result from the transfer of matter from the companion star to the WD. The transferred hydrogen-rich matter accumulates on the surface of the WD until hydrogen ignition starts a thermonuclear runaway in the degenerate matter of the WD envelope. The resulting expansion of the hot envelope causes the brightness of the WD to rise by a typical outburst amplitude of $\\sim$ 12 magnitudes within a few days, and mass to be ejected at high velocities \\citep[see][and references therein]{2005ASPC..330..265H,1995cvs..book.....W}. However, a fraction of the hot envelope can remain in steady hydrogen burning on the surface of the WD \\citep{1974ApJS...28..247S,2005A&A...439.1061S}, powering a supersoft X-ray source (SSS) that can be observed directly once the ejected envelope becomes sufficiently transparent \\citep{1989clno.conf...39S,2002AIPC..637..345K}. In this paper, we define the term ``appearance of the SSS'' from an observational point of view as the time when the SSS becomes visible to us, due to the decreasing opacity of the ejected material. This is not the same as the actual turn-on of the H-burning on the WD surface at the beginning of the thermonuclear runaway. The turn off of the SSS, on the other hand, does not suffer from extinction effects and clearly marks the actual turn off of the H-burning in the WD atmosphere and the disappearance of the SSS.  The duration of the supersoft phase is related to the amount of H-rich matter that is {\\it not} ejected and on the luminosity of the white dwarf. More massive WDs need to accrete less matter to initiate the thermonuclear runaway, because of their higher surface gravity \\citep{1998ApJ...494..680J}. In general, more massive WDs retain less accreted mass after the explosion, although this also depends on the accretion rate, and reach higher luminosities \\citep{2005ApJ...623..398Y}. Thus, the duration of the SSS state is inversely related to the mass of the WD. In turn, the transparency requirement mentioned above implies that the appearance time is determined by the fraction of mass ejected in the outburst \\citep[see][]{2005A&A...439.1061S,1998ApJ...503..381T,2006ApJS..167...59H}. X-ray light curves therefore provide important clues to the physics of the nova outburst, addressing the key question of whether a WD accumulates matter over time to become a potential progenitor for a type Ia supernova (SN-Ia). The duration of the SSS state provides the only direct indicator of the post-outburst envelope mass in CNe. For massive WDs, the expected SSS phase is very brief ($<$ 100 d) \\citep{2005A&A...439.1061S,1998ApJ...503..381T} and could have easily been missed in previous surveys \\citep[e.g.][]{2001A&A...373...63S,2004ApJ...609..735W,2005A&A...434..483P}.  There are only a few novae with very short SSS states known. One of them is the well-studied galactic nova RS~Oph. RS~Oph is a recurrent nova (RN), which underwent its last outburst in 2006. X-ray observations documented a rapid decline in supersoft emission at $\\sim 90$ days after the optical outburst and a duration of the SSS phase of $\\sim 60$ days \\citep{2006ATel..838....1O,2007ApJ...659L.153H}. From theoretical models, \\citet{2007ApJ...659L.153H} estimated a WD mass of $(1.35\\pm0.01)$\\msun. RNe are classified by the observational fact that they have had more than one recorded outburst. These objects are very good candidates for SN-Ia progenitors, as RNe are believed to contain massive WDs. However, one of the problems with this RNe $-$ SN-Ia connection was their rare occurrence in optical surveys \\citep{1996ApJ...473..240D}. But hypothetical previous outbursts may have been missed, and then these objects were just classified as CNe in optical surveys. Interestingly, a very short SSS stage of a CN may indicate a very massive WD, which is capable of frequent outbursts, and thus is a potential RN.  Another example for a nova with short SSS state is the CN Nova LMC 2000, for which \\citet{2003A&A...405..703G}, based on non-detections of supersoft emission, suggest an SSS phase shorter than seven weeks. The work of \\citet{2007A&A...465..375P} already discusses nova M31N~2004-11f as an object with a very short SSS state of 35--55 days. This nova is an RN.  This work presents the results of a new monitoring strategy tailored to CNe with short SSS states. We carried out combined \\xmmk/\\chandra monitoring of the central region of \\m31 \\citep[distance 780 kpc,][]{1998AJ....115.1916H,1998ApJ...503L.131S} from November 2007 to mid-February 2008. Individual observations were separated by just 10 days. Section \\ref{sec:obs} provides detailed information on the X-ray and optical data sets. Results are presented in Sect.\\,\\ref{sec:results} and discussed in Sect.\\,\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss} From the positional and temporal agreement and from the HRC-I hardness ratios, we infer that \\source and \\nova are the same object.  The ejected mass in the nova outburst can be estimated from the start date of the SSS phase and from the expansion velocity of the ejected material. The decrease in the optical thickness of the expanding ejecta is responsible for the rise in the X-ray light curve of the post-outburst novae, as shown for V1974 Cyg \\citep{1996ApJ...463L..21S,1996ApJ...456..788K}. Assuming that the material ejected by the nova explosion forms a spherical, homogeneous shell expanding at constant velocity $v$, the hydrogen mass density of the shell will evolve in time $t$ as $\\rho=M^{ej}_{H} / (\\frac{4}{3}\\pi v^{3}t^{3})$ where $M^{ej}_{H}$ is the ejected hydrogen mass \\citep{1996ApJ...456..788K}. The column density of hydrogen evolves with time as $N_{H} ({\\rm cm}^{-2})= M^{ej}_{H} / (\\frac{4}{3}\\pi u v^{2}t^{2})$, where $u=1.673\\times10^{-24}$ g is the atomic mass unit. We assume a typical value for the expansion velocity of 2000 km s$^{-1}$, since we do not have an optical spectrum for \\nova, and that the SSS turns on when the absorbing hydrogen column density decreases to $\\sim10^{21}$ cm$^{-2}$. With this, the appearance time of the SSS source (between 6 and 16 days) constrains the ejected mass to the range $(0.4-3)\\times10^{-7}$\\msun. This is about two orders of magnitude lower than typical ejected masses determined in a similar way for M31 novae \\citep{2007A&A...465..375P}. However, the unknown expansion velocity error adds additional uncertainties to the estimate of $M^{ej}_{H}$. Figure\\,\\ref{fig:param} shows the range of $M^{ej}_{H}$ for different expansion velocities as a function of the SSS appearance time.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{1581fig2.eps}} \\caption{Relation of ejected mass (in solar units) and SSS appearance time for different expansion velocities $v$ (see Sect.\\,\\ref{sec:discuss}).} \\label{fig:param} \\end{figure}  The turn-off time of the SSS phase allows us to estimate the amount of hydrogen-rich material burned on the WD surface, $M^{burn}=(L\\Delta t) / (X_H \\epsilon)$, where $L$ is the bolometric luminosity, $\\Delta t$ the duration of the SSS phase, $X_H$ the hydrogen fraction of the burned material, and $\\epsilon=5.98\\times10^{18}$ erg g$^{-1}$. We compute the burned mass during the short SSS phase of \\nova, assuming a bolometric luminosity of $3\\times10^4L_{\\odot}$ and a hydrogen mass fraction $X_H=0.5$. The SSS turn off at 45--58 days after the optical outburst constrains the burned mass to the range $(8-10)\\times10^{-8}$\\msun. It is noteworthy that this is within a factor 2--3 the same as the ejected mass. We used typical values for $L$ and $X_H$, since we do not know the actual values for \\nova.  \\nova has the shortest SSS phase observed so far, which can be caused either by a WD with $M_{WD} >$ 1.1 \\msun for a standard hydrogen fraction in the envelope or by very hydrogen-poor envelope in a WD with $M_{WD} \\sim$ 1.0 \\msun. The mass of the stable H-burning envelope decreases with increasing WD mass, with massive WDs showing the shortest SSS phases. Massive WDs also eject less massive shells during outburst, leading to a faster rise in the SSS emission.  From the optical light curve we find a decline rate of $\\sim 0.6$ mag/d. Extrapolating the magnitude to the day before the first observation (MJD = 54405.3) leads to an estimated maximum magnitude of $m_R \\sim 16.1$, which implies an absolute magnitude at maximum of $M_R \\sim -8.4 \\pm 0.4$ at the distance of \\m31. Within the uncertainties, this value matches the maximum magnitude versus rate of decline (MMRD) relationship obtained by \\citet{1995ApJ...452..704D}. Assuming that the strength of a nova outburst (for a given temperature of the WD and accretion rate) mostly depends on the mass of the WD \\citep[see e.g.] [and references therein]{1992ApJ...393..516L,1995ApJ...452..704D,2002AIPC..637..443D}, one can estimate WD masses from optical nova light curves. However, \\citet{1992ApJ...393..516L} mentions that great caution should be exercised when applying his maximum magnitude vs. WD mass relation to individual novae, since there are weaker dependences of the nova luminosity on the magnetic field of the WD and the enrichment of the envelope by heavy elements, which are neglected in his formula. Using our optical light curve, we compute that\\nova likely originated from a WD with $M_{WD} \\sim$ 0.9--1.15 \\msun. This result is consistent with the X-ray observations.  As mentioned before, many of the novae with known short SSS phases turned out to be RNe. To check for the possibility that \\nova is an RN, we looked for historical novae with positions close to \\novak. In the compiled on-line nova catalog of \\citet{2007A&A...465..375P}\\footnote{http://www.mpe.mpg.de/$\\sim$m31novae/opt/m31/index.php} there are two historical novae within a radius of $\\sim 20 \\arcsec$ around \\novak. \\citet{2000AstL...26..433S} found a nova (M31N~1998-07n) at a distance of about $15 \\arcsec$, and \\citet{1929ApJ....69..103H} reported a nova (M31N~1920-10a) at a distance of about $20 \\arcsec$ to \\novak. The J2000 coordinates of M31N~1920-10a were taken from the Combined General Catalogue of Variable Stars \\citep{2004yCat.2250....0S}. From the finding chart given by \\citet{2000AstL...26..433S}, nova M31N~1998-07n and \\nova do not seem to be identical. This agrees with the position accuracy of $0\\,\\farcs5$ given by \\citet{2000AstL...26..433S}. For M31N~1920-10a no finding chart exists and the position given by \\citet{1929ApJ....69..103H} has a minimum accuracy of $0\\,\\farcm1$. The actual unknown position accuracy of historical novae leaves the (small) possibility open that both novae in fact originated on the same WD. A relatively large position error for a historical nova could be possible, since e.g. the re-analysis of photographic plates by \\citet{2008A&A...477...67H} gives several examples for historical novae with position errors of $\\sim 20 \\arcsec$ or more. Clarification of this point could only come from re-examination of the original photographic plates used in the work of \\citet{1929ApJ....69..103H}."
    },
    "0812/0812.1033_arXiv.txt": {
        "abstract": "{ The existence of strong lensing systems with Einstein radii covering the full mass spectrum, from $\\sim 1-2\\arcsec$ (produced by galaxy scale dark matter haloes) to $>10\\arcsec$ (produced by galaxy cluster scale haloes) have long been predicted. Many lenses with Einstein radii around $1-2\\arcsec$ and above $10\\arcsec$ have been reported but very few in between. In this article, we present a sample of 13 strong lensing systems with Einstein radii in the range $3\\arcsec\\,-\\,8\\arcsec$ (or image separations in the range $6\\arcsec\\,-\\,16\\arcsec$), i.e. systems produced by \\emph{galaxy group scale dark matter haloes}. This group sample spans a redshift range from 0.3 to 0.8. This opens a new window of exploration in the mass spectrum, around 10$^{13}$-\\,10$^{14}$\\,M$_{\\sun}$, a crucial range for understanding the transition between galaxies and galaxy clusters, and a range that have not been extensively probed with lensing techniques. These systems constitute a subsample of the Strong Lensing Legacy Survey (SL2S), which aims to discover strong lensing systems in the Canada France Hawaii Telescope Legacy Survey (CFHTLS). The sample is based on a search over 100 square degrees, implying a number density of $\\sim$\\,0.13 groups per square degree. Our analysis is based on multi-colour CFHTLS images complemented with \\emph{Hubble Space Telescope} imaging and ground based spectroscopy. Large scale properties are derived from both the light distribution of elliptical galaxies group members and weak lensing of the faint background galaxy population. On small scales, the strong lensing analysis yields Einstein radii between 2.5$\\arcsec$ and 8$\\arcsec$. On larger scales, strong lens centres coincide with peaks of light distribution, suggesting that light traces mass. Most of the luminosity maps have complicated shapes, implying that these intermediate mass structures may be dynamically young. A weak lensing signal is detected for 6 groups and upper limits are provided for 6 others. Fitting the reduced shear with a Singular Isothermal Sphere, we find $\\sigma_{\\rm SIS}\\,\\sim$ 500\\,km\\,s$^{-1}$ with large error bars and an upper limit of $\\sim$\\,900\\,km\\,s$^{-1}$ for the whole sample (except for the highest redshift structure whose velocity dispersion is consistent with that of a galaxy cluster). The mass-to-light ratio for the sample is found to be M/L$_i$\\,$\\sim$ 250 (solar units, corrected for evolution), with an upper limit of 500. This compares with mass-to-light ratios of small groups (with $\\sigma_{\\rm SIS} \\sim$ 300 km\\,s$^{-1}$) and galaxy clusters (with $\\sigma_{\\rm SIS} >$ 1\\,000 km\\,s$^{-1}$), thus bridging the gap between these mass scales. The group sample released in this paper will be complemented with other observations, providing a unique sample to study this important intermediate mass range in further detail. } ",
        "introduction": "\\subsection{Galaxy Groups \\& Cosmology} Galaxy groups are believed to play a key role in the formation and evolution of structures in the Universe. They contain the majority of the galaxies in the Universe, and within a hierarchical framework, they span the regime between individual galaxies and massive galaxy clusters, making them cosmologically significant. They are also more varied in their properties than galaxy clusters, as demonstrated by comparisons between various scaling relations in galaxy groups to those in galaxy clusters. This indicates that galaxy groups are probably not a homogeneous class of objects, e.g. they cannot be considered simple scaled-down versions of galaxy clusters. Detailed studies of this intermediate regime of the mass spectrum ($\\sim$10$^{13}$-\\,10$^{14}$ \\Msun) are needed to understand better the physical processes behind formation of galaxy groups and how galaxy groups participate to structure formation and evolution.  Properties of low and intermediate redshift galaxy groups ($z<0.5$) have been studied using X-ray and optical tracers \\citep[\\emph{e.g.}][and references therein]{group1,group2,group3,group4,group5,group6,mamongroup,group7,fabio,group8,group9,sun}. Galaxy groups have also been studied numerically \\citep[\\emph{e.g.}][and references therein]{elena05,jespergroup}.  \\subsection{Strong Lens Statistics} Fig.~\\ref{Fig9oguri06} shows the predicted contributions of different types of haloes in the image separation distributions \\citep[from][Fig.~9]{oguri06} extracted from N-body simulations, where $\\theta$ corresponds to about twice the Einstein radius of the lens\\footnote{In this paper, we will use the Einstein radius \\R\\, to characterise the lenses.}. This plot shows that there are large overlaps among the halo types, but also that some scales are dominated by different regimes: on small scales, the distribution is dominated by galaxy-scale dark-matter haloes, with \\R\\,$<$\\,2$\\arcsec$, and on larger scales, by cluster-scale dark-matter haloes, with \\R\\,$>$\\,10$\\arcsec$. In between, the distribution is dominated by haloes generating strong lensing deflection of radii between $\\sim$\\,3$\\arcsec$ and $\\sim$\\,8$\\arcsec$, whose mass is in the range 10$^{13}$-\\,10$^{14}$\\,M$_{\\sun}$, i.e. by \\emph{group-scale dark-matter haloes}.  Hence according to simulations, strong lenses from intermediate mass systems should exist in a $\\Lambda$CDM Universe. Until now, only a few cases have been detected observationally in the SDSS, which we discuss in Section~\\ref{cassowary}.  \\paragraph{Strong Lensing Deflector:} In this paper, we refer to a ``deflector'' as the foreground mass distribution giving rise to multiple images, and we focus on its central part. Specifically, a strong lensing deflector refers to the area of the foreground mass distribution enclosed by the multiple images. This region can be populated by more than a single galaxy on scales smaller than \\R, as it is the case for some of the galaxy groups reported in this work. In order to avoid confusion, we use the term deflector instead of lensing galaxy. This is particularly important for group scale deflectors since their mass may no longer be associated to a single galaxy. This definition is well adapted to most of the groups studied in this paper, where the multiple images identification is clear and therefore the region enclosed by these images is well defined. However, this is not the case for three of the groups presented in this paper.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.65]{11473fg1.jpg} \\caption{The contributions of different types of haloes on the image separation distribution, from \\citet{oguri06}: galaxy (dashed); groups (dot) and clusters (dot-dashed). The sum of distributions of three types is shown by thick solid line.} \\label{Fig9oguri06} \\end{center} \\end{figure}  \\subsection{Strong Lensing \\emph{in} and \\emph{by} a Galaxy Group} We shall emphasise the distinction between strong lensing \\emph{in} a galaxy group and \\emph{by} a galaxy group. Strong lensing in a galaxy group corresponds to the observation of a strong lensing feature, with a typical Einstein radius of $\\sim2\\arcsec$, generated by a galaxy group member. Since the high-density environment is likely to enhance the strong galaxy-galaxy lensing cross section \\citep[see][but see \\citet{cecile} in the \\textsc{cosmos} field]{kovner,ole,fassnacht06,oguri05,king1689,slacsclustering}, many strong lensing systems in groups (or clusters) have been reported \\citep[\\emph{e.g.}][]{kundic97a,fassnacht02,morgan05,williams06,momcheva,auger07,ring1689,auger08,grillo12,treuaroundslacs,inada}. If the creation of multiple images is boosted by external shear and convergence from a smooth group component, these strong lensing systems are, to first order, generated by a galaxy-scale dark-matter halo populated by a single galaxy \\citep[most of the time an elliptical galaxy, but note that efforts are currently underway to find lenses generated by haloes populated by spiral galaxies, see][]{chloe,trott}.  On the other hand, strong lensing by a galaxy group corresponds to the observation of strongly lensed features with a typical Einstein radius between $\\sim 3\\arcsec$ and $\\sim 8\\arcsec$. This translates into a projected mass enclosed within this radius around $\\sim$ 10$^{12.5}$\\,-\\,10$^{13}$\\,\\Msun, i.e. characteristic of a \\emph{group-scale dark-matter halo}. Note that the highest limit of 8$\\arcsec$ adopted here is somewhat arbitrary and may be considered by others as a characteristic size of a poor cluster.  \\subsection{Galaxy Groups in the Strong Lensing Legacy Survey}  A straightforward method for finding lens systems is to identify them directly within existing multi-band wide-field imaging surveys. For this purpose, we have started to explore systematically the Canada France Hawaii Telescope Legacy Survey (CFHTLS), using a dedicated automatic search procedure optimised for the detection of arcs. This subsample of the CFHTLS forms the Strong Lensing Legacy Survey \\citep[SL2S,][]{sl2s}. This sample is not biased towards high masses by any prior X-ray selection and because of the wide area covered, cosmic variance is neglected and the resulting sample of strong lenses is assumed to be representative of what can be found in the Universe. However, selection biases inherent to our detection technique do exist and will be addressed in a forthcoming publication using numerical simulations. The SL2S sample probes a wide range of masses and a wide range of redshifts for the lenses. Thanks to unsurpassed combined depth, area and image quality of the CFHTLS, we have uncovered a new population of lenses, with Einstein radii between $\\sim 3\\arcsec$ and $\\sim 8\\arcsec$, i.e. generated by \\emph{galaxy-group scale-dark matter haloes}. This new population effectively bridges the gap between single galaxies and massive clusters, opening a \\emph{new window of exploration in the mass spectrum}. The purpose of this paper is to present the first representative sample of such lenses. Note that this sample is not statistically complete since the survey is not complete yet and that we do not understand properly our selection function.  This paper is organised as follows: data used in this work are presented in Section 2. Sections 3, 4 and 5 present the methodology of the systematic analysis we apply on each group. Results are presented in Section 6 and discussed in Section 7. In the Appendix, each group is described individually, with some images. All our results are scaled to a flat, $\\Lambda$CDM cosmology with $\\Omega_{\\rm{M}} = 0.3, \\ \\Omega_\\Lambda = 0.7$ and a Hubble constant \\textsc{H}$_0 = 70$ km\\,s$^{-1}$ Mpc$^{-1}$. All images are aligned with WCS coordinates, i.e. north is up, east is left. Magnitudes are given in the AB system. ",
        "conclusions": "In this Section, we discuss the results obtained on the group sample. Our statements are clearly limited by the large error bars obtained on the quantities of interest, in particular on the weak-lensing mass estimates. Therefore, part of the following discussion is qualitative only.  \\subsection{Mass is Traced by Light} The strong lensing system is always found to coincide with the peak of the light distribution, as can be seen in Fig.~\\ref{all} (see, however, the case of SL2S\\,J08544-0121 discussed in the Appendix). Because a strong lensing system is always associated with a significant mass concentration (occurrence of strong lensing implies $\\Sigma\\,>\\,\\Sigma_{\\rm critic}$), this suggests that the light is a good tracer of the mass and this justifies the choice of the strong lensing system as the group centre when we perform a weak lensing analysis. We note that this statement is limited to the fact that the centres of the luminosity and mass distributions are located approximately at the same position. We cannot say much at this point because of the limitations of the weak lensing analysis (Section~6.3).  \\subsection{A Homogeneous Sample ?}  The sample is homogeneous in term of luminosities ranging between 1 and 4 10$^{12}$\\,L$_{\\sun}$ (Fig.~\\ref{reslum2}). However, Fig.~\\ref{all} shows that the luminosity contours have very different shapes from one group to another and that all but two (SL2S\\,J09413-1100 at $z=0.38$ and SL2S\\,J02254-0737 at $z=0.51$) have complicated shapes, suggesting that these intermediate mass structures may be dynamically young.  \\subsection{Evolution with Redshift ?}  We looked at possible evolution of the mass-to-light ratio with redshift in Fig.~\\ref{msurl}. However, given the large error bars on this quantity, we cannot reliably infer any trend.  \\subsection{Comparison to Other Studies} There have been at least two different estimations of the mass-to-light ratio of galaxy groups from weak lensing. The first measurement of the average mass-to-light ratios was performed by \\citet{hoekstra01} in the CNOC2 survey. They stacked the weak lensing signal obtained for 50 groups that have been selected on the basis of a dynamical analysis of group candidates. Their weak lensing detection points towards much less massive systems than the one presented in this paper: their best-fit singular isothermal sphere yields a velocity dispersion of 274$^{+48}_{-59}$ km\\,s$^{-1}$. However, they find a mass-to-light ratio of 254 $\\pm$ 110\\,$h$ (solar units, B band rest frame corrected for evolution).  Subsequently, another similar study on a sample of 116 CNOC2 galaxy groups was performed by \\citet{parker05}. Given more statistics, they were able to split the group sample into subsets of poor and rich groups. Their poor galaxy groups were found to have an average velocity dispersion of 193 $\\pm$ 38 km\\,s$^{-1}$ and a mass-to-light ratio of 134 $\\pm$ 26\\,$h$ (solar units, B band rest frame), while their rich galaxy groups have a velocity dispersion of 270 $\\pm$ 39 km\\,s$^{-1}$ and a mass-to-light ratio of 278 $\\pm$ 42\\,$h$, thus finding evidence for a steep increase in the mass-to-light ratio as a function of mass.  The groups presented in the present work are more massive than the rich galaxy group sample from \\citet{parker05}. Given the error bars obtained on the mass-to-light ratios, it is difficult to reliably compare our findings to those from \\citet{parker05}. However, considering the 6 groups for which we achieved a weak lensing detection, we get a mean velocity dispersion of 623 km\\,s$^{-1}$ and a mean mass-to-light ratio of 242, consistent with the results of \\citet{parker05}. Note that these values are close to mass-to-light ratios of galaxy clusters \\citep[\\emph{e.g.}][]{carlberg97,bardeau07}.  Our results suggest that after the steep increase in mass-to-light ratio as a function of mass pointed out by \\citet{parker05} which seems to occur at velocity dispersions around 200\\,km\\,s$^{-1}$ (or mass scale of order $\\sim$10$^{13}$ M$_{\\sun}$), mass-to-light ratios reach a plateau and do not evolve much from the mass scale of groups with velocity dispersions around 300\\,km\\,s$^{-1}$ up to rich galaxy clusters with velocity dispersion larger than 1\\,000\\,km\\,s$^{-1}$. The group sample studied in this article therefore helps to bridge the gap between estimates of the mass-to-light ratio of galaxy groups with velocity dispersions in the range 200\\,-\\,300 km\\,s$^{-1}$ and for galaxy clusters.\\\\  \\subsection{The Naked Cusp Sample} Eight galaxy groups reported in this article exhibit typical cusp configurations, with a clear tangential arc on one side of the deflector. For 5 of these groups, we have not been able to find the expected counter image which is usually located on the other side of the deflector, even when HST imaging was available. This incomplete cusp configuration is called 'naked cusp' \\citep[see, \\emph{e.g.}][]{bible,oguri04}. In two cases, SL2S\\,J02215-0647 and SL2S\\,J08544-0121, the counter image of the cusp is clearly identified on the HST images. In SL2S\\,J22133+0048, the detection of the counter image is tentative only. Moreover, note that the case of SL2S\\,J08544-0121 is a bit special since this cusp configuration is not located at the centre of the total light distribution, whereas the other cusp groups are. To summarise, we find that most of the cusp groups are naked. It may be that the current observations did not allow us to find the counter images, but this naked state of a cusp can be linked with lens statistic issue. If CDM haloes have central density slopes $<$~1.5, \\citet{oguri04} predict that a significant fraction ($>20$\\%) of large separation lenses should have naked cusp image configurations. Therefore, our results could suggest that the groups studied in this work may have shallow central density slopes.  \\subsection{Other Intermediate Separation Lenses} \\label{cassowary} We claim this work to be the first representative sample of strong lensing systems generated by group scale dark matter haloes, with intermediate image separation (in the range $\\sim 3\\arcsec-\\,8\\arcsec$). However, at least five intermediate separation lenses have been reported so far, even though the authors did not always consider the group-scale haloes as being responsible for the image separation, and are interpreting the results in terms of massive galactic scale haloes. A remarkable example is the so called Cosmic Horseshoe \\citep{belokurov07,dye08}, an almost complete Einstein ring of radius 5$\\arcsec$ from a luminous red galaxy at $z=0.44$. The authors argued that this is the most massive galaxy lens discovered so far, with a mass enclosed within its Einstein radius of 5\\,10$^{12}$M$_{\\sun}$, making this galaxy as massive as the entire Local Group. We rather suggest that this lensing system is simply produced by a \\emph{group-scale} dark-matter halo, whose centre is populated by the luminous red galaxy. Indeed, this galaxy is the brightest object in the group of $\\sim$\\,26 members as revealed by the SDSS photometry. As suggested by the nearly perfect circle outlined by the ring, no external shear was required in the model of the system, which is not surprising because the deflector traces the centre of the group halo. This important discovery led to a more systematic data mining looking for large separation lenses in the Sloan Digital Sky Survey: the CAmbridge Sloan Survey Of Wide ARcs in the skY (CASSOWARY\\footnote{http://www.ast.cam.ac.uk/research/cassowary/}). Recently, \\citet{belokurov08} reported two new large separation lenses, with Einstein radius equal to 4$\\arcsec$ and 8$\\arcsec$. More recently, \\citet{lin}, in a systematic search for bright arcs in the SDSS, reported a strongly lensed $z=2$ galaxy by a deflector whose Einstein Radius equals 3.82$\\arcsec$. We also mention an intermediate separation lens (\\R $\\sim$ 2.5$\\arcsec$) found in the CLASS survey, B2108+213 \\citep{mckean,anu08}.  We expect to find new intermediate separation lenses on completion of the CFHTLS, which, in combination with group scale lenses from other promising dedicated searches in surveys such as the SDSS\\footnote{one month after this work was submitted, \\citet{kubogroups} reported the finding of five new group scale lenses in the SDSS}, will increase the statistics of intermediate separation lenses, thus allowing us to get insights on this intermediate mass range."
    },
    "0812/0812.3941_arXiv.txt": {
        "abstract": "We analyze the effect of weak field gravitational waves on the timing of pulsars, with particular attention to gauge invariance, that is, to the effects that are independent of the choice of coordinates. We find: (i)~the Doppler shift cannot be separated into gauge invariant gravitational wave and kinetic contributions; (ii)~a gauge invariant separation {\\it can} be made for the time derivative of the Doppler shift in which the gravitational wave contribution is directly related to the Riemann tensor, and the kinetic contribution is that for special relativity; (iii) the gauge dependent effects in the Doppler shift play no role in the program of gravitational wave detection via pulsar timing. The direct connection shown between pulsar timing and the Riemann tensor of the gravitational waves will be of importance in discussions of gravitational waves from alternative (non-Einsteinian) theories of gravitation. ",
        "introduction": "\\label{sec:intro}   The possibility of detecting gravitational waves through pulsar timing, first suggested independently by \\cite{sazhin78} and \\cite{Detweiler79}, is of increasing interest as part of the birth of gravitational wave astronomy, and the details of this technique continue to be advanced by several researchers \\citep{HD83,backerhellingsAnnRev86,hellingsAJ86,ktr94,LB01,JHLM05,Hobbs08}. The basis of this detection technique is the effect that a gravitational wave has on the arrival times of pulsar signals.  In the analysis of this effect \\citep{hellingsPRD81,backerhellingsAnnRev86}, the gravitational wave has been described in terms of perturbations $h_{\\mu\\nu} \\equiv g_{\\mu\\nu}-\\eta_{\\mu\\nu} $ of the spacetime metric $g_{\\mu\\nu}$. This description of gravitational waves is analogous to describing electromagnetic waves using the potentials $\\{\\Phi,{\\bf A} \\}$. In particular, coordinate transformations are analogous to the gauge transformations of electromagnetic theory, and -- at least in linearized general relativity -- are also called gauge transformations.  If electromagnetic fields are to be described in terms of potentials, then one must choose a particular gauge for the description. Similarly in relativistic gravity, coordinate conditions (``gauge fixing'') must be added.  The need for a gauge choice gives rise to the possibility of confusing a gauge effect with a physical effect.  This confusion, for relativistic gravity, kept it a controversial question for many years whether gravitational waves were physical, or were coordinate waves \\citep{KEN07}. A further source of confusion has been the distinction between the Doppler shift due to gravitational waves, and the Doppler shift due to relative motion of the emitting pulsar and the receiving telescope.  This has the potential to be particularly confusing, since the relative motion is a manifestation of spacetime curvature, as is the gravitational wave.  Our purpose here is to clarify the description of pulsar timing, by analyzing pulsar timing with particular attention to what aspects of the pulsar Doppler shift are gauge dependent and what aspects are physical. What we will show is that pulsar timing can be analyzed in a way in which there is an incontrovertible contribution due to gravitational waves, and a contribution due to relative motion of emitter and receiver. The result is couched completely in terms of quantities that are intuitively appealing, as well as mathematically well defined. The gravitational wave contribution will be based on the Riemann tensor which is gauge invariant very much as are the electric and magnetic fields in classical electromagnetism.  We will see, moreover, that in the literature there are claims about Doppler shift effects that are not gauge invariant, but that these statements refer to intermediate steps in the analysis of pulsar timing and do not affect the overall program of pulsar timing. The technique of gravitational wave detection by pulsar timing is physically valid; it is cleanly distinguishable both from coordinate effects and from source/receiver motions.  Besides clarifying what is and is not physically meaningful in pulsar timing, this paper helps to establish a firm background for studying the nonstandard polarization modes of alternative theories of gravitation, modes for which pulsar-timing detection may be particularly useful. Such modes are most easily described in terms of components of the Riemann tensor, not in terms of metric perturbations.   The paper is organized with Sec.~\\ref{sec:gaugeinv} as its mathematical heart.  In that section we derive the expressions that form the basis of our analysis, and that show clearly what is and what is not gauge invariant in pulsar timing.  To help focus on the distinction between kinematic and gravitational effects, in Sec.~\\ref{sec:gaugeinv} we assume that the motions of the emitter and the receiver of pulses are driven by nongravitational forces. We remedy this unphysical assumption in Sec.~\\ref{sec:pngravity}, where we add gravity as the source of the astrophysical motions of the emitter and receiver.  In Sec.~\\ref{sec:conc} we discuss the implications of these results for the way in which pulsar-timing gravitational wave detection is to be carried out, and is to be viewed. In the Appendix we fill in some details of Sec.~\\ref{sec:gaugeinv} that we thought would divert too much attention from the main point of that section.   We have chosen to present these results with a minimum of unnecessary generality that would make the mathematics more elegant, but more obscure.  Except as noted, we use the notation and conventions of the text by \\cite{MTW}. ",
        "conclusions": "\\label{sec:conc}  A physical measurement, such as the Doppler shift of a photon, is not changed by the changes in our mathematical choices, so the gauge invariance of the expression in Eq.~(\\ref{DoppwithU}) is simply a check of consistency. What {\\it is} subject to our mathematical choices is the interpretation of the terms in Eq.~(\\ref{DoppwithU}), in particular the interpretation of the integral term as the gravitational wave contribution. This interpretation is clearly incorrect since we could, for example, choose coordinates that makes the integral vanish for any photon, and put all of the Doppler shift into the kinetic term.  Though integrals like that in Eq.~(\\ref{DoppwithU}) have appeared in the literature \\citep{Detweiler79,MAS82,BCR83} in discussions of pulsar timing, in practice, this causes no real difficulty.  The Doppler shift can be understood to be the time integral of the gauge invariant expression in Eq.~(\\ref{dDoppdt}).  That time integral will contain an integration constant that cannot be meaningfully separated into gravity and kinematics. The observed phase of the arriving pulses, the raw data of pulsar timing, will be the next time integral, and will contain a term $A+Bt $, where $A$ and $B$ are such integration constants. These integration-constant terms play no role in the actual analysis of timing residuals. The program of gravity wave detection by pulsar timing, therefore, has a solid physical foundation based on the Riemann tensor.   This relationship of timing residuals to the Riemann tensor has a useful secondary benefit. As pointed out in the introduction, pulsar timing has the potential to be a particularly sensitive probe of non-Einsteinian polarization modes of gravitational waves \\citep{leejenetprice}. The description of these modes uses the components of the Riemann tensor. The expression in Eq.~(\\ref{dDoppdt}) gives a clear and unequivocal description of how these nonstandard gravitational waves affect pulsar timing residuals."
    },
    "0812/0812.4085_arXiv.txt": {
        "abstract": "{The mechanical and radiative feedback that exists in the star formation history affects the subsequent star formation rate. } {After considering the effects of negative feedback on the process of star formation, we explore the relationship between star formation process and the associated feedback, by investigating how the mechanical feedback from supernovae(SNe) and radiative feedback from luminous objects regulate the star formation rate and therefore affect the cosmic reionization.} {Based on our present knowledge of the negative feedback theory and some numerical simulations, we construct an analytic model in the framework of the Lambda cold dark matter model. In certain parameter regions, our model can explain some observational results properly.} { In large halos($T_{\\rm vir}>10^4K$), both mechanical and radiative feedback have a similar behavior: the relative strength of negative feedback reduces as the redshift decreases. In contrast, in small halos ($T_{\\rm vir}<10^4K$) that are thought to breed the first stars at early time, the radiative feedback gets stronger when the redshift decreases. And the star formation rate in these small halos depends very weakly on the star-formation efficiency.} {Our results show that the radiative feedback is important for the early generation stars. It can suppress the star formation rate considerably. But the mechanical feedback from the SNe explosions is not able to affect the early star formation significantly. The early star formation in small-halo objects is likely to be self-regulated. The radiative and mechanical feedback dominates the star formation rate of the PopII/I stars all along. The feedback from first generation stars is very strong and should not be neglected. However, their effects on the cosmic reionization are not significant, which results in a small contribution to the optical depth of Thomson scattering.} ",
        "introduction": "In the past decades, the so-called ``bottom-up\" hierarchical scenario for the large-scale structure formation in the Lambda cold dark matter ($\\Lambda$CDM) cosmogonies has been getting decisive support from more and more observations (e.g. from HST, WMAP, SDSS) and from high-resolution N-body/hydrodynamics numerical simulations. Owing to the complexity of the baryonic evolution in the radiative background and gravitational field, the galaxy formation and evolution in dark matter halos and the reionization history of IGM have not been fully understood yet. The difficulties stem from our lack of knowledge of those early formed objects that are unobservable at present. The formation of the first objects is relevant to the cosmic reionization history. Recent detections of the Gunn-Peterson trough\\cite{gunn1965} in the spectra of QSOs with $z\\simeq 6$ indicate less 50\\% neutral hydrogen at $z \\sim 6.5$ \\cite{fan,loeb}. The ongoing observations by the WMAP satellite of cosmic microwave background (CMB) and the highest redshift QSOs put very tight constraints on the reionization history of the Universe. The WMAP five-year observation manifests the Thomson scattering optical depth, $\\tau_{\\rm e} = 0.084^{+0.016}_{-0.016}$ \\cite{Komatsu08}, which suggests that our Universe might be reionized during the period of redshift $9.4 \\le z_{\\rm re} \\le 12.2$.  It is well-known that chemical elements heavier than lithium are produced exclusively through stellar nucleosynthesis. Some of the first generation stars (the so-called population III stars, hereafter, PopIII) die as SNe explosions, which can expel heavier elements into the intergalactic medium (IGM). When the metal elements in IGM are enriched to a certain threshold $Z_{\\rm crit}$, the population II/I stars (hereafter PopII/I) will form and take the place of the first stars to light the universe. Thence, the SNe from PopIII stars determine the transition from PopIII to PopII/I. The existence of PopIII stars can help in explaining the metal enrichment from $Z\\simeq 10^{-12}-10^{-10}$ to the lowest metallicity of PopII stars $Z\\simeq 10^{-4}-10^{-3}$, the formation of massive black holes, the reionization of the universe, the starting engine for the formation of the first galaxies and the G-dwarf, and so on\\cite{cf05}. But there is still no hope in observing the first generation stars until the launch of the $James\\ Webb\\ Space\\ Telescope$ (JWST), the successor of the $Hubble\\ Space\\ Telescope$ (HST) \\cite{bl01}. JWST is expected to find the pair instability SNe (PISNe) from massive PopIII stars \\cite{wise05}.  Due to the absence of the observational data of very high-redshift ($z\\ge 10$) objects, theoretical investigations are mainly based on numerical simulations. Some works concentrated on the effects of the first-generation SNe explosions\\cite{Yoshida03,KY05,Greif07}, while the others focused on the strong stellar and galactic winds from PopIII stars\\cite{MEM06,Ricotti08}. Based on these works, we put forward a model to describe the global effects from the PopIII stars at an early time, such as the SFR density, the IGM reionization, and so on.  In this paper, we study the effects of the mechanical feedback from SNe and the radiative feedback from stars and UV background, especially the negative feedback effects on the SFH and the cosmic reionization history. The radiative feedback from PopIII stars includes the ISM photoevaporation and IGM reionization. The outline of this paper is as follows. In section 2, we describe the evolution of dark matter halos in the $\\Lambda$CDM model. Within this framework, the SFR in a halo can be expressed as an analytic formula with some free parameters. Mechanical feedback from stars is studied in section 3. In section 4, we introduce an analytic model to deal with the radiative feedback from PopIII. Finally, our discussion and conclusions are presented in the last section. Throughout this paper, we adopt the cosmological parameters consistent with the 5 years WMAP data:$\\Omega_{\\rm m}h^{\\rm 2} = 0.1358$, $\\Omega_\\Lambda = 0.726$, $\\Omega_{\\rm b}=0.0456$, $h = 0.7$, $\\sigma_8=0.812$\\cite{Komatsu08}. ",
        "conclusions": "In this paper we develop an analytic model of the mechanical and radiative feedback of star formation based on some results of numerical simulations. Basically, our results can fit the observational data. However, this model also depends on some initial parameters (most of them are listed in $\\hbox{Table\\ 1}$) and working assumptions.  $(i)$ It is possible to fit the reionization result from the highest QSOs observation by providing another pair of $f^{\\rm {II}}_{*}$ and $f^{\\rm {II}}_{\\rm esc}$. This is because the IGM reionization mainly depends on $n^{\\rm II}_\\gamma$ and $f^{\\rm {II}}_{*}\\times f^{\\rm {II}}_{\\rm esc}$. The value of $n^{\\rm II}_\\gamma$ is nearly fixed under certain conditions (e.g. metallicity, IMF, and stellar mass range), but lacking the SFRD observational data constraint on $f^{\\rm {II}}_{*}$, one can find many pairs of $f^{\\rm {II}}_{*}$ and $f^{\\rm {II}}_{\\rm esc}$ by keeping $f^{\\rm {II}}_{*}\\times f^{\\rm {II}}_{\\rm esc}$ invariant to guarantee $\\tau_{\\rm e} = 0.084^{+0.016}_{-0.016}$.  $(ii)$ As for the assumption on $F_{\\rm III}(z)$, there are a lot of similar functions. This is an uncertainty factor in this model. Hopefully, it may be determined by the numerical simulations and even the future JWST observation.  $(iii)$ We only consider the negative feedback in this work, the neglect of some positive feedback, such as the X-ray background or dense shell (from SNe explosions) fragmentation, results in the underestimation of the SFRD, $Q_{\\rm HII}$, $\\tau_e$, and so on. This is what we will investigate in the next work.  $(iv)$ Although we can fit the optical depth to Thomson scattering in this work, it does not mean the X-ray pre-ionization is not important. Obviously, our results are close to the low fringe of the error bar. Perhaps, the left space in the error bar is waiting for the contribution from X-ray pre-ionization.  In the following, we summarize the results of this work in more detail.  1. The SNe feedback can suppress the star formation in large halos ($T_{\\rm vir}>10^4 K$). And the relative feedback strength of SNe feedback [$\\dot{\\rho}^{\\rm FB}_{\\rm III}(z)/\\dot{\\rho}_{\\rm III}(z)$ or $\\dot{\\rho}^{\\rm FB}_{\\rm II/I}(z)/\\dot{\\rho}_{\\rm II/I}(z)$] reduces as the redshift decreases, because massive halos appear abundant at low redshift and massive halos whose binding energy is much higher than before can reduce the feedback effect from SNe(see Eq.\\ref{snfb} and Eq.\\ref{snfb3}).  2. The radiative feedback mechanism can affect the SFR in both small ($T_{\\rm vir}<10^4 K$) and large($T_{\\rm vir}>10^4 K$) PopIII halos. In small halos, the shallow gravitational well and the poor cooling-efficiency cannot prevent the hot gas from escaping due to the strong photoevaporation from the massive OB stars. The radiative feedback can suppress the PopIII SFR considerably. In large halos, the ionizing background enhances the Jeans mass and makes the SFR suppressed more or less.  3. The SFRD in small PopIII halos($T_{\\rm vir}<10^4 K$) is sensitive to $\\tau_{\\rm FB}$, which depends on $f^{\\rm III}_{\\rm esc}$ and $n^{\\rm III}_{\\gamma}$ seriously. If $f^{\\rm III}_{\\rm esc}$ and $n^{\\rm III}_{\\gamma}$ increases, $\\tau_{\\rm FB}$ will decreases and so does the SFRD. Furthermore, our results support this viewpoint: early star formation is likely to be self-regulated \\cite{Ricotti02a,Ricotti02b,Yoshida03,cf05,KY05,Ricotti08}.  4. Although PopIII stars have high value of $f^{\\rm III}_{\\rm esc}$ and $n^{\\rm III}_{\\gamma}$, they are not able to reionize the universe considerably (about $20\\%$ at $z\\sim 14$). And the optical depth to Thomson scattering provided by PopIII stars is less than $22\\%$.  In general, we find the radiative feedback is important to the formation of the early generation stars. It suppresses the star formation considerably. But the mechanical feedback from the SNe explosions is not able to affect the early star formation significantly. The radiative and mechanical feedback dominates the star formation rate of the PopII/I stars. The feedback on SFRD from first generation stars is very strong and should not be neglected. However, their effect on the cosmic reionization is not significant, which results in a small contribution to the optical depth of electrons $\\tau_e$."
    },
    "0812/0812.1944_arXiv.txt": {
        "abstract": "{Motivated by the possibility of inflation in the cosmic landscape, which may be approximated by a complicated potential, we study the density perturbations in multi-field inflation with a random potential. The random potential causes the inflaton to undergo a Brownian-like motion with a drift in the $D$-dimensional field space, allowing entropic perturbation modes to continuously and randomly feed into the adiabatic mode. To quantify such an effect, we employ a stochastic approach to evaluate the two-point and three-point functions of primordial perturbations. We find that in the weakly random scenario where the stochastic scatterings are frequent but mild, the resulting power spectrum resembles that of the single field slow-roll case, with up to 2\\% more red tilt. The strongly random scenario, in which the coarse-grained motion of the inflaton is significantly slowed down by the scatterings, leads to rich phenomenologies. The power spectrum exhibits primordial fluctuations on all angular scales. Such features may already be hiding in the error bars of observed CMB $TT$ (as well as $TE$ and $EE$) power spectrum and have been smoothed out by binning of data points. With more data coming in the future, we expect these features can be detected or falsified. On the other hand the tensor power spectrum itself is free of fluctuations and the tensor to scalar ratio is enhanced by the large ratio of the Brownian-like motion speed over the drift speed. In addition a large negative running of the power spectral index is possible. Non-Gaussianity is generically suppressed by the growth of adiabatic perturbations on super-horizon scales, and is negligible in the weakly random scenario. However, non-Gaussianity can possibly  be enhanced by resonant effects in the strongly random scenario or arise from the entropic perturbations during the onset of (p)reheating if the background inflaton trajectory exhibits particular properties. The formalism developed in this paper can be applied to a wide class of multi-field inflation models including, e.g. the N-flation scenario. }  \\begin{document} ",
        "introduction": "The study of the inflationary universe is mainly on the slow-roll model with a single inflaton field. However, motivated by a variety of reasons, including those from superstring theory, multi-field inflation has received a lot of attention recently. An important reason to focus on the single field case is simplicity. A typical multi-field inflationary scenario is clearly very complicated and not predictive. There are special multi-field models which are easy to analyze while still predictive: models where the fields do not interact with each other during inflation were analyzed first in Ref.\\cite{Starobinsky:1986fxa}.  A special case of this class of models is when all inflaton fields have the same potential, such as assisted inflation\\cite{assisted}. While in more general cases, e.g. N-inflation\\cite{Ninflation}, and similar scenarios\\cite{Battefeld:2008py, Misra:2007cq}, each inflaton can have its own potential.  If the inflaton fields couple to each other during inflation and reheating, the problem seems rather intractable. Motivated by the realization of inflation in superstring theory, we expect a multi-field inflationary scenario. This is particularly true if inflation takes place in the stringy cosmic landscape\\cite{Kachru:2003aw}, where the number of moduli is probably as large as hundreds. Even if not all of them participate in inflation, it is likely that a large subset of moduli contribute directly or indirectly to the inflationary scenario. Their interaction can be rather complicated (maybe strongly interacting as well), generating a cosmic landscape that is random (aperiodic) in some directions and quasi-periodic (within certain field range) in others.  A realistic illustration could be the following potential for moduli $\\rho_i$ and axions $\\phi_i$\\cite{axionf}, \\begin{equation} V_{T} (\\rho_i, \\phi_i) =  V(\\rho_j) - \\alpha_i \\cos \\left(\\frac{\\phi_i}{f_i} \\right)  + \\beta_{ij}\\cos \\left(\\frac{\\phi_i}{f_i} - \\frac{\\phi_j}{f_j}\\right) + \\dots  + U(\\rho_i, \\phi_i) ~. \\end{equation} $f_i$'s are the axion decay constants and $\\alpha_i =  M_i^4 e^{-S_{\\mathrm{inst}}^{i}}  \\gg \\beta_{ij} $, where $M_i$ is a natural mass (say, the string) scale, and the $i$th instanton has action $S_{\\mathrm{inst}}^{i}$. $V(\\rho_j)$ contains the potential coming from the moduli and includes $D$-term contributions. Presumably, it contains the dominant contribution to the vacuum energy density in $V_{T}(\\rho_i, \\phi_i)$, while $U(\\rho_i, \\phi_i)$ expresses the couplings of the moduli and the axions. In the region where $V(\\rho_j)$ is relatively flat or relatively localized, extensive inflation may take place, and the inflaton is represented by the collection of $\\phi_i$ axions. With very small $\\alpha_i$, the wavefunction is either not trapped (because binding in higher dimension typically requires a deeper attractive potential) or tunneling out of any classical vacuum site is very rapid. This potential is periodic except for the interaction term $U(\\rho_i, \\phi_i)$, which is perturbative in general, but it can also introduce relatively strong interactions at times. Here, $U(\\rho_i, \\phi_i)$ plays the role analogous to impurities in a crystal or doping in condensed matter physics.  On a more general ground, the complicated scalar potential could be described using the random matrices\\cite{Aazami:2005jf}. Generically, the potential probably looks rugged at short distances but smoother at coarse-grained level. We shall assume the randomness in the potential comes in different scales. Roughly speaking, we have in mind the following scenario: the small scale randomness of the potential causes the inflaton to undergo Brownian-like motion in the $D$-dimensional field space, allowing the entropic perturbations to feed into the adiabatic perturbations during inflation. The large scale behavior of the potential allows the inflaton to slowly move down in the landscape and eventually end the inflationary epoch. The ending can be brought about by the appearance of tachyons, fast tunneling to a low point in the potential, or a quick roll down when the inflaton suddenly reaches a steep slope.  It is important to point out\\cite{Tye:2006tg} that the number of scalar fields relevant for inflation should be bigger than 2. Otherwise, the inflaton wavefunction will be trapped (i.e., localized) in the inflaton potential. If so, eternal inflation is unavoidable and inflation does not end. Our scenario implicitly assumes there is no eternal inflation and this assumption becomes more likely as $D$, the number of scalar fields participating in inflation, increases beyond 2. Here we have in mind a relatively large $D$ (say, dozens to hundreds).  In slow-roll inflation, the large number $N_e$ of e-folds of inflation follows from the flatness of the inflaton potential. In general, this requires a fine tuning. Here, the large scale structure of the inflaton potential has to be reasonably flat too. However, the large scale flatness constraint is substantially relaxed compared to that in the slow-roll case, since the inflaton exhibits a random walk like motion and scatterings tend to slow down its overall motion. For this reason, we consider multi-field inflation in the landscape to be quite natural.  Motivated by this picture, it is reasonable to ask if one can make any predictions about such an inflationary scenario. Although the analysis is rather general under certain conditions, it helps to have a concrete example in mind. As an illustration, we could imagine an extended version of brane inflation, where a $D$3-brane is moving in the cosmic landscape. Its position during inflation includes both its position in the moduli space and its position within the 6-dimensional compactified bulk. Its motion includes classical rolling and scattering, which may result in percolation. Quantum fluctuation leads to density perturbation. Presumably,  quantum diffusion may be included as well. One may consider this scenario as a generalization of chain inflation\\cite{chain}, which includes only repeated fast tunnelings, and related scenarios in Ref.\\cite{Davoudiasl:2006ax}.  Note that, in the absence of quantum fluctuation, no cosmological perturbation is generated in the classical but random trajectory of the inflaton. Even for a convoluted random looking trajectory, the classical path of the inflaton will be the same at every causal patch of the universe during inflation and so no cosmological perturbation is  generated in the absence of quantum effects. So the cosmological perturbation is entirely sourced by the quantum fluctuation. As is well known, in multi-field inflation, the quantum perturbation results in the adiabatic mode and the entropic modes of the cosmological perturbation. If the classical trajectory is a straight line in the field space, the entropic modes and the adiabatic mode remain decoupled, and the density perturbation comes entirely from the adiabatic mode. The randomness of the classical background allows for the random mixing of adiabatic and entropic modes, which is how the randomness of the classical background path enters into the cosmological perturbations, and how the resulting power spectrum incorporates the randomness of the inflaton potential. In general, a random potential will lead to a combination of classical percolation and quantum diffusion. However, knowing that the density perturbation is very small, i.e. $\\delta \\rho/\\rho  \\sim 10^{-5}$, it should be a very good approximation to treat the quantum diffusion as a quantum fluctuation around the classical (but complicated) background trajectory. In this case, the data will never be able to reveal individual bumps. Our analysis in this paper  is carried out in this approximation. A full analysis of quantum diffusion will be interesting but is beyond the scope of this paper. (Of course, one may take the point of view that, since $\\delta \\rho/\\rho$ is so small, such an analysis may not be necessary.)  Here we would like to show that if the potential is random enough, some definitive statements about such a multi-field inflationary scenario are possible. In fact, the predictions can be quite specific and detail independent.  \\begin{itemize} \\item Even if the inflaton moves relatively fast, its drift velocity can be small, dictated by the slope of the potential over large field scales. Inflation naturally lasts for many e-folds.  \\item One may easily recover the usual slow-roll, almost scale-invariant power spectrum, with a slight additional red tilt (up to 2\\% when scatterings are frequent but mild) in the power spectrum index.  \\item The running of the power spectral index can be large, because the conversion of entropic modes into the adiabatic mode introduces additional scale dependence.  \\item The tensor to scalar ratio $r$ can be large. On one hand, it is suppressed by the amplified scalar power spectrum; on the other hand, it gets enhanced by a factor equal to the square of the ratio of the Brownian-like motion speed to the drift speed, which we expect to be much greater than one in the strongly random scenario. The net effect may still be an enhancement in $r$.  \\item Non-Gaussianity is generically suppressed by the growth of the adiabatic mode after horizon exit due to the presence of extra light fields. However, it may get enhanced by resonant effects\\cite{Chen:2008wn}. On the other hand, the large number of strongly coupled fields can enhanced the non-Gaussianity induced during the onset of (p)reheating.  \\item We generically expect random fluctuations in the scalar power spectrum due to scatterings of the inflaton. Such fluctuation may appear with random amplitudes and periodicity. The precise shape of the power spectrum is not quite informative here, but the variance of such fluctuations is well determined by the underlying microscopic properties of the scalar potential (landscape). The currently observed $TT$ power spectrum, looking at the resolution of each multiple moments, is not smooth. The error bar of each data points (roughly $10\\%$) is comparable to the fluctuation of data itself (for $100 \\lesssim l \\lesssim 800$)\\cite{Nolta:2008ih}, so it is hard to claim any fluctuating features in the currently observed power spectrum. Furthermore, current analysis based on binned data points from a few tens of multiple moments may have smoothed out such fluctuations.  However, with more observational data coming in the future, we expect that analyzing the power spectrum at the resolution of single multiple moments is possible and the fluctuations in the power spectrum can certainly be detected or falsified if the error bars become substantially smaller than the primordial fluctuations. Furthermore, if such fluctuations are detected in the $TT$ power spectrum, one expects to see the same pattern in $TE$ and $EE$ correlations as well.  \\end{itemize}  The paper is organized as follows. The relevant properties of the multi-field inflationary scenario are reviewed in Section \\ref{sec_rev}. In Section \\ref{sec_infmf}, we consider the model where the randomness in the potential leads to a Brownian-like motion for the inflaton. That is, only classical percolation is included. The power spectrum is discussed in Section \\ref{sec_pzeta} and the non-Gaussianity is discussed in Section \\ref{sec_fnl}. In Section \\ref{sec_dis}, we discuss the more general situation where the random scattering is also caused by quantum diffusion. We attempt to argue that the generic results obtained in Section \\ref{sec_pzeta} and Section \\ref{sec_fnl} may hold even in the presence of such quantum diffusion. Section \\ref{sec_sum} contains the summary and some remarks. Some details are relegated to the two appendices. ",
        "conclusions": "\\label{sec_dis}  So far in this paper, we have considered a random potential where scatterings in the potential lead to inflaton motion resembling that of a random walk. We analyze the impact of this classical percolation property on the density perturbation sourced by quantum fluctuations. In a more realistic situation, we expect quantum diffusion effects on the motion of the inflaton as well. By this, we mean fast tunnelling, resonance scatterings as well as quantum hopping. Here, quantum hopping is hopping over a low barrier, due to the presence of the Gibbons-Hawking temperature $T=H/2 \\pi$, so one may treat this as a thermal effect.  In this scenario, we assume that the inflaton is not trapped at any classically stable site that will lead to eternal inflation. Rather, the inflaton is mobile in the random potential and can move freely, via classical percolation or quantum diffusion. The argument for this property is based on the multi-field feature and discussed in Ref.\\cite{Tye:2006tg}. There, the multi-field feature ($D \\gg 1$) is crucial for the mobility of the inflaton. Here we see that the resulting magnitude of the density perturbation requires many scatterings during a single e-fold of inflation, consistent with the overall picture that emerges.  Instead of the decay rate per unit volume $\\Gamma$, it is convenient to introduce the tunneling probability in one Hubble time, the dimensionless parameter $\\gamma$ given by $\\gamma=\\Gamma/H^4$, so $\\gamma$ is a measure of the nucleation rate relative to the expansion rate of the universe. We are interested in situations where $\\gamma > 1$. Let the probability of an arbitrary point remaining in a false de Sitter vacuum at time $t$ be $P(t)$. Then $P(t)$ behaves as $ P(t)\\sim \\exp ({-{4\\pi\\over 3}\\gamma Ht})$. Thus the lifetime of the field in a false vacuum is estimated as \\begin{equation} t_F\\simeq {3\\over 4\\pi H\\gamma} = {3 H^3\\over 4\\pi \\Gamma} \\label{tf} \\end{equation} If $\\gamma \\gg 1$, then $Ht_F \\ll 1$. Let us assume that inflation is dominated by repeated fast tunneling events. Following Ref.\\cite{Guth:1982ec, Starobinsky:1982ee}, or in the $\\delta N$ formalism, we can easily estimate the amplitude of density perturbation \\begin{equation} \\zeta = \\delta N \\simeq H \\delta t \\sim Ht_F \\label{adp} \\end{equation} COBE normalization is $\\zeta \\sim 10^{-5}$ and thus $t_F\\sim 10^{-5}H^{-1}$ which implies that there are roughly $10^5$ steps during one e-fold. This is the key idea of chain inflation\\cite{chain}. We do like to note that an axionic potential with an exponentially small barrier height and a small decay constant do have the right feature for such a scenario.  Here, we assume that each site has a dominant decay channel for fast tunneling, so that the inflaton path is uniquely determined. Since $\\gamma \\gg 1$, the phase transition completes rapidly (that is, the universe percolates). So the inflaton path mimics a classical path. This puts a tight constraint on the inflaton potential. Presumably, this condition can be relaxed somewhat; that is, the intermediate steps can lead to different paths, though towards the end of inflation, the inflaton ends up in the same vacuum site in the potential. Such deviation in paths can lead to a new type of density perturbation, which must stay small to satisfy the observational bound. It will be very interesting to study this density perturbation and what it tells us about the inflaton potential.  Even with fast tunneling, a locally attractive site will trap the inflaton even for a little while. In high dimensions, it is harder to trap a wavefunction. So in some cases, the inflaton is not trapped by a local attractive site, but rather it scatters as a resonance over that site. In both cases, $\\epsilon$ or $\\dot \\sigma$ will decrease and so $\\zeta$ will increase. It is also likely that $\\zeta$ fluctuates, with vales much bigger than its average $\\sqrt{\\la\\zeta^2\\ra} \\sim 10^{-5}$. As long as such fluctuations happen at a scale too small to be picked up by the CMB data, (say, $10^3$ steps per e-fold), it will be difficult to detect this effect.  Such a scattering (or fast tunneling) from a meta-stable site A to a nearby site B can release a substantial fraction of the vacuum energy to the kinetic energy of the inflaton. Since the next step happens rapidly, there is no time for the Hubble expansion or the inflaton decay to dilute its kinetic energy, so next step scattering or tunneling can go to a site C with vacuum energy density comparable to that of A and above that of B. This means repeated scatterings and/or fast tunnelings do not have to go downhill. This leads to the Brownian-like motion. However, over some number of steps, the Hubble damping will take place and the inflaton will move downwards. This leads to the drift velocity. So at the coarse-grained scale, the resulting picture can easily mimic the slow-roll scenario.  Even though the lifetime is much shorter than the Hubble time, the motion of the scalar field does not really roll down along a smooth potential, but instead a bumpy potential. Usually the small bumps do not affect the amplitude of the power spectrum, but it may impart a large distinctive non-Gaussianity\\cite{Chen:2008wn}. In our case the tunneling/percolation time is much shorter than the Hubble time, which means that the ``period'' of the bumps is very short and the non-Gaussian features from bumps may be too fluctuating to be picked up observationally. However, by the very nature of randomness in the potential, it is entirely possible that a big bump (or dip) in the potential or a particularly strong scattering can cause an observable feature in the CMB. This might be the origin of the feature at $l\\sim 20$ in the WMAP data. Such fluctuations may appear in the power spectrum as a function of the wave number (or $l$). They are imprinted primordially, and should not be confused with the statistical fluctuations in data analysis.  One possible subtle issue we have not discussed in this paper is the effect of isocurvature perturbations after reheating. When the universe is reheated through the decay of a large number of scalar fields, the abundance of the created particles may depend on the initial value of the inflaton fields. The isocurvature perturbations can be characterized as the perturbation on the ratio of the particle number densities between different species, e.g. between photons and dark matter particles. Recent observations from WMAP+BAO+SN have put a stringent bound on the amount of isocurvature perturbations allowed in the primordial power spectrum\\cite{Komatsu:2008hk}. The isocurvature power spectrum cannot exceed 6.7\\% of its adiabatic counterpart. In our scenario, the power spectrum of adiabatic perturbation at the end of inflation is proportional to $\\vartheta\\Neff \\sim (D-1)\\Theta^2\\Neff$, while we also have entropic perturbations by themselves with total power spectrum proportional to $(D-1)$, so the ratio of entropy perturbations over adiabatic perturbations is $1/(\\Neff\\Theta^2)$ which is about a few percent with $\\Theta \\lesssim 1$. The amount of isocurvature perturbations left after reheating could be well within the current observational bound. On the other hand, the reheating process may result in additional power in the curvature perturbation and possibly additional non-Gaussianity. This is certainly an interesting effect that deserves future study."
    },
    "0812/0812.5083_arXiv.txt": {
        "abstract": "We study the impact of catastrophic errors occurring in the photometric redshifts of galaxies on cosmological parameter estimates with cosmic shear tomography. We consider a fiducial survey with 9-filter set and perform photo-z measurement simulations. It is found that a fraction of $1\\%$ galaxies at $z_{\\rm spec}\\sim0.4$ is misidentified to be at $z_{\\rm phot}\\sim3.5$. We then employ both $\\chi^2$ fitting method and the extension of Fisher matrix formalism to evaluate the bias on the equation of state parameters of dark energy, $w_0$ and $w_a$, induced by those catastrophic outliers. By comparing the results from both methods, we verify that the estimation of $w_0$ and $w_a$ from the fiducial 5-bin tomographic analyses can be significantly biased. To minimize the impact of this bias, two strategies can be followed: (A) the cosmic shear analysis is restricted to $0.5<z<2.5$ where catastrophic redshift errors are expected to be insignificant; (B) a spectroscopic survey is conducted for galaxies with $3<z_{\\rm phot}<4$. We find that the number of spectroscopic redshifts needed scales as $N_{\\rm spec} \\propto f_{\\rm cata}\\times A$ where $f_{\\rm cata}=1\\%$ is the fraction of catastrophic redshift errors (assuming a 9-filter photometric survey) and A is the survey area. For $A=1000 \\hbox{ deg}^2$, we find that $N_{\\rm spec}>320$ and 860 respectively in order to reduce the joint bias in $(w_0,w_a)$ to be smaller than $2\\sigma$ and $1\\sigma$. This spectroscopic survey (option B) will improve the Figure of Merit of option A by a factor $\\times 1.5$ thus making such a survey strongly desirable. ",
        "introduction": "The weak gravitational lensing effect by large scale structures, namely the cosmic shear, is believed to be a powerful tool to measure the cosmological parameters including the equation of state of dark energy. Over the last ten years, weak lensing observations have already provided increasingly better constraints on the matter density $\\Omega_M$ and the amplitude of mass fluctuations often represented by $\\sigma_8$, the $rms$ amplitude of the smoothed and linearly extrapolated mass density perturbations on the scale of $8\\hbox{ Mpc}\\ h^{-1}$ (Hoekstra, Yee \\& Gladders 2002, Jarvis et al. 2003, Rhodes et al. 2004, Heymans et al. 2004, Hoekstra et al. 2006, Semboloni et al. 2006,  Massey et al. 2007, Benjamin et al. 2007, Fu et al. 2008). Their contributions to the constraints on the total neutrino mass have also been investigated (Tereno et al. 2008, Gong et al. 2008, Ichiki et al. 2008, Li et al. 2008). Several more ambitious weak lensing surveys targeting at hundred to thousand times larger areas, such as DES\\footnote{http://www.darkenergysurvey.org}, LSST\\footnote{http://www.lsst.org}, JDEM\\footnote{http://jdem.gsfc.nasa.gov} and EUCLID\\footnote{http://sci.esa.int/science-e/www/object/index.cfm?fobjectid=42266}, are being planned aiming at exploring the properties of dark energy in great detail.  Future large surveys are expected to be able to control statistical errors to insignificant levels. Therefore different systematics can become dominant sources of contaminations in future weak lensing studies, and should be investigated carefully. These include the effect of the point spread function (PSF) (e.g., Amara \\& Refregier 2008), the methodology in the shear measurement (Bacon et al. 2001, Erben et al. 2001, Hirata \\& Seljak 2003, Heymans et al. 2006, Massey et al. 2007, Miller et al. 2007,  Kitching et al. 2008, Bridle et al. 2008, Ngan et al. 2008, Semboloni et al. 2008), redshift errors (Ma et al. 2006, Huterer et al. 2006), intrinsic alignments of source galaxies (see Heavens 2001 for a review, Bridle \\& King 2007, Fan 2007), uncertainties in the computation of the non-linear power spectrum (see discussion in Huterer 2002, Van Waerbeke et al. 2002), and so on. In this paper, we focus on the impact of the photometric redshift errors on cosmological studies.   The necessity of measuring vast amount of galaxy redshifts in cosmic shear observations makes the multi-color photometric determination of redshifts the only feasible way in obtaining the redshift information of source galaxies. However the inaccuracy of photometric redshifts can contaminate considerably the cosmological information embedded in cosmic shears (Huterer et al. 2006; Ma et al. 2006; Amara \\& Refregier 2007; Wittman et al. 2007; Abdalla et al. 2008; Margoniner \\& Wittman 2008; Stabenau et al. 2008). Taking into account the bias and the scatter of $z_{\\rm phot}$ with respect to $z_{\\rm spec}$, Ma, Hu \\& Huterer (2006) analyze the corresponding degradation of the constraints on dark energy. It is found that these errors can completely erase the additional information from tomographic lensing observations. In order to make full use of the advantage of the tomographic binning, both the bias and scatter of $z_{\\rm phot}$ have to be known to better than 0.003-0.01. Amara \\& Refregier (2007) and Abdalla et al. (2008) pay special attention to the so-called catastrophic errors in $z_{\\rm phot}$, the outliers with $|z_{\\rm spec}-z_{\\rm phot}|>1$, and explore their effects on the figure of merit of dark energy parameter constraints.   While most of the existing studies concentrate on the error range of parameter constraints, the systematic bias in the determination of cosmological parameters due to the photometric redshift errors should also be carefully analyzed. A large systematic bias can lead to a completely wrong conclusion on the values of the cosmological parameters although the statistical errors may be small (Amara \\& Refregier 2008). Our study presented in this paper emphasizes precisely this aspect of effects on dark energy parameter estimates arising from catastrophic failures in photo-z measurements. We focus on a survey with similar design as described in Aldering et al. (2004), with its 9-filter set described also in Dahlen et al. (2008) and Jouvel et al. (2009 in prep.). We analyze in particular the bias on dark energy parameter estimates induced by the catastrophic failures occurred around $(z_{\\rm spec}, z_{\\rm phot})\\sim (0.4, 3.5)$. And we investigate in detail the requirements of the number of spectroscopic redshift measurements to retrieve the correct estimates.  The outline of the paper is as follows. In \\S 2, we introduce the methods to describe the lensing observables and to evaluate the biases. We present the results of bias in \\S 3. In \\S 4, we discuss the calibration requirements to remove the biases due to the catastrophic failures. Discussions and conclusions are given in \\S 5. ",
        "conclusions": "\\label{sect:dis}  In this paper, we concentrate on the island part of the galaxies in the $(z_{\\rm spec},z_{\\rm phot})$ plot where $z_{\\rm phot} $ is completely mis-estimated. This type of large catastrophic failures in the photometric redshift measurements arise from the confusion of different characteristics of galaxies' SED due to the limited wavelength coverage. The most notable features used in the photo-z measurement of galaxies are the Lyman and $4000 \\AA$ breaks. The mis-identification of the two results in the islands at $(z_{\\rm spec}, z_{\\rm phot})\\sim (0.4, 3.5)$ and $(z_{\\rm spec}, z_{\\rm phot})\\sim (3, 0.2)$ seen in Figure 1. We particularly analyze the systematic bias induced by such islands on cosmological parameter estimations. We find that a fraction of $f_{\\rm cata}=1 \\%$ catastrophic failures due to galaxies at $z_{\\rm spec}\\sim 0.4$ being misidentified to be at $z_{\\rm phot}\\sim3.5$, which is the dominant case of catastrophic failure occurrence in our photo-z simulations, can bring systematic biases on $w_0$ and $w_a$ at a level far exceeding the statistical errors in $5$-tomographic-bin analysis.   It has been realized that spectroscopic redshift calibration may be necessary for future weak lensing surveys in order to reduce the uncertainty of photometric redshift errors. The large bias shown in our analyses reinforces such an argument. We investigate the requirements of the spectroscopic calibration to retrieve the unbiased parameter estimates. We find that for $f_{\\rm cata}=1\\%$, and $f_{\\rm sky}=0.025$, about 400-1000 spectral redshift measurements for galaxies with $z_{\\rm phot}=[3,\\ 4]$ are needed. We further present a scaling relation of $N_{\\rm spec}\\propto {f_{\\rm cata}\\times f_{\\rm sky}}$. Thus for a future survey with a survey area of $10,000 \\hbox{ deg}^2$, the size of the calibration sample in the redshift range $z_{\\rm phot}=[3,\\ 4]$ should be on the order of $10^4$ if $f_{\\rm cata}$ remains to be $\\sim 1\\%$. In analyzing the error degradation from general photometric errors, people have argued that the inaccuracy of the high redshift photometric measurements can be compensated by increasing the measurement accuracy for galaxies at relatively low redshifts. Therefore much less spectroscopic calibration samples at high redshifts are needed (e.g., Ma et al. 2006). Concerning the catastrophic-error-induced bias discussed in this paper, however, there is no such a compensation. Thus the required $\\sim 10^4$ spectral calibrations in $z_{\\rm phot}=[3,\\ 4]$ may be regarded as challenging tasks.  It is also argued that the inclusion of both U and near-IR filters is essential for the accuracy of photo-z and the minimization of the catastrophic outliers (e.g., Dahlen et al. 2008, Abdalla et al. 2008), as the Lyman and $4000 \\AA$ breaks can be effectively distinguished. Given that our simulations with $9$-filter set already include the near-IR bands, we consider the improvements from U band coverage. Extending the wavelength coverage to $3200 \\AA$, our simulations show that the catastrophic failure fraction decreases dramatically to the level around $0.1\\%$ (Jouvel et al. 2009, in prep.). In this case, the bias should no longer be significant compared to the statistical errors. It is thus important to investigate the possibility to include $<3500 \\AA$ sensitivity for the next generation space dark energy mission.  It has also been proposed to filter out those problematic galaxies, such as late type galaxies, low signal-to-noise galaxies, or very low/high redshift (e.g., $z<0.5$/$z>2.5$) galaxies whose photo-zs are most suspicious, in the shear analysis to reduce the impact of catastrophic failures (Abdalla et al. 2008; Margoniner \\& Wittman 2008; Heavens, Kitching \\& Taylor 2006). In such cases, the statistical power of a survey is inevitably lowered to certain extents depending on the fraction of galaxies discarded. In Figure \\ref{Fig:cut}, we show the degradation of the constraints if those galaxies at $z<0.5$ and $z>2.5$ (thus 3 bins left) are excluded in the analyses, compared to the case with calibrated (known) $f_{\\rm cata}$. We see that the Figure of Merit (denoted as F.o.M)\\footnote{F.o.M = $\\sqrt {\\rm det (F)}$, where F is the reduced Fisher matrix for $w_0$ and $w_a$ after marginalization over other parameters.} decreases from 41 to 24, i.e., almost a factor of two degradation. If we re-divide the galaxies with $0.5\\le z\\le 2.5$ into $5$ finer bins (right panel of Figure \\ref{Fig:cut}), the F.o.M = 29, or a degradation factor $\\sim 1.5$ (see also table 1 for a summary.). All these calculations are done with our fiducial survey conditions and Gaussian priors $\\sigma(\\Omega_b)=0.01$ for $\\Omega_b$ and $\\sigma(p_i)=0.05$ are applied upon all other parameters.  In our current analyses, we isolate the island galaxies appearing around $(z_{\\rm spec}$,$z_{\\rm phot})=(0.4,3.5)$ to see clearly their effects on cosmic shear analysis. On the other hand, the scatters and small bias of $z_{\\rm phot}$ around $z_{\\rm spec}$ have been modeled and studied in, for example, Ma et al. (2006). With $\\sigma_z$ and $z_{\\rm bias}$ as well as their uncertainties $\\Delta\\sigma_z$ and $\\Delta z_{\\rm bias}$ taken into account, the statistical errors of the constrained cosmological parameters will be degraded depending on the values of ($\\sigma_z$, $z_{\\rm bias}$, $\\Delta\\sigma_z$, $\\Delta z_{\\rm bias}$). For instance, considering ($\\sigma_z$, $z_{\\rm bias}$, $\\Delta\\sigma_z$, $\\Delta z_{\\rm bias}$) $=(0.05(1+z), 0, 0.003, 0.003)$, the degradation factor for $w_0$ and $w_a$ is about $1.5$ (Ma et al. 2006). Then a crude estimate based on our analyses gives the needed $N_{\\rm spec}$ at $z\\sim[3,\\ 4]$ to be about $8600/(1.5)^2\\sim 4000$ to control the bias down to $1\\sigma$ level for $f_{\\rm cata}=1\\%$ and $f_{\\rm sky}=0.25$. More complete analyses should be done for both the degradation and the bias simultaneously taking into account both the small and the large photo-z errors."
    },
    "0812/0812.4200_arXiv.txt": {
        "abstract": "We argue that we may be able to sort out dark matter models in which electrons are generated through the annihilation and/or decay of dark matter, by using a fact that the initial energy spectrum is reflected in the cosmic-ray electron flux observed at the Earth even after propagation through the galactic magnetic field.  To illustrate our idea we focus on three representative initial spectra: (i)monochromatic (ii)flat and (iii)double-peak ones.  We find that those three cases result in significantly different energy spectra, which may be probed by the Fermi satellite in operation or an up-coming cosmic-ray detector such as CALET. ",
        "introduction": "\\label{sec:1} The presence of dark matter has been firmly established by numerous observational data, although we have not yet understood what dark matter is made of. Recent cosmic-ray measurements aiming for indirect dark matter detection may be providing us with insights into dark matter.  The PAMELA data~\\cite{Adriani:2008zr} showed that the positron fraction starts to deviate from a theoretically expected value for secondary positrons around $10$ GeV, and continues to increase up to about $100$\\,GeV. The ATIC collaboration~\\cite{ATIC-new} has recently released the data, showing a clear excess in the total flux of electrons plus positrons peaked around $600 - 700$\\,GeV, in consistent with the PPB-BETS observation~\\cite{Torii:2008xu}.  The excess may be explained by astrophysical sources like pulsars~\\cite{Aharonian(1995),Hooper:2008kg} or microquasars~\\cite{Heinz:2002qj}, although it is not easy to account for the electron flux with a sharp drop-off observed by ATIC~\\footnote{ The nearby pulsars may be able to explain the PAMELA data, though~\\cite{Hooper:2008kg}.  }.  An alternative explanation is the annihilation and/or decay of dark matter.  Indeed, the exciting PAMELA and ATIC/PPB-BETS data has stimulated new directions in dark-matter model building~\\cite{DM-models}. In this letter we take a step further toward sorting out those dark matter models.  One important constraint on the dark matter models comes from the absence of any excess in the anti-proton flux~\\cite{Adriani:2008zq,Yamamoto:2008zz}.  Also, the PAMELA and ATIC/PPB-BETS anomalies in the electrons and positrons suggest that the initial energy spectrum of the electrons and positrons should be hard.  Those observational evidences suggest that electrons (or muons) must be directly produced from dark matter, with the hadronic branch being suppressed.  The models proposed so far are broadly divided into two categories concerning how to suppress the antiproton production. One category is such that the dark matter particle mainly annihilates or decays into leptons. For instance, the dark matter may be a hidden $U(1)_H$ gauge boson decaying into the standard model particles through a kinetic mixing with a $U(1)_{B-L}$ gauge boson; the smallness of quark's quantum number under the $U(1)_{B-L}$ naturally suppresses the anti-proton production~\\cite{Chen:2008yi}. Perhaps the dark matter particle has a lepton number~\\cite{Asaka:2005cn,Chen:2008dh}, or the lepton number as well as a discrete symmetry, which is responsible for the longevity of dark matter, may be slightly broken altogether~\\cite{Takayama:2000uz,Buchmuller:2007ui,Ibarra:2008qg}.  The other category introduces a light particle in the dark sector so that the dark matter particle annihilates or decays into the light particles, which then decay into the standard model particles. If the mass of the light particle is lighter than $1$\\,GeV, the hadronic branch will be suppressed~\\cite{Cholis:2008vb}.  In addition, the presence of such light particle may enhance the annihilation rate to account for the relatively large positron production rate suggested by the PAMELA and ATIC/PPB-BETS data.  Interestingly, the initial energy spectrum of electrons and positrons are quite different in the above two classes of the dark matter models. Such difference in the initial source spectrum may persist in the cosmic-ray electron spectrum observed at the Earth. This will be an important clue to distinguish the dark matter models, since we expect to measure the energy spectrum more precisely in the near future.  For instance, the Fermi satellite~\\cite{FGST} can measure electrons with an energy resolution of about $5$\\% at $20$ GeV to $20$\\% at $1000$ GeV~\\cite{Moiseev:2007js}.  There is also a dedicated experiment proposed to measure the electron spectrum, CALET~\\cite{Torii:2006qb}, which is an instrument to observe very high energy electrons and gamma rays on the Japanese Experiment module Exposure Facility (JEM-EF) of International Space Station (ISS).  The CALET detector has a sensitivity to electrons from $1$\\,GeV to $10$\\,TeV with an energy resolution better than a few \\% for energies greater than 100 GeV. Those measurements will have much more events than the current ATIC/PPB-BETS data. Thus, those promising cosmic-ray electron measurements may help us to distinguish different dark matter models if the ATIC/PPB-BETS excess is indeed due to the dark matter.         In this letter we study the energy spectrum of electrons generated through the annihilation and/or decay of dark matter, particularly paying attention to differences in the energy spectra measured at the solar system for different initial energy spectra. To illustrate our idea we will consider the following initial energy spectra of the electrons: (i) monochromatic (ii) flat and (iii) double-peak ones. We normalize the production rate of the electrons and positrons so as to account for the ATIC/PPB-BETS anomaly.  As we will see, the three cases result in quite different energy spectra at the solar system even after long propagation through the galaxy. We will also discuss whether we can distinguish different dark matter models in the experiments such as Fermi and CALET.  The energy resolution is essential to identify the origin of  the electron and position excess coming from dark matter. While the discontinuity of the spectrum expected for the monochromatic electron spectrum can be identified by Fermi, the other spectra (ii) and (iii) are less prominent, which may leave room for astrophysical explanation since the electron spectrum from supernova remnant will also drop significantly with a certain energy cutoff~\\footnote{ It was discussed in Ref.~\\cite{Hall:2008qu} whether we can distinguish between dark matter and pulsar origins with atmospheric Cherenkov Telescopes. }.  We will however see that the end point of the distribution will be clearly seen with the resolution of a few \\%,  and that it will be possible to distinguish the two models (ii) and (iii) at more than $5 \\,\\sigma$ C.L.  for the expected statistics at CALET. ",
        "conclusions": "\\label{sec:4} In this letter we have estimated the electron energy spectrum at the solar system for the three different initial energy spectra, $dN_e/dE$, given by (\\ref{eq:ini_spectrum1}) - (\\ref{eq:ini_spectrum}), in both decaying and annihilating dark matter scenarios, varying the diffusion model parameters.  We have found that the difference in the initial spectra are reflected in the electron flux measured at the Earth even after long propagation through the galaxy. We have explicitly shown that such behavior is robust against changing diffusion parameters as long as we are concerned with the electron flux in the high energy region (say $E \\gtrsim 400 - 500$\\,GeV).  This observation will enable us to sort out dark matter models by precisely measuring the electron energy spectrum in a future observation such as CALET or the Fermi satellite in operation.  We have also shown that the discontinuity predicted by the monochromatic initial spectrum (i) can be identified with an energy resolution of Fermi, while the other spectra (ii) and (iii) are less prominent as the electron spectrum from supernova remnant will also drop significantly with a certain energy cutoff.  Also we have studied if we can distinguish the flat and double-peak initial spectra, which result in relatively similar energy spectra at the solar system.  We have seen that the end point of the electron spectrum will be clearly seen with the resolution of about  a few \\% assuming the statistics consistent with the ATIC anomaly, and that it will be possible to distinguish the two models at more than $5 \\sigma$ C.L.  for the expected statistics at CALET."
    },
    "0812/0812.1563_arXiv.txt": {
        "abstract": "The first paper in this series explored the effects of altering the chemical mixture of the stellar population on an element by element basis on stellar evolutionary tracks and isochrones to the end of the red giant branch.  This paper extends the discussion by incorporating the fully consistent synthetic stellar spectra with those isochrone models in predicting integrated colors, Lick indices, and synthetic spectra. Older populations display element ratio effects in their spectra at higher amplitude than younger populations. In addition, spectral effects in the photospheres of stars tend to dominate over effects from isochrone temperatures and lifetimes, but, further, the isochrone-based effects that are present tend to fall along the age-metallicity degeneracy vector, while the direct stellar spectral effects usually show considerable orthogonality. ",
        "introduction": "Little is known about the influence of individual element abundances on the integrated flux of a stellar population. As present, it is impossible to derive an accurate age within 10\\% from the integrated light of a cluster-like single-age and single-abundance stellar population, and the primary reason is the complication due to abundance ratio effects \\citep{wor98}.  The secondary reason is that the input ingredients and choices made in even the most sophisticated stellar evolution calculations induce scatter in the results \\citep{char96}.  The tertiary reason is the set of difficulties associated with stellar flux knowledge, such as colors, line strengths, or spectra, that are needed at each stellar evolutionary phase to represent the component stars.  Additional uncertainties exist, such as dust extinction, stellar rotation and activity, the blue straggler frequency, and the effects of close binaries.  Paper I in this series, \\citet{dot07b}, explored the effects of 12 chemical mixtures on the $L$, $T_{\\rm eff}$, and lifetimes of stars along stellar evolutionary isochrones and upon the opacities needed to calculate the stellar models. The mixtures explored were solar, $\\alpha$-element enhanced, and ten cases where only one element at a time was enhanced, for elements C, N, O, Ne, Mg, Si, S, Ca, Ti, and Fe. The mixtures were re-scaled so that the mass fraction of heavy elements, $Z$, was constant. \\footnote{Added to these 12 mixtures were three at non-constant $Z$ and variable $X$ and $Y$ but constant [Fe/H] with 0.2 dex more Carbon, 0.3 dex more Nitrogen and 0.3 dex more Oxygen (see \\citealt{dot07b}).}  The conclusions from that paper could be summarized by splitting the elements that were investigated into three categories. ``Displacers'' include O, C, N, and Ne. These are elements that are abundant but supply less opacity per unit mass than heavier elements. At fixed $Z$, boosting a displacer element will therefore decrease the opacity. This leads to shorter stellar lifetimes and hotter stars. ``Boosters'' include Mg and Si. These elements are good opacity sources, so that boosting a booster will increase opacity and make cooler stars that live somewhat longer. ``Oddball'' elements defy these trends, and include Ca and Ti, both of which make cooler stars, that, nevertheless, have shorter lifetimes. Another oddball is S, which decreases low-temperature opacity but increases high-temperature opacity so that its ultimate effect on temperatures and lifetimes is small. An element worth special mention is Fe, because its high-temperature opacity has considerable structure. One of its opacity peaks corresponds to the temperatures characteristic of the outer edge of the convective core that develops in stars slightly more massive than the sun. Increasing the Fe abundance therefore strongly couples to the onset of convective core overshooting effects, and leads to a strong luminosity effect in the region of the main sequence turnoff. The temperature, luminosity, and lifetime effects are illustrated in Paper I, but further consequences will be illuminated in this paper.  Besides isochrones, the other major ingredient in any population synthesis effort is some representation of stellar flux for each star in the isochrone. The effects of individual element abundances on the spectra of stars have been studied for years, and enormous progress has been made in finding the structure of stellar atmospheres and calculating the emergent flux. The two most-cited programs for constructing model atmospheres of stars are ATLAS \\citep{kur70} and MARCS \\citep{gus75} and we use results from these code in this work.  For spectral synthesis of the emergent flux we use SYNTHE \\citep{kur70}, FANTOM \\citep{coe05,cay91}, and SSG \\citep{bell94}. For ordinary population synthesis, empirical spectra or colors can be used \\citep{vaz99}, but for investigating element-by-element effects, using synthetic spectra is clearly the way forward. Investigations into the effects of individual elements on stellar spectra, but with the intent of applying the results to galaxy spectra, include \\citet{trip95} and \\citet{korn05}, who gauged the effects of ten individual elements on the Lick system of 25 pseudo-equivalent width indices \\citep{wor94b,wo97}. \\citet{serv05} investigated 24 elements in spectra that ran from 3500 \\AA\\ to 9000 \\AA\\ with velocity smoothing appropriate for dynamically hot systems such as elliptical galaxies.  The present paper combines the new isochrones described in Paper I with greatly extended and updated synthetic spectra in order to obtain {\\em ab initio} population synthesis models as a function of {\\em individual} element abundances\\footnote{\\citet{coe07} did a similar investigation with a flat enhancement of all the $\\alpha$-elements over a range of metallicity.}.  All parts of the models, including high temperature and low temperature opacities, stellar evolutionary models, and synthetic spectra, include the altered abundances in the same way.  At present the grid includes only solar $Z$, but calculations are underway to extend to many different abundances.  This paper describes the new spectra and their color index results in $\\S$ 2, Lick index results in $\\S$ 3, concluding with a discussion and summary section. ",
        "conclusions": "We have commenced this project to make stellar population models which incorporate flexible chemistry, so that almost any interesting chemical mixture can be interpolated.  Paper I dealt with the effects on the stellar evolution models, examining the temperature and luminosity changes due to the altered opacites when ten chemical elements are individually tweaked to the end of the red giant branch. In this paper we combine those isochrone effects with the stellar spectral effects in order to investigate their mixed effects on the integrated spectrophotometric indices as well as in their integrated spectra themselves.  We again emphasize here that the models in this study are {\\em incomplete in terms of inclusion of all stellar evolutionary phases and should not be used blindly when comparing to real stellar populations} until the helium-burning phases are properly incorporated.  A version with a full range of metallicity and with horizontal-branch and asymptotic giant branch stars included is planned.  Comparison of our models with observations of Virgo cluster galaxies will be presented as well (Serven et al. in preparation).  Within our spectral coverage (3000 \\AA\\ to 10000 \\AA), we have investigated the broadband color behaviors in the $UBVR_CI_C$ filter set.  As one would expect, older populations show larger spectral effects due to element-by-element abundance changes.  This mostly reflects temperature effects.  We have also confirmed that Mg is the most important $\\alpha$-element that shapes the $U-B$ and $B-V$ colors as already depicted in \\citet{cas04}.  From our investigation of Lick indices using the \\emph{integrated} spectra, we find that (1) CN$_2$ is a useful nitrogen indicator once we have good carbon abundances from G4300 and C$_2$4668, but good silicon (and also titanium for the C$_2$4668) abundances are also needed, (2) Ca4227 is a robust calcium indicator with some good constraints of C, N, and O, (3) Fe4531 and Fe5015 are very useful titanium indicators where an independent iron abundances is provided, (4) Mg $b$ and Fe5406 are good magnesium and iron indicator, respectively, and (5) the variation of individual elements affects the Balmer lines.  We defer the investigation of NaD and TiO$_1$ and TiO$_2$ indices until we have the Na-enhanced isochrones and can model TiO molecular effects with confidence.  Below, we illuminate some points that we have described here with the help of full spectra.  Fe4531 and Fe5015: From Figure 9 we see that both Fe4531 and Fe5015 prove to be good titanium indicators.  They are almost equally sensitive to titanium and iron. Figures 21--23 show the SAURON \\citep{bac01} spectral range, 4810 \\AA\\ to 5350 \\AA, that includes H$\\beta$, Fe5015, Mg$_1$, Mg $b$, Fe5270 and a part of Fe5335.  Here the integrated spectra of solar-scaled, Mg-, Fe-, and Ti-enhanced cases are shown at 2 Gyr with 300 km/sec velocity dispersion normalized at 4750 \\AA.  H$\\beta$ is not influenced much by the enhancement of these elements at this age, as Figure 10 also shows. Mg $b$ is mostly sensitive to Mg, as Figure 21 shows, while Fe5015, Fe5270, and Fe5335 are notably sensitive to Fe-enhancement, although Fe5015 is equally sensitive to titanium, as Figures 22 and 23 display.  Moreover, we have found that the Fe-enhanced and Ti-enhanced spectra show their centroids at different wavelengths in the case of Fe4531, opposite sides from the scaled-solar spectra, because of different locations of iron and titanium lines in their index bandpass.  Figure \\ref{fig24} displays this case\\footnote{ This phenomenon at the stellar level with higher resolution can be seen at http://astro.wsu.edu/hclee/Fe45$\\_$4000mp00gp25$\\_$ti.pdf  and http://astro.wsu.edu/hclee/Fe45$\\_$4000mp00gp25$\\_$fe.pdf}. We have preliminary hints from high-S/N, high-resolution spectra of Virgo cluster galaxies taken at the Kitt Peak 4-m telescope (Serven et al. in preparation) that galaxies with the higher velocity dispersion tend to have their centroids at the longer wavelengths, indicating higher Ti-enhancements.  If confirmed, this would indicate a Ti-$\\sigma$ relation, similar to the well-known Mg-$\\sigma$ relation (see also \\citealt{mil00}).  Balmer lines:  H$\\beta$: The bottom right panel of Figure 10 shows that, in an $\\alpha$-enhanced mixture at the fixed [Fe/H] (Alpha+), H$\\beta$ becomes weaker.  This sounds a bit different from the conclusion of LW05 where H$\\beta$ is $\\alpha$-element insensitive.  As we can see from Table 5, the effects that we see from the bottom right panel of Figure 10 are mostly isochrone effects, especially at young ages.  The H$\\beta$ index becomes weaker because of the decrease of temperature of stellar isochrones from the $\\alpha$-enhanced case at fixed [Fe/H] (see also \\citealt{coe07}; \\citealt{sch07}).  LW05 used only solar-scaled isochrones at solar and super-solar metallicities, although these authors employed modified stellar spectra because of the $\\alpha$-enhancement.  Moreover, it is worth emphasizing that $\\alpha$-enhanced isochrones at fixed $Z$ are now no longer significantly hotter than the solar-scaled ones.  This has been revealed by Paper I, in its figure 11, after incorporating the latest development of low-temperature opacities by \\citet{fer05}.  If this behavior is indeed true, it will certainly affect the conclusions of earlier works (e.g., \\citealt{tm03}) that are based upon the \\citet{sal00} isochrones, which were later shown by \\citet{wei06} to present problems.  H$\\gamma$ and H$\\delta$: Fe-enhancement makes these indices significantly weaker, especially for the broader indices, H$\\gamma_A$ and H$\\delta_A$ (see also \\citealt{gs08}).  Figures 17 and 18 show that this is because of Fe-lines that inhabit the continuum region.  It is suggested therefore that the narrower indices (H$\\gamma_F$ and H$\\delta_F$) may serve as better age indicators in case the stellar systems are of wildly different chemical mixtures.  Our goals are to narrow down the uncertainties of mean age estimations to 10\\% when derived from a single integrated light spectrum.  We believe that we can reach these goals as we are able to take into account the effects from individual chemical elements.  The ultimate goal then obviously would be the understanding of the chemical enrichment and star formation histories of galaxies, stepping away from the study of mean ages and mean metallicities to begin to tackle the more realistic problem of multi-age, multi-metallicity stellar populations as they exist in galaxies.  We believe that several efforts such as SPECKMAP \\citep{ocv06}, MOPED \\citep{pan07}, and STARLIGHT \\citep{cid08} will be even more useful once they use sophisticated stellar population models with flexible chemistry as inputs."
    },
    "0812/0812.2756_arXiv.txt": {
        "abstract": "This paper studies the formation and evolution of binary supermassive black holes (SMBHs) in rotating galactic nuclei, focusing on the role of stellar dynamics. We present the first $N$-body simulations that follow the evolution of the SMBHs from kiloparsec separations all the way to their final relativistic coalescence, and that can robustly be scaled to real galaxies. The $N$-body code includes post-Newtonian ($\\PN$\\,) corrections to the binary equations of motion up to order $2.5$; we show that the evolution of the massive binary is only correctly reproduced if the conservative $1 \\PN$ and $2 \\PN$ terms are included. The orbital eccentricities of the massive binaries in our  simulations are often found to remain large until shortly before coalescence. This directly affects not only their orbital evolution rates, but has important consequences as well for the gravitational waveforms emitted during the relativistic inspiral. We estimate gravitational wave amplitudes when the frequencies fall inside the band of the (planned) Laser Interferometer Space Antennae (LISA). We find significant contributions --- well above the LISA sensitivity curve --- from the higher-order harmonics. ",
        "introduction": "Supermassive black holes (SMBHs) are commonly observed at the centers of nearby galaxies, and the existence of quasars at redshifts $z \\approx 6$ implies that many of these SMBHs reached nearly their current masses at very early times. Bright galaxies are believed to form via hierarchical merging of smaller galaxies, and galaxy mergers lead inevitably to the formation of binary SMBHs (Begelman, Blandford \\& Rees 1980). At their formation, such SMBH binaries typically have separations $a \\approx a_{\\mathrm h}$, where  \\begin{equation} a_{\\mathrm h} \\equiv \\frac{G \\mu}{4 \\sigma^2} \\approx 3.3\\,{\\mathrm{pc}}~\\frac{q}{(1+q)^2} \\left( \\frac{m_1+m_2}{10^8\\, {\\mathrm M}_\\odot} \\right)^{0.6}\\  , \\end{equation}   \\noindent \\citep[see, e.g.,][]{MM05}.  This ``hard binary separation'' $a_{\\mathrm h}$ is defined in the standard way as the value of the binary semi-major axis at which passing stars are ejected with velocities high enough to unbind them from the SMBHs; $a_{\\mathrm h}$ is a function of the binary reduced mass $\\mu \\equiv m_1\\,m_2/(m_1+m_2)$ and the ambient stellar velocity dispersion $\\sigma$ ($q \\equiv m_2/m_1 \\le 1$). At least one binary SMBH has probably been observed, with a projected separation of about $7$\\,pc \\citep{Rod06}. The quasi-periodic outbursts of the quasar OJ287 have also plausibly been modeled as a binary system with separation $\\sim 0.05$ pc and eccentricity $\\sim 0.6$ \\citep{Val07}. A few other galactic systems are known to contain dual SMBHs with separations of a few kiloparsec, presumably the precursors of binary SMBHs \\citep{kom2003,bia08}.  Binary SMBHs may eventually coalesce, but only after stellar- and/or gas-dynamical processes have first brought the two SMBHs to separations small enough ($\\sim 10^{-3}$\\,pc) that gravitational wave emission is effective.  Whether Nature typically succeeds in overcoming this ``final parsec problem,'' or whether long-lived binary SMBHs are the norm, is currently unknown. Persistence of the binaries would have a number of potentially important consequences, including ejection of SMBHs from galaxy nuclei via three-body interactions and a reduction in the mean ratio of SMBH mass to galaxy mass \\citep[Volonteri, Haardt \\& Madau 2003;][]{Madau04,Libes06}. In addition, plans to detect gravitational waves from the final inspiral of binary SMBHs by space-based observatories such the Laser Interferometer Space Antennae (LISA) (Sesana, Volonteri \\& Haardt 2007) might need to be re-considered if evolution of the binaries typically stalls at large separations.  In a spherical, gas-free galaxy, the evolution of a binary SMBH slows down  at separations of $\\sim a_{\\mathrm h}$ because stars on orbits that intersect the binary -- so-called ``loss-cone orbits'' -- are depleted in just a few galaxy crossing times (e.g., Begelman et al. 1980, Makino et al. 1993, Makino 1998, Berczik, Spurzem \\& Merritt 2005; Merritt 2006). A massive binary can continue to evolve in such a galaxy only if the loss-cone orbits are somehow repopulated. Gravitational scattering is one possible mechanism for loss-cone repopulation, but it is only effective in low-luminosity galaxies with central relaxation times below $\\sim 10$\\,Gyr \\citep{Yu02, MMS07}. Dense concentrations of gas can substantially accelerate the evolution of a massive binary by increasing the drag on the individual SMBHs \\citep{Escala04,Escala05,Dotti07}. The plausibility of such gas accumulations, with masses comparable to the masses of the SMBHs, is unclear however, particularly in the most massive galaxies. Galaxy merger simulations including gas have followed the two SMBHs until separations of a few tens of parsecs, roughly the hard binary separation \\citep{Kazant05,Mayer07}. These simulations contain useful information about the formation and early evolution of SMBH binaries but -- at least so far -- do not have the resolution needed to address the final parsec problem.  An alternative pathway exists for binary SMBHs to evolve beyond $a \\approx a_{\\mathrm h}$, even in gas-free galaxies with long relaxation times.  If the galaxy potential is non-axisymmetric, many of the stars will be on centrophilic orbits, i.e., orbits that pass near the center of the galaxy once per orbital period \\citep{GB85}. If even a small fraction of a galaxy's mass is associated with such orbits, the ``feeding'' rate of a central binary can be enormously enhanced compared with the rates in spherical galaxies \\citep{MPoon04}. \\citet{BMSB06} explored this pathway in $N$-body simulations containing particle numbers up to one million; the initial galaxy models were rotating, and in some cases the rotation was rapid enough to induce the formation of a triaxial bar. Evolution of the SMBH binary was observed not to stall at $a \\approx a_{\\mathrm h}$ in the triaxial models; furthermore the evolution rates exhibited no discernible dependence on particle number, as predicted if the mechanism of loss-cone refilling is collisionless. To date, the \\citet{BMSB06} simulations are the only ones that have successfully followed the evolution of a binary SMBH all the way from galactic to sub-parsec scales, and that can robustly be scaled to real galaxies due to the lack of an appreciable $N$-dependence of the evolution rates.  The final binary separation in the \\citet{BMSB06} simulations was $\\sim 0.05 a_{\\mathrm h}$, just small enough that gravitational wave emission would induce coalescence in $10$\\,Gyr (assuming a scaling that assigns a mass $10^8 {\\mathrm M}_\\odot$ to the binary SMBH). In this paper, we repeat the Berczik et al. simulations using a new $N$-body code that incorporates the post-Newtonian ($\\PN$~) corrections to the equations of motion of the SMBHs. This allows us to follow the evolution of the binary all the way to final coalescence. In addition, we are able to estimate the strength of the gravitational waves emitted during the final inspiral. A key result is that the eccentricity of the binary can remain high until shortly before coalescence. Both the strength of the emitted gravitational waves, and the timescale for coalescence, are strongly dependent on the binary eccentricity. Highly eccentric black hole binaries would represent appropriate candidates for forthcoming verification of gravitational radiation through the planned mission of LISA.  The outline of this paper is as follows: in Section~2 we give a description of the numerical methods and initial models used in this work. In Section~3 we then present the results of a set of $N$-body simulations using both the classical, i.e., Newtonian gravity (Sec.~3.1) and the relativistic, i.e., post-Newtonian gravity (Sec.~3.2). In Section~4 we discuss our results in a broader astrophysical context, focusing on the SMBH merger as sources of gravitational wave emission. In Section~5 we finish with the conclusions. ",
        "conclusions": "We present the first $N$-body models which self-con\\-sis\\-tently follow the evolution of binary supermassive black holes in merged galaxies from kiloparsec separations down to gravitational-wave-induced coalescence, and in which the early evolution of the binary is driven by collisionless loss-cone repopulation, allowing the results to robustly be scaled to real galaxies. Our simulations include post-Newtonian corrections to the equations of motion of the SMBH binary up to order $2.5 \\PN$; we show that inclusion of the energy-conserving $1 \\PN$ and $2 \\PN$ terms is also crucial for obtaining the correct time dependence of the binary orbital parameters. We identify evolutionary phases in which the binary is still evolving due to stellar encounters but in which relativistic corrections to its two-body motion are also important, thereby showing that the consideration of all $\\PN$ orders {\\em in} the $N$-body simulations is necessary for an accurate prediction of its orbital elements evolution. The SMBH binaries in our simulations often form with large eccentricities, and these high eccentricities are maintained during the Newtonian phases of the evolution. We show that the gravitational wave signal measured by an interferometer like LISA would contain significant contributions from high-order harmonics from GWs emitted from binaries at cosmological distances,  and that the nature of the signal can depend strongly on the dynamical history of the binary prior to its entering the gravitational wave dominated regime.  We have shown that supermassive black hole binaries in galactic nuclei can overcome the stalling barrier and will reach the relativistic coalescence phase in a timescale shorter than the age of the universe. A gravitational wave signal expected for the LISA satellite from these SMBH binaries is expected, in particular due to the high eccentricity of the SMBH binary when entering the relativistic coalescence phase. We present for the first time a comprehensive set of models which cover self-consistently the transition from the Newtonian dynamics, dynamical friction phase (with yet unbound SMBH binary) to the situation when relativistic, post-Newtonian corrections start to influence the relative SMBH motion. After the shrinking time scale becomes very short the binary decouples from the rest of the galactic nucleus and can be treated as a relativistic two-body problem. We follow this evolution formally to the coalescence of the two black holes using $\\PN$ terms of up to order $2.5 \\PN$ and determine the gravitational wave emission in different modes relative to the LISA sensitivity curve. We find that the orbital parameters of a SMBH binary, when entering the LISA band, depends on the previous dynamical history - in particular there exists a phase where the SMBH binary is still partially coupled to the stellar environment via three-body encounters, but relativistic corrections to its two-body motion already play a role.  Since purely Newtonian models of SMBH binaries in rotating galactic nuclei \\citep{BMSB06} have shown a {\\color{black} quasi-$N$-independence of the stellar-dynamical driven hardening rates, we conclude that the results obtained in this paper regarding the time required until relativistic merger holds for galactic nuclei with realistic parameters.} In this work we limited ourselves to models with particle numbers up to $50$k only, however, due to the independence of the results in the Newtonian phase from $N$ in the current work and in \\citet{BMSB06}, we do not expect any significant changes in our main results for models with $N>50$k.  It should be noted, however, that our results show a strong dependence on initial conditions, and that our initial model is a very simple approximation to a post-merger galactic nucleus. Therefore our conclusion from this work can only be that it is possible to reach relativistic coalescence in a reasonable time ($10^8$ years). Any prediction of event rates for LISA would, however, require a more careful estimate of the distribution of parameters for a realistic set of mergers, e.g. by taking data from semi-analytic merger trees, similar to \\cite{Sesana05}. This is the subject of presently ongoing work.  We have also examined the effect of using $2.5 \\PN$ alone or {\\em full} $\\PN$ corrections. Simple two-body $\\PN$ experiments show significant differences of the BHs trajectories, and we find generally a dependence of the merging time in the two cases. Different orbital parameters of the SMBH binary in the final phase would have a major impact on the waveforms of the GWs. However, the computation of the latter also requires simulations with higher $\\PN$ corrections, at least up to order ${\\cal O}({c}^{-6})$, i.e., $3 \\PN$ in the equations of motion. This is beyond the scope of the current paper.  Results on whether our simulated SMBH binary will fall into the LISA frequency band and at which frequencies will be discussed in a cosmological context in a forthcoming paper."
    },
    "0812/0812.2610_arXiv.txt": {
        "abstract": "I review  several issues related to statistical description of gravitating systems  in both static and expanding backgrounds. After briefly reviewing the results for the static background, I concentrate on gravitational clustering\\index{gravitational clustering}  of collision-less particles in an expanding universe. In particular, I describe (a) how the non linear mode-mode coupling transfers power from one scale to another in the Fourier space if the initial power spectrum is sharply peaked at a given scale and (b) what are the asymptotic characteristics  of gravitational clustering that are independent of the initial conditions.  Numerical simulations as well as analytic work shows that  power transfer leads to a universal power spectrum at late times, somewhat reminiscent of the existence of Kolmogorov spectrum in fluid turbulence. ",
        "introduction": "\\label{intro}  The statistical mechanics of systems dominated by gravity  has close connections with areas of condensed matter physics, fluid mechanics, re-normalization group etc. and poses an incredible challenge as regards the basic formulation.  The ideas also find application in many different areas of astrophysics and cosmology, especially in the study of globular clusters, galaxies and gravitational clustering in the expanding universe. (For a  overall review of statistical mechanics of gravitating systems, see \\cite{tppr}, \\cite{textone}, \\cite{texttwo}; review of gravitational clustering in expanding background is available in \\cite{tpiran} and in several textbooks in cosmology \\cite{cosmotext};  for a sample of different attempts to understand these phenomena by different groups; see \\cite{chavanis}, \\cite{sanchez}, \\cite{vala}, \\cite{fola}, \\cite{botta}, \\cite{roman}  and the references cited therein.) It will be useful to begin with a broad overview and a  description of the issues which will be addressed in this article.  In Newtonian theory, the gravitational force can be described as a gradient of a scalar potential and the evolution of a set of particles under the action of gravitational forces can be described the equations \\begin{equation} \\ddot{\\bf x}_i = - \\nabla \\phi ({\\bf x}_i, t); \\quad \\nabla^2 \\phi = 4 \\pi G \\sum_i m_i \\delta_D ({\\bf x} -{\\bf x}_i) \\end{equation} where  ${\\bf x}_i$ is the position of the $i-$th particle, $m_i$ is its mass.  For an isolated system with sufficiently large number of particles, it is useful to investigate whether some kind of statistical description  of such a system is possible. Such a description, however,  is complicated by  the long range, \\index{long range} unscreened, nature of gravitational force. If a self gravitating system is divided into two parts, the total energy of the system cannot be expressed as the sum of the gravitational energy of the components. The conventional results in statistical physics are based on the extensivity of the energy \\index{extensivity of the energy} which is clearly invalid for gravitating systems. To construct the statistical description of such a system, one must begin with the construction of the micro-canonical ensemble describing such a system. If the Hamiltonian of the system is $H(p_i, q_i)$ then the volume $g(E)$ of the constant energy surface  $H(p_i, q_i) =E$ will be of primary importance in the  micro-canonical ensemble\\index{micro-canonical ensemble}. The logarithm of this function will give the entropy $S(E) = \\ln g(E) $ and the temperature of the system will be $T(E)\\equiv \\beta(E)^{-1} = (\\pder{S}{E})^{-1}$.  Systems for which a description based on canonical ensemble is possible, the Laplace transform of $g(E)$ with respect to a variable $\\beta$ will give the partition function $Z(\\beta)$. It is, however, trivial to show that gravitating systems of interest in astrophysics cannot be described by a canonical ensemble \\cite{tppr}, \\cite{dlbone}, \\cite{dlbtwo}. Virial theorem holds for such systems  and we have $(2K+U) =0$ where $K$ and $U$ are the total kinetic and potential energies of the system. This leads to $E=K+U= -K$; since the temperature of the system is proportional to the total kinetic energy, the  specific heat\\index{specific heat} will be negative: $C_V \\equiv (\\pder{E}{T})_V \\propto (\\pder{E}{K}) < 0$.   On the other hand, the specific heat of any system described by a canonical ensemble $C_V = \\beta^2 <(\\Delta E)^2>$ will be positive definite. Thus one cannot describe self gravitating systems of the kind we are interested in by  canonical ensemble\\index{canonical ensemble}.  One may attempt to find the equilibrium configuration for self gravitating systems by maximizing the entropy $S(E)$ or the phase volume $g(E)$. It is again easy to show that no global maximum for the entropy exists for classical point particles interacting via Newtonian gravity. To prove this, we only need to construct a configuration with   arbitrarily high entropy which can be achieved as follows: Consider a system of $N$ particles initially occupying a region of finite volume in phase space and total energy $E$. We now move  a small number of these particles (in fact, a pair of them, say, particles 1 and 2  will do) arbitrarily close to each other. The potential energy of interaction of  these two particles, $-Gm_1m_2/r_{12}$, will become arbitrarily high as $r_{12}\\to 0$. Transferring some of this energy to the rest of the particles, we can increase their kinetic energy without limit. This will clearly increase the phase volume occupied by the system without bound. This argument can be made more formal by dividing the original system into a small, compact core and a large diffuse halo and allowing the core to collapse and transfer the energy to   the halo.   The absence of the global maximum for entropy --- as argued above --- depends on the idealization that there is no short distance cut-off in the interaction of the particles, so that we could take the limit $r_{12}\\to 0$. If we assume, instead, that each particle has a minimum radius $a$, then the typical lower bound to the gravitational potential energy contributed by a pair of particles will be $-Gm_1m_2/2a$. This will put an upper bound on the amount of energy that can be made available to the rest of the system.  We have also assumed that part of the system can expand without limit --- in the sense that any particle with sufficiently large energy can move to arbitrarily large distances. In real life, no system is completely isolated and eventually one has to assume that the meandering particle is better treated as a member of another system. One way of obtaining a truly isolated system is to confine the system  inside a spherical region of radius $R$ with, say, reflecting wall. (Most of our discussion is confined to 3-dimensions and the situation is diffrent in 2-dimensions; see e.g \\cite{two-d})  The two cut-offs $a$ and $R$ will make the upper bound on the entropy finite, but even with  the two cut-offs, the primary nature of gravitational instability cannot be avoided. The basic phenomenon described  above  (namely, the formation of a  compact core\\index{compact core} and a  diffuse halo\\index{diffuse halo}) will still occur since this is the direction of increasing entropy. Particles in the hot diffuse component will permeate the entire spherical cavity, bouncing off the walls and having a kinetic energy which is significantly larger than the potential energy. The compact core will exist as a gravitationally bound system with very little kinetic energy. A more formal way of understanding this phenomena is based on  the virial theorem for a system with a short distance cut-off confined to a sphere of volume $V$. In this case, the virial theorem will read as \\cite{textone} \\begin{equation} 2T + U = 3PV + \\Phi \\label{modvirial} \\end{equation} where $P$ is the pressure on the walls and $\\Phi$ is the correction to the potential energy arising from the short distance cut-off. This equation can be satisfied in essentially three different ways. If $T$ and $U$ are significantly higher than $3PV$, then we have $2T + U \\approx 0$ which describes a self gravitating systems in standard virial equilibrium but not in the state of maximum entropy. If $T \\gg U$ and $3PV\\gg \\Phi$, one can have $2T \\approx 3PV$ which describes an ideal gas with no potential energy confined to a container of volume $V$; this will describe the hot diffuse component at late times. If  $T \\ll U$ and $3PV\\ll \\Phi$, then one can have $U\\approx \\Phi$ describing the compact potential energy dominated core at late times. In general, the evolution of the system will lead to the production of the core and the halo and each component will satisfy the virial theorem in the form (\\ref{modvirial}). Such an asymptotic state with two distinct phases  \\cite{aaron}  is quite different from what would have been expected for systems with only short range interaction. Considering its importance, I shall briefly describe in section \\ref{sec:phases} a toy model which captures the essential physics of the above system in an exactly solvable context.  The above discussion focussed on the existence of global maximum to the entropy and we proved that it does not exist in the absence of two cut-offs. It is, however, possible to have {\\it local} extrema of entropy which are not global maxima. Intuitively, one would have expected the distribution of matter in  the configuration which is a local extrema of entropy to be described by a Boltzmann distribution, with the density given by $\\rho ({\\bf x}) \\propto \\exp[-\\beta \\phi({\\bf x})]$ where $\\phi$ is the gravitational potential related to $\\rho$ by Poisson equation. This is indeed true; for a formal proof see \\cite{tppr}.  This configuration is usually called the  isothermal sphere\\index{isothermal sphere} (because it can be shown that, among all solutions to this equation, the one with spherical symmetry maximizes the entropy) and  it is a local maximum of entropy.  The second (functional)  derivative of the entropy with respect to the configuration variables will determine whether the local extremum of entropy is a local maximum or a saddle point \\cite{antonov}, \\cite{tpapjs}.  The relevance of the long range of gravity in all the above phenomena can be understood by studying model systems with an attractive potential varying as $r^{-\\alpha}$ with different values for $\\alpha$. Such studies confirm the results and interpretation given above; (see \\cite{ispo} and references cited therein).  Let us now consider the situation in the context of an expanding background. There is considerable amount of observational evidence to suggest that one of the dominant energy densities in the universe is contributed by self gravitating point particles. The smooth average energy density of these particles drive the expansion of the universe while any small deviation from the homogeneous energy density will cluster gravitationally. [For a review of cosmology from a contemperorary perspective, see e.g., \\cite{cosmoreview}] One of the central problems in cosmology is to describe the non linear phases of this gravitational clustering starting from a initial spectrum of density fluctuations. It is often enough (and necessary) to use a statistical description and  relate different statistical indicators (like the  power spectra\\index{power spectra}, $n$th order  correlation functions\\index{correlation functions} ....) of the resulting density distribution to the statistical parameters (usually the power spectrum) of the initial distribution. The relevant scales at which gravitational clustering is non linear are less than about 10 Mpc (where 1 Mpc = $3\\times 10^{24}$ cm is the typical separation between galaxies in the universe) while the expansion of the universe has a characteristic scale of about few thousand Mpc. Hence, non linear gravitational clustering in an expanding universe can be adequately described by Newtonian gravity provided the rescaling of lengths due to the background expansion is taken into account. This is easily done by introducing a  {\\it proper} coordinate\\index{proper coordinate} for the $i-$th particle ${\\bf r}_i$, related to the  {\\it comoving} coordinate\\index{comoving coordinate}  ${\\bf x}_i$, by ${\\bf r}_i = a(t) {\\bf x}_i$ with $a(t)$ describing the stretching of length scales due to cosmic expansion. The Newtonian dynamics works with the proper coordinates ${\\bf r}_i$ which can be translated to the behaviour of the comoving coordinate ${\\bf x}_i$ by this rescaling. [This implies that, for all practical purposes, we are still in the domain of Newtonian gravity. There is a far deeper connection between thermodynamics and gravity \\cite{grav-thermo} in the general relativistic domain which we will not discuss in these lectures.]   As to be expected, cosmological expansion  completely changes the nature of the problem because of several new factors which come in: (a) The problem has now become time dependent and it will be pointless to look for equilibrium solutions in the conventional sense of the word. (b) On the other hand, the expansion of the universe has a civilizing influence on the particles and  acts counter to the tendency of gravity to make systems unstable. (c) In any small local region of the universe, one would assume that the conclusions describing a finite gravitating system will still hold true approximately. In that case, particles in any small sub region will be driven towards configurations of local extrema of entropy (say, isothermal spheres) and towards global maxima of entropy (say, core-halo configurations).  An extra feature comes into play as regards the expanding halo from any sub region. The expansion of the universe acts as a damping term in the equations of motion and drains the particles of their kinetic energy --- which is essentially the lowering of temperature of any system participating in cosmic expansion.  This, in turn, helps gravitational clustering since the potential wells of nearby sub regions can capture particles in the expanding halo of one region when the kinetic energy of the expanding halo  has been sufficiently reduced.  The actual behaviour of the system will, of course, depend on the form of $a(t)$. However, for understanding the nature of clustering, one can take $a(t) \\propto t^{2/3}$ which describes a matter dominated universe with critical density. Such a power law has the advantage that there is no intrinsic scale in the problem. Since Newtonian gravitational force is also scale free, one would expect some scaling relations to exist in the pattern of gravitational clustering. Converting this intuitive idea into a concrete mathematical statement turns out to be non trivial and difficult.   It is clear that cosmological expansion  introduces several new factors into the problem when compared with the study of statistical mechanics of isolated gravitating systems. (For a general review of statistical mechanics of gravitating systems, see \\cite{tppr}. For a sample of different approaches, see \\cite{sample} and the references cited therein. Review of gravitational clustering in expanding background is also available in several textbooks in cosmology \\cite{cosmotext,lssu}.) Though this problem can be tackled in a `practical' manner using high resolution numerical simulations (for a review, see \\cite{numsim}), such an approach hides the physical principles which govern the behaviour of the system. To understand the physics, it is necessary to attack the problem from several directions using analytic and semi analytic methods. Several such attempts exist in the literature based on Zeldovich(like) approximations \\cite{za}, path integral and perturbative techniques \\cite{pi}, nonlinear scaling relations \\cite{nsr} and many others. In spite of all these it is probably fair to say that we still do not have a clear analytic grasp of this problem, mainly because each of these approximations have different domains of validity and do not arise from a central paradigm.  I propose to attack the problem from a different angle, which has not received much attention in the past. The approach begins from the dynamical equation for the the density contrast in the Fourier space  and casts it as an integro-differential equation. Though this equation is known in the literature  (see, e.g. \\cite{lssu}), it has received very little attention  because it is not `closed' mathematically; that is, it involves variables which are not natural to the formalism and thus further progress is difficult. I will, however, argue that there exists a natural closure condition for this equation  based on Zeldovich approximation thereby allowing us to write down a \\textit{closed} integro-differential equation for the gravitational potential in the Fourier space.  It turns out that this equation can form the basis for several further investigations some of which are described in ref. \\cite{gc1} and in the second reference in \\cite{tppr}. Here I will concentrate on just two specific features, centered around the following issues:   \\begin{itemize} \\item If the initial power spectrum is sharply peaked in a narrow band of wavelengths, how does the evolution transfer the power to other scales? In particular, does the non linear evolution in the case of gravitational interactions lead to a universal power spectrum (like the Kolmogorov spectrum in fluid turbulence)?  \\item What is the nature of the time evolution at late stages? Does the gravitational clustering at late stages wipe out the memory of initial conditions and evolve in a universal manner? \\end{itemize} Fair amount of progress can be made as regards these questions using the integro-differential equation mentioned above and some of these aspects will be discussed in detail. ",
        "conclusions": ""
    },
    "0812/0812.2929_arXiv.txt": {
        "abstract": "We  use SDSS-DR4 photometric  and spectroscopic  data out  to redshift $z\\sim0.1$ combined with ROSAT All  Sky Survey X-ray data to produce a sample of  twenty-five fossil groups  (FGs), defined as  bound systems dominated by a single,  luminous elliptical galaxy with extended X-ray emission.   We  examine  possible  biases introduced  by  varying  the parameters  used  to define  the  sample  and  the main  pitfalls  are discussed.  The  spatial density of  FGs, estimated via  the $V/V_{\\rm MAX}$ test, is 2.83 $\\times$10$^{-6} h_{75}^3$ Mpc$^{-3}$ for L$_{X} > 0.89 \\! \\times \\!  10^{42} h_{75}^{-2}$ erg s$^{-1}$ consistent with Vikhlinin  et   al.   (1999),  who  examined   an  X-ray  overluminous elliptical galaxy sample (OLEG).  We compare the general properties of FGs  identified  here  with  a  sample  of  bright  field  ellipticals generated  from   the  same  dataset.   These  two   samples  show  no differences in  the distribution  of neighboring faint  galaxy density excess, distance from the red sequence in the color-magnitude diagram, and  structural   parameters  such  as  a$_{4}$   and  internal  color gradients.  Furthermore, examination of stellar populations shows that our   twenty-five   FGs   have   similar  ages,   metallicities,   and $\\alpha$-enhancement as the  bright field ellipticals, undermining the idea that these systems represent fossils of a physical mechanism that occurred at  high redshift.  Our  study reveals no  difference between FGs and field ellipticals, suggesting that FGs might not be a distinct family of true fossils, but rather the final stage of mass assembly in the Universe. ",
        "introduction": "Searching for  preserved remnants of physical  processes that occurred in the  cosmic past is one  of cosmologists' main  tools in developing our  understanding of how  galaxies, groups,  and clusters were formed and evolved, and ultimately came to define the large scale structure we observe. Recent studies have shown that there seems to be a  class of  systems with  exceptional preservation,  known  as fossil groups (FGs, Ponman  et al. 1994). These systems  consist of isolated, luminous early-type  galaxies embedded in an extended  X-ray halo. The main motivation for seeking such  galactic systems stems from the fact that the merger  time scales for $\\rm L > L^*$  group galaxies is much shorter than the cooling time  scales for the hot gas component within which these bright galaxies are  embedded (e.g. Barnes 1989, Ponman \\& Bertram 1993). In  this simplistic view FGs are  a natural consequence of the merging  process in normal groups and therefore  can be used to trace the history of coalescence in the not-so-high-redshift Universe.  From the observational viewpoint, FGs are seen in the optical as large elliptical galaxies,  but with X-ray luminosities  comparable to those of an entire group of galaxies ($L_X > 4.4 \\times 10^{42} ergs/s$; see Jones et  al.  2003,  hereafter JO03).  To  date, no clear  picture of their origin has emerged from  the collected data.  Two main scenarios have been suggested: FGs result  from the complete merging of galaxies that constituted  a loose group  in the past,  collapsing at an early epoch, but never becoming incorporated into clusters (Hausman \\& Ostriker 1978;  Ponman et al.  1994; Jones et al.   2000; Khosroshahi, Jones \\& Ponman  2004); or they originate from  a region that inhibits formation of  $L^*$ galaxies  in these groups  leading to  an atypical galaxy luminosity function (LF, Mulchaey \\& Zabludoff 1999).  These  systems  have  been  detected  out to  redshifts  of  at  least $z\\sim0.6$ (Ulmer et al.~2005),  have high $M/L$ ratios suggesting low star formation efficiency (Vikhlinin et al. 1999) and represent 8\\% to 20\\%  of all  systems with  similar X-ray  luminosities.  Ten  FGs are already known and  well studied from optical and  X-ray data (Mulchaey \\& Zabludoff 1999, Romer et al. 2000, hereafter R00, Jones et al 2003, Khosroshahi,  Jones, \\&  Ponman~2004, Sun  et al.   2004,  Yoshioka et al. 2004, Ulmer et al.  2005).  A recent paper by Santos et al.~(2007) presents a  sample of 34  FG candidates based  on SDSS data.   In this work, we  use optical (SDSS) and  X-ray (ROSAT All Sky  Survey -- RASS) data to  define a sample of  FG candidates. We select  FGs following a strategy similar  to Jones et  al.  (2003), using  spectroscopic (age, metallicity, $\\rm \\alpha/Fe$  enhancement) and photometric ($\\rm a_4$, color  gradient) parameters  to  further constrain  what  should be  a FG. The main  goal is to establish a link  between FGs, compact groups and isolated  ellipticals based on their global  properties and simple expectations,  if  mergers are  the  dominant  events determining  the evolution of a FG.  This  paper is  organized as  follows. Section  2 presents  the galaxy catalog  used to  search for  FGs, describing  the parameters  and the X-ray measurements that define the search. The actual selection of FGs is  then detailed.  Section  3  presents the  control  sample used  to compare  FG properties  to  those of  \"normal\"  ellipticals, which  in principle  did  not  result  from  merging. Section  4  discusses  the magnitude distribution of the  brightest and second brightest galaxies of FGs,  while in  Section 5 we  characterize the distribution  of the faint galaxies around  FGs.  Since FGs are thought  to be old systems, in Section 6 we examine the colors of FGs relative to the red sequence of early type  galaxies.  Section 7 deals with  the internal structure of  seed galaxies  of  FGs,  as  measured by  their  structural parameters  and internal  color gradients.  In Section  8,  we analyze stellar population properties of  FGs, by looking at the distributions of age, metallicity, and  $\\alpha$ enhancement.  Finally, in Section 9, we summarize  and discuss the  main results.  Throughout the  paper we use a cosmology  with $\\Omega_{\\rm m} = 0.3$,  $\\Omega_\\Lambda = 0.7$, and $\\rm H_0 = 75 \\, \\rm km \\, s^{-1} \\, Mpc^{-1}$. ",
        "conclusions": "\\label{sec:DISCUSSION}  In the current hierarchical model of galaxy formation, the basic idea is that massive galaxies form  later as the result of the merging history  of their  dark matter  halos,  but their  stars form  earlier (downsizing).   In this sense  special attention  is being  focused on massive  (luminous) galactic systems,  whether they  reside in  low or high  density locations.   We  may find  fossils  at different  cosmic epochs,  and they  provide important  clues that  help  us distinguish between nature and nurture.  Our main results from this paper are:  We found  578 candidates following our  optical photometric definition of FGs in  the redshift range 0.05 $<$ z  $<$ 0.095.  After discarding incorrect morphological  classifications and AGN we are  left with 102 FGs for which there is a significant X-ray detection in ROSAT (maximum total flux of 10$^{-14} \\rm erg \\, cm^{-2} s^{-1}$).  Only twenty-five FGs  were  selected  with  likely  extended  X-ray  emission  and  not associated with  rich clusters. These  constitute our final  sample of FGs. Applying  the V/V$_{\\rm max}$ test  we find a  spatial density of 2.83  $\\times$10$^{-6}  \\rm  {h_{75}}^{3}$ Mpc$^{-3}$,  comparable  to three other  independent observational  studies (V99, RO00,  and JO03) and to the recent analysis  of the Millennium Simulation by Dariush et al. (2007).  As shown  in Figure 12,  the brighter FGs  always have $\\Delta  \\rm M$ greater  than $\\sim  3.0$ mag  while fainter  FGs are  limited  by the completeness of  the SDSS  catalog. The lack  of systems in  the lower left part of the figure,  at magnitudes brighter than $^{0.1}M_{\\rm r} \\sim -23$,  might be  caused by  selection effects as  well as  by the shape  of the  galaxy LF.   Thus,  it is  not clear  from the  present analysis if the absence of systems in the lower left part of Fig.12 is related to  some physical characteristics  of the sample. It  would be reasonable to expect  that as mergers progress the  LF of such systems will  have  a larger  gap  between the  first  and  the second  ranked galaxies.  The density excess around FGs is fully consistent with what we measure for field galaxies, implying that the environments around both systems are  statistically   identical  when  considering   the  faint  galaxy population. This result seems to  indicate that FGs and field galaxies originate  from  similarly  overdense   regions  of  the  Large  Scale Structure.  By studying the  distribution of color offsets from  the red sequence, we find  that FGs  occupy a similar  region to field  ellipticals.  We also find  that the mean color  gradients in both samples  (FG and FS) are statistically the same.  However,  the scatter around the mean for FGs  is  significantly (1.6  $\\sigma$)  smaller  than  that for  field ellipticals, which may indicate a  more regular process in the buildup of  FGs, such  as mergers  of L$^{*}$  galaxies.  The  distribution of $a_{4}$,  which  measures  isophotal  shape  deviations  from  a  pure ellipses, indicates that FGs and field ellipticals are similar.  Studying the stellar populations  in the twenty-five FGs and seventeen field galaxies we find that there is no difference in age, metallicity or $\\alpha$-enhancement, indicating that the star formation history of fossil groups seems to be  analogous to that of field ellipticals.  We further examined  elliptical galaxies in compact groups  studied by de la Rosa  et al.   (2007).  As already  established in  previous works, elliptical galaxies  in CGs are  older and more metal-poor  than field ellipticals (Proctor et al.~2004; de  la Rosa et al.  2007) and fossil groups.  Particularly, when we look  at the relation between [Z/H] and log$\\sigma$, the mass-metallicity relation  is evident (see Lee et al. 2006; Kobayashi 2005).   More massive galaxies (in both  FGs and field galaxies) retain their metals more  efficiently than the low mass ones (i.e., ellipticals in CGs).   The $\\alpha$-enhancement diagram shows a clear   trend   between   $\\alpha/F_{e}$   and   velocity   dispersion (mass). This can be interpreted as a manifestation of downsizing - the less  massive galaxies have  a more  extended star  formation history. Based  on  these  results  we  see  that  FGs  are  similar  to  field ellipticals but  cannot be formed by  dry mergers of  ellipticals in a CG. However, a wet merger with a gas rich disk system may also explain the  observed  relations.   Here, the  luminosity-weighted  parameters (Age,  [Z/H],   [$\\alpha$/Fe])  are  more  sensitive   to  the  recent star-formed-population, generating lower ages, higher [Z/H] and higher [Mg/Fe].  This  assumes that the merger-starburst is  soon followed by SN/AGN feedback which depletes the gas content.  McCarthy  et al.~(2004)  have found  a significant  population  of red galaxies at high redshift ($1.3<z<2.2$), and claim that they cannot be descendants of  the $z \\sim  3$ Lyman-break galaxies.  Based  on their spectroscopic data  they also  conclude that most  of the  present day massive  galaxies   had  an  early   ($z_{\\rm  f}=  2.4$)   and  rapid formation. If we restrict ourselves to the twenty-five FGs, their last main star formation  episode must have occurred at  $z>0.3$, much more recently than in the high-$z$  galaxies. It appears that a more recent burst  occurred in  the  FGs  which would  explain  their lower  ages. Therefore, we may be seeing  massive ellipticals that accreted a small galaxy at $z>0.3$.  Our  results indicate that  FGs are  not significantly  different from bright ellipticals  found in low-density  regions of the  universe. As claimed by  Mulchaey \\&  Zabludoff (1999) in  a detailed study  of NGC 1132, FGs  could simply be \"failed\"  groups - a  few relatively bright galaxies merged forming the dominant  system we see today, but without enough surrounding  matter to form additional  nearby bright galaxies, resulting in  an atypical LF.  This  is also in  agreement with recent results    obtained    by   D'Onghia    et    al.~(2005)   based    on N-body/hydrodynamical  cosmological simulations  where they  find that FGs are formed by the infall of $L \\sim L_{*}$ systems along filaments with impact parameters as small as 5 kpc.  Our  findings  suggest  that  the objects  meeting  the  observational criteria expected for fossil groups  at low redshift might not be true fossils.   The extensive  study presented  here whereby  we  select FG candidates  based  on optical  data,  and  then  examining their  X-ray counterparts   reveals   no   difference   between   FGs   and   field ellipticals. We  note that  fossils can be  created anytime  in cosmic history,  and these  systems may  represent  the final  stage of  mass assembly as suggested in the  analysis of the Millennium simulation by Dariush et al. (2007), instead of forming a distinct class.          \\appendix"
    },
    "0812/0812.2742_arXiv.txt": {
        "abstract": "The Be/X-ray binary system \\SO\\ underwent a giant outburst in May-June 2005, followed by a dimmer outburst in August-September 2005. This increased intensity provided an opportunity to search for redshifted neutron-capture lines from the surface of the neutron star. If discovered, such lines would constrain the neutron star equation of state, providing the motivation of this search. The spectrometer (SPI) on board the \\emph{INTEGRAL} satellite observed the dimmer outburst and provided the data for this research. We have not detected a line with enough significance, with the width-dependent upper limits on the broadened and redshifted neutron capture line in the range of (2 -- 11) $\\times$ 10$^{-4}$ photons cm$^{-2}$ s$^{-1}$. To our knowledge, these are the strongest upper limits on the redshifted 2.2 MeV emission from an accreting neutron star. Our analysis of the transparency of the neutron star surface for 2.2 MeV photons shows that photons have a small but finite chance of leaving the atmosphere unscattered, which diminishes the possibility of detection. ",
        "introduction": "\\label{sec:intro}  Matter that is being accreted onto a neutron star has such large energies that it can undergo nuclear reactions by colliding with particles in the star's atmosphere. Nuclei that are heavier than Hydrogen will therefore break apart and generate free neutrons. The nuclear spallation process of $^{4}$He, discussed in detail in \\citet{Bildsten93}(hereafter, \\emph{BSW}), begins with the $^{4}$He atom entering the star's atmosphere with a very high kinetic energy. The atom undergoes many Coulomb collisions with atmospheric electrons and slows down. The atom may hit an atmospheric proton and splits into a neutron and $^{3}$He. After this initial spallation, the $^{3}$He travels deeper into the atmosphere and thermalizes at a high column density. The thermal $^{3}$He can do one of three things: a) absorb a neutron (the most possible outcome), b) collide with a fast proton and be destroyed, or c) participate in a fusion reaction. The neutron that was liberated during the initial spallation undergoes elastic scatterings with atmospheric protons, thermalizes and drifts down due to gravity. Since most of the $^{3}$He atoms will capture the free neutrons, the rate of production of 2.2 MeV photons will depend on the amount of excess neutrons. Another restraining factor is the depth to which neutrons travel before recombining with protons. If recombination (and therefore 2.2 MeV photon production) occurs deep in the atmosphere, the photon will have to undergo many Compton scatters before it can leave the atmosphere. \\citet{Bildsten93} have shown that for a neutron star with moderate magnetic field strength ($\\sim10^{9}$ G), a thermal neutron drifts approximately 1 Compton length before recombining and have calculated that only ~10\\% of the 2.2 MeV photons will escape unscattered.  These photons, emitted with an intrinsic energy \\emph{E$_{0}$} from the atmosphere of a neutron star of radius \\emph{R} and mass \\emph{M}, experience a gravitational redshift for an observer at infinity to an energy \\emph{E} given by  \\[ \\frac {R}{M} = \\frac {2G}{c^{2}(1-\\frac{E}{E_{0}})} \\]  \\citep{Ozel03}. Measuring \\emph{E/E$_{0}$} would allow a direct measurement of the gravitational redshift and place significant constraints on the underlying equation of state. Using the \\emph{P}-$\\rho$ relation for a given model, the equation of hydrostatic equilibrium can be numerically integrated to obtain the global mass-radius (\\emph{M}-\\emph{R}) relation for that equation of state, defining a single curve in the neutron star \\emph{M}-\\emph{R} diagram \\citep{Lattimer04}. While the ratio \\emph{M}/\\emph{R} cannot identify the nuclear equation of state uniquely, it can rule out several models and place constraints on the remaining models. If the mass can be measured independently, using other means, then the nuclear equation of state can be uniquely identified.  When calculating the 2.2 MeV photon production mechanism, BSW have assumed a moderate magnetic field  and that matter arrives radially onto the neutron star. The magnetic field of \\SO\\ has recently been measured as $\\sim 4 \\times 10^{12}$G \\citep{Caballero07}. In such high magnetic fields, the ability of such line photons to escape the neutron star is contingent upon the optical depths for Compton scattering and also for magnetic pair creation being less than unity. The Compton scattering constraint amounts to one on proximity of the emission point to the upper layers of the stellar atmosphere; as mentioned above, it was addressed in BSW, using non-magnetic Compton scattering cross sections.  The pair production constraint is more involved, and will be addressed at some length in the Discussion below. The bottom line is that pair creation transparency of the magnetosphere to 2.2 MeV photons will be achieved if their origin is at polar magnetic colatitudes less than around 20 degrees {\\it and} they are beamed within around 15--20 degrees relative to their local surface magnetic field vector.  These two requirements limit the expectation for visibility of a 2.2 MeV emission line to small, but non-negligible probabilities.  The 2.2 MeV line is expected to be both redshifted and broadened, due to gravitational and relativistic effects \\citep{Ozel03}. Since we don't know a priori how much the line will be redshifted, we scanned the entire 1 MeV - 2.2 MeV region - wide enough to contain the redshifted line according to several different equations of state \\citep{Shapiro83} given in Table 1. The width of the redshifted 2.2 MeV line is also unknown, which led us to scan the 1 - 2 MeV region using different energy bands: 5keV, 11keV(${\\Delta}$E/E$_{0}$ = 0.5\\%), 22 keV (${\\Delta}$E/E$_{0}$ = 1\\%), 44 keV (${\\Delta}$E/E$_{0}$ = 2\\%), 100 keV, 150 keV and 170 keV. We expect the line broadening to be minimal, since it is not a rapidly rotating neutron star with a pulsation period of 103 seconds \\citep{Clark98}.  \\begin{table}[t] \\caption{\\label{table:EOSs} Mass and Radius Values for Common Equations of State for Neutron Stars$^{a}$ and the Expected Energies of the 2.2 MeV Line} \\begin{minipage}{\\linewidth} \\begin{center} \\begin{tabular}{c|c|c|c} \\hline \\hline Equation of State$^{b}$ & $M_{max}$/$M_{\\odot}$ & R (km) & E (MeV) \\\\ \\hline $\\pi$ & 1.5 & 7.5 & 1.41\\\\ R & 1.6 & 8.5 & 1.47 \\\\ BJ & 1.9 & 10.5 & 1.50 \\\\ TNI & 2.0 & 9 & 1.30 \\\\ TI & 2.0 & 12-15 & 1.56-1.72 \\\\ MF & 2.7 & 13.5 & 1.41 \\\\ \\hline \\hline \\end{tabular}\\\\ {$^{a}$ \\cite{Shapiro83}  $^{b}$ $\\pi$: Pion Condensate, R:Reid's, BJ: Bethe-Johnson, TNI: 3-Nucleon Interaction Approximation, TI: Tensor Interaction, MF: Mean Field Approximation } \\end{center} \\end{minipage} \\end{table}   The redshifted 2.2 MeV line from \\SO was previously studied by \\citet{Boggs06b}, using \\emph{RHESSI} data. The \\emph{RHESSI} data was gathered during \\SO's giant outburst (May 18 - June 25), whereas we used data from a smaller outburst that took place approximately 3 months later \\citep[see][for a detailed discussion of \\emph{INTEGRAL} and \\emph{RXTE} analysis of the source] {Caballero07}. The \\emph{RHESSI} instrument uses Germanium crystals as detectors like SPI, but its area is smaller, and there is no shielding. It is specifically designed to study solar flares, and due to the close proximity of the Sun to this source during the giant outburst, while \\emph{INTEGRAL} wasn't able to study this event, \\emph{RHESSI} had an unique opportunity to observe the entire outburst. \\citet{Boggs06b} subsequently calculated flux upper limits for \\SO\\, which we used for comparison. ",
        "conclusions": "\\label{sec:discussion}  Despite the excellent energy resolution and relatively high area of \\emph{SPI} for high energy photons, the redshifted 2.2 MeV line from \\SO\\ was not detected significantly. However, we were able to obtain stronger constraints in the flux than that of \\emph{RHESSI} instrument. The non detection is not surprising, as the transparency of the neutron star, detailed below, is greatly diminished for high magnetic fields, reducing the yield of 2.2 MeV photons.  The visibility of any atmospheric 2.2 MeV neutron capture line is subject to potential attenuation in the neutron star magnetosphere.  The leading process for such absorption is magnetic pair production, $\\gamma\\to e^+e^-$, the process that underpins standard radio and polar cap gamma-ray pulsar models (e.g. \\citet{Sturrock71}, \\citet{Daugherty96}). This mechanism is the highest order photon absorption interaction in strong-field QED, and is prolific at the field strengths inferred for A0535+262.  It has a threshold (for the $\\parallel$ polarization state) of $2 m_ec^2/\\sin\\theta_{kB}$, where $\\theta_{kB}$ is the angle the photon momentum vector makes to the local field line, above which the interaction rate precipitously rises (e.g. \\citet{Daugherty83}, \\citet{Baring08}. At energies only modestly above threshold, the attenuation lengths in local fields greater than around $10^{12}$ Gauss are considerably less than 1cm. Accordingly, once the threshold is crossed, 2.2 MeV photons can be expected to be destroyed in A0535+262, provided this occurs at altitudes less than around one stellar radius above the surface.  This accessibility of the threshold defines the broad criterion for transparency versus opacity of the magnetosphere for the line photons, the outcome of which is critically dependent on geometry.  Attenuation phase space can readily be assessed as functions of the surface polar field, the surface emission locale, and the emission direction of the photons.  Such a detailed analysis was offered by \\citet{Baring01} at much higher energies that are germane to pulsar and magnetar continuum emission.  It should be remarked that familiar two photon pair creation is negligible in comparison.  The transparency of the magnetosphere for 2.2 MeV line photons was investigated here principally via the salient threshold criterion, to quickly ascertain the pertinent approximate phase space.  Our analysis was performed in flat spacetime, sufficient for a general indication of transparency, for which rather simple geometrical relations for photon transport through the magnetosphere can be prescribed, along the lines of the approach that \\citet{Baring07} used for resonant Compton upscattering in magnetars.  The involved mathematical details of this analysis will be omitted here, for the sake of brevity, and only the principal conclusions outlined.  Photons at 2.2 MeV that are emitted at the atmospheric surface are below magnetic pair creation threshold when they initially have $\\sin\\theta_{kB} 2 m_ec^2/(2.2 MeV) \\Rightarrow \\theta_{kB}\\lesssim 27^\\circ$. Note that at the field strengths of $\\lesssim 4\\times 10^{12}$ Gauss appropriate for A0535+262, the pair threshold for the $\\perp$ photon polarization state is extremely close in energy to that for the $\\parallel$ state.  For equatorial emission locales, where the radius of field curvature is comparatively small, even photons that are initially beamed close to field lines rapidly acquire sufficient angles with respect to {\\bf B} to precipitate pair attenuation. Hence the magnetosphere is necessarily opaque to $\\gamma\\to e^+e^-$ for emission colatitudes $\\theta_e$ generally greater than around $30^\\circ$. Near the pole, if $\\theta_{kB}\\lesssim 27^\\circ$ initially, transport of photons to significant altitudes is required to access pair threshold. There, the field line curvature radius is $\\rho_c\\approx 4R_{ns}/(3\\sin\\theta_e)$ for a neutron star radius of $R_{ns}$. Accordingly, photons that initially are beamed close to the field, acquire an angle of $\\theta_{kB}\\gtrsim 27^\\circ$ when reaching an altitude $h\\sim 4R_{ns}\\sin 27^\\circ/(3\\sin\\theta_e)$.  This altitude must generally be more than around a stellar radius for the field decline to render the pair creation rate insufficient to sustain photon attenuation, as opposed to free escape (see, e.g. \\citet{Baring01} for a detailed examination). To summarize our analysis, here we find that pair creation transparency of the magnetosphere to 2.2 MeV photons will occur if their origin is at polar magnetic colatitudes $\\theta_e \\lesssim 20^\\circ$ {\\it and} they are initially beamed within around 15--20 degrees relative to their local surface magnetic field vector. These two requirements limit the expectation for visibility of any 2.2 MeV emission line to modest or small, but non-negligible, surface areas and solid angles.  With a more powerful tool for background rejection and a dramatical improvement in sensitivity (broad line sensitivity of 1.2 $\\times$ 10$^{-6}$ photons cm$^{-2}$ s$^{-1}$ and a spectral resolution of 0.2 -- 1 $\\%$) \\citep{Boggs06_act}, the Advanced Compton Telescope could be a good opportunity to follow up on the search for the broadened and redshifted 2.2 MeV line."
    },
    "0812/0812.4813_arXiv.txt": {
        "abstract": "Collapsars are fast-spinning, massive stars, whose core collapse liberates an energy, that can be channeled in the form of ultrarelativistic jets. These jets transport the energy from the collapsed core to large distances, where it is dissipated in the form of long-duration gamma-ray bursts. In this paper we study the dynamics of ultrarelativistic jets produced in collapsars. Also we extrapolate our results to infer the angular energy distribution of the produced outflows in the afterglow phase. Our main focus is to look for global energetical properties which can be imprinted by the different structure of different progenitor stars. Thus, we employ a number of pre-supernova, stellar models (with distinct masses and metallicities), and inject in all of them jets with fixed initial conditions. We assume that at the injection nozzle, the jet is mildly relativistic (Lorentz factor $\\sim 5$), has a finite half-opening angle ($5^\\circ$), and carries a power of $10^{51}\\,$erg\\,s$^{-1}$. In all cases, well collimated jets propagate through the progenitor, blowing a high pressure and high temperature cocoon. These jets arrive intact to the stellar surface and break out of it.  A large Lorentz factor region $\\Gamma\\simmore 100$ develops well before the jet reaches the surface of the star, in the unshocked part of the beam, located between the injection nozzle and the first recollimation shock. These high values of $\\Gamma$ are possible because the finite opening angle of the jet allows for free expansion towards the radial direction.  We find a strong correlation between the angular energy distribution of the jet, after its eruption from the progenitor surface, and the mass of the progenitors.  The angular energy distribution of the jets from light progenitor models is steeper than that of the jets injected in more massive progenitor stars. This trend is also imprinted in the angular distribution of isotropic equivalent energy. ",
        "introduction": "Recent observations of gamma-ray bursts (GRBs) suggest that long duration GRBs and type Ib/c supernova (SN) explosions are tightly connected.  For example, SN1998bw was observed in the positional error box of GRB980425 \\citep{Galama98}.  In this case the GRB/SN association was based on the spatial and temporal coincidence of both events.  The most remarkable example of long GRB/SN link came in 2003, when the spectra of both the GRB030329 afterglow and of the SN2003dh were measured, since the burst happened closeby and it was quite bright.  The supernova spectrum, which includes many complex lines, gradually appeared from the decaying afterglow spectrum after a few tens of days from the burst.  The spectrum of SN2003dh after about a month from the explosion is quite similar with that of SN1998bw at the same stage.  Both SN1998bw and SN2003dh are type Ic supernovae \\citep{Iwamoto98,Stanek03} whose progenitor had lost the hydrogen and the helium envelopes during the pre-supernova stage.  They are also categorized within a special class of supernova explosions, so-called hypernovae, whose explosion energy is about ten times higher, i.e., $\\sim 10^{52}\\ \\mbox{erg}$, than that of ordinary supernova. Indeed, our common view is that most long-lasting GRBs are produced by core-collapse supernovae akin to SN1998bw.  GRBs are not exclusively linked to hypernovae. For instance, the long lasting burst ($t_d\\sim 2\\times 10^3$ s) GRB060218 (or XRF060218) was associated with the type Ic \\object{SN2006aj}, which is not a hypernova.  The explosion energy falls within the regular range for type Ic events \\citep{Campana06}.  \\cite{Mazzali06} argued that the progenitor star may not form a black hole but a neutron star, since the estimated mass of the progenitor during the main sequence is $\\sim 20 M_\\odot$.  On the other hand, there are recent examples of cases of long-lasting GRBs (\\object{GRB060505} and \\object{GRB060614}) where no supernova signature was observed at all even if, considering the distance, the observational trace of a supernova should have been detected \\citep{Fynbo06,DellaValle06,GalYam06}.  Though a variety of long duration GRBs which have a strong connection with SNs are observed, we still lack of a complete picture of the processes by which a sizable fraction of the energy involved in a type Ib/c SN explosion is tapped in a relatively narrow channel, and produces a GRB at a large distance from the original site of generation.  \\citet{Rees92} proposed that the death of massive stars can be an origin of GRBs.  \\cite{Woosley93} introduced the collapsar model to account for the progenitor system of long GRBs.  According to this model, a non-spherical outflow could be formed from the deep inside of the progenitor where a black hole or a proto-neutron star is born as a result of the collapse of the iron core \\citep{MacFadyen99}. If the specific angular momentum of the iron core is sufficiently large, an accretion torus may develop around the central compact object.  The system formed by a central object and an accretion disk has the potential to launch a bipolar outflow.  The mechanisms proposed to extract energy out of such central engines are basically two \\citep[e.g.,][]{AO07}: thermal or hydromagnetic.  Thermal mechanisms rely on depositing a considerable amount of thermal energy in the vicinity of the rotation axis of the system, just above the poles of the central compact object, where a low density funnel has develop in the course of the evolution. The accretion energy of the hot torus is converted into a copious flux of neutrinos $\\nu$ and anti-neutrinos ${\\bar \\nu}$. From the $\\nu{\\bar \\nu}$-annihilation a hot $e^+e^-$-plasma results. In its turn, $e^+e^-$-pairs annihilate yielding a {\\em fireball} of high energy photons. The conversion of the thermal energy of the fireball into kinetic energy partly determines the subsequent evolution of the plasma.  It accelerates to ultrarelativistic speeds (reaching Lorentz factors of $\\sim 50$ \\citealt{Aloy00}) while, at the same time, interacts with the progenitor system. Alternatively, MHD process may tap a fraction of the rotational energy of the BH or of the accretion disk to form an outflow \\citep{Proga03,Mizuno04,McKinney06,Nagataki07,Komissarov07,Takiwaki07}.  A number of works have dealt with the hydrodynamic properties of the outflows generated in collapsar progenitors as they propagate through the progenitor system (and in some cases beyond the surface of the progenitor star).  The problem is addressed by means of numerical relativistic hydrodynamic simulations with different degrees of complexity (\\citealt{Zhang03,Zhang04,Mizuta06,Morsony07,Tominaga07}), which assume that a quasi-steady momentum flux has been produced at a certain distance from the region where the energy is released (independent of which is the actual energy extraction mechanism -MHD or thermal-). Therefore, such numerical works assume the existence of a nozzle through which the injection of a supersonic jet is produced, and put their focus on the modification of the morphology and of the dynamics of collimated outflows as they travel through the progenitor star. Depending on the exact inflow conditions, a variety of different outflows result. For instance, \\citet{Mizuta06} finds a whole spectrum of outflows ranging from collimated, relativistic jets to poorly collimated expanding winds. Thus, \\citeauthor{Mizuta06} argue that such a variety of resulting outflows supports the idea that the same collapsar scenario can yield a number of different phenomena (in agreement with the previous ideas of unification of high energy transients, e.g., \\citealp{Ramirez-Ruiz02b}), such as, GRBs, X-ray rich GRBs, X-ray Flashes, and normal supernovae.  In the prompt GRB phase we observe the emission of high energy photons from an ultrarelativistic outflow, which is generated at a distance $10^{13-15}$\\,cm, very far from the central engine.  Due to the large optical thickness of the outflow at scales comparable to that of the progenitor, we have not observed any electromagnetic emission directly from the progenitor so far, except, perhaps, some precursor activity \\citep{Woosley00,Ramirez-Ruiz02}.  Since the direct detection of progenitors of GRBs is nowadays impossible (if they happen at cosmological distances), the observation of the association between GRBs and SNe probably provides the best clue to understand the progenitors of the GRBs.  Though more than 100 GRBs per year are identified, most of them occur far away.  Thus, it is technically impossible to identify their accompanying supernovae.  Another clue regarding the nature of the GRB progenitors comes from the environment and the host galaxies in which the burst takes place. Both clues indicate that long-duration GRBs are associated with the death of the most massive stars \\citep{Fruchter06}.  The typical hosts of long GRBs are star-forming, low metallicity galaxies (with an star formation rate $\\sim 10^3\\,M_\\odot$y$^{-1}$; \\citealt{Berger01,Frail02}) but bluer than typical starburst galaxies, with little dust \\citep{LeFloch04}, and lower masses than current ellipticals, i.e., they correspond to the typical environments of formation of massive stars.  \\cite{Fruchter06} also conclude, that the host galaxies of the GRBs are significantly fainter and more irregular than the hosts of core-collapse supernovae.  A very interesting question is weather it is possible to get any information on the nature of a GRB progenitor from the observation of the afterglow emission.  One possibility is to look at the angular distribution of integrated energy per unit of solid angle, as observed in the afterglow phase of the burst.  \\citet{Lazzati05} estimated theoretically such an angular distribution assuming, that the kinetic energy of the jet is converted to thermal energy in the cocoon, till the head of the jet reaches the progenitor surface.  The cocoon originates from jet material which crosses through the terminal strong shock of the collimated outflow and moves away from the center of the progenitor surrounding the beam of the jet.  A fat cocoon develops for light jets, i.e., jets whose rest-mass density is much lower than that of the ambient gas into which the jet propagates.  When the jet breaks out the progenitor surface, the thermal energy is released to a low-density inter stellar medium (ISM).  \\cite{Lazzati05} concluded, that the energy distribution per solid angle ($dE/d\\Omega$) of the jet displays a $\\theta^{-2}$ dependence with the viewing angle after its eruption through the progenitor surface.  Recently, \\citet{Morsony07} have used hydrodynamic simulations to test the theoretical prediction of \\citeauthor{Lazzati05}, and found that their numerical models do not follow the inferred theoretical angular energy distribution.  In this paper we also try to verify the analytic relation for the angular dependence of the energy with the polar angle that was proposed by \\cite{Lazzati05}. We explore a parameter space different from that of \\cite{Morsony07} in order to compute the dependence of the angular energy distribution on the structure of the progenitor. The progenitor models are built upon the pre-supernova models of \\cite{Woosley06}.  Along the way, we also characterize the hydrodynamic properties of relativistic jets propagating through different progenitor stars.  The paper is organized as follows. We describe our physical model, the choice of stellar progenitors, and relevant numerical details in Sect.~\\ref{model}. In the appendix, we provide a study of a selected sample of models in order to justify our choice of numerical resolution and the effects it has on our conclussions. The dynamics of the injected bipolar outflows, and the extrapolated angular energy distribution in the afterglow phase is considered in Sec.~\\ref{results}. Finally, we discuss our results and write down the conclusions of this work in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we explore the relation of the dynamical properties of ultrarelativistic jets generated in collapsars with the properties of the progenitor stars in which such jets propagate. We have particularly focused on the correlations that exist between the angular profile of the energy per solid angle (as seen in the afterglow phase) and the properties of the progenitors. Along the way, we have pointed out which is the relevance of the fact that our numerical models are set up with a finite injection half-opening angle.  Using a non-zero injection angle affects the way in which the conversion of internal-to-kinetic energy takes place. If the flow is injected parallel to the polar axis, the development of reconfinement shocks happens closer to the injection nozzle than if the flow is injected radially within a cone of finite half-opening angle. When the recollimation shock occurs far away from the nozzle, the unshocked beam flow accelerates along a larger distance in a rarefaction that precedes such shock. Thus, the beam reaches there larger Lorentz factors than if the jet is injected parallel to the polar axis. The dissipation in cross shocks acts by simply recycling part of the kinetic energy of the outflow into thermal energy. The thermal energy is not lost, since the jet propagation is roughly adiabatic inside of the progenitor (radiation losses are negligible there). Instead, this thermal energy can be further converted into kinetic energy at larger distances. This process may happen several times before the outflow becomes transparent and radiation can freely escape. This explains why, in spite of the differences in the beam dynamics, all models (independent of the injection half-opening angle; c.f., \\citealt{Mizuta06}) develop a roughly similar propagation speed and, by the time they reach the head of the jet, the gross properties of the outflow are similar.   We have estimated the angular distribution of energy per solid angle in the afterglow phase by extrapolation of the state of our models when they reach a distance of $\\sim 10^{11}\\,$cm. This extrapolation relies on the fact that the jets develop a rough self-similar behavior soon after they emerge from the progenitor surface. Our results show that the equivalent isotropic energy per solid angle $dE/d\\Omega|_{\\rm iso}$ is only partly consistent with that of \\cite{Lazzati05}. However, the results in this paper do appear to be consistent with previous numerical simulations such as those of \\cite{Zhang04} and \\cite{Morsony07}. \\cite{Lazzati05} obtain that the angular energy distribution displays a relatively flat core which is flanked by a region where $dE/d\\Omega|_{\\rm iso} \\propto \\theta^{-2}$. Our results show that the core of the distribution (close to $\\theta=0^\\circ$) is not flat (but decays as $\\theta^{-1}$) and that the energy per solid angle decays much faster than $\\theta^{-2}$ (it does it as $\\theta^{\\kappa}$, with a value of $\\kappa\\simless -3.6$ depending on the mass of the progenitor; see below). We can fit the $dE/d\\Omega|_{\\rm iso}$-data with SBP functions up to $\\theta\\simless 3.4^\\circ$. At smaller latitudes, a simple power-law with a slope close to $-2.6$ fits better the data. In this region ($4.5^\\circ < \\theta < 8^\\circ$), the cocoon contribution is the dominant (at smaller values of $\\theta$, the beam of the jet dominates the energetics), and we find there the best consistency with the model of \\cite{Lazzati05}.  We have correlated the properties of the angular $dE/d\\Omega|_{\\rm iso}$ distribution of the jet with the fundamental parameters of the progenitor star in which the jet has propagated. We find that the shape of the distribution is mostly influenced by the mass of the progenitor. When the mass of the progenitor is small ($M\\sim 5 \\Msun$), because of the occurrence of large mass losses due to winds in the latest stages of the star's evolution, then $dE/d\\Omega|_{\\rm iso}$ decays faster with $\\theta$ than if the mass of the progenitor is large ($M\\sim 15 \\Msun$). We find that the reason for this behavior is that the average density of the progenitors tends to grow (approximately) with the mass. This means that the average jet propagation speed inside of the star is smaller, the larger is the mass. A smaller jet propagation speed results into thicker and hotter cocoons, since we fix the same mass, momentum and energy fluxes at the nozzle for all models. Also the beam of the jet in the more massive progenitors is wider. This is the reason why low mass progenitors develop narrower $dE/d\\Omega|_{\\rm iso}$ profiles than high mass ones. The difference in the collimation of the energy distributions resulting from low and high mass progenitors has a direct influence on the number of observed events. We expect to see more events produced in heavy progenitors than in low mass ones.  One could question whether the correlation that we have found between the mass of the progenitor and the width of the $dE/d\\Omega|_{\\rm iso}$ is an artifact of our numerical set up. We are fixing the luminosity of the jet to be the same independent of the progenitor mass. However, progenitors of different mass may develop central engines which release different power. A good proxy of the power of the central engine is the mass of the iron core of the progenitor. One may expect that the collapse of more massive iron cores results into larger central compact objects. If the power released by the central engine is dominated by the size of the central compact object, then models with more massive iron cores could release a larger power than models with low mass cores. Then, how do we justify our numerical assumption that the power injected in the jet is roughly independent of the progenitor's mass?. In support of our point we argue that, first, according to \\cite{Woosley06} data, there is not a one to one correlation between the mass of the iron core and the mass of the progenitor and, second, the iron mass varies by less than a 30\\% in all the models considered here, while the total mass can be different by a factor of 3. Thus, within the simplifications we do in our models, the assumption of a common luminosity independent of the mass of the progenitor is justified.  Irrespective of these two arguments, if heavier progenitors would result into more luminous central engines, we shall point out that this trend will also result into wider jets and cocoons. This is a result early pointed out by \\cite{Aloy00}, were it was shown that, taking the same progenitor, but increasing the luminosity of the central engine by a factor of 10, results into thicker cocoons than in cases in which the luminosity is more moderate. The reason being that the release of a large power triggers large amplitude Kelvin-Helmholtz instabilities at the basis of the outflow, which transfer a large fraction of the momentum of the beam to a thicker shear layer between the beam and the cocoon. Effectively, this process widens the cross sectional area of the beam, reducing its propagation speed and, consistently inflating larger cocoons.  We have also found, that comparing progenitors of similar mass, the metallicity of the star has a small impact on the extrapolated $dE/d\\Omega|_{\\rm iso}$ profiles. The collimation of the jet is similar regardless of the stellar metallicity. However, the cocoon is more narrowly collimated in low-metallicity stars, because of the large density drop close to their surfaces. This reduced density makes the jets in low-metallicity stars much more ballistic, once they break out the stellar surface, than in solar-metallicity progenitors. Unfortunately, this difference might not be observable, since it happens in regions where the energy per solid angle is much smaller than at the jet core (unless orphan afterglows could be detected; see, e.g., \\citealp{Totani02,Rossi08}).  Finally, we have found significant differences between the $dE/d\\Omega|_{\\rm iso}$ profiles of collapsar-jets and those of jets produced in merger remnants (i.e., between the angular energy distribution of jets associated to long and to short GRBs). Collapsar-jets are more narrowly collimated than jets of merger remnants, and the decay of $dE/d\\Omega|_{\\rm iso}$ beyond the central flat core is much steeper in the latter than in the former. This intrinsic difference manifest itself as a larger chance of detectability of jets from merger remnants at viewing angles up to $\\theta \\sim 12^\\circ$, while, on the other hand, collapsar-jets could hardly be seen off-axis."
    },
    "0812/0812.2815_arXiv.txt": {
        "abstract": "We report on optical imaging of the X-ray binary \\sax\\ with the 8-m Gemini South Telescope. The binary, containing an accretion-powered millisecond pulsar, appears to have a large periodic modulation in its quiescent optical emission.  In order to clarify the origin of this modulation, we obtained three time-resolved $r'$-band light curves (LCs) of the source in five days.  The LCs can be described by a sinusoid, and the long time-span between them allows us to determine optical period $P=7251.9$~s and phase 0.671 at MJD 54599.0 (TDB; phase 0.0 corresponds to the ascending node of the pulsar orbit), with uncertainties of 2.8~s and 0.008 (90\\% confidence), respectively. This periodicity is highly consistent with the X-ray orbital ephemeris. By considering this consistency and the sinusoidal shape of the LCs, we rule out the possibility of the modulation arising from the accretion disk. Our study supports the previous suggestion that the X-ray pulsar becomes rotationally powered in quiescence, with its energy output irradiating the companion star, causing the optical modulation. While it has also been suggested that the accretion disk would be evaporated by the pulsar, we argue that the disk exists and gives rise to the persistent optical emission. The existence of the disk can be verified by long-term, multi-wavelength optical monitoring of the source in quiescence, as an increasing flux and spectral changes from the source would be expected based on the standard disk instability model. ",
        "introduction": "While it had long been believed that neutron star (NS) X-ray binaries (XRBs) are progenitors of the recycled millisecond radio pulsars \\citep{bv91}, it was the discovery of coherent pulsations from the transient XRB \\sax\\ (hereafter \\saxs) during its X-ray outburst in 1998 that first and finally confirmed the connection between the two systems \\citep{wv98}: in this binary, the accreting NS is a 2.49 ms X-ray pulsar. As the first example of accretion-powered millisecond pulsar systems, \\saxs\\ has been extensively studied, with various interesting properties revealed (see \\citealt{har+08} and references therein).  In this paper, we focus on the optical periodic modulation seen in this binary and report on our observational study of the modulation.  The orbital period of \\saxs\\ is $P_{\\rm orb}\\simeq$7249.157 s ($\\simeq$2.01 hr), accurately known to one part in $10^{10}$ from Doppler modulations of the millisecond pulsations \\citep{cm98,har+08}. Combined with the derived mass function of 3.8$\\times 10^{-5} M_{\\sun}$, the period implies that the mass-transferring companion could be a 0.17 $M_{\\sun}$ low-mass main-sequence star, but more likely a $\\sim$0.05 $M_{\\sun}$ brown dwarf \\citep{bc01}. At a distance of $D=3.5$ kpc \\citep{gc06}, the optical counterpart in quiescence is several magnitudes brighter ($V=20.7$, $L_{V}\\simeq 3.0\\times 10^{32}$ ergs s$^{-1}$ assuming isotropic emission and extrinction $A_V=0.73$; see \\S~2 and \\S~4) than the possible types of stars suggested as the companion, probably indicating that the optical emission arises from the accretion disk in the binary \\citep{hccv01}. However in the quiescent state, 10--40\\% sinusoidal-like modulations in the source's optical light curves (LCs) have been reported (\\citealt{hccv01,cam+04}), and this is puzzling because the quiescent X-ray luminosity is approximately $L_{\\rm X}\\simeq 5\\times 10^{31}$ ergs s$^{-1}$ (e.g., \\citealt{hei+07}), two orders of magnitude lower than that required to account for the modulation \\citep{bur+03}. Typically in a low-mass X-ray binary (LMXB), sinusoidal optical modulation arises from X-ray heating of the companion star by the central X-ray source: the visible area of the heated face varies as a function of orbital phase (e.g., \\citealt{ak93}).  In \\saxs, depending on the companion's star types, only 0.5--1.4\\% [estimated by $(R_2/D_b)^2/4$, where $R_2$ is radius of the companion and $D_b$ is the separation distance between the NS and companion] of the total energy flux from the central NS would be received by the companion for isotropic emission. This leads to the suggestion that in quiescence, the NS might switch to be a rotation-powered pulsar so that the rotational energy would be the energy source that heats half surface of the companion star and causes the modulation \\citep{bur+03}.  However, there are other possibilities that do not require a rotation-powered pulsar, and we have considered whether or not the accretion disk could give rise to the modulation. It has been known that ``superhumps\", which are commonly seen in short-period cataclysmic variables (CV; \\citealt{war95}), also appear in LMXBs (e.g., \\citealt{has+01}). These periodic modulations have periods a few percent longer than the orbital periods and can be sinusoidal-like with an amplitude of $\\sim$10\\%, arising from a precessing, eccentric accretion disk (e.g., \\citealt{wk91}). Indeed, it has been suggested that those NS LMXBs with $P_{\\rm orb} <4.2$ hr are potential superhump sources \\citep{has+01}. In addition, several parts of an accretion disk could contribute significantly to optical modulation (e.g., \\citealt{mc82}). It has also been suggested that for an X-ray transient, its quiescent optical emission may come from a bright spot on the accretion disk \\citep{mc01}.  In particular, the superhump possibility was suggested by the X-ray LC obtained in the source's 2002 outburst. As shown in Figure~\\ref{fig:xlc}, the LC exhibits a $\\sim$5-day periodic modulation at the end of the outburst. If this indicates the precession periodicity ($P_{\\rm prec}\\simeq 5$ days) of the accretion disk, it would imply a superhump period of $P_{\\rm sh}= 7373$~s ($1/P_{\\rm sh}=1/P_{\\rm orb}-1/P_{\\rm prec}$) and superhump excess $\\epsilon =0.017$ [$\\epsilon =(P_{\\rm sh}-P_{\\rm orb})/P_{\\rm orb}$] \\citep{pat01}. The excess value is consistent with those obtained for cataclysmic variables and LMXBs (\\citealt{pat+05, has+01}). Furthermore, a mass ratio of $q\\simeq 0.08$ could be estimated from the relation $\\epsilon =0.18q+0.29q^2$ \\citep{pat+05}, implying a companion mass of 0.11 $M_{\\sun}$ for 1.4 $M_{\\sun}$ NS mass. This companion mass is within the range implied by the mass function.  Previously, time-resolved imaging observations over a small period of time (covering only $\\sim$1.5 orbital period of the binary) were made. However, these observations were carried out either with a small telescope (\\citealt{hccv01}) or under very poor observing conditions \\citep{cam+04}, resulting in large uncertainties in the obtained LCs. In order to study the optical emission from \\saxs, and particularly to probe whether it could be a superhump source, we have obtained high quality optical LCs of the source in its quiescent state through time-resolved photometry. The observations were made with the 8-m Gemini South Telescope over five days, allowing us to determine the period and phase of the optical modulation accurately. We note that \\citet{hei+08} (see also \\citealt{del+08}) recently observed the source simultaneously at X-rays and optical $g'i'$ wavelengths, and from the observations they confirmed the inconsistency between the large amplitude optical modulation and low X-ray luminosity. \\begin{figure} \\plotone{f1.eps} \\figcaption{X-ray light curve of \\saxs\\ during its 2002 outburst, obtained with the Proportional Counter Array (PCA; 2--60 keV energy range) on board the \\textit{Rossi X-ray Timing Explorer} satellite. Gaps in the light curve were due to Earth occultations of the source. At the end of the outburst, a $\\sim$5 day periodicity is tentatively suggested. \\label{fig:xlc}} \\end{figure} ",
        "conclusions": "Using the 8-m Gemini South Telescope, we have obtained, for the first time, well-determined LCs from \\saxs\\ in its quiescent state over a time span of four days. From the above studies of the LCs, we find that the optical period and phase are consistent with the X-ray ephemeris, indicating that the optical modulation is orbital in origin. In studies of several tens of LMXBs at optical wavelengths (e.g., \\citealt{vm95}), in no instance has there been an accretion disk giving rise to a sinusoidal modulation at the orbital period. In addition, the sinusoidal maximum must correspond to superior conjunction of the companion star. Because of these, we rule out the possible disk origin for the modulation that we have suspected. However, the source in outburst could still be a superhumper, which might have been hinted in the X-ray LC (Figure~\\ref{fig:xlc}). As the outward extension of accretion disks in outburst has both been observed and reproduced in disk instability simulations \\citep{osa96, dhl01}, it would not be unexpected for the accretion disk in \\saxs\\ to have extended to the resonance zone during the 2002 outburst, developing into an eccentric form due to the tidal instability \\citep{wk91}. In fact, superhumps have been seen in outbursts of both black-hole and NS LMXB systems \\citep{oc96, ele+08}. In order to determine this possibility for the long periodicity seen in the 2002 outburst, time-resolved imaging observations, like ours, of the source in outburst are needed. Since the source will be as bright as $\\sim$17 mag in an outburst (e.g., \\citealt{ghg99}; \\citealt{wan+01}), a search for superhump modulation will be feasible even with a small telescope.  Based on the current observational studies of LMXBs (e.g., \\citealt{vm95}), it seems extremely unlikely that the observed optical modulation would arise from a source other than the companion star. Thus far, pulsar wind heating of the companion is the only model that has been suggested \\citep{bur+03, cam+04}. The long term spin-down rate of the pulsar has been measured, indicating a rotational energy loss rate of 9$\\times 10^{33}$ ergs s$^{-1}$ \\citep{har+08}. This energy output, presumably in the form of a pulsar wind, would illuminate the companion star. Assuming isotropic emission and a brown dwarf companion \\citep{bc01}, the fraction of the total energy received by the companion is $\\sim$0.005$\\eta_{\\ast}(R_2/0.13\\ R_{\\sun})^2$, where $\\eta_{\\ast}$ is the fraction of the received energy absorbed by the companion. Following \\citet{ak93}, the companion's heated face would have temperature $\\sim$$7430 \\eta_{\\ast}^{1/4}$ K, due to pulsar wind heating by the putative rotation-powered pulsar.  Using such a hot face that varies following a function of $[1+\\sin i\\sin (2\\pi t/P)]$, where $i$ is the inclination angle of the binary, and also including a constant flux component $F_C$, we tested whether we could re-generate the averaged LCs of \\saxs. The distance and extinction to the source were fixed at 3.5 kpc and $A_V=0.73$, respectively, where the extinction value is estimated from $A_V=N_{\\rm H}/0.179\\times 10^{22}$~cm$^{-2}$ by assuming hydrogen column density to the source $N_{\\rm H}=0.13\\times 10^{22}$ cm$^{-2}$ (\\citealt{dl90}; \\citealt{hei+07}). The extinction law for Sloan filters given by \\citet{sfd98} was used. We found that the parameter values of $i\\simeq 63\\arcdeg$ ($M_2\\simeq 0.049 M_{\\sun}$), $\\eta_{\\ast}\\simeq 0.46$, and $F_C\\simeq 19\\ \\mu$Jy can provide the observed modulation (the resulting $\\chi^2\\simeq 2100$, with no systematic uncertainties considered). Although we used a very simple model, these derived parameter values are consistent with its known properties. In addition to the fact that the companion is likely a $\\sim$0.05 $M_{\\sun}$ star, the source shows no X-ray eclipses or dips, implying $i\\leq 70\\arcdeg$. The obtained $\\eta_{\\ast}$ values are within the range found for two binary radio pulsars \\citep{sta+01,rey+07}, in which it is known that the companion is irradiated by the pulsar wind. Therefore, it is plausible that the NS in \\saxs\\ does turn into a rotation-powered pulsar in quiescence, giving rise to the optical modulation. We note that very recently, \\citet{del+08} used an advanced model to fit their $g'i'$ light curves, and also found that the required heating energy should be $\\sim 10^{34}$ ergs s$^{-1}$, consistent with the derived spin-down luminosity (which has 30\\% uncertainty; \\citealt{har+08}).  The origin of the persistent optical emission is not clear. \\citet{hccv01} tried explaining the emission from an X-ray irradiated disk around the pulsar, but it may not be appropriate to use a steady thin disk model to describe a disk in the thermally stable cold state (lower cold branch of the standard thermal equilibrium $S$-curve; e.g., \\citealt{las01}), since a disk temperature profile in the cold state can be drastically different from the hot state (the steady disk case). \\citet{cam+04} used a shock front, arising from the interaction between the companion star and pulsar wind, and the irradiated companion to account for the emission. Here we argue that the accretion disk in quiescence exists, against the suggestion that the disk would be evaporated by the pulsar \\citep{bur+03, hei+08}, and this can be tested by monitoring \\saxs\\ at optical wavelengths.  According to the standard disk instability model (DIM; e.g., \\citealt{osa96,las01}), while the mass accretion rate to the NS in \\saxs\\  is very low during quiescence, $\\dot{M}_{\\rm acc}\\leq 6.2\\times 10^{-15}\\ M_{\\sun}$ yr$^{-1}$ (estimated from the observed X-ray flux), the average mass transfer rate from the companion to the accretion disk is as high as $\\sim$10$^{-11} M_{\\sun}$ yr$^{-1}$ (estimated from the X-ray fluence in each outburst; \\citealt{gal08}). The transferred mass is stored in the disk, building up the surface density for triggering the next outburst. The average persistent $r'$ flux from \\saxs\\ in our observations is estimated to be $F_C=19\\ \\mu$Jy, corresponding to a disk luminosity of $L_{r'}= 2\\pi D^2 F_C/\\cos i \\simeq 1.5\\times 10^{32}$ ergs s$^{-1}$ ($i=63\\arcdeg$ is assumed). There is plenty of gravitational energy available to power this emission as matter moves inwards through the outer disk. At the time average accretion rate of $\\dot M\\sim 10^{-11}\\ M_\\odot \\ {\\rm yr^{-1}}$, matter falling into a radius of 4000 km releases gravitational energy at a rate that matches the observed luminosity. This radius is far larger than those that are suggested for the inner radius $r_{\\rm in}$ of the disk. Generally, $r_{\\rm in}$ would be close to the Alfv\\'{e}n radius, $r_{\\rm in}\\simeq 56\\ {\\rm km}$ $(\\dot{M_{\\rm in}}/10^{-11} M_{\\sun} {\\rm yr}^{-1})^{-2/7} (\\mu/10^{26}\\ {\\rm G\\ cm} ^{3})^{4/7}$, where $\\dot{M_{\\rm in}}$ is the mass accretion rate in the inner edge of the disk and $\\mu$ is the magnetic moment, $\\mu\\simeq 10^{26}$ G cm$^3$ for \\saxs\\ (\\citealt{har+08}). In the cold state, $\\dot{M}_{\\rm in}$ would be lower than $\\dot{M}$, and we note that for $\\dot{M}_{\\rm in}= 0.1 \\dot{M}$ \\citep{dhl01}, $r_{\\rm in}$ is 110 km. However, since a radio pulsar presumably would have no interactions with a surrounding disk, $r_{\\rm in}$ would be larger than the light cylinder radius of the pulsar, which is 120 km. As can be seen, it is possible that in quiescence, the disk in \\saxs\\ would be outside of the light cyliner. In addition, the disk temperature profile in quiescence may be described by a constant, at least right after an outburst (e.g., \\citealt{osa96, dhl01}). For \\saxs, we find that an effective temperature of 4600~K for the disk can give rise to the persistent $r'$ flux, where the disk is assumed to be cut off at the tidal radius 3.7$\\times 10^{10}$ cm ($\\simeq 0.9R_1$, where $R_1$ is the NS's Roche lobe radius). This temperature value is consistent with those typically considered in the DIM (\\citealt{las01}; the critical effective temperature for having an outburst is $\\sim 6000$ K).  In order to verify our suggestion that the persistent optical emission arises from the disk, long-term, multi-wavelength optical monitoring of the source in its quiescent period is required. From such observations, we might expect to see an increasing flux from the source. Moreover, since in the DIM the temperature profile as a function of disk radius is predicted to be changing, turning from a constant right after an outburst to a power-law--like function prior to the next outburst (e.g., \\citealt{dhl01}), we would also see flux spectrum changes. This type of well-behaved changes would not be expected from the pulsar wind shock model \\citep{cam+04}, thus allowing to determine the origin of the persistent emission.  If the companion star is irradiated by the pulsar wind, there is no reason to think that the disk is not. It has been suggested that the disk in quiescence might be evaporated by the pulsar (e.g., \\citealt{bur+03, hei+08}), but according to the recent calculations by \\citet{jon07}, a pulsar wind may only be effective in heating a disk. Basically, as X-rays from a NS would ionize the surface of a disk, the Poynting flux, which is dominant in a wind when it is not far from the light cylinder of the pulsar, would interact with the ionized particles, converting energy into disk heating. Using equation~(16) in \\citet{jon07}, we estimate that the baryon loss rate of the disk at the inner radius is approximately 3$\\times 10^{21}(r_{\\rm in}/120\\ {\\rm km})^{-3}$ cm$^{-2}$ s$^{-1}$, only 0.05\\% of the surface density ($\\sim$10--100 g cm$^{-2}$) that is generally considered in the accretion disk models (e.g., \\citealt{dhl01}). This suggests that the disk in \\saxs\\ could exist and might be irradiated by the pulsar wind. However, using the model provided by \\citet{jon08}, the flux due to pulsar wind heating would be 2 $\\mu$Jy for parameter $\\zeta=0.3$ (0.03$\\lesssim\\zeta\\lesssim$0.3 and a larger $\\zeta$ value corresponds to a higher disk effective temperature; see details in \\citealt{jon08}). The flux would be 10\\% of the average $r'$ flux, which would suggest a weak pulsar-wind heating effect in \\saxs.  Finally, it will be of great interest if \\saxs\\ can be determined to become rotation-powered during quiescence. We note that the source could be very similar to PSR~J2051$-$0827 \\citep{sta+96}, a binary millisecond pulsar system. For example, the latter has an orbital period of 2.38 hr and a mass function of 1.0$\\times 10^{-5}$ $M_{\\sun}$, and the pulsar has a spin-down luminosity of 6$\\times 10^{33}$ ergs s$^{-1}$. However, searches for pulsed radio emission from \\saxs\\ have not been successful (e.g., \\citealt{burgay+03}). Here we suggest that the source might be identified by searching for its pulsed $\\gamma$-ray emission. Observations of millisecond pulsars suggest that their efficiency at $\\gamma$-ray energies may be as high as $\\sim$7\\% \\citep{kui+00}. This implies a $\\gamma$-ray flux of $\\sim$5$\\times 10^{-13}$ ergs cm$^{-2}$ s$^{-1}$ for \\saxs, possibly detectable by deep observations with \\textit{Fermi Gamma-Ray Space Telescope}."
    },
    "0812/0812.0021_arXiv.txt": {
        "abstract": "We show that the excellent optical and gamma-ray data available for GRB~080319B rule out the internal shock model for the prompt emission. The data instead point to a model in which the observed radiation was produced close to the deceleration radius ($\\sim10^{17}$\\,cm) by a turbulent source with random Lorentz factors $\\sim10$ in the comoving frame. The optical radiation was produced by synchrotron emission from relativistic electrons, and the gamma-rays by inverse Compton scattering of the synchrotron photons.  The gamma-ray emission originated both in eddies and in an inter-eddy medium, whereas the optical radiation was mostly from the latter.  Therefore, the gamma-ray emission was highly variable whereas the optical was much less variable.  The model explains all the observed features in the prompt optical and gamma-ray data of GRB~080319B. We are unable to determine with confidence whether the energy of the explosion was carried outward primarily by particles (kinetic energy) or magnetic fields.  Consequently, we cannot tell whether the turbulent medium was located in the reverse shock (we can rule out the forward shock) or in a Poynting-dominated jet. ",
        "introduction": "Two major unsolved questions in the field of gamma-ray bursts are: (i) the mechanism by which the energy in relativistic jets is converted to random particle kinetic energy, and (ii) the radiation process by which the particle energy is converted to gamma-ray photons.  A number of ideas have been proposed for the conversion of jet energy to particle energy (cf. Piran 1999, 2005; M{\\'e}sz{\\'a}ros 2002; Thompson 1994, 2006; Lyutikov \\& Blandford 2003; Zhang 2007).  Most popular among these is the so-called internal shock model (Piran, Shemi \\& Narayan 1993; Rees \\& Meszaros 1994; Katz 1994), in which different parts of the relativistic GRB jet travel at different speeds.  Faster segments collide with slower segments in shocks, and a fraction of the jet kinetic energy is converted to thermal energy. Gamma-rays are then produced by either the synchrotron process or the synchrotron-self-Compton (SSC) process (e.g., Piran 2005, M{\\'e}sz{\\'a}ros 2002, Zhang 2007).  Another model is the external shock model (Dermer 1999) in which gamma-rays are produced via the synchrotron process in an external shock driven into the circumstellar medium by the GRB jet.  Difficulties with this model have been pointed out by a number of authors (cf. Piran 1999).  The excellent data obtained by the Swift and Konus satellites for GRB 080319B -- dubbed the ``the naked eye burst\" -- has provided a new opportunity to investigate the viability of these models and to understand the fundamental nature of GRBs.  We summarize here the main observational properties of this burst (details may be found in Racusin et al.  2008).  GRB 080319B lasted for about 50~s and had a burst fluence in the 20~keV~--~7~MeV band of $5.7 \\pm 0.1 \\times 10^{-4}$~erg~cm$^{-2}$, which corresponds to an isotropic energy release of $E_\\gamma = 1.3 \\times 10^{54}$~erg (Golenetskii et al. 2008) for a redshift $z=0.937$ (Vreeswijk et al 2008; Cucchiara \\& Fox 2008).  The time-averaged gamma-ray spectrum during the burst had a peak at around 650~keV.  The maximum flux at this energy was $\\sim7$~mJy, and the time averaged flux was $\\sim 3$~mJy.  The time-averaged gamma-ray spectrum was measured by Konus-Wind (Racusin et al. 2008) to be $F_\\nu \\propto \\nu^{0.18 \\pm 0.01}$ for photon energies below 650~keV and $F_\\nu \\propto \\nu^{-2.87 \\pm 0.44}$ at higher energies.  (The spectrum evolved during the course of the burst, as we discuss in \\S2, and this provides additional information on the radiation process.)  In the optical band, at photon energies around 2~eV, the peak flux of GRB 080319B was $V=5.4$~mag or 20~Jy, and the time-averaged flux was about 10~Jy (Karpov et al 2008). The optical lightcurve varied on a time scale of about 5~s, while the $\\gamma$-ray flux varied on time scales of $\\sim0.5$~s.  GRB 080319B has seriously challenged our understanding of gamma-ray bursts.  The problems posed by this burst are described in Kumar \\& Panaitescu (2008) and discussed in further detail in the present paper.  We show that the observations cannot be explained with any of the standard versions of the internal shock model.  We find, however, that a consistent model is possible if we give up the idea of internal shocks and instead postulate that the gamma-ray source is relativistically turbulent (see Narayan \\& Kumar 2008).  We begin in \\S2 by arguing that the radiation in GRB 080319B must have been produced by the SSC mechanism.  Following this, we derive in \\S3 the basic equations describing a GRB that radiates via SSC.  In \\S4, we combine these equations with the internal shock model and attempt to explain the data on GRB 080319B.  We find that no consistent model is possible.  In \\S5, we consider a relativistically turbulent model, again with SSC radiation, and show that in this case it is possible to obtain a consistent model of GRB 080319B.  We summarize the main conclusions in \\S6. The Appendix discusses the effects of source inhomogeneity on our calculation of the synchrotron self-absorption frequency, and the synchrotron and inverse-Compton fluxes. ",
        "conclusions": "We have shown in this paper that the gamma-ray and optical data for GRB 080319B rule out the popular internal shock model for generation of the prompt radiation.  According to this model, the duration ($\\lta\\ 1$\\,s) of spikes in the gamma-ray lightcurve sets an upper bound on the quantity $R/(2c\\Gamma^2)$, where $R$ is the radius of the source relative to the center of the explosion and $\\Gamma$ is the bulk Lorentz factor.  When we apply this condition, we find that it is impossible to fit the observed optical and gamma-ray flux simultaneously.  Specifically, any model that fits the optical flux under predicts the gamma-ray flux by nearly two orders of magnitude (Fig.~1).  This is an unacceptably large discrepancy which cannot be eliminated with any reasonable modification of the internal shock model.  An equally powerful qualitative argument against the internal shock model is the fact that we find the radius $R$ of the source to be constrained quite tightly by the observations.  The energy required in magnetic field increases very rapidly as we decrease $R$, $E_B\\propto R^{-22/7}$, whereas the energy in particles increases rapidly, $E_e\\propto R^{17/7}$.  Also, the cooling time of electrons becomes too short to be compatible with observations if $R<3\\times10^{16}$\\,cm.\\footnote{Collisions at a smaller radius would produce a weak optical flash with flux decreasing roughly as R$^{11/12}$.  The electrons would undergo very rapid cooling and produce a low energy spectrum in the gamma-ray band of $f_\\nu\\propto\\nu^{-1/2}$. Prompt optical observations of GRB 080319B show variations in the optical flux by less than a factor two for much of the 50\\,s duration of the burst except at the beginning and the end.  Moreover, the low energy spectral index for the gamma-ray emission was greater than 0 throughout the burst.}  All of these factors together constrain the location where the prompt emission in GRB~080319B was produced to lie within a narrow range of radius: $4\\times10^{16}\\ \\lta\\ R\\ \\lta\\ 8\\times10^{16}$\\,cm (see Fig.~1). This is problematic for the internal shock model.  According to this model, there is a large number of internal shocks among independent ejecta, with a separate shock producing each of the $\\sim50$ spikes in the gamma-ray lightcurve of GRB 080319B.  Why would all the ejecta collide within such a narrow range of radius?  Moreover, why should the radius be so close to the deceleration radius $R_d\\sim 10^{17}$\\,cm, where the ejecta meet the external medium and begin to slow down?  This coincidence is suspicious.  All of these problems are eliminated if we give up the internal shock model and consider instead a model in which the variability in the gamma-ray lightcurve is produced by relativistic turbulence in the source with random eddy Lorentz factors $\\gamma_t\\sim10$.  In this model, the quantity $R/(2c\\Gamma^2)$ is no longer constrained to be less than 1\\,s, but only needs to be comparable to the burst duration $\\sim50$\\,s (Narayan \\& Kumar 2008).  With this modification, we find that we obtain a remarkably consistent model of GRB 080318B (see Fig.~2) in which the prompt optical emission was produced by synchrotron emission and the gamma-rays were the result of inverse Compton scattering.  The predicted gamma-ray flux is perfectly compatible with observations.  Also, estimates of various quantities such as the total energy, cooling time, Lorentz factor, etc. are all very reasonable and consistent (\\S5.1).  The radius of the source is calculated to be in the range $6\\times10^{16}<R<3\\times10^{17}$\\,cm; if we select a nominal value $R\\sim10^{17}$\\,cm, we obtain an excellent fit to all the observations.  In the context of a physical model, the picture that emerges from this model is that the energy of the relativistic jet in GRB~080319B was converted to optical \\& $\\gamma$-ray radiation either via a relativistic reverse shock when ejecta (composed of $p^+$s and $e^-$s) ran into the circumstellar medium or that much of the jet energy was in magnetic field that was dissipated close to the deceleration radius.  Theoretically, it is difficult to understand how a reverse shock might produce relativistic turbulence with $\\gamma_t\\sim10$ (\\S5.2.3).  Also, it is easier to understand the optical data for the time period $10^2\\ \\lta\\ t\\ \\lta\\ 10^3$\\,s if we assume a Poynting jet (\\S5.2.4).  For these reasons we have a mild preference for the Poynting-dominated jet model.  A potential problem for the Poynting jet model is that the ratio of magnetic to particle kinetic energy is about 0.1 for our best solution (Fig.~2). However, this ratio is similar to that inferred for the pulsar wind termination shock for the Crab pulsar (Kennel \\& Coroniti, 1984). Moreover, this ratio of 0.1 does not rule out the Poynting model for another reason which is that, in all of our formulae, $B$ is the projection of the magnetic field perpendicular to the electron momentum vector.  Thus, if electron momenta are preferentially parallel to the magnetic field, then the true $E_B$ would be larger than our estimate (easily by a factor 10 compared to the value given in eq. \\ref{eb08b}), and we can have $E_B/E_{e}\\sim1$.  Note that electrons are accelerated parallel to the magnetic field in reconnection regions and so this possibility is not as arbitrary as it might appear.  In the relativistic turbulence model, fluctuations in the observed gamma-ray lightcurve are produced as a result of random relativistic variations in the velocity field of the source, with turbulent Lorentz factor $\\gamma_t\\sim 10$.  The model predicts that there should be $\\sim\\gamma_t^2\\sim100$ spikes in the gamma-ray light curve (Narayan \\& Kumar 2008), which is consistent with the $\\sim50$ spikes seen in GRB~080319B.  The optical synchrotron flux is dominated by the inter-eddy medium rather than eddies.  Therefore, we expect much less variability in the optical flux, as was indeed observed.  The model can explain the sharp rise in the optical flux of GRB~080319B at the beginning of the burst.  For this, we must postulate that the synchrotron peak frequency increased from $\\sim0.5$\\,eV to $\\sim1.5$\\,eV during the first $\\sim15$s.  Since the synchrotron peak frequency $\\nu_i$ was below the optical band (2\\,eV), the optical flux was smaller than the peak synchrotron flux by a factor $(4/\\nu_i)^{2.8}$ (the spectral index above the peak is known from gamma-ray observations).  A modest increase in $\\nu_i$ by a factor of 3 early in the burst would thus produce a factor of 20 increase in the observed optical flux.  The reason that the gamma-ray flux increased by a much smaller factor during the same time is that the peak IC flux is proportional to $\\tau_e N_e B$, which would change little during this time period.  The end of the gamma-ray prompt emission phase occurred $\\sim 8$\\,s after the prompt optical in GRB~080319B.  This is probably a result of inverse Compton emission from turbulent eddies lying a bit outside of the primary $1/\\Gamma$ cone, but which happened to point toward us because of a fortuitous alignment of their turbulent velocity (the probability for this happening is of order unity).  Since much of the synchrotron emission comes from non-turbulent fluid in between eddies, the optical flux would not have a similar effect."
    },
    "0812/0812.0810_arXiv.txt": {
        "abstract": "We present the results of concurrent X-ray and optical monitoring of the Seyfert~1 galaxy \\mkn79 over a period of more than five years. We find that on short to medium time-scales (days to a few tens of days) the 2--10 keV X-ray and optical $u$ and $V$ band fluxes are significantly correlated, with a delay between the bands consistent with zero days. We show that most of these variations may be well reproduced by a model where the short-term optical variations originate from reprocessing of X-rays by an optically thick accretion disc. The optical light curves, however, also display long time-scale variations over thousands of days, which are not present in the X-ray light curve. These optical variations must originate from an independent variability mechanism and we show that they can be produced by variations in the (geometrically) thin disc accretion rate as well as by varying reprocessed fractions through changes in the location of the X-ray corona. ",
        "introduction": "\\label{sec:intro}  Active Galactic Nuclei (AGN) emit across almost the entire electromagnetic spectrum. In most AGN, the bulk of the emission is radiated in the UV/optical region in what is known as the `big blue bump'. The blue bump is generally accepted to be thermal emission, originating in an optically thick accretion disc surrounding the central black hole. For a review, see \\citet{koratkarblaes99}. AGN have been known to exhibit optical variability since shortly after their discovery \\citep[e.g.][]{smithhoffleit63} but the origin of the variability is still unclear. The two main possibilities are that either the optical variations arise from intrinsic variations of the disc emission, perhaps due to accretion rate variations, or that the variations arise from reprocessing of the variable X-ray emission, boosting the intrinsic thermal emission from the disc.  If the optical variations arise from reprocessing of X-rays, we expect them to lag behind the X-ray variations. As the X-rays originate from close to the central black hole and the longer wavelength emission comes from further out in the disc, light travel time arguments for a standard optically thick disc predict that the lag will increase with wavelength as $\\tau\\propto\\lambda^{4/3}$ \\citep{collieretal99}. For a typical AGN we expect short delays, of order a day or less, between adjacent photometric bands. Such wavelength-dependent delays have been seen in many AGN \\citep{wandersetal97, collier98, krissetal00, collieretal01, oknyanskij03, sergeevetal05, doroshenkoetal06}. For a detailed discussion of the disc reprocessing model, see \\citet{cackettetal07}. \\citet{gaskell07} suggested that the wavelength-dependent delays in the UV-optical may also be due to a contribution from reprocessing material much further out in the torus.  On the other hand, if the optical variations arise from variations in the intrinsic disc emission as accretion rate perturbations propagate inwards, we expect the optical variations to precede the X-ray variations. By the same argument, the longer wavelength optical emission will also lead those at shorter wavelengths.  The variations will propagate through the disc on a viscous, rather than light-travel time-scale, so we expect much longer delays ($\\sim$months--years) between wavebands than in the reprocessing scenario.   Coordinated X-ray and optical monitoring observations provide a method for determining the sign of the delay between the optical and X-ray variations and hence for determining the dominant process underlying optical variability. However, previous studies have yielded varied and confusing results, with some sources showing well correlated emission and others no correlation at all. For example, \\citet{mason02} report a 0.14~day lag of the UV emission relative to the X-rays in NGC~4051, but on longer time-scales \\citet{shemmer03} find a 2.4 day optical lead. \\citet{peterson00} find a correlation between the X-rays and optical on time-scales of a few weeks in the same source, but little agreement on time-scales of days.  For NGC~3516, \\citet{maoz00} reported a possible 100~day X-ray to optical lag after their 1.5~year monitoring of this source. However, with the addition of 3.5~years more data \\citep{maoz02}, they were unable to confirm this lag, or measure any other statistically significant correlation at any lag in the light curves. For 3C~390.3, the relationship between passbands is also complicated: X-ray flares sometime lead and sometimes follow optical/UV events \\citep{gaskell08}. \\citet{uttley03} found the X-ray and optical continuum in NGC~5548 to be well correlated on month time-scales, and also that the amplitude of the optical variations was larger than that of the X-rays on these long time-scales. On short time-scales \\citet{suganumaetal06} report a clear 1.6~day optical lag to the X-ray variations and an optical amplitude much smaller than the X-ray amplitude, suggesting that, unlike on the longer time-scales, reprocessing is the main driver of the optical variability  on time-scales of days. Intensive monitoring by \\citet{gaskell06}, however, fails to show this short time-scale correlation.  Besides simply measuring the optical--X-ray lag, coordinated X-ray and optical monitoring observations enable us to put other useful constraints on the X-ray and optical emission processes and on the geometry of the emitting regions. For example, the greater fractional variability of the X-rays on short ($\\sim$day) time-scales indicates that, in a Compton scattering scenario for the X-ray production, at least part of the X-ray variability is not produced by variable up-scattered seed photons, but by variations in the scattering corona itself. The relative amplitude of optical and X-ray variability is also important because if the optical variations, which represent variations in the bolometrically dominant blue bump, exceed those in the X-rays, it is hard to explain the optical variations as purely due to reprocessing, e.g. as in NGC~5548 \\citep{uttley03} and MR2251-178 \\citep{arevalo08}. In these AGN the large amplitude optical variations occur on time-scales of years and may represent accretion rate fluctuations propagating through the emission region.  It has been suggested \\cite[e.g.][]{uttley03} that the strength of the correlation may be dependent on the mass of the central black hole. Scaled in terms of the gravitational radius, AGN of higher black hole mass or lower accretion rate, with cooler discs, will have the thermal optical emitting region closer to the black hole than the less massive/high accretion rate objects. We may expect stronger correlations in the high mass systems, as the accretion rate variations will be better connected between the X-ray and optical emitting regions. Conversely, in the low mass/high accretion rate systems, the very different local viscous time-scales will allow the emission regions to vary independently, causing poorer correlations on long time-scales.  In order to try and understand the optical variability of AGN, and the relative importance of the two proposed variability mechanisms, we have set up a long term monitoring programme of a sample of AGN of differing mass and fractional accretion rate $\\dot{m}=\\dot{M}/\\dot{M}_{\\mbox{\\small{Edd}}}$. Previously \\citep{uttley03} we have reported on observations of NGC~5548 [$M\\sim 7\\times10^{7}M_\\odot$; $\\dot{m}\\sim0.08$ \\citep{peterson04,woourry02}] and MR2251-178 [$M\\sim10^{9}M_\\odot$;  $\\dot{m}\\sim 0.04$] \\citep{arevalo08} and here we report on the observations of \\mkn79.  \\mkn79 (UGC~3973, MCG+08--14--033) is a nearby ($z=0.022$) Seyfert~1.2 galaxy with a black hole mass of $(5.24\\pm1.44)\\times 10^7 M_\\odot$ \\citep{peterson04} and an average $\\dot{M}$ of $\\sim6\\%$ of its Eddington accretion rate \\citep{kaspi00}. A $V$ band image of the galaxy is shown in Figure~\\ref{fig:pic}. It is a well-known X-ray source and its optical variability is well established \\citep{peterson98,webbmalkan00}. We present the combined optical light curves from four independent monitoring programs, together spanning more than five years, as well as the X-ray monitoring light curve as obtained by the \\rxte~satellite for the same time. Data collection, construction of the light curves and basic variability properties are discussed in Section~2. We present the cross-correlation analysis in Section~3 and in Section~4 we show that the medium term fluctuations are well reproduced by a reprocessing model and that the long time-scale trends can be accounted for by allowing the disc and/or corona geometry to vary over these time-scales. We conclude in Section~5.   \\begin{figure} \\psfig{file=pic_negative.ps,width=8cm} \\caption{A combined $V$ band image of the Seyfert galaxy Markarian 79 (marked with a red circle) as obtained with the Liverpool Telescope. The bar at the bottom of the image respesents a scale of 1 arcminute. The comparison star used for the optical photometry is marked with a blue square, and the summed flux from the other two bright stars in the image was used to confirm that the comparison star did not vary.} \\label{fig:pic} \\end{figure} ",
        "conclusions": "We have presented the long-term optical ($u$ and $V$ band) and X-ray light curves of \\mkn79, obtained by combining data from several monitoring programs of this source. The $u$ band light curve is the result of 2.5 years of monitoring with the Liverpool Telescope (LT) and the $V$ band light curve was constructed from the combined data from six ground based telescopes, including the LT. The X-ray observations, which were made every $\\sim2$ days, are from our continuing monitoring program with the \\rxte~satellite. We further included the results from a 3 month multiwavelength program on the Faulkes Telescope in $u$, $B$, $V$, $R$ and $i$.  We summarise our findings as follows: \\begin{enumerate} \\item The light curves are variable on time-scales as short as days (X-rays) and as long as thousands of days (optical), with highly correlated variability on time-scales of tens of days or months. \\item Long time-scale ($\\sim$years) variations are clearly seen in the optical light curve, but are not present in the X-ray light curve. \\item Although there is more variability in the X-rays than in the optical bands on short ($\\sim$day) time-scales, the fractional variability of the $V$ band light curve on long time-scales is greater when corrected for host galaxy flux, due to the large amplitude of these long-term optical variations. \\item The lags measured between the X-rays and long-term optical bands, are consistent with zero days for the data set as a whole. Formally we find that the $V$ band lags behind the X-rays by $0.00^{+1.91}_{-4.00}$ days and the $u$ band lags by $-1.95^{+1.95}_{-3.89}$ days. \\item While some strong events on a time-scale of a few days are consistent with the zero days lag, there are also time periods where the optical and the X-rays match up poorly and events are only seen in one passband. \\item The variations between the short multicolour light curves are also well correlated. The measured centroid lags, $\\tau$, are consistent with the $\\tau\\propto\\lambda^{4/3}$ prediction of X-ray reprocessing by an optically thick accretion disc. However, as the multicolour light curves are quite short, with few prominent peaks and troughs to constrain the lags, the error bars on the lag measurements are large and it is not possible to constrain the reprocessing model particularly tightly. We also cannot exclude the possibility that the measured delays across the optical bands are affected by a contribution from light reprocessed in the dusty torus further out. \\item Optical variations on time-scales of $\\sim$months are well reproduced by a reprocessing model, but longer term variations ($\\sim$years) cannot be accounted for in this way. \\end{enumerate}   In order to reproduce the long time-scale optical trends mentioned above, we are forced to include a second mechanism for the production of optical variations. This mechanism could be a variation in the geometry of either the purported accretion disc or the irradiating X-ray source. Changing the height of the X-ray source changes the fraction of the disc illuminated and hence the amount of optical emission produced by reprocessing. The intrinsic disc flux produced by viscous dissipation, is governed by $\\dot m$ and changes in the inner truncation radius $R_{\\rm in}$ affects the amount of optical emission produced. We have considered the effect of changing $h_X$, $\\dot m$ and $R_{\\rm in}$ individually, in steps between the three sections of the optical light curves (Figure \\ref{rep_tests}). We find that when we vary $h_X$, we are able to reproduce the large and small fluctuations of the observed light curves slightly better than when we vary the disc parameters. It is not possible, however, at this level of modelling, to rule out the other two possibilities. A crude model like this cannot reproduce the observed light curves exactly. In reality the parameters will probably all vary together, rather than one at a time as we have done here, and also gradually, rather than in a step-wise fashion. The description of the X-ray source as a point source above the centre of the disc, although providing a reasonable approximation to an extended hemispherical corona, may also not be correct. However, the requirement for an additional source of optical variation, in addition to X-ray reprocessing, is robust.  Despite the good correlation on short time-scales, there are some short time-scale variations which appear to occur only in one passband. This mismatch needs to be addressed by future studies and may be a consequence of anisotropic X-rays emission \\citep[see, e.g.][]{gaskell06}.   Within the standard accretion disc model \\citep{shakura}, the high black hole mass and low accretion rate of \\mkn79 implies that it has a cool disc. In such a disc, the X-ray and optical emitting regions are closer together, in terms of gravitational radii, than in a system of low black hole mass and/or high accretion rate. Hence both reprocessing and intrinsic variations are expected to contribute to the observed optical variability. Reprocessing benefits from the large solid angle subtended by the disc in the inner regions and intrinsic variations propagate on observable time-scales, because of the shorter viscous time-scale at smaller disc radii. We have previously reported this behaviour in the quasar MR2251-178 \\citep{arevalo08}. This system has a very massive central black hole ($M\\sim10^9M_\\odot$) and hence a cool disc, with the X-ray and optical emitting regions even closer together than in \\mkn79, in terms of gravitational radii. When  accretion rate fluctuations propagate through the disc, the X-ray and optical emitting regions are modulated together, leading to a good correlation on long time-scales. Reprocessing of X-rays imprints small rapid fluctuations onto the optical light curve, resulting in a well correlated behaviour on short time-scales as well. The strong X-ray--optical correlation we have reported here is in agreement with this picture.   For \\mkn79 we expect that 50\\% of the $V$ band emission should originate within 70$R_g$ of the black hole. Taking a typical value for the viscosity parameter, $\\alpha\\sim 0.1$, in an optically thick, standard accretion disc \\citep{shakura}, with thickness $(h/R)^2 \\sim 0.01$, we find that the viscous time-scale at 70$R_g$ is $t_{\\mbox{\\small{visc}}}\\sim 5$ years. In a geometrically thin disc, this time-scale is even longer.  This (approximate) time-scale, on which we would expect accretion rate perturbations to propagate inwards to the X-ray emitting region, is longer than the time over which we have observed the steep decrease in optical flux and hence  consistent with the fact that we have not yet seen a decrease in long term average X-ray luminosity. However the connection between the disc and the corona is not well understood and it is not clear how accretion rate variations, which may affect the optical disc emission, will be translated to the X-ray emission. If this interpretation is correct, the X-ray light curve should decline significantly over the next 5 or more years.   In subsequent papers we will compare the optical and X-ray variations in a number of other AGN of differing black hole mass and accretion rate, and hence of differing disc structure, in order to determine how disc structure affects both the long and short term optical variability of AGN."
    },
    "0812/0812.2354_arXiv.txt": {
        "abstract": "The Gaia space project, planned for launch in 2011, is one of the ESA cornerstone missions, and will provide astrometric, photometric and spectroscopic data of very high quality for about one billion stars brighter than V=20. This will allow to reach an unprecedented level of information and knowledge on several of the most fundamental astrophysical issues, such as mapping of the Milky Way, stellar physics (classification and parameterization), Galactic kinematics and dynamics, study of the resolved stellar populations in the Local Group, distance scale and age of the Universe, dark matter distribution (potential tracers), reference frame (quasars, astrometry), planet detection, fundamental physics, Solar physics, Solar system science.\\\\ I will present a description of the instrument and its main characteristics, and discuss a few specific science cases where Gaia data promise to contribute fundamental improvement within the scope of this Symposium. ",
        "introduction": "\\label{s:intro}  The idea of measuring the position of stars on the sky systematically and for a scientific purpose dates back to the 2nd Century BC, when the Greek astronomer Hipparchus measured about 1000 stars naked eye. This has been repeated several times during the following centuries, with steadily increasing power and accuracy, until the most recent Hipparcos satellite (1989-1993) that measured $\\sim$ 120,000 sources with $\\sim$ 1 mas accuracy and produced a Catalogue (\\cite[Perryman et al. 1997]{per97}) later revised by \\cite[van Leeuwen (2007)]{fvl07}.  Gaia represents the natural follow up of the Hipparcos mission, with huge improvement in terms of: i) measurement accuracy, ii) limiting magnitude and hence number of observed objects, and iii) the combination of nearly simultaneous astrometric, photometric and spectroscopic observations.  Gaia is a cornerstone mission of the ESA Space Program, that will perform an all-sky survey and produce accurate astrometry and photometry for about 1.5$\\times$10$^9$ objects down to a limiting magnitude of 20~mag, and additional spectroscopy for objects brighter than V=16-17. This will allow to obtain a stereoscopic and kinematic view of the Galaxy, and to address key questions of modern astrophysics regarding the formation and evolution of the Milky Way. Such an observational effort has been compared to the mapping of the human genome for the impact that it will have in Galactic astrophysics.  In addition, Gaia will provide a fundamental contribution in a much broader range of scientific areas (see Sect. \\ref{s:science}).  More information on the Gaia mission and its science can be found in the \\cite[{\\it Gaia Concept and Technology Study Report} (2000)]{ctsr}, the Proceedings of the Symposium \\cite[``The 3-Dimensional Universe with Gaia'' (2005)]{3D}, and at {\\bf http://www.rssd.esa.int/Gaia}.  In Table \\ref{t:compare} Gaia and Hipparcos characteristics and performances are compared.  \\begin{table} \\begin{center} \\caption{Comparison of Hipparcos and Gaia characteristics and performance} \\label{t:compare} {\\scriptsize \\begin{tabular}{|l|c|c|} \\hline & {\\bf Hipparcos } & {\\bf Gaia}  \\\\ \\hline Completeness to ...   & V $\\sim$ 9 & V $\\sim$ 20 (blue) - 22 (red) \\\\ \\hline Magnitude limit & V$\\sim$12.4     &  $\\sim$ 1 mag. fainter than completeness \\\\ \\hline N. Sources      & $\\sim$1.2 10$^5$ & $\\sim$1.5 10$^9$  \\\\ ~~~~~ Quasars   & 0          & $\\sim$10$^6$ \\\\ ~~~~~ galaxies  & 0          & $\\sim$10$^7$ \\\\ \\hline Astrometric accuracy         & $\\sim$ 1 mas & $\\le$ 7 $\\mu$as at V$\\le$10 \\\\ &              & 12(red)-25(blue) $\\mu$as at V=15 \\\\ &              & 100(red)-300(blue) $\\mu$as at V=20  \\\\ \\hline Photometry                   & 2 bands      & ~~~low-res spectrophotometry, 330-1050 nm ~~~\\\\ \\hline Spectroscopy (R$\\sim$ 11,000)~~~  & none    & 1-10 km s$^{-1}$ at V $\\le$ 16(blue)-17(red) \\\\ \\hline Target selection     & ~~~input catalogue~~~ & real-time onboard selection \\\\ \\hline \\end{tabular} } \\end{center} \\end{table} ",
        "conclusions": ""
    },
    "0812/0812.0398_arXiv.txt": {
        "abstract": "Following novel development and adaptation of the Metric Space Technique (MST), a multi-scale morphological analysis of the Sloan Digital Sky Survey (SDSS) Data Release 5 (DR5) was performed. The technique was adapted to perform a space-scale morphological analysis by filtering the galaxy point distributions with a smoothing Gaussian function, thus giving quantitative structural information on all size scales between 5 and 250 Mpc. The analysis was performed on a dozen slices of a volume of space containing many newly measured galaxies from the SDSS DR5 survey. Using the MST, observational data were compared to galaxy samples taken from \\emph{N}-body simulations with current best estimates of cosmological parameters and from random catalogs. By using the maximal ranking method among MST output functions we also develop a way to quantify the overall similarity of the observed samples with the simulated samples. ",
        "introduction": "From redshift surveys such as the Sloan Digital Sky Survey (SDSS; \\citealt{york00}) and the Two-Micron All Sky Survey (2MASS; \\citealt{skr00}), the local universe shows intricate patterns with clusters, filaments, bubbles, sheet-like structures and the so-called voids. For a review of the structural analysis of the universe, see \\citet{wei05}. At the same time, Lambda Cold Dark Matter $\\lambda$CDM models have been developed, see \\citet{gill04} and \\citet{dol08}. Several simulations have been created, such as the Millennium Simulation \\citep{spr05} done by \\citet{cro05} and another \\emph{N}-body simulation by \\citet{ber06}. These models describe a universe that consists mainly of dark energy and dark matter and calculate the evolution of the universe from a short time after the big bang to the present time. Work has been done to verify the similarity between the real universe and simulated universe \\citep{spr05,ber06} and they agree well, based on the comparative techniques used in these studies.   To supplement the widely used correlation function and power spectrum, alternatives have been proposed to quantify structure in the galaxy distribution, such as the genus curve \\citep{zel82a}, percolation statistics \\citep{zel82a,sha83,sah97}, Rhombic Cell analysis \\citep{wu04}, void probability functions \\citep{white79}, high-order correlation function \\citep{peeb80}, and multi-fractal measures \\citep{saar07}. However, all of these consider a single map as a space. Here we generalize the Metric Space Technique (MST), a tool used to analyze and classify astrophysical maps \\citep{ada92}, to perform a multi-scale analysis. Key facets of the MST approach are consideration of any given map as an element in the space of all such maps and definitions of a distance function to make the spaces of all maps into a topological space. Moreover, the other methods focus on summary statistics that convey little of the geometric and topological properties of the galaxy distribution. The MST method gives desired quantitative summary statistics of the difference between maps. However, a primary benefit of our method is that the output functions, such as filamentation, number of components,  density, volume and pixels, are straightforward and simple to understand and particularly useful in maps comparisons. Finally, MST is based on the use of threshold values, which will make sure that we can unambiguously define a space on the map with an interesting topology \\citep{ada92}.  The MST allows an objective and quantitative comparison of any two images. All such images are considered to be elements of a metric space, where, instead of comparing images on a pixel-to-pixel basis, the comparison is made by considering the metric distance between two images' output functions. The MST was first used to analyze Galactic molecular cloud data \\citep{ada94,wis94}. Several mathematical and technical improvements to the technique were presented in \\citet{kha04a} (for more details, see \\citet{kha04b}, where the updated formalism was used to analyze Galactic atomic hydrogen gas regions from the Canadian Galactic Plane Survey \\citep{tay03}. For both studies the output functions were applied to two-dimensional gray-scale images which described a smooth density field. But as originally suggested by \\citet{ada92}, one can choose to smooth point distribution data (e.g., galaxy distribution) in order to obtain gray-scale data from which the output functions can be calculated. The first application to point distribution is done by \\citet{wu08}. Importantly, however, the smoothing level becomes critical in creating the density field \\citep{don88,sil81}. Efforts have concentrated on determining the best density estimate from optimal smoothing length \\citep{cole95,vic05}. % In this paper, we consider a wide range of smoothing levels for multi-scale filtering \\citep{kha06}. By varying the size of the smoothing function over a range of scales, exactly like the wavelet transform, a complete multi-scale description of galaxy distributions in metric space becomes possible.  The goal of this paper is therefore to use the multi-scale MST to quantify morphological differences between the SDSS observational data and two sets of simulation sample data and then illustrate the use of those differences to understand degrees and types of structure in the galaxy distribution. Here the point (galaxy) distribution data was filtered by a smoothing function over a continuous range of scales. Using this novel approach, the MST not only informs us, quantitatively, about the structure information of the universe and which mock sample most resembles the observational data, but also how the information and resemblance vary over size scales. ",
        "conclusions": "We have used a slightly modified MST of \\citet{ada94} on multiple scale to study the morphology of galaxy distributions. The technique gives a detailed morphological description of galaxy distributions in metric space, on scales from about 5 Mpc to about 250 Mpc, with five output functions showing strong statistical differences. We also find that the filament output function values are high for the observations at small filtering scales but at around 50 Mpc the function approaches a lower stable value. Considering that most voids in SDSS galaxies are around 30--50 Mpc \\citep{gott05}, this seems a likely signature of those voids.    The key motivation for this work is to supplement traditional tools with a more informative way of quantifying the similarity in the ``visual'' morphological properties between simulations and the observed universe. We use the ``metric distance'' as the parameter to describe that similarity through multiple measures by calculating the value of each of the MST output functions. We combine the values of each of the output functions into one ``final'' parameter for each simulation by the maximal ranking method. In Table 1 and Fig. \\ref{distance_scale}, it was demonstrated that two \\emph{N}-body simulations have done a similar job of approximating our universe and that NYUr is more close to the observed sample than MPAr. From the analysis for Fig. \\ref{physical_space}, we surprisingly found that MPAp is more closely matched than the NYUp to the observed sample in redshift space, with the implication that velocity determinations for simulation galaxies is a major contributor to the relatively poorer match of the MPA simulation. The velocities of satellite galaxies in NYU simulation halos are assigned randomly from the dark matter particles within the haloes \\citep{ber06}. However, in the MPA simulation, even satellite galaxies have interpolated velocities (taken from the subhalo) rather than just randomly assigned ones  \\citep{cro08}. It is very likely that the mechanism for producing the velocity of satellite galaxies in MPA simulation contributes noticeably to the relative shortcomings of MPAr in Fig. \\ref{distance_scale}.  While the MST yields a single statistic for comparison of structure maps, in a way similar to other measures of large scale structure, we submit that its greater utility is in providing multiple intermediate outputs that convey insight into the physical differences between samples that lead to the statistical result. Of the many topological characteristics, threshold values, and scale samplings that the MST aggregates into a final result, we highlight a few examples of the specific physical differences that the technique reveals.  First, we have the expected result that the random sample is much different from all other samples at virtually all scales for all output functions. We have chosen only to use that case for normalization in the maximal ranking step.  Now, for our much more meaningful comparisons among MPA, NYU, and observed samples, we see that for the density output function, MPA has more high-density pixels  (at about the 5$\\sigma$ level) than NYU sample, and the NYU sample has more high-density pixels than observed sample (1-2 $\\sigma$). The volume and pixels output functions show similar trends as the density output function with the implication that those high-density pixels are also accompanied by large area regions of pixels above the various thresholds. The components and filament output functions are more complicated and fluctuate with the increasing scale. Simply speaking, for scales less than 50 Mpc, NYU and the observed sample are close to each other (around 1 $\\sigma$), but MPA clearly has many more sizeable clumps (greater than 3$\\sigma$) and is also more filamentary (greater than 3$\\sigma$). For scales more than 50 Mpc, MPA shows more filamentary structure (greater than 3$\\sigma$) than NYU sample, which is a little more filamentary (0.5-2 $\\sigma$) than the observed sample. And observed sample has more clumps (1$\\sigma$) than both mock samples.  Our next step is to apply the MST to the full three-dimensional galaxy distribution, which will require redevelopment/extension of output functions."
    },
    "0812/0812.2448_arXiv.txt": {
        "abstract": "We have used the Gemini Near-infrared Integral Field Spectrograph (NIFS) to map the emission-line intensity distributions and ratios in the Narrow-Line Region (NLR) of the Seyfert galaxy NGC\\,4151 in the Z, J, H and K bands at a resolving power $\\ge$\\,5000, covering the inner  $\\approx$ 200\\,$\\times$\\,300\\,pc of the galaxy at a spatial resolution of $\\approx$8\\,pc. We present intensity distributions in 14 emission lines, which show three distinct behaviours. (1) Most of the ionized gas intensity distributions are extended to $\\approx$100\\,pc from the nucleus along the region covered by  the known biconical outflow (position angle PA=60/240$\\degr$ -- NE--SW), consistent with an origin in the outflow; while the recombination lines show intensity profiles which decrease  with distance $r$ from the nucleus as $I\\,\\propto\\,r^{-1}$, most of the forbidden lines present a flat intensity profile ($I\\,\\propto\\,r^0$) or even increasing with distance from the nucleus towards the border of the NLR.  (2) The H$_2$ emission lines show completely distinct intensity distributions, which avoid the region of the bicone, extending from $\\approx$10\\,pc to $\\approx$\\,60\\,pc  from the nucleus approximately along the large scale bar, almost perpendicular to the bicone axis. This morphology supports an origin for the H$_2$-emitting gas in the galaxy plane. (3) The coronal lines show a steep intensity profile, described by $I\\,\\propto\\,r^{-2}$; the emission  is clearly resolved only in the case of [Si\\,{\\sc vii}], consistent with an origin in the inner NLR.  Using the line-ratio maps [Fe{\\sc\\,ii}]\\,1.644/1.257 and Pa\\,$\\beta$/Br\\,$\\gamma$ we obtain an average reddening of E(B-V)\\,$\\approx$\\,0.5 along the NLR and E(B-V)\\,$\\ge$\\,1 at the nucleus. Our line-ratio map [Fe{\\sc\\,ii}]\\,1.257$\\mu$m/[P{\\sc\\,ii}]\\,1.189$\\mu$m of the NLR of NGC\\,4151 is the first such map of an extragalactic source. Together with the [Fe{\\sc\\,ii}]/Pa\\,$\\beta$ map, these line ratios correlate with the radio intensity distribution, mapping the effects of shocks produced by the radio jet on the NLR. These shocks probably release the Fe locked in grains and produce an enhancement of the [Fe{\\sc\\,ii}] emission at $\\approx$\\,1\\arcsec\\ from the nucleus. At these regions, we obtain electron densities N$_e \\approx\\,4000$\\,cm$^{-3}$ and temperatures T$_e\\,\\approx\\,15000$\\,K for the [Fe{\\sc\\,ii}]-emitting gas. For the H$_2$-emitting gas we obtain much lower temperatures of T$_{exc}\\,\\approx$\\,2100\\,K and conclude that the gas is in thermal equilibrium. The heating necessary to excite the molecule may be due to X-rays escaping perpendicular to the cone (through the nuclear torus, if there is one) or to shocks probably produced by the accretion flow previously observed along the large scale bar.  The  distinct  intensity distributions and physical properties of the ionized and molecular gas, as well as their locations, the former along the outflowing cone, and the latter in the galaxy plane surrounding the nucleus, suggest that the H$_2$-emitting gas traces the AGN {\\it feeding}, while the ionized gas traces its {\\it feedback}. ",
        "introduction": "NGC\\,4151 is the nearest and apparently brightest Seyfert 1 galaxy (or Seyfert 1.5, according to \\citet{ost76}), and thus harbours one of the best studied active galactic nuclei (hereafter AGN) \\citep{crenshaw07}. As pointed out by \\citet{mundell99b}, from old optical observations, NGC\\,4151 was believed to be a small spiral galaxy with major axis extending by $\\approx$2$\\farcm$5 along position angle (hereafter PA) PA$\\approx$130$\\degr$. But this is just the inner part of the galaxy, comprising an oval distortion, or  ``weak fat bar'' \\citep{mundell99a}, beyond which there is a much larger disk, weak in the optical, but bright in H{\\sc\\,i} emission, with spiral arms extending up to 6\\,arcmin from the nucleus \\citep{davies73}. From the H{\\sc\\,i} kinematics, it was concluded that the galaxy has a small inclination of $i\\approx$21$\\degr$, and a major axis PA$\\approx$22$\\degr$ \\citep{pedlar92,davies73,mundell99b,das05}. Its radial velocity is $cz=997$\\,km\\,s$^{-1}$ \\citep{pedlar92}, the Hubble type is (R')SAB(rs)ab, and we will adopt in this paper a distance of 13.3\\,Mpc, corresponding to a scale at the galaxy of 65\\,pc\\,arcsec$^{-1}$ \\citep{mundell03}.  In radio continuum observations NGC\\,4151 shows a linear structure comprising several knots elongated over $\\sim 3\\farcs$5 along PA$\\approx$77$\\degr$, which is embedded in a diffuse emission extending over $\\sim 10\\farcs$5 \\citep{pedlar93,mundell95}. More recent high resolution radio images \\citep{mundell03} reveal a faint jet underlying the discrete components which seem to be shocklike features produced by interactions of the jet with gas clouds in the galaxy, as well as neutral gas absorption consistent with being due to an obscuring torus.  In the optical, the narrow-line region (hereafter NLR) of NGC\\,4151 has been found to have a biconical morphology (both in [O{\\sc\\,iii}]$\\lambda$5007 and H$\\alpha$ emission lines), with the line-of-sight outside but close to the edge of the cones \\citep{evans93,hut98}. The projected opening angle of the cones is $\\sim$75$\\degr$ and the projected axis is oriented along a position angle PA$\\sim$60/240$\\degr$. Optical spectroscopy reveals outflows along the cones, with the approaching side to the SW \\citep{med95,evans93}. The outflows have been modelled in a number of studies \\citep[e.g.][]{das05,crenshaw00,hutchings99}, in which it is also argued that the AGN should be the primary ionization source of the bicone as no clear correlation with the radio jet is observed. The radio jet may be nevertheless associated with high velocity clouds observed in the NLR \\citep{winge97}.  Combining X-ray, UV and optical spectra, \\citet{kra05,kra06} and \\citet{crenshaw07} where able to characterize outflows estimated to be at only $\\sim$0.1\\,pc from the nucleus, obtaining a high mass outflow rate ($\\sim$0.16\\,M$_\\odot$\\,yr$^{-1}$) which is about 10 times the accretion rate necessary to feed the AGN in NGC\\,4151. They sugest that this outflow originates in an accretion disk wind.  In the near-infrared, \\citet{thompson95} obtained spectra from 0.87$\\mu$m to 2.5$\\mu$m, and concluded that, from the high emission-line ratios of [Fe{\\sc\\,ii}]/H{\\sc\\,i}, the majority of the Fe should be in gaseous form, thus implying grain destruction to release a significant fraction of the iron usually tied up in dust. He obtained an electron temperature $T_e\\approx\\,10^4$K and density $N_e\\approx\\,10^4$\\,cm$^{-3}$ for the region emitting  [Fe{\\sc\\,ii}]  lines. Long-slit J-band observations of the kinematics of the [Fe{\\sc\\, ii}]$\\lambda$1.2570$\\mu$m and Pa$\\beta$ emission lines showed similar velocity structure to that observed in [O{\\sc iii}] \\citep{knop96}, which was later confirmed with two-dimensional Integral Field Unit (hereafter IFU) observations \\citep{turner02}. This similarity supports the same origin for [O{\\sc\\,iii}] and [Fe{\\sc\\,ii}] emission, namely photoionization by the AGN in the NLR. Nevertheless,  broadening of the [Fe{\\sc\\,ii}] emission relative to Pa$\\beta$, which is observed along the NLR \\citep{knop96} suggests additional processes contributing to the [Fe{\\sc\\,ii}] emission, such as shocks from  an AGN wind or jet.  Despite being a well-studied galaxy \\citep{ulrich00}, the dynamics and excitation of the NLR, as well as the role of the radio jet, are  not yet fully understood. In the present paper we use the Gemini Near-infrared Integral Field Spectrograph (NIFS) equipped  with the adaptive optics module ALTAIR -- to map the NLR gas distribution and excitation in the inner $\\approx 200$\\,pc$\\times$400\\,pc, at a spatial resolution of $\\approx$7\\,pc at the galaxy. The data cover the wavelength range 0.95--2.51~$\\mu$m at a spectral resolving power over 5000. In a companion paper \\citep{sl08} (hereafter Paper II), we use these data to obtain the gas kinematics and present emission-line channel maps of the NLR obtained by slicing the strongest emission-line profiles in velocity bins of 60\\,km\\,s$^{-1}$.  The present paper is organized as follows. In section 2 we describe the observations and reductions. In section 3 we present the results, which include flux measurements of 55 emission lines and 14 intensity distribution maps as well as line ratio maps. In section 4 we derive physical parameters for the NLR and discuss the origin of the [Fe{\\sc\\,ii}] and H$_2$ emission, and in section 5 we present our conclusions. ",
        "conclusions": "We have mapped the emitting gas intensity distributions, reddening and excitation in the NLR of NGC\\,4151 using Gemini NIFS observations of the inner $\\approx 200 \\times$300\\,pc$^2$ of the galaxy, covering the near-IR Z, J, H and K spectral bands, at a resolving power R$\\ge$5000 and spatial resolution of $\\approx$8\\,pc. The main results of this paper are: \\begin{itemize} \\item The intensity distributions in the recombination lines of H and He as well as in  the  [S{\\sc\\,iii}],  [P{\\sc\\,ii}], [ [Fe{\\sc\\,ii}] and [S{\\sc\\,iii}] emission lines, similarly to that in the optical  [O{\\sc\\,iii}], are most extended along PA=$60/240\\degr$ -- the axis of the previously known bicone. The emitting region is somewhat brighter and reaches a larger projected distance of $\\approx$130\\,pc from the nucleus to the SW (the near side of the bicone) than to the NE (the far side of the bicone), where it reaches a distance of $\\approx$100\\,pc. The fluxes in the recombination lines decrease with distance $r$ from the nucleus as $\\propto\\,r^{-1}$, while those on the other emission lines above  remain constant or even increase with distance from the nucleus along the NLR.  \\item The H$_2$ intensity distribution is completely different from that of the ionized gas, avoiding the region of the bicone, probably due to the destruction of the H$_2$ molecule by the strong ionizing flux along the bicone. Most of the H$_2$ emission seems to be coming not from the outflowing gas, but from gas which is in the plane of the galaxy. This is supported by its kinematics (Paper II). The concentration of the  H$_2$ emission approximately perpendicular to the bicone indicates that there is attenuation of the ionizing flux at these locations, precluding the destruction of the H$_2$ molecule. This attenuation may be due to an obscuring torus and/or to the bottom part of the biconical outflow (as suggested by \\citet{kra08}).   \\item The intensity distribution in the coronal line [Si{\\sc\\,vii}] is resolved and that in the [Ca{\\sc\\,viii}] is marginally resolved, consistent with an origin  in the inner NLR.  The  intensity distributions in these and the other coronal lines --  [S{\\sc\\,viii}] and [S{\\sc\\,ix}] -- decrease steeply with distance as $\\propto\\,r^{-2}$, similarly to those of the calibration stars. \\item The line ratios [Fe{\\sc\\,ii}]$\\lambda$1.257/1.644 and Pa$\\beta$/Br$\\gamma$ were used to map the reddening along the NLR, which ranges from E(B-V)=0 to E(B-V)=0.5. Within 0$\\farcs$5 from the nucleus, the reddening increases up to E(B-V)$\\ge$1.  \\item The [S{\\sc\\,iii}]/Pa\\,$\\beta$ line ratio has a value of $\\approx$11 to the SW (typical for the NLR of Seyfert galaxies), where we are looking at the inner wall of the near cone, and is $\\approx$ 30\\% lower to the NE, where we are looking at the outer wall of the far cone. This difference seems not to be due to reddening and  indicates  lower excitation to the NE.  \\item The line-ratio map [Fe{\\sc\\,ii}]\\,1.257/[P{\\sc\\,ii}]\\,1.187 of the NGC\\,4151 NLR is the first 2D such map of an extragalactic source, and, similarly to  the [Fe{\\sc\\,ii}]\\,1.257/Pa\\,$\\beta$ ratio shows larger values at the locations where the [Fe{\\sc\\,ii}] emission is enhanced at $\\approx$\\,1\\arcsec\\ from the nucleus. The increase in these line ratios maps the shocks produced by the radio jet in the NLR, which release the Fe{\\sc\\,ii} usually tied up in dust grains. This is confirmed by the correlation between the linhe-ratio maps and the outer parts of the radio map. From the many emission lines observed at the locations where the [Fe{\\sc\\,ii}] emission is enhanced, we have obtained the gas density, N$_e\\,\\approx$\\,4000\\,cm$^{-3}$ and temperature T$_e\\,\\approx$\\,15000$\\pm$5000K.   \\item From the fluxes of 10 H$_2$ emission lines, we have concluded that the H$_2$ emitting gas is in thermal equilibrium at the excitation temperature $T_{exc}=2155\\,K$, ruling out any significant contribution from fluorescence to the excitation of the H$_2$ molecule. The thermal excitation may be due to X-rays from the AGN escaping perpendicular to the bicone axis or else to shocks produced by the accretion flow observed along the bar \\citep{mundell95} when the gas reaches the nuclear region.   \\item We have calculated the mass of the  ionized and molecular gas, obtaining for the former $M_{HII}\\approx\\,2.4\\,\\times10^6\\,{\\rm M_\\odot}$ and for the latter only M$_{H_2}\\,\\approx$\\,240\\,M$_\\odot$. This small  mass is  nevertheless only that  of the ``hot skin''  of a probably much larger molecular gas mass.  \\item The distinct intensity distribution and smaller temperature (as well as distinct kinematics; see Paper II) of the H$_2$ emission when compared with those of the ionized gas, supports a distinct origin for the emitting gas.  The H$_2$ emission is probably a tracer of a large molecular gas reservoir,  built up from the accretion flow observed along the large scale bar, which may  be the source of fuel to the AGN. The H$_2$ emission can thus be considered a tracer of the {\\it feeding} of the AGN, while the ionized gas emission, which maps the outflowing gas, is a tracer of the {\\it feedback} from the AGN.   \\end{itemize}"
    },
    "0812/0812.2481_arXiv.txt": {
        "abstract": "We present a table of redshifts for 2907 galaxies and stars in the 145~arcmin$^2$ {\\em HST\\/} ACS GOODS-North, making this the most spectroscopically complete redshift sample obtained to date in a field of this size. We also include the redshifts, where available, in a table containing just under 7000 galaxies from the ACS area with $K_{s,{\\rm AB}}<24.5$ measured from a deep $K_s$ image obtained with WIRCam on the CFHT, as well as in a table containing 1016 sources with NUV$_{\\rm AB}<25$ and 478 sources with FUV$_{\\rm AB}<25.5$ (there is considerable overlap) measured from the deep GALEX images in the ACS area. Finally, we include the redshifts, where available, in a table containing the 1199 $24~\\mu$m sources to 80~$\\mu$Jy measured from the wider-area {\\em Spitzer\\/} GOODS-North. The redshift identifications are greater than 90\\% complete to magnitudes of ${\\rm F435W\\/}_{\\rm AB}=24.5$, ${\\rm F850LP\\/}_{\\rm AB}=23.3$, and $K_{s,{\\rm AB}}=21.5$ and to 24~$\\mu$m fluxes of 250~$\\mu$Jy. An extensive analysis of these data will appear in a parallel paper, but here we determine how efficient color-selection techniques are at identifying high-redshift galaxies and Active Galactic Nuclei. We also examine the feasibility of doing tomography of the intergalactic medium with a 30~m telescope. ",
        "introduction": "\\label{secintro}  Substantial progress has been made in tracing the formation and evolution of galaxies through extensive multiwavelength observations of extragalactic survey fields. However, the most challenging and time-consuming task is spectroscopically identifying the sources. This is especially true when one is trying to identify complete samples to faint magnitude limits. Although with existing multiwavelength data sets we can construct spectral energy distributions (SEDs) and measure photometric redshifts for galaxies,  spectra provide an enormous amount of additional information on the sources. We have therefore carried out a very high quality, nearly spectroscopically complete, magnitude limited redshift survey of the Great Observatories Origins Deep Survey-North (GOODS-N; Giavalisco et al.\\ 2004) field. The GOODS-N is one of the most intensively studied regions in the sky, and in many bandpasses it has the deepest images ever obtained. It has also been the target of extensive spectroscopic observations over the years (e.g., Cohen et al.\\ 2000; Wirth et al.\\ 2004; Cowie et al.\\ 2004). Thus, it is nearly ideal for the present study.  In a parallel paper (Cowie \\& Barger 2008) we use this unique spectral database to place galaxies into the overall evolutionary scheme by combining spectral line and break diagnostics with colors and by analyzing metal evolution. Here we present the data catalogs and use the redshift information to test the efficiency of using color-selection techniques to identify populations of high-redshift galaxies and Active Galactic Nuclei (AGNs). We also determine which AGNs can be distinguished by their colors. We then use the data set to examine the feasibility of doing intergalactic medium (IGM) tomography on a 30~m telescope.  Techniques have been developed to pre-select high-redshift candidates for spectroscopic confirmation using simple photometric selection criteria. The Lyman Break Galaxy (LBG) dropout technique (e.g., Cowie et al.\\ 1988; Songaila et al. 2000; Lilly et al. 2001; Steidel \\& Hamilton 1993; Steidel et al.\\ 1995) is the most heavily used of these and has really opened up the field to the study of $z>3$ star-forming galaxies. However, with the advent of wide-area near-infrared (NIR) arrays, other selections, including the Distant Red Galaxy (DRG) color selection (Franx et al.\\ 2003; van Dokkum et al.\\ 2003; $z\\gtrsim 2$) and the $BzK$ color selection (Daddi et al.\\ 2004; $1.4\\lesssim z\\lesssim 2.5$), have recently become popular. The above selections are, of course, limited to sources that can be observed in optical and NIR surveys, but they have produced a substantial sample of spectroscopically confirmed high-redshift galaxies and AGNs. H$^{-}$ methods using the mid-infrared (MIR) spectral shapes have also been suggested (e.g., Simpson \\& Eisenhardt 1999; Sawicki 2002). Although these techniques have not been heavily used, they are of considerable interest for optically faint sources.  With all of these avenues for finding high-redshift sources, we would hope that there are not still significant populations that remain unidentified. Recently, however, Le F{\\`e}vre et al.\\ (2005) and Paltani et al.\\ (2007) have challenged this view using the VIMOS Very Large Telescope Deep Survey (VVDS), a purely $I$ flux-selected sample (targets have magnitudes between $I_{AB}=17.5$ and $I_{AB}=24$) of $9295$ galaxies (``first epoch'') with measured redshifts in the range $0\\le z \\le 5$. These authors claim surface densities consistently larger than those of Steidel et al.\\ (1996, 1999, 2004) by factors of $1.6-6.2$ at $z\\approx3$, with the largest factor applying to the brightest magnitudes. In calculating the surface densities they corrected for the fraction of galaxies they did not observe ($\\sim 76$\\%) and for their estimate of the fraction of galaxies with incorrect redshifts. The need for the latter correction arises from their use of redshifts with all four of their reliability flags (flags 3 and 4 have very reliable measurements; flag 2 has a reasonably reliable measurement, but there is a non-negligible probability that the redshift is wrong; flag 1 has a tentative measurement, but there is a significant probability that the redshift is wrong). They concluded that the LBG selection techniques and photometric redshift studies are not able to identify the full population of high-redshift galaxies found through spectroscopic observations of pure magnitude selected samples. This claim requires verification from a highly spectroscopically complete sample with very reliable redshift identifications.  A 30~m telescope will revolutionize wide-field spectroscopy, making sources that are currently 10 times too faint for moderate-to-high dispersion spectroscopy in the optical accessible and mapping the distribution of galaxies over the redshift range $2\\le z \\le 4$. Of most interest, however, is whether a 30~m telescope will allow dense sampling of the IGM through the use of background galaxies. This would be an enormous improvement over the one-dimensional information now obtained through the study of rare high-redshift quasars, whose surface density on the sky at magnitudes accessible with Keck is very low. To trace the three-dimensional distribution of diffuse gas at high redshifts would require a high density of background probes, but it needs to be determined whether the necessary surface density of such probes exists.  The structure of the paper is as follows. In \\S\\ref{secsurvey1} and \\S\\ref{secsurvey2} we describe, respectively, our photometry and our spectroscopy. In \\S\\ref{seccolor} we examine how well different color selection techniques do, including LBG selection in \\S\\ref{secLBG}, $BzK$ selection in \\S\\ref{secK}, and H$^-$ selection in \\S\\ref{sechminus}. In \\S\\ref{sectom} we consider whether IGM tomography is feasible with a 30~m telescope, and in \\S\\ref{secsummary} we summarize our results.  We assume $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$ throughout. All magnitudes are given in the AB magnitude system, where an AB magnitude is defined by $m_{AB}=-2.5\\log f_\\nu - 48.60$. Here $f_\\nu$ is the flux of the source in units of ergs~cm$^{-2}$~s$^{-1}$~Hz$^{-1}$. ",
        "conclusions": "\\label{secsummary}  In this paper we presented deep $K_s$-band imaging and the most spectroscopically complete redshift sample obtained to date of the 145~arcmin$^2$ {\\em HST\\/} ACS GOODS-N region. We provided the data in a variety of tables as a function of magnitude from the NUV to the MIR. The redshift identifications are greater than 90\\% complete to magnitudes of ${\\rm F435W\\/}_{\\rm AB}=24.5$, ${\\rm F850LP\\/}_{\\rm AB}=23.3$, and $K_{s,{\\rm AB}}=21.5$ and to 24~$\\mu$m fluxes of 250~$\\mu$Jy.  We used the data to analyze various color selection techniques that have been proposed to find high-redshift galaxies. The LBG technique still appears to be the most robust method for providing a high selection completeness with low contamination. We do not confirm the presence of additional luminous galaxies not picked out by the LBG selection, as proposed by Le F{\\`e}vre et al.\\ (2005). $BzK$, H$^-$, and IRAC color techniques can be highly efficient at selecting high-redshift ($z>1.4$, $z>1.3$, and $z>1.6$, respectively) galaxies and AGNs but at the expense of a fairly high degree of contamination by lower redshift galaxies. Samples selected using these techniques therefore need to be spectroscopically followed up in order to obtain clean high-redshift samples.  Finally, we have used the ACS images to make detailed calculations of the S/N which can be obtained with spectral resolution 6000 seeing-limited spectra on a 30~m telescope. We find that spectra with S/N$\\ge 30$ per resolution element in 10~hr integrations can only be obtained for galaxies brighter than ${\\rm F606W}_{\\rm AB}=23.7$. We have also computed the observed cumulative surface densities of galaxies and AGNs as a function of magnitude in the $z=2-3$ and $z=3-4$ ranges, as well as the maximum surface densities which could lie in these intervals when unidentified sources are included. These surface densities suggest that moderate spectral resolution (6000) observations with S/N$\\ge 30$ on a 30~m telescope will still primarily be of the AGN population and that the samples will be sparse with average separations greater than 1~arcmin at $z=2-4$."
    },
    "0812/0812.3164_arXiv.txt": {
        "abstract": "We present three 43~GHz images and a single 86~GHz image of Markarian 501 from VLBA observations in 2005. The 86~GHz image shows a partially resolved core with a flux density of about 200~mJy and a size of about 300 Schwarzschild radii, similar to recent results by Giroletti et al. Extreme limb-brightening is found in the inner parsec of the jet in the 43~GHz images, providing strong observational support for a `spine-layer' structure at this distance from the core. The jet is well resolved transverse to its axis, allowing Gaussian model components to be fit to each limb of the jet. The spine-layer brightness ratio and relative sizes, the jet opening angle, and a tentative detection of superluminal motion in the layer are all discussed. ",
        "introduction": "\\label{intro} Markarian 501 is among the best known and most prominent of the TeV blazars. It was the second such source to be discovered (Quinn et al. 1996), and during flaring episodes in 2005 it displayed TeV variability on timescales of several minutes (Albert et al. 2007). The remarkably fast variability found for this source (and for PKS~2155$-$304) may require very high Lorentz factors ($\\Gamma>50$) in the gamma-ray production region (Begelman et al. 2008).  In contrast to the TeV observations, VLBI studies of jet kinematics in the TeV blazars (including Mrk 501, Edwards \\& Piner [2002]) have suggested much lower Lorentz factors ($\\Gamma\\sim3$) in their parsec-scale VLBI jets (Piner et al. 2008). A number of explanations for this discrepancy have been discussed in the literature, with one possibility being that the jet is structured transverse to its axis into a fast `spine' that produces the high-energy emission, and a slower `layer' responsible for the radio emission (e.g., Ghisellini et al. 2005). Transverse jet structures are also expected on theoretical grounds (e.g., Zakamska et al. 2008). Limb-brightened VLBI jets that confirm such transverse structures have been seen in M87 (Ly et al. 2007; Kovalev et al. 2007) and in Mrk~501 itself at frequencies below 22~GHz (Giroletti et al. 2004, hereafter G04), and tentatively at high frequencies (Giroletti et al. 2008, hereafter G08).  In this paper, we report the detection of significant limb-brightening in the inner parsec of Mrk~501, based on a sequence of three 43~GHz VLBA images from 2005. We use cosmological parameters $H_{0}=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}=0.27$, and $\\Omega_{\\Lambda}=0.73$. At the redshift of Mrk~501 (0.034), an angular separation of 1 milliarcsecond (mas) corresponds to a projected linear distance of 0.66~pc. We assume a black hole mass of $10^{9} M_{\\odot}$ (Wagner 2008), consistent with that used by G08. ",
        "conclusions": "The high resolution of the 43~GHz VLBA observations, combined with the sensitivity of their 512 Mbps bandwidth, has enabled us to resolve the spine-layer structure of the jet in the inner parsec of Mrk~501 with unprecedented clarity. For the first time, our data resolve the transverse structure well enough that separate Gaussian models can be fit to each limb, which is crucial for measurements of the structure deconvolved from the observing beam. Below we enumerate the major results on the core and limb-brightened jet obtained in this Letter: \\begin{enumerate} \\item{The $3\\sigma$ lower limit on the layer/spine brightness ratio measured from the Gaussian model fits at 0.5 pc (projected) from the core is $>$10:1.} \\item{A brightness peak in the layer appears to move outward with an apparent pattern speed of $3.3\\pm0.3~c$.} \\item{The jet is approximately cylindrical from 0.2 to 0.7 pc (projected) from the core, and the apparent opening half angle decreases from about $30\\arcdeg$ to about $10\\arcdeg$ over this region (Figure~4).} \\item{An estimate of the relative sizes of the layer and spine (from an epoch where the minor axes of the ellipses fit to the layer are resolved) shows that the volume of the layer is about twice that of the spine.} \\item{The detection of the 86~GHz core constrains the most compact radio emission in Mrk~501 to come from a size less than 300x200 Schwarzschild radii, the same size upper limit found by G08.} \\end{enumerate} These results can all aid in constraining fundamental properties of the limb-brightened structure; like the Lorentz factor, viewing angle, and Doppler factor of the spine and layer, which are important for modeling the radio to TeV gamma-ray emission from the two transverse regions of the jet in this source. Further observations currently in the VLBA's dynamic queue should reveal more about the nature of the inner jet structure of Mrk~501, during the first year of Fermi (GLAST) observations.  \\vspace{-0.1in}"
    },
    "0812/0812.1161_arXiv.txt": {
        "abstract": "We report the identification of 32 transiting-planet candidates in HATNet field G205.  We describe the procedures that we have used to follow up these candidates with spectroscopic and photometric observations, and we present a status report on our interpretation of the 28 candidates for which we have follow-up observations. Eight are eclipsing binaries with orbital solutions whose periods are consistent with their photometric ephemerides; two of these spectroscopic orbits are singled-lined and six are double-lined. For one of the candidates, a nearby but fainter eclipsing binary proved to be the source for the HATNet light curve, due to blending in the HATNet images.  Four of the candidates were found to be rotating more rapidly than $v \\sin i = 50$ \\kms\\ and were not pursued further. Thirteen of the candidates showed no significant velocity variation at the level of 0.5 to 1.0 \\kms.  Seven of these were eventually withdrawn as photometric false alarms based on an independent reanalysis using more sophisticated tools.  Of the remaining six, one was put aside because a close visual companion proved to be a spectroscopic binary, and two were not followed up because the host stars were judged to be too large.  Two of the remaining candidates are members of a visual binary, one of which was previously confirmed as the first HATNet transiting planet, HAT-P-1b.  In this paper we confirm that the last of this set of candidates is also a a transiting planet, which we designate \\hatcur{b}, with mass $M_{\\rm p} = \\hatcurPPm\\,\\mjup$, radius $R_{\\rm p} = \\hatcurPPr\\,\\rjup$, and photometric period $P = \\hatcurLCP$ days. \\hatcur{b} has an inflated radius for its mass, and a large mass for its period.  The host star is a solar-metallicity F dwarf, with mass $\\mstar = \\hatcurYYm\\,\\msun$ and $\\rstar = \\hatcurYYr \\rsun$. ",
        "introduction": "\\label{sec:introduction}  Photometric surveys are now finding many hundreds of stars that exhibit periodic dimmings that look like they might be due to transiting planets, but so far the vast majority of these candidates have turned out to be stellar systems involving eclipsing binaries and not planets.  To make a convincing case that the light curve is due to a transiting planet it is necessary, but not sufficient, to show that the detailed shape of the dimming, including depth and duration, can be accurately modeled as a planet.  The proof of the pudding is then an orbit for the host star that matches the photometric ephemeris and implies a planetary mass for the companion.  For a spectroscopic orbit it is also necessary to demonstrate that the measured velocity variations are not due to astrophysical effects other than orbital motion.  As of 1 July 2008, spectroscopic orbits had been properly published for 30 stars hosting transiting planets.  In five cases (HD 209458, HD 149026, HD 189733, Gliese 436, and HD 17156) the orbit came first from Doppler surveys, and transits were subsequently identified.  These five have been the source of the most interesting exoplanet astrophysics that has emerged so far, primarily because they are the nearest and brightest.  In contrast, the initial excitement of confirming a few planets from the OGLE photometric surveys has now cooled because of the difficulty of making follow-up observations for such faint systems.  Wide-angle ground-based photometric surveys utilizing small telescopes are now the main source of new transiting planets, with 19 confirmed and published planets from TrES, HAT, XO, and WASP, and many more on the way.  Although generally fainter than the five from Doppler surveys, some of these are bright enough for interesting follow-up work, and the growing numbers are fleshing out the extraordinary population of close-in giant planets.  At the Harvard-Smithsonian Center for Astrophysics (CfA), some of us have been following up transiting planet candidates for almost ten years, starting with targets supplied by the Vulcan team in 1999.  We soon learned that the sample was dominated by systems involving eclipsing binaries \\citep{Latham:03}.  Indeed, we have not yet been able to confirm that any of the 66 Vulcan candidates that we observed are actually planets.  More recently, two other wide-angle surveys, TrES and HATNet, have been the source of a large number of transiting-planet candidates that have been followed up initially at CfA.  As of 1 July 2008 a total of 811 candidates from these two surveys had been observed or scheduled for observations, but as of 1 July 2008 only 11 of these had led to proper publication as confirmed planets.  Over the years our procedures for reducing and analyzing HATNet photometry and for identifying good transiting-planet candidates have evolved considerably, and we have encountered a variety of stellar systems that can mimic transiting planets.  Thus one of the goals of this paper is to report our present procedures and some of the lessons learned.  A second, and perhaps more important goal is to document the false positives that we have identified in the process of our work to follow up candidates identified in HATNet field G205.  As more of the sky gets covered by wide-angle surveys, it is inevitable that some of the same candidates will be identified by more than one survey.  Ways must be found to document the follow-up work that has already been done, to avoid wasteful duplication of effort.  Last, but not least, we confirm and characterize \\hatcur{b}, a second planet to join with HAT-P-1b in field G205. ",
        "conclusions": "\\label{sec:discussion}  In this paper we describe the procedures and observations that we use to follow up transiting-planet candidates identified by wide-angle ground-based photometric surveys, and we report the status of our efforts to follow up 32 candidates identified in HATNet field G205.  Initial spectroscopic observations obtained with the CfA Digital Speedometers disclosed that twelve of the candidates showed variable radial velocities with amplitudes of at least a few \\kms, indicating that in each case the companion responsible for the velocity variations must be too massive to be a planet.  We accumulated enough observations for ten of these spectroscopic binaries to allow the derivation of orbital solutions.  Eight of these orbits have parameters that are consistent with the photometric ephemerides, thus proving that the stellar companion is the source of the transit-like light curve.  For all but two of these eight eclipsing binaries the original HATNet photometric period turned out to be half of the true orbital period, presumably because the primary and secondary eclipses were not distinguishable in the HATNet photometry, implying that the two stars must be similar.  Given the shallowness of the observed dips, the eclipses must be grazing, a conclusion that is supported by the fact that we were able to detect the spectrum of the secondary and derive double-lined orbits for six of the seven.  For the two eclipsing binaries where the orbital period matches the original HATNet photometric period, the companions are M dwarfs, small enough to produce a transit-like light curve with full eclipses of the F-star primary but no detectable secondary eclipses \\citep{Beatty:07}.  \\begin{figure}[!ht] \\plotone{exoplanet_m_r_without_lowmass.eps} \\caption{ Mass--radius diagram for published transiting extrasolar planets.  The data are taken from \\citet{Torres:08}, and www.exoplanet.eu.  Filled circles mark HATNet discoveries, and the large filled circle marks \\hatcur{b}. Overlaid are \\citet{Baraffe:03} (zero insolation planetary) isochrones for ages of 0.5~Gyr (upper, dotted line) and 5~Gyr (lower dashed-dotted line), respectively. \\hatcur{b} lies well above the main stream for Hot Jupiters. \\label{fig:mr}} \\end{figure}   Two of the single-lined spectroscopic binaries have orbital periods considerably longer than the HATNet photometric periods.  Thus the unseen companion responsible for the orbital motion can not be the source of the transit-like light curve.  Possible explanations include hierarchical triple systems with an eclipsing companion to either one of the stars in the observed outer orbit, or a quadruple system with an eclipsing binary in a much wider orbit than the one observed.  The eclipsing binary could even be an accidental alignment, unresolved in our images.  There is also the possibility that the original HATNet detection was a photometric false alarm.  We also used the spectra from the CfA Digital Speedometers to estimate the effective temperature and surface gravity of the candidate stars. In seven cases our spectroscopic gravities implied that the stars are giants, too large to allow detection of a transiting planet with HATNet, let alone that with such short periods, such putative planets would orbit inside the giant star. The possible explanations for these systems are similar to those in the previous paragraph. Either the giant is diluting the light of an eclipsing binary, or the photometric detection is a false alarm.  Indeed, one of the giants is also the primary in one of the two spectroscopic binaries discussed in the previous paragraph.  Four of the candidates are rotating too rapidly to allow very precise radial-velocity measurements, and two of these also showed large velocity variations based on the CfA spectra.  Thus these four candidates were not pursued further.  Four of the candidates have not yet been followed up with initial reconnaissance spectroscopy because they are fainter than $V=13.0$ mag and difficult targets for the old CfA Digital Speedometers.  We plan to follow them up in the near future with the Tillinghast Reflector Echelle Spectrograph, a modern fiber-fed echelle spectrometer that has recently come into operation at FLWO.  Spectra obtained with the CfA Digital Speedometers failed to show velocity variations at the level of about 0.5 to 1.0 \\kms\\ for seven of the candidates.  Two of these were subsequently withdrawn as photometric false alarms.  One is quite hot and has not yet been followed up with very precise radial velocities. One is a member of a close visual binary where the companion is a spectroscopic binary that is likely to be the source of the dip in the HATNet light curve.  Two are the components of another visual binary, one of which was previously confirmed as the host of the transiting planet HAT-P-1b.  In this paper we report that the seventh of these candidates is also a planet, \\hatcur{b}.  Our newly confirmed planet has mass $M_{\\rm p} = \\hatcurPPm\\,\\mjup$\\ and radius $R_{\\rm p} = \\hatcurPPr\\,\\rjup$, which places it among the most inflated of the transiting Hot Jupiters currently known (see Figure~\\ref{fig:mr}).  The 3.3\\,\\lsun\\ luminosity of the host star and the $a=\\hatcurPParel$\\,AU semi-major axis correspond to an equivalent semi-major axis of $a=0.026$\\,AU if this object orbited our Sun and received the same flux. (Note, however, that the spectrum of the solar flux would be different).  The theoretical radius from \\citet{Fortney:08} would be 1.32\\rjup, 1.25\\rjup\\ and 1.17\\rjup, respectively, for ages of 300\\,Myr, 1\\,Gyr and 4.5\\,Gyr, assuming coreless models, and 0.02\\,AU orbital distance. \\hatcur{b} with its \\hatcurPPr\\,\\rjup\\ radius is thus inflated at the 2-3-$\\sigma$ level. This is also consistent with \\citet{liu:08}, where the 1.5\\,\\mjup\\ mass and 0.026\\,AU equivalent solar distance yields such a large equilibrium radius only if the ratio between core heating power and insolation power is as high as $2\\times10^{-3}$. The equilibrium time-scale for the radius evolution would be very short, approx 20\\,Myr. For comparison, WASP-5b \\citep{Anderson:08} has a similar mass (1.58\\,\\mjup), but its radius is only 1.09\\,\\rjup.  \\begin{figure}[!ht] \\plotone{exoplanet_P_g.eps} \\caption{ The period vs.~surface gravity diagram for published transiting extrasolar planets. The data are taken from \\citet{Torres:08}, and www.exoplanet.eu. Filled circles mark HATNet discoveries, and the large filled circle marks \\hatcur{b}, which lies on the main stream for Hot Jupiters, if the supermassive planets, such as HAT-P-2b and XO-3b, are ignored. \\label{fig:pg}} \\end{figure}   \\hatcur{b} is also offset to the high-mass side of the period--mass relation. Planets with orbital periods similar to \\hatcur{b} tend to have smaller masses, around 0.8\\,\\mjup\\ as compared to the \\hatcurPPmshort\\,\\mjup\\ mass of \\hatcur{b}.  However, with the discovery of more massive planets, such as HAT-P-2b \\citep{Bakos:07a}, XO-3b \\citep{Johns-Krull:08}, and Corot-Exo-3b \\citep{Deleuil:08} the period--mass relation has become less clear-cut.  Although the mass and radius of \\hatcur{b} are larger than usual, together they imply an unremarkable surface gravity, $\\log g = \\hatcurPPlogg$ (cgs), and \\hatcur{b} falls tightly on the period--surface gravity relations reported by \\citet{Southworth:07} and by \\citet{Torres:08}.  The Safronov number of \\hatcurb\\ is $\\Theta = \\hatcurPPtheta$, indicating that it belongs to the so-called Class-I planets as defined by \\citet{Hansen:07}. Curiously, \\hatcur{b} lies at the very hot end of the distribution in the equilibrium temperature vs.~Safronov number diagram. All the hotter transiting planets are in Class II\\@. Furthermore, while ``inflated'' planets mostly belong to Class II, \\hatcur{b} has an inflated radius, yet belongs to Class I. With the recent discovery of several new transiting exoplanets, the Safronov dichotomy is becoming less pronounced, e.g.~WASP-11b/HAT-P-10b \\citep{West:08, Bakos:08} falls in between the two groups, and the high- and low-mass exoplanets also seem to fall outside these groups.    The incident flux on \\hatcur{b} is $1.91 \\times 10^{9}$ erg s$^{-1}$ cm$^{-2}$, placing it in the pM class, as defined by \\citet{Fortney:08}. This implies that \\hatcur{b} has significant opacity due to the absorption by molecular TiO and ViO in its atmosphere, leading to a temperature inversion and a hot stratosphere. This could be tested with future {\\em Spitzer Space Telescope} \\/observations. Based on its mass and incident flux, \\hatcur{b} is similar to TrES-3 \\citep{O'Donovan:07}.    The Transiting Exoplanet Survey Satellite (TESS) is under study for a NASA SMEX mission to survey the entire sky for nearby bright transiting planets.  The TESS cameras are similar to the HATNet cameras, with a similar pixel size and similar image quality.  Thus TESS will share some of the same challenges that HATNet has faced in following up transiting-planet candidates.  TESS will observe targets selected from an input catalog, with the goal of avoiding evolved stars.  This should eliminate one source of false positives that has plagued the magnitude-limited ground-based photometric surveys.  In the case of HATNet field G205, 8 out of 28 candidates proved to be evolved stars.  In addition, the most interesting targets for TESS will be the nearest and brightest stars, for example roughly 50,000 Main Sequence stars from the Hipparcos Catalogue.  Because these stars have been much better studied over the years than the fainter stars typical of the HATNet candidates, much more will be known about them ahead of time.  TESS will achieve roughly ten times better limiting photometric precision than HATNet, and the frequency of grazing eclipsing binaries masquerading as transiting planets will be much smaller for light curves with shallower dips, compared to the 7 grazing eclipsing binaries out of 28 candidates found in HATNet field G205.  On the other hand, some of the shallower dips will be due to fainter eclipsing binaries that contaminate the light of the primary targets.  These will pose a difficult challenge in the follow-up of TESS candidates."
    },
    "0812/0812.3955_arXiv.txt": {
        "abstract": "The physics of neutron star crusts is vast, involving many different research fields, from nuclear and condensed matter physics to general relativity. This review summarizes the progress, which has been achieved over the last few years, in modeling neutron star crusts, both at the microscopic and macroscopic levels. The confrontation of these theoretical models with observations is also briefly discussed. ",
        "introduction": "\\label{sect.intro}  Constructing  models of neutron stars requires knowledge of the physics of matter with a density significantly exceeding the density of atomic nuclei. The simplest picture of the atomic nucleus is a drop of highly incompressible nuclear matter. Analysis of nuclear masses tells us that nuclear matter at saturation (i.e.\\ at the minimum of the energy per nucleon) has the density  $\\rho_0= 2.8\\times 10^{14}~\\mdens$, often called \\emph{normal nuclear density}. It corresponds to $n_0=0.16$ nucleons per fermi cubed. The density in the cores of massive neutron stars is expected to be as large as $\\sim 5\\mbox{\\,--\\,}10\\rho_0$ and in spite of decades of observations of neutron stars and intense theoretical studies, the structure of the matter in neutron star cores and in particular its equation of state remain the well-kept secret of neutron stars (for a recent review, see the book by Haensel, Potekhin and Yakovlev~\\cite{haensel-06}). The physics of matter with $\\rho\\sim 5\\mbox{\\,--\\,}10\\rho_0$ is a huge challenge to theorists, with observations of neutron stars being crucial for selecting a correct dense-matter model. Up to now, progress has been slow and based overwhelmingly on scant observation~\\cite{haensel-06}.  The outer layer of neutron stars with density $\\rho<\\rho_0$ -- the neutron star crust --  which is the subject of the present review, represents very different theoretical challenges and observational opportunities. The elementary constituents of the  matter are neutrons, protons, and electrons -- like in the atomic matter around us. The density is ``subnuclear'', so that the methods developed and successfully applied in the last decades to terrestrial nuclear physics can be applied to neutron star crusts. Of course, the physical conditions are extreme and far from terrestrial ones. The compression of matter by gravity crushes atoms and forces, through electron captures, the neutronization of the matter. This effect of huge pressure was already predicted in the 1930s (Sterne~\\cite{sterne-33}, Hund~\\cite{hund-36, hund-36b}). At densities $\\rho\\gtrsim 4\\times 10^{11}~\\mdens$ a fraction of the neutrons is unbound and forms a gas around the nuclei. For a density approaching $10^{14}~\\mdens$, some 90\\% of nucleons are neutrons while nuclei are represented by proton clusters with a small neutron fraction. How far we are taken from terrestrial nuclei with a moderate neutron excess! Finally, somewhat above $10^{14}~\\mdens$ nuclei can no longer exist -- they coalesce into a uniform plasma of nearly-pure neutron matter, with a few percent admixture of protons and electrons: we reach the bottom of the neutron star crust.  The crust contains only a small percentage of a neutron star's mass, but it is crucial for many astrophysical phenomena involving neutron stars. It contains matter at subnuclear density, and therefore  there is no excuse for the theoretical physicists, at least in principle: the interactions are known, and many-body theory techniques are available. Neutron star crusts are wonderful cosmic laboratories in which the full power of theoretical physics can be demonstrated and hopefully confronted with neutron star observations.  To construct neutron star crust models we have to employ atomic and plasma physics, as well as the theory of condensed matter, the physics of matter in strong magnetic fields, the theory of nuclear structure, nuclear reactions, the nuclear many-body problem,  superfluidity, physical kinetics, hydrodynamics, the physics of liquid crystals, and the theory of elasticity. Theories have to be applied under extreme physical conditions, very far from the domains where they were originally developed and tested. Therefore, caution is a must!  Most of this review is devoted to theoretical descriptions of various aspects of neutron star crusts. The plural ``crusts'' in the title is well justified; depending on the scenario of their formations, the crusts may be very different in their composition and structure as sketched in Figure~\\ref{fig.sect.intro.surfaces}. In Section~\\ref{sect.plasma} we briefly describe the basic plasma parameters relevant to crust physics and delineate various plasma regimes in the temperature-density plane. We also address the important question of the magnetic field.   \\epubtkImage{neutron_star_surfaces.png}{% \\begin{figure}[htbp] \\centerline{\\includegraphics[scale=0.8]{neutron_star_surfaces}} \\caption{Schematic pictures of various neutron star surfaces.} \\label{fig.sect.intro.surfaces} \\end{figure}}   The ground state of the crust (one of the possible ``crusts'') is reviewed in Section~\\ref{sect.groundstate}. We also discuss there the uncertainties concerning the densest bottom layers of the crust and we mention possible deviations from the ground state. As we describe in Section~\\ref{sect.accretion}, a crust formed via accretion is expected to be very different from that formed during the aftermath of a supernova explosion. We show how it can be a site for nuclear reactions. We study its thermal structure during accretion, and briefly review the phenomenon of X-ray bursts. We quantitatively analyze the phenomenon of deep crustal heating.  To construct a neutron star model one needs the equation of state (EoS) of the crust, reviewed in Section~\\ref{sect.eos}. We consider separately the ground-state crust, and the accreted crust. For the sake of comparison, we also describe another EoS of matter at subnuclear density -- that relevant for the collapsing type~II supernova core.  Section~\\ref{sect.structure} is devoted to the stellar-structure aspects of neutron star crusts. We start with the simplest case of a spherically-symmetric static neutron star and derive approximate formulae for crust mass and moment of inertia. Then we study the deformation of the crust in a rotating star. Finally, we consider the effects of magnetic fields on the crust structure. Apart from isotropic stress resulting from pressure, a solid crust can support an elastic strain. Elastic properties are reviewed in Section~\\ref{sect.elast}. A separate subsection is devoted to the elastic parameters of the so-called ``pasta'' layers, which behave like liquid crystals. The inner crust is permeated by a neutron superfluid. Various aspects of crustal superfluidity are reviewed in Section~\\ref{sect.super}. After a brief introduction to superconductivity and its relevance for neutron star crusts, we start with the static properties of neutron superfluidity, considering first a uniform neutron gas, and then discussing the effects of the presence of the nuclear crystal lattice. In the following, we consider superfluid hydrodynamics. We stress those points, which have been raised only recently. We consider also the important problem of the critical velocity above which superfluid flow breaks down. The interplay of superfluid flow and vortices is reviewed. The section ends with a discussion of entrainment effects. Transport phenomena are reviewed in Section~\\ref{sect.cond}. We present calculational methods and results for electrical and thermal conductivity, and shear viscosity of neutron star crusts. Differences between accreted and ground-state crusts, and the potential  role of impurities are illustrated by examples. Finally, we discuss the very important effects of magnetic fields on transport parameters. In Section~\\ref{sect.hydro}, we review macroscopic models of the crust, and we describe in particular a two-fluid model, which takes into account the stratification of crust layers, as well as the presence of a neutron superfluid. We show how entrainment effects between the superfluid and the charged components can be included using the variational approach developed by Brandon Carter~\\cite{carter-89}. Section~\\ref{sect.neutrino} is devoted to a description of the wealth of neutrino emission processes associated with crusts. We limit ourselves to the basic mechanisms, which, according to existing calculations, are the most important ones at subsequent stages of neutron-star cooling.  The confrontation of theory with observations is presented in Section~\\ref{sect.obs}. Neutron stars are born very hot, and we briefly describe in Section~\\ref{sect.obs.supernova} the present status of the theory of hot dense matter at subnuclear densities; this layer of the proto-neutron star will eventually become the neutron star crust. The crust is crucial for neutron star cooling, as observed by a distant observer. Namely, the crust separates the neutron star core from its surface, where the observed X-ray radiation is produced. The relation between crust physics and observations of cooling neutron stars is studied in Section~\\ref{sect.obs.cooling}.  In Section~\\ref{sect.obs.rproc} we briefly consider possible r-processes associated with the ejection and subsequent decompression of the neutron star crusts. Pulsar glitches are thought to originate in neutron star crust and glitch models are confronted with observations of glitches in pulsar timing data in Section~\\ref{sect.obs.glitches}. The asteroseismology of neutron stars from their gravitational wave radiation is discussed in Section~\\ref{sect.obs.gw}. Due to its elasticity, the solid crust can support mountains and shear (torsional) oscillations, both associated with gravitational wave emission. The crust-core interface can be crucial for the damping of r-modes, which, if unstable, could be a promising source of gravitational waves from rotating neutron stars. Observations of oscillations in the giant flares from Soft Gamma Repeaters are confronted with models of torsional oscillations of crusts in Section~\\ref{sect.obs.sgr}. As discussed in Section~\\ref{sect.obs.LMXB}, the modeling of phenomena associated with low-mass X-ray binaries (LMXB)  requires a rather detailed knowledge of the physics of neutron star crusts. New phenomena discovered in the last decade (and some very recently) necessitate realistic physics of accreted neutron star crusts, including deep crustal heating  and the correct degree of purity. All aspects of accreted crusts, relevant for soft X-ray transients, X-ray superbursts, and persistent X-ray transients, are discussed in this section.  \\newpage ",
        "conclusions": "\\label{sect.conclusion}  The conditions prevailing inside the crusts of neutron stars are not so extreme as those encountered in the dense core. Nonetheless, they are still far beyond those accessible in terrestrial laboratories. The matter in neutron star crust is subject to very high pressures, as well as huge magnetic fields, which can attain up to $10^{14}\\mbox{\\,--\\,}10^{15}$~G in magnetars. For comparison, it is worth reminding ourselves that the strongest (explosive) magnetic fields ever produced on the Earth reach ``only'' $3\\times10^7$~G~\\cite{motokowa-04}. The description of such environments requires the interplay of many different branches of physics, from nuclear physics to condensed matter and plasma physics.  Considerable progress in the microscopic modeling of neutron star crusts has been achieved during the last few years. Yet the structure and properties of the crust remain difficult to predict, depending on the formation and subsequent cooling of the star. Even in the ground state approximation, the structure of the neutron star crust is only well determined at $\\rho\\lesssim 10^{11}~\\mdens$, for which experimental data are available. Although all theoretical calculations predict the same large-scale picture of the denser layers of the crust, they do not quantitatively agree, reflecting the uncertainties in the properties of very exotic nuclei and uniform highly-asymmetric nuclear matter. The inner crust is expected to be formed of a lattice of neutron-rich nuclear clusters coexisting with a degenerate relativistic electron gas and a neutron liquid. The structure of the inner crust, its composition and the shape of the clusters are model dependent, especially in the bottom layers at densities $\\sim 10^{14}~\\mdens$. However, the structure of the inner crust is crucial for calculating many properties, such as neutrino emissivities, as well as transport properties like electric and thermal conductivities. Superfluidity of unbound neutrons in the crust seems to be well established, both observationally and theoretically. However, much remains to be done to understand its properties in detail and, in particular, the effects of the nuclear lattice.  The interpretation of many observed neutron star phenomena, like pulsar glitches, X-ray bursts in low-mass X-ray binaries, initial cooling in soft X-ray transients, or quasi-periodic oscillations in soft gamma repeaters, can potentially shed light on the microscopic properties of the crust. However, their description requires  consistent models of the crust from the nuclear scale up to the macroscopic scale. In particular, understanding the evolution of the magnetic field, the thermal relaxation of the star, the formation of mountains, the occurrence of starquakes and the propagation of seismic waves, requires the development of global models combining general relativity, elasticity, magnetohydrodynamics and superfluid hydrodynamics. Theoretical modeling of neutron star crusts is very challenging, but not out of reach. A confrontation of these models with observations could hopefully help us to unveil the intimate nature of dense matter at subnuclear densities. The improvement of observational techniques, as well as the development of gravitational wave astronomy in the near future, open very exciting perspectives.  We are aware that some aspects of the physics of neutron star crusts have not been dealt with in this review. Our intention was not to write an exhaustive monograph, but to give a glimpse of the great variety of topics that are necessarily addressed by different communities of scientists. We hope that this review will be useful to the reader for his/her own research.   \\newpage"
    },
    "0812/0812.1027_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\bigskip \\noindent A pivotal part in the ongoing search for extra-solar planets is the quest to identify planetary habitability, i.e., the principal possibility of life. In a previous paper, \\cite{kast93} presented a one-dimensional climate model to define a zone of habitability (HZ) around the Sun and other main-sequence stars that assumed as basic premise an Earth-like model planet with a CO$_2$/H$_2$O/N$_2$ atmosphere and that habitability requires the presence of liquid water on the planetary surface.  In the meantime, other definitions of habitable zones have been proposed such as the Galactic HZ, the UV-HZ, and the photosynthesis-sustaining HZ (pHZ).  The Galactic HZ \\citep{line04} caters to the requirement that a sufficient amount of heavy elements (notably those contained in carbon and silicate compounds) must be present for the build-up of planets and life, a condition easily met in the solar neighborhood.  The UV-HZ \\citep{bucc06,cunt08} is based on the premise that no lethal amounts of stellar UV flux is produced (regarding life forms assuming carbon-based biochemistry), a condition that tends to favor the environment of old main-sequence stars and giants \\citep{guin02} as well as planets with appreciable atmospheres, notable significant ozone layers \\citep{segu03}.  Another definition of habitability first introduced by \\cite{fran00a,fran00b} is associated with the photosynthetic activity of the planet, which critically depends on the planetary atmospheric CO$_2$ concentration. This type of habitability is thus strongly influenced by the planetary geodynamics, encompassing climatological, biogeochemical, and geodynamical processes (``Integrated System Approach'').  This concept has previously been used in studies of fictitious planets around 47~UMa \\citep{cunt03,fran03} and 55~Cnc \\citep{bloh03}, as well as detailed studies of observed super-Earth planets in the Gliese 581 system \\citep{bloh07b}.  The latter investigation showed that Gliese 581c is clearly outside the habitable zone, since it is too close to the star, whereas Gliese 581d located near the outer edge of the habitable zone is probably habitable, at least for certain types of primitive life forms \\citep[see also][]{sels07}.  Moreover, \\cite{bloh07a} have used this type of model to compile a detailed ranking of known star-planet systems regarding the principal possibility of life, which by the way led to the conclusion that the Solar System is not the top-tier system (``Principle of Mediocrity'').  In case of Earth-mass planets (1~$M_\\oplus$), a detailed investigation of geodynamic habitability was presented by \\cite{fran00b} with respect to the Sun as well as stars of somewhat lower and higher mass as central stars.  \\citeauthor{fran00b} found that Earth is rendered uninhabitable after 6.5~Gyr as a result of plate tectonics, notably the growth of the continental area (enhanced loss of atmospheric CO$_2$ by the increased weathering surface) and the dwindling spreading rate (diminishing CO$_2$ output from the solid Earth).  This implies that there is no merit in investigating the future habitability of Earth during the post--main-sequence evolution of the Sun, as in the framework of pHZ models, the lifetime of habitability is limited by terrestrial geodynamic processes.  However, this situation is expected to be significantly different for super-Earth planets due to inherent differences compared to Earth-mass planets \\citep[e.g.,][]{vale07a}. A further motivation for this type of work stems from the ongoing discovery of super-Earths in the solar neighborhood with the Gliese 876 \\citep{rive05} and Gliese 581 \\citep{udry07} systems as prime examples.  In the following, we discuss the definition of the photosynthesis-sustaining habitable zone, including the relevant geodynamic assumptions.  Next, we describe the most recent model of solar evolution that is used as basis for our study.  Thereafter, we present our results including comparisons to previous work.  Finally, we present our summary and conclusions. ",
        "conclusions": "We studied the habitability of super-Earth planets based on the integrated system approach that has previously been used in various theoretical planetary studies \\cite[e.g.,][]{fran00b,cunt03,bloh03,fran03,bloh07a,bloh07b}.  This work is motivated by the quest to identify habitability outside the Solar System as well as the ongoing discovery of super-Earths in the solar neighborhood with the Gliese 876 \\citep{rive05} and Gliese 581 \\citep{udry07} systems as prime examples.  In agreement with previous studies, it is found that photosynthesis-sustaining habitability strongly depends on the planetary characteristics.  For planets of a distinct size, the most important factor is the relative continental area. Habitability was found most likely for water worlds, i.e., planets with a relatively small continental area.  For planets at a distinct distance from the central star, we identified maximum durations of the transit of the pHZ.  A comparison of planets with different masses revealed that the maximum duration of the transit increases with planetary mass.  Therefore, the upper limit for the duration of the transit for any kind of Earth-type planet is found for most massive super-Earth planets, i.e., 10~$M_\\oplus$, rather than 1~$M_\\oplus$ planets, which are rendered uninhabitable after 6.5~Gyr, as previously pointed out by \\cite{fran00b}.  Our study forwards a thermal evolution model for a 10~$M_\\oplus$ super-Earth orbiting a star akin to the Sun.  The calculations consider updated models of solar evolution obtained by \\cite{schr08} with a detailed mass loss description provided by \\cite{schr05}.  The latter is relevant for the change of luminosity along the Red Giant Branch as well as the increase of the orbital distances of any putative planets during that phase.  By employing the integrated system approach, we were able to identify the sources and sinks of atmospheric carbon dioxide on the planet, allowing us to describe the photosynthesis-sustaining habitable zone (pHZ) determined by the limits of biological productivity on the planetary surface.  Concerning the pHZ, we identified the following properties:  \\smallskip \\noindent (1) Geodynamic solutions are identified for different solar ages, including the RGB phase.  The pHZ increases in width over time and moves outward.  For example, for ages of 11.0, 11.5, 12.0, and 12.1 Gyr, the pHZ is found to extend from 1.41 to 2.60, 1.58 to 2.60, 4.03 to 6.03, and 6.35 to 9.35 AU, respectively.  \\smallskip \\noindent (2) Habitable solutions at large ages, especially for the subgiant and giant phase, are water worlds.  This also means that the possibility of water worlds in principle results in an extension of the outer edge of habitability.  The reason is that planets with a considerable continental area have higher weathering rates that provide the main sink of atmospheric CO$_2$.  Therefore, such planets, contrary to water worlds, are unable to build up CO$_2$-rich atmospheres that prevent the planet from freezing or allowing photosynthesis-based life.  \\smallskip \\noindent (3) The total duration of the transit of the habitable zone is similar to the predictions by \\cite{lope05} based on the conservative limits of the climatic HZ obtained by \\cite{kast93}.  For water worlds with $c=0.1$, the transit times of the pHZ at 2, 3, and 5~AU are obtained as 3.7, 0.25, and 0.10~Gyr, respectively, whereas for planets at 10 and 20~AU, much smaller transit times are found.  \\medskip  Our results are a further motivation to consider super-Earth planets in upcoming or proposed planet search missions such as Kepler, TPF or Darwin. Moreover, our results can also be viewed as a reminder to seriously contemplate the possibility of habitable planets around red giants, as previously pointed out by \\cite{lope05} and others.  For central stars with a higher mass than the Sun, a more rapid evolution will occur that will also affect the temporal and spatial constraints on planetary habitability when the central stars have reached the RGB.  \\noindent {\\bf Acknowledgments}  We would like to thank Norman Sleep and an anonymous referee for their helpful comments which allowed us improving the paper.  \\pagebreak"
    },
    "0812/0812.2139_arXiv.txt": {
        "abstract": "We have carried out a photometric and spectroscopic survey of bright high-amplitude $\\delta$~Scuti (HADS) stars. The aim was to detect binarity and multiperiodicity (or both) in order to explore the possibility of combining binary star astrophysics with stellar oscillations. Here we present the first results for ten, predominantly southern, HADS variables. We detected the orbital motion of RS~Gru with a semi-amplitude of $\\sim$6.5~\\kms\\ and 11.5~days period. The companion is inferred to be a low-mass dwarf star in a close orbit around RS~Gru. We found multiperiodicity in RY~Lep both from photometric and radial velocity data and detected orbital motion in the radial velocities with hints of a possible period of 500--700~days.  The data also revealed that the amplitude of the secondary frequency is variable on the time-scale of a few years, whereas the dominant mode is stable. Radial velocities of AD~CMi revealed cycle-to-cycle variations which might be due to non-radial pulsations. We confirmed the multiperiodic nature of BQ~Ind, while we obtained the first radial velocity curves of ZZ~Mic and BE Lyn. The radial velocity curve and the O--C diagram of CY~Aqr are  consistent with the long-period binary hypothesis. We took new time series photometry on XX~Cyg, DY~Her and DY~Peg, with which we updated their O--C diagrams. ",
        "introduction": "$\\delta$~Scuti stars are short-period pulsating variables of A-F spectral types, located at the intersection of the main sequence and the classical instability strip in the Hertzsprung-Russell diagram. Typical periods are in the order of a few hours with amplitudes less than 1~mag. A prominent group within the family comprises the high-amplitude $\\delta$~Scuti stars (HADS), which have V-band amplitudes larger than 0.3~mag. Population~II members of the group are also known as SX~Phoenicis stars, often found in globular clusters (for a review of $\\delta$~Scuti stars see e.g. \\citet{rod01}).  HADS's are the short-period counterparts of the classical Cepheids, excited by the $\\kappa$-mechanism and pulsating in one or two radial modes, usually in the fundamental and first-overtone modes \\citep{mcn00}. The data in \\citet{rod01} show that there is no rapidly rotating HADS ($v\\sin{i}\\leq40$~\\kms), suggesting an intimate relationship between the rotational state and the excitation of pulsations. In recent years, the number of HADS known with multimode oscillations has rapidly grown, hinting a new potential for asteroseismic studies of these objects \\citep[e.g.][]{por03,por05}. Several investigations suggested that some of the stars may have non-radial pulsation modes present \\citep{mcn00,por05}. However, an important parameter in modeling stellar oscillations is the mass of the star, which is usually constrained from evolutionary models. This will inevitably lead to great uncertainties in any kind of modeling attempts, so that independent mass estimates could be of paramount importance. Thus, binary $\\delta$~Scutis may play a key role in understanding oscillations of these stars.  There has been a great interest recently in $\\delta$~Scuti stars that reside in eclipsing binary systems \\citep[e.g.][]{kim03,mkr06,soy06,pig07,chr07}. Currently, we know about 40 such $\\delta$~Scutis, of which only one belongs to the HADS group \\citep{chr07}. There are also a handful of non-eclipsing binary HADS \\citep{rod01}, usually deduced from the apparent cyclic period changes but with a few exceptions, like the single-lined spectroscopic binary SZ~Lyn. The low number of binary $\\delta$~Scutis suggests a strong observational bias, because one needs accurate observations with a long time span to detect multiplicity unambiguously \\citep{rod01}.  In this paper we present the results of our investigations into binarity and multiple periodicity in bright HADS variables. The sample was initially selected from the variable star catalogue of the Hipparcos satellite, which contains 21 ``SX~Phe'' type stars \\citep{per97b}. Of these, here we discuss 9 stars (and RY~Lep in addition), i.e. our present sample contains almost half of all known bright HADS. Using a wide range of telescopes and instruments, we have been monitoring the target stars over the last five years, extending the earlier studies by our group \\citep{kis95,kis02,der03,sza08}.  The paper is organized as follows. The observations and the data analysis are described in Sect. 2. The main discussion is in Sect. 3, in which the results for individual stars are presented. A brief summary is given in Sect. 4. ",
        "conclusions": "We have carried out multicolor photometry and medium- and high-resolution spectroscopy of ten bright high-amplitude $\\delta$ Scuti stars over 5 years. Our aim was to detect binarity and/or multiperiodicity in HADS variables in order to deepen our knowledge of interaction between oscillations and binarity.  To put our binary targets in a broader context, we have compiled a complete list of binary HADS variables, presented in Table\\ \\ref{hadscmass}. How do RS Gru and RY Lep, the two newly confirmed spectroscopic binaries, compare with other known systems? Looking at the 8 stars in Table\\ \\ref{hadscmass}, we can see three distinct groups with markedly different orbital periods. RS~Gru is one of the shortest-period binary and it is interesting to note that neither UNSW-V-500 nor RS~Gru show evidence of multimode pulsations. RY~Lep is similar to SZ~Lyn both in the orbital period and the reasonably large mass of the companion. These intermediate-period systems are also promising for detecting spectral features of the companion in the ultraviolet or infrared region, thus allowing a full dynamical mass determination. To be able to detect spectroscopically the binary nature of the long-period systems, will require very high-precision spectroscopy, since the expected v$_{\\gamma}$ change is in the range of 1~\\kms.  \\begin{table} \\begin{center} \\caption{\\label{hadscmass} Summary of the estimated masses of the companions in the known HADS binary systems. Sources for $P_{\\rm orb}$ and masses are: (1) \\citet{chr07}, (2) \\citet{mof88}, (3) \\citet{fu08} (4) \\citet{fu99}, (5) \\citet{hur07}, (6) \\citet{fu04}.} \\label{hadscmass} \\begin{tabular}{|lrrc|} \\hline\\hline Star & $P_{\\rm \torb}$ & $m_{\\rm comp} (M_\\odot)$ & Refs. \\\\ \\hline UNSW-V-500 & 5.35~d & $\\sim$0.3 & 1 \\\\ RS Gru & 11.5~d & 0.1--0.2 & present paper \\\\ RY Lep &  730~d: & $\\sim$1.1 & present paper \\\\ SZ Lyn & 3.2~yr & 0.7--1.6 & 2 \\\\ KZ Hya & 26.8~yr & 0.83--3.4 & 3 \\\\ BS Aqr & 31.7~yr & 0.1--0.33 & 4 \\\\ AD CMi & 42.9~yr & 0.15--1.0 & 5 \\\\ CY Aqr & 52.5~yr & 0.1--0.76 & 6 \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table}  To summarize, the main results of this paper are the follows:   \\begin{enumerate}   \\item We monitored RS~Gru spectroscopically on 17 nights in order to measure the orbital period. We derived the orbital period as 11.5~days.  \\item We confirmed the multimode pulsation of RY~Lep from CCD photometry, detecting and refining the frequencies of two independent modes. Spectroscopic measurements also show the multimode pulsation. We detected the orbital motion in the radial velocity curve, confirming the preliminary results of \\citet{lan02} on the binary nature of RY~Lep. Our 700~day-long dataset is in good qualitative agreement with \\citet{lan02}. The limits on the orbital period and RV amplitude suggest a binary companion of about 1~M$_{\\odot}$, possibly a white dwarf star.  \\item The radial velocity curve of AD~CMi shows cycle-to-cycle variations that support the presence of a low frequency mode pulsation reported by \\citet{kho07}. The center-of-mass velocity is in good agreement with the previous measurement by \\citet{bal83} and does not contradict the binary hypothesis of the star, since the predicted $\\gamma$-velocity change is around 1~\\kms.  \\item We obtained the first spectroscopic measurements for BE~Lyn. The RV curve has an amplitude of $\\sim$34~\\kms\\ and the center-of-mass velocity is 3.4~\\kms.  \\item We confirmed the double-mode nature of BQ~Ind, corresponding to the fundamental and first overtone modes.  \\item We detected a low-amplitude secondary period in the photometry of ZZ~Mic but further observations are needed to confirm its validity. The RV curve has a full amplitude of 22~\\kms.  \\item We updated the O--C diagram for CY~Aqr, corroborating binarity found by \\citet{zho99} and \\citet{fu03}. Our radial velocity data are in a good agreement with previous observations by \\citet{str49} and \\citet{fer87} but have better accuracy.  \\item We obtained new time series photometry on XX Cyg, DY Her and DY Peg and updated their O--C diagrams with new times of maximum. DY Her and DY Peg show continuous period decrease, while XX Cyg has continuous period increase.   \\end{enumerate}  \\noindent Further analysis of the photometric and spectroscopic observations (e.g. determination of physical parameters) will be presented in a subsequent paper."
    },
    "0812/0812.0599_arXiv.txt": {
        "abstract": "{Photometric observations for the OGLE-II microlens monitoring campaign have been taken in the period 1997$-$2000. All light curves of this campaign have recently been made public. Our analysis of these data has revealed 13 low-amplitude transiting objects among $\\sim$15700 stars in three Carina fields towards the galactic disk. One of these objects, OGLE2-TR-L9 (P$\\sim$2.5 days), turned out to be an excellent transiting planet candidate.} {In this paper we report on our investigation of the true nature of OGLE2-TR-L9, by re-observing the photometric transit with the aim to determine the transit parameters at high precision, and by spectroscopic observations, to estimate the properties of the host star, and to determine the mass of the transiting object through radial velocity measurements.} {High precision photometric observations have been obtained in $g'$, $r'$, $i'$, and $z'$ band simultaneously, using the new GROND detector, mounted on the MPI/ESO 2.2m telescope at La Silla. Eight epochs of high-dispersion spectroscopic observations were obtained using the fiber-fed FLAMES/UVES Echelle spectrograph, mounted on ESO's Very Large Telescope at Paranal.} {The photometric transit, now more than 7 years after the last OGLE-II observations, was re-discovered only $\\sim$8 minutes from its predicted time. The primary object is a fast rotating F3 star, with $v$sin$i$=39.33$\\pm$0.38 km/s, T=6933$\\pm$58 K, log g = 4.25$\\pm$0.01, and [Fe/H] = $-$0.05$\\pm$0.20. The transiting object is an extrasolar planet with M$_{\\rm{p}}$=4.5$\\pm$1.5 M$_{\\rm{Jup}}$ and R$_{\\rm{p}}$=1.61$\\pm$0.04R$_{\\rm{Jup}}$. Since this is the first planet found to orbit a fast rotating star, the uncertainties in the radial velocity measurements and in the planetary mass are larger than for most other planets discovered to date. The rejection of possible blend scenarios was based on a quantitative analysis of the multi-color photometric data. A stellar blend scenario of an early F star with a faint eclipsing binary system is excluded, due to the combination of 1) the consistency between the spectroscopic parameters of the star and the mean density of the transited object as determined from the photometry, and 2) the excellent agreement between the transit signal as observed at four different wavelengths.}{} ",
        "introduction": "Transiting extrasolar planets allow direct measurements of their fundamental parameters, such as planet mass, radius, and mean density. Furthermore, their atmospheres can be probed through secondary eclipse observations (e.g. Charbonneau et al. 2005; Deming et al. 2005; Knutson et al. 2007), and atmospheric transmission spectroscopy (e.g. Charbonneau et al. 2002; Tinetti et al. 2007; Snellen et al. 2008). This makes transiting exoplanets of great scientific value.  Many photometric monitoring surveys are currently underway. Several of these surveys are very successful, such as the transit campaigns of the Optical Gravitational Lens Experiment (OGLE-III; 7 planets; e.g. Udalski et al. 2008), the Trans-Atlantic Exoplanet Survey (TrES; 4 planets; e.g. Mandushev et al. 2007 ), the Hungarian Automated Telescope Network (HATNet; 9 planets; e.g. Shporer et al. 2008), the XO survey (5 planets; e.g. Burke et al. 2007), and the Wide Area Search for Planets (SuperWASP; 15 planets; e.g. Anderson et al. 2008). In addition, CoRoT is targeting planet transits from space (4 planets; e.g. Aigrain et al. 2008), ultimately aimed at finding Earth or super-Earth size planets.  In this paper we present a new transiting extrasolar planet, OGLE2-TR-L9\\,b. The system has an I-band magnitude of I=13.97, and an orbital period of $\\sim$2.5 days. The host star is the fastest rotating and hottest (main sequence) star around which an orbiting extrasolar planet has been detected to date. The transit system was the prime planetary candidate from a sample of thirteen objects presented by Snellen et al. (2007; S07). They were drawn from the online database of the second phase of the OGLE project, a campaign primarily aimed at finding microlensing events, conducted between 1997 and 2000 (OGLE-II; Szymanski 2005; Udalski et al. 1997). The low amplitude transits were discovered among the light curves of $\\sim$15700 stars, with 13.0$<$I$<$16.0, located in three Carina fields towards the galactic plane. Note that the light curve data have a different cadence than usual for transit surveys, with 1$-$2 photometric points taken per night, totalling 500$-$600 epochs over 4 years. This work shows that such a data set is indeed sensitive to low-amplitude transits, and can yield transiting extrasolar planets.  In section 2 of this paper we present the analysis of new transit photometry of OGLE2-TR-L9, taken with the GROND instrument mounted on the MPI/ESO 2.2m telescope at La Silla. The particular goals of these observations are the re-discovery of the transit, improved transit parameters and orbital ephemeris. In section 3 the spectroscopic observations with the FLAMES/UVES multi-fiber spectrograph on ESO's Very Large Telescope are described, including a description of the analysis of the radial velocity variations, of the bisector span, and the spectroscopic parameters of the host star. In section 4 and 5 the stellar and planetary parameters with their uncertainties are determined, and arguments against a stellar blend scenario laid out. The results are discussed in section 6. ",
        "conclusions": "More than seven years and $>$1000 orbital periods after the last observations of OGLE2-TR-L9, the transit signal was rediscovered only 8 minutes from its predicted time (from S07). It not only shows that an observing campaign with large time intervals between measurements can produce reliable light curves, it also shows it produces extremely accurate orbital periods.  OGLE2-TR-L9b is the first extrasolar planet discovered transiting a fast rotating (v$sin$i=39 km/s) F star. OGLE2-TR-L9 is also the star with the highest surface temperature (T=6933 K) of all main sequence stars that host an exoplanet known to date. It is therefore not surprising that the uncertainties in the radial velocity variations are higher than for most other transiting exoplanets presented in the literature. Only due to the high mass of OGLE2-TR-L9b, we were able to detect its radial velocity signature. Note however that, since a blend scenario can be rejected at high significance, an upper limit to the mass of OGLE2-TR-L9b would have been sufficient to claim the presence of a transiting extrasolar planet, although with an unknown mass. Similar arguments may have to be used in the case of future detection of transits of Earth-size planets from Kepler or CoRoT, since their radial velocity signature may be too small to measure.  OGLE2-TR-L9b has a significantly larger radius than expected for a planet of about 4.5 times the mass of Jupiter, even if it is assumed that 0.5\\% of the incoming stellar luminosity is dissipated at the planet's center (Fressin et al. 2007). However, it is not the only planet found to be too large (e.g. CoRoT-exo-2b, TrES-4b, and XO-3b). Several mechanisms have been proposed to explain these 'bloated' radii, such as more significant core heating and/or orbital tidal heating (see Liu, Burrows \\& Ibgui 2008 for a recent detailed discussion).  The measured v$sin$i and estimated stellar radius combine to a rotation period of the host star of $\\sim$1.97$\\pm$0.04 days. It means that the rotation of the star is not locked to the orbital period of OGLE2-TR-L9b. A v$sin$i of 39 km/sec is within the normal range for stars of this spectral type. The mean v{\\sl sin}i of  F5 to F0 stars in the solar neighbourhood range from 10$^{2}$ to 10$^{3}$ km/sec respectively. Note that the v{\\sl sin}i of OGLE2-TR-L9a is only $\\sim$9\\% of the expected break-up velocity for a star of this mass and radius. Assuming the general Roche model for a rotating star (e.g. Seidov, 2004), the ratio of polar to equatorial radius of OGLE2-TR-L9a will be on the order of, $1-\\frac{1}{2}(v/v_{\\rm{max}})^2\\sim$~0.996. Thus, the rotational flattening of the host star is not expected to significantly influence the transit shape. OGLE2-TR-L9 is expected to exhibit a strong Rossiter-McLaughlin effect. Simulation using a segmented stellar surface predict an amplitude of 230 m/sec."
    },
    "0812/0812.2249_arXiv.txt": {
        "abstract": "Although the MiniBooNE experiment has severely restricted the possible existence of light sterile neutrinos, a few anomalies persist in oscillation data, and the possibility of extra light species contributing as a subdominant hot (or warm) component is still interesting. In many models, this species would be in thermal equilibrium in the early universe and share the same temperature as active neutrinos, but this is not necessarily the case.  In this work, we fit up-to-date cosmological data with an extended $\\Lambda$CDM model, including light relics with a mass typically in the range 0.1--10~eV. We provide, first, some nearly model-independent constraints on their current density and velocity dispersion, and second, some constraints on their mass, assuming that they consist either in early decoupled thermal relics, or in non-resonantly produced sterile neutrinos.  Our results can be used for constraining most particle-physics-motivated models with three active neutrinos and one extra light species. For instance, we find that at the $3\\sigma$ confidence level, a sterile neutrino with mass $m_s = 2$ eV can be accommodated with the data provided that it is thermally distributed with $T_s/T_\\nu^{\\rm id} \\lesssim 0.8$, or non-resonantly produced with $\\Delta N_{\\text{eff}} \\lesssim 0.5$. The bounds become dramatically tighter when the mass increases. For $m_s \\lesssim 0.9$ eV and at the same confidence level, the data is still compatible with a standard thermalized neutrino. ",
        "introduction": "Neutrino oscillation is a well studied phenomenon, confirmed by strong experimental evidences. Most experimental results are well explained with a three-neutrino oscillation model, involving two independent and well-measured square-mass differences: $\\Delta m^2_{sol} = (7.59 \\pm 0.21) \\times 10^{-5}$ eV$^2$ \\cite{Abe:2008ee} and $\\Delta m_{atm} = (2.74^{+0.44}_{-0.26}) \\times 10^{-3}$ eV$^2$ \\cite{Adamson:2007gu}. However, some other experiments have shown some anomalies which do not fit in this hypothesis (LSND \\cite{Aguilar:2001ty}, Gallium experiments \\cite{Abdurashitov:2005tb}, MiniBooNE low energy anomaly \\cite{AguilarArevalo:2007it}). These anomalous results might be due to unknown systematic effects, but all attempts to identify such systematics have failed until now. Otherwise, they could be interpreted as exotic neutrino physics.  In Ref. \\cite{Giunti:2007xv}, the MiniBooNE anomaly was explained through a renormalization of the absolute neutrino flux and a simultaneous disappearance of electron neutrinos oscillating into sterile neutrinos (with $P_{\\nu_e \\to \\nu_e} = 0.64 ^{+0.08}_{-0.07}$). The LSND and Gallium radioactive source experiment \\cite{Hampel:1997fc, Abdurashitov:1998ne, Abdurashitov:2005ax} anomalies have been studied in Ref. \\cite{Giunti:2006bj}, where is it claimed that all these anomalies could be interpreted as an indication of the presence of, at least, one sterile neutrino with rather large mass (few eV's). Ref. \\cite{Acero:2007su} also studied the compatibility of the Gallium results with the Bugey \\cite{Declais:1994su} and Chooz \\cite{Apollonio:2002gd} reactor experimental data, concluding that such a sterile neutrino should have a mass between one and two eV's. Finally, the MiniBooNE collaboration performed global fits of MiniBooNE, LSND, KARMEN2, and Bugey experiments in presence of a fourth sterile neutrino \\cite{AguilarArevalo:2008sk} (assuming no renormalization issue for MiniBooNe unlike Ref.~\\cite{Giunti:2007xv}). When all four experiments are combined, the compatibility between them is found to be very low (4\\%); however, when only three of them are included, the compatibility level is usually reasonable (the largest tension being found between LSND and Bugey). In this analysis, the preferred value of the sterile neutrino is usually smaller than 1eV, but still of possible cosmological relevance (for instance, for all four experiments, the best fit corresponds to $\\Delta m^2 \\sim 0.2 - 0.3$~eV$^2$).  These various developments suggest that it is important to scrutinize cosmological bounds on scenarios with one light sterile neutrino, which could help ruling them out, given that current bounds on the total neutrino mass assuming just three active neutrinos are as low as $\\sum m_{\\nu} < 0.61$eV (using WMAP5, BAO and SN data \\cite{Komatsu:2008hk}). This result cannot be readily applied to the models which we consider here. Indeed, scenarios with extra neutrinos require a specific cosmological analysis, for the simple reason that besides affecting the total neutrino mass, additional neutrinos also increase the abundance of relativistic particles in the early universe.  From the point of view of Cosmology, there have been many works constraining simultaneously the sum of neutrino masses and the contribution to the relativistic energy density component of the Universe, parametrized as the effective number of neutrinos, $N_{\\text{eff}}$ (see for example \\cite{Dolgov:2002wy, Crotty:2004gm, Cuoco:2005qr, Lesgourgues:2006nd, Hannestad:2006mi, Bell:2005dr}). Most of these works assume either that the heaviest neutrino (and hence the most relevant one from the point of view of free-streaming) has a thermal distribution, sharing the same value of temperature as active neutrinos, or that all neutrinos are degenerate in mass. However, the results of Refs.~\\cite{Cirelli:2006kt, Dodelson:2005tp} can also be applied to the case of very light active neutrinos plus one heavier, non-necessarily thermal sterile neutrino, which is the most interesting case for explaining oscillation anomalies. In terms of physical motivations, it is very likely that the light sterile neutrino required by the LSND anomaly acquires a thermal distribution in the early universe, through oscillations with active neutrinos in presence of a large mixing angle \\cite{DiBari:2001ua}. On the other hand, there are some proposals to avoid these contrains (for a list of some scenarios, see \\cite{Abazajian:2002bj}). One of such possibilities is based on a low reheating temperature ($T_R$) Universe \\cite{Giudice:2000dp, Gelmini:2004ah, Gelmini:2008fq, Yaguna:2007wi}, in which, for a sufficiently low $T_R$, the sterile neutrinos could be non-thermal \\cite{Kawasaki:2000en} and its production would be suppressed \\cite{Gelmini:2008fq}, such that usual cosmological bounds are evaded. In fact, in these models, sterile neutrinos are allowed to have a large mass without entering in conflict with other experimental results, while $T_{\\text{dec}} \\lesssim T_R \\lesssim 10$MeV ($T_{\\text{dec}}$ being the temperature of the cosmic plasma at neutrino decoupling).  In absence of thermalization, cosmological bounds on the sterile neutrino mass become potentially weaker. Hence, it is interesting to study the compatibility of recently proposed scenarios with a light sterile neutrino with the most recent cosmological data, keeping in mind the possibility of a non-thermal distribution. The goal of this paper is hence to study the compatibility of cosmological experimental data (WMAP5 plus small-scale CMB data, SDSS LRG data, SNIa data from SNLS and conservative Lya data from VHS) with the hypothesis of a sterile neutrino with the characteristics sketched above, i.e., with a mass roughly of the order of the electron-Volt, and a contribution to $N_{\\text{eff}}$ smaller than one. ",
        "conclusions": "In this work, we studied the compatibility of cosmological experimental data with the hypothesis of a non-thermal sterile neutrino with a mass in the range $0.1-10$ eV (or more), and a contribution to $N_{\\text{eff}}$ smaller than one. We computed Bayesian confidence limits on different sets of parameters, adapted to the case of thermal relics (section \\ref{TH}), of non-resonantly produced sterile neutrinos {\\it \\`a la} Dodelson \\& Widrow (DW, section \\ref{DW}), or of generic parameters leading to nearly model-independent results (section \\ref{GEN}). In each case, we performed a specific parameter extraction from scratch, in order to obtain reliable results assuming flat priors on the displayed parameters. For simplicity, we assumed that the masses of the three active neutrinos are negligible with respect to that of the sterile neutrino.  For a cosmological data set consisting in recent CMB and LSS data, as well as older but very conservative Lyman-$\\alpha$ data, we found the conditional probability e.g. on the mass of a thermal relic given its temperature, or on the mass of a DW neutrino given its density suppression factor, etc. These proabilities are such that if the fourth neutrino is a standard one (with $\\Delta N_{\\rm eff}=1$), it should have a mass $m_s \\lesssim 0.4$ eV ($2\\sigma$ C.L.) or $m_s \\lesssim 0.9$ eV ($3\\sigma$ C.L.).  At the $3\\sigma$ C.L., a mass $m_s = 1$ eV can be accommodated with the data provided that this neutrino is thermally distributed with $T_s/T_\\nu^{\\rm id} \\lesssim 0.97$, or non-resonantly produced with $\\Delta N_{\\text{eff}} \\lesssim 0.9$. The bounds become dramatically tighter when the mass increases.  At the same confidence level, a mass of just $m_s = 2$~eV requires either $T_s/T_\\nu^{\\rm id} \\lesssim 0.8$ or $\\Delta N_{\\text{eff}} \\lesssim 0.5$, while a mass $m_s = 5$~eV requires $T_s/T_\\nu^{\\rm id} \\lesssim 0.6$ or $\\Delta N_{\\text{eff}} \\lesssim 0.2$.  Our bounds can hopefully be used for constraining particle-physics-motivated models with three active and one sterile neutrinos, as those investigated recently in order to explain possible anomalies in neutrino oscillation data. Many of these models can be immediately localized in our figures \\ref{figB} or \\ref{figC}. For sterile neutrinos or other particles which do not fall in the thermal or DW category, a good approximation consists in computing their velocity dispersion and localizing the model in our figure \\ref{figD} \\footnote{However, this approximation could be not so good when the distribution $p^2 f(p)$ of the non-thermal relic peaks near $p=0$, as if part of these relics were actually cold, see \\cite{Boyarsky:2008prep}.}.  Future neutrino oscillation experiments are expected to test the self-consistency of the standard three-neutrino scenario with increasing accuracy. If anomalies and indications for sterile neutrinos tend to persist, it will be particularly useful to perform joint analysis of oscillation and cosmological data, using the lines of this work for the latter part."
    },
    "0812/0812.2763_arXiv.txt": {
        "abstract": "An overview on the present observational status and phenomenological understanding of cosmic rays above 10$^{16}$~\\eV{} is given. Above these energies the cosmic ray flux is expected to be gradually dominated by an extra-galactic component. In order to investigate the nature of this transition, current experimental activities focus on the measurement of the cosmic ray flux and composition at the 'ankle' or 'dip' feature at several~\\EeV.  At the ultra high energy end of the spectrum, the flux suppression above 50~\\EeV{} is now well established by the measurements of HiRes and the Pierre Auger Observatory and we may enter the era of charged particle astronomy. ",
        "introduction": "\\bstctlcite{IEEEexample:BSTcontrol}  The all particle spectrum of cosmic rays is known to follow a power law, $dN/dE\\varpropto E^{-\\gamma}$ over many orders of magnitude. However, at the highest energies, shown in Fig.~\\ref{fig_pdg}, it exhibits three remarkable features.  The {\\itshape knee}~\\cite{kulikov1958}, a steepening of the flux by $\\Delta\\gamma\\approx 0.5$, at a few~\\PeV{} followed by a flattening called the {\\itshape ankle}~\\cite{linsley1963} at several~\\EeV{} and a flux suppression at ultra high energies~\\cite{Greisen:1966jv,*Zatsepin:1966jv}.  The first two features are suspected to be an indication of the end of the galactic cosmic ray spectrum and the transition to an extra-galactic component.  At energies above several hundreds of \\TeV{} the particle fluxes are too low to allow for a direct measurement of the properties of cosmic rays.  Instead, as will be explained in Sec.~\\ref{sec:airshowers}, the analysis of {\\itshape air showers} plays a crucial role to measure their flux and composition. These two observables are essential to study the transition from galactic to extra-galactic cosmic rays and to distinguish between the various models put forward to explain the ankle (see Sec.\\ref{sec:galegal} and~\\ref{sec:compo}).  At the same hour this talk was given, the first proton beams were injected to the Large Hadron Collider~\\cite{cernbeams}, that will eventually be able to accelerate protons up to $7\\cdot10^{12}$~\\eV. The ultra-high energy frontier of physics is however beyond $10^{20}$~\\eV, where several cosmic rays have already been detected\\cite{Linsley:1963km,Bird:1994uy,Matthews:2005ve}.  At these extreme energies, particles are expected to suffer significant energy losses during their propagation to earth. The corresponding flux suppression was predicted over forty years ago~\\cite{Greisen:1966jv,*Zatsepin:1966jv} and it is only now, that experiments gathered enough statistics to study it carefully (see Sec.~\\ref{sec:gzk}).  The astrophysical sources that are able to accelerate particles to such tremendous energies are still unknown.  Their unambiguous identification requires to study the arrival directions of cosmic rays, i.e. to do {\\itshape particle astronomy}.  Whereas the experimental knowledge on cosmic rays made a major leap forward in the current {\\itshape hybrid era}~\\cite{Watson:2008wv}, new projects aim to accumulate more data below 1~\\EeV{} and at ultra high energies.  Moreover, new data on hadronic interactions at man-made particle accelerators are needed to facilitate the interpretation of air shower data~\\cite{Engel:2002id}.  These efforts will be described in Sec.~\\ref{sec:outlook}. \\begin{figure} \\includegraphics[width=\\linewidth]{pics/cr_fig9_eV_07.eps} \\caption[flux]{All particle flux of cosmic rays (\\cite{Amsler:2008zz} and references therein)} \\label{fig_pdg} \\end{figure} ",
        "conclusions": ""
    },
    "0812/0812.0285_arXiv.txt": {
        "abstract": "{Herbig Ae/Be stars (HAeBe) are pre-main sequence objects in the mass range 2~M$_{\\odot} < M_* < $~8~M$_{\\odot}$. Their X-ray properties are uncertain and, as yet, unexplained.}{We want to elucidate the X-ray generating mechanism in HAeBes.}{We present a \\emph{XMM-Newton} observation of the HAeBe HD~163296. We analyse the light curve, the broad band and the grating spectra, fit emission measures and abundances and apply models for accretion and wind shocks.}{We find three temperature components ranging from 0.2~keV to 2.7~keV. The \\ion{O}{vii} He-like triplet indicates a X-ray formation region in a low density environment with a weak UV photon field, i.~e. above the stellar surface. This makes an origin in an accretion shock unlikely, instead we suggest a shock at the base of the jet for the soft component and a coronal origin for the hot component. A mass outflow of $\\dot M_{\\rm shock} \\approx 10^{-10} M_{\\sun}$~yr$^{-1}$ is sufficient to power the soft X-rays.}{ HD~163296 is thought to be single, so this data represent genuine HAeBe X-ray emission. HD~163296 might be prototypical for its class.} ",
        "introduction": "\\label{introduction} The phase between protostar and main-sequence object is a key stage for planet formation and, furthermore, the properties of a star are determined precisely during that phase for the rest of its life. It is therefore important to study objects of this age in order to understand the origin of planetary systems. Herbig Ae/Be stars (HAeBes) are thought to be the predecessors of main-sequence (MS) stars in the mass range 2~M$_{\\odot} < M_* < $~8~M$_{\\odot}$, although the empirically defined object class of HAeBes may in fact encompass a wider range of evolutionary stages such as stars with magnetospheric accretion (Ae) and boundary layer objects (Be). HAeBes are certainly young, they do have clear signatures of surrounding disks and thus must be considered to be the more massive brothers of the better studied classical T Tauri stars (CTTS), which are young ($<10\\;\\mathrm{Myr}$), low mass ($M_*<3M_{\\odot}$), pre-main sequence stars exhibiting strong H$\\alpha$ emission. CTTS are also surrounded by disks, as evidenced by a strong infrared excess, and are actively accreting material.  Strong X-ray emission is a characteristic property of most young stellar objects. A review of the observational situation before \\emph{XMM-Newton} and \\emph{Chandra} for low mass stars is given by \\citet{1999ARA&A..37..363F}. \\citet{1994A&A...292..152Z} were the first to carry out a systematic X-ray survey of HAeBes. Restricting attention to those stars located within 500~pc, 70~\\% of all HAeBes were found to produce detectable X-ray emission with X-ray luminosities ranging between a few $10^{29}$~erg~s$^{-1}$ and $10^{31}$~erg~s$^{-1}$, thus exhibiting a much larger detection rate than found for field A-type stars (10-15~\\%, according to \\citet{2007A&A...475..677S}). Also, \\citet{2004A&A...413..669G} report the detection of a large X-ray flare (with \\emph{XMM-Newton}) from the Herbig Ae star \\object{V892 Tau}, and present strong arguments that the flare in fact originated from the Ae star in the V892 Tau system; note, however, that \\citet{2005A&A...431..307S} show V892~Tau to be a close binary. Until today it is unclear whether the observed X-ray emission originates from the HAeBes themselves or from unresolved companions, which by necessity would have to be low-mass, young, and active stars. \\citet{2004ApJ...614..221S} studied a sample of ten close HAeBes and showed that the X-ray emission probably originates from a magnetically confined plasma, although it remains unclear if this plasma is associated with a companion. In a sample of 17~HAeBes  \\citet{2006A&A...457..223S} find X-ray emission in about 80~\\% of the objects, more than can be reasonably expected from late-type companions, although it is possible that their sample is biased towards known X-ray sources. In a follow-up paper \\citet{BeateHAeBe} detect every object, but again many of them show spectral characteristics compatible with low-mass companions.  X-ray astronomy has been revolutionised by the availability of high resolution grating spectroscopy in the current generation of X-ray satellites. Grating spectra of HAeBes published so far are taken from \\object{AB Aur} \\citep{ABAur} and the spectroscopic binary \\object{HD 104237} \\citep{2008ApJ...687..579T}, where the main component is a Herbig~Ae star with a CTTS companion of spectral type K3. More observations are available from CTTS; their high resolution X-ray spectra typically show a soft component (which may be hidden by large absorption columns) and  unusually low f/i-ratios in the He-like triplets of Ne and O. For CTTS the most promising explanation for these phenomena is a magnetically funnelled infall model. The inner disk is truncated at the corotation radius and ionised material is loaded onto the field, flowing along the magnetic field lines and hitting the stellar surface close to free-fall velocities \\citep[e.g.][]{1994ApJ...429..781S}. In the photosphere an accretion shock develops, which produces the soft X-ray component. Because the plasma has relatively high densities, the f/i-ratios in the \\ion{Ne}{ix} and \\ion{O}{vii} triplets are small. This model has been successfully applied in explaining the X-ray emission from the CTTS \\object{TW Hya} in terms of an accretion shock model plus a hot corona \\citep{acc_model}.  A question unanswered to date is whether the accretion disks surrounding HAeBes lead to similar phenomena as those surrounding CTTS or if other modes of X-ray generation operate in the more massive stars. This could be, e.g., magnetically confined winds  colliding in the equatorial plane as suggested for \\object{IQ Aur} by \\citet{1997A&A...323..121B}, or internal shocks in unstable winds as in the CTTS \\object{DG Tau} \\citep{2005ApJ...626L..53G,2008A&A...478..797G,Schneider,dgtau}.  All these models require magnetic fields, which are only weak in evolved A and B stars, because of the absence of an outer convection zone required for a solar-type $\\alpha-\\Omega-$dynamo. Nevertheless the HAeBes could have magnetic fields, produced by the compression of a primordial field of the proto-stellar cloud. It has been suggested that they are the progenitors of the magnetic Ap/Bp stars \\citep{2005A&A...442L..31W}, which comprise about 5\\% of the total A star population.  In order to enlighten the origin of the X-rays from HAeBes we performed X-ray observations of \\object{HD 163296} with \\emph{XMM-Newton}, whose stellar properties we explain in Sect.~\\ref{stellarproperties} before presenting the observations in Sect.~\\ref{observations}. We show the results of our analysis in Sect.~\\ref{results} and discuss their implications in Sect.~\\ref{discussion}. Sect.~\\ref{conclusion} gives our conclusions. ",
        "conclusions": "\\label{conclusion} We present the observation of HD~163296 with \\emph{XMM-Newton}. This is after AB~Aur the second grating spectrum of a HAeBe, without -as far as we know- an unresolved companion. The signal-to-noise ratio (SNR) is significantly better than in the first observation. Congruently both stars show soft spectra with little temporal variability. In all ana\\-lysed He-like triplets the line ratios point to an emission origin in regions of low densities and weak UV photon field, that is at a few stellar radii. On HD~163296 we find a FIP abundance pattern typical for inactive stars.  Both the spectral shape and data from other instruments make an origin of the X-ray in an unknown companion to HS~163296 very unlikely. Therefore we believe to present genuine X-ray data for HAeBe which can be representative for the whole class of this objects for all we know today.  The optically observed mass flux rates are much larger than required by the energetics of the X-ray radiation in the case of accretion shocks or emission from outflows, but the large $f/i$-ratios rule out an accretion shock scenario on the stellar surface. We favour a model where the main contribution to the soft X-ray emission comes from shock heated plasma in the jet or wind.  Possible geometries include shocks in unsteady winds or collimation shocks on the base of the jet. The collimation of the jets is likely related to a magnetic field, whose origin is unclear. This field can be a frozen-in primordial magnetic field or the star generates one in a thin convective layer. We find no other convincing explanation for the hard emission than heating by energetic events from the magnetic field. Thus we predict that with longer observation times we should observe flare-like behaviour on HAeBes."
    },
    "0812/0812.0766_arXiv.txt": {
        "abstract": "{LS~5039\\,/\\,RX~J1826.2-1450 is one of the few High Mass X-ray binary systems from which radio and high energy TeV emission has been observed. Moreover, variability of the TeV emission with orbital period was detected.} {We investigate the hard X-ray ($25-200$~keV) spectral and timing properties of the source with the monitoring IBIS/ISGRI instrument on-board the \\textsl{INTEGRAL} satellite. } {We present the analysis of \\textsl{INTEGRAL} observations for a total of about 3~Msec exposure time, including both public data and data from the Key Programme. We search for flux and spectral variability related to the orbital phase.} {The source is observed to emit from 25 up to 200~keV and the emission is concentrated around inferior conjunction. Orbital variability in the hard X-ray band is detected and established to be in phase with the orbitally modulated TeV emission observed with H.E.S.S. For this energy range we determine an average flux for the inferior conjunction phase interval of $(3.54 \\pm 2.30) \\times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$, and a flux upper limit for the superior conjunction phase interval of $1.45 \\times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$ ~(90\\% conf. level respectively). The spectrum for the inferior conjunction phase interval follows a power law with an index $\\Gamma = 2.0^{+0.2}_{-0.2} $~(90\\% conf. level).} {} ",
        "introduction": "LS~5039\\,/\\,RX~J1826.2-1450 is a High Mass X-ray Binary (HMXRB) system, which was discovered in the \\emph{ROSAT} all-sky survey and identified with optical follow-up observations by \\citet{motch:97}. The compact companion is estimated to have a mass of $M\\sim3.7M_\\odot$, orbiting a bright, $V=11.2$, O6.5V((f)) star with an orbital period $P_{\\rm orb}\\sim3.9060 \\pm 0.00017$~d, an eccentricity $e=0.35\\pm0.04$ and a distance of $d=2.5\\pm0.1$~kpc \\citep{casares:05}. Based on the observation of persistent asymmetric milliarcsecond radio outflows (extending up to $\\sim$1000 AU), \\citet{marti:98} and \\citet{paredes:00} suggested the presence of a mildly relativistic ($v\\sim0.2~c$) jet. \\citet{paredes:02} and \\citet{ribo:08} have therefore proposed the source to be a microquasar.  LS~5039 has been observed by several X-ray satellites. A summary of all observations can be found in \\citet{bosch-ramon:05a}. Flux variations on timescales of days as well as miniflares on shorter timescales have been observed. RXTE observations, covering the whole orbital period, show variability with a maximum of emission at orbital phase $\\phi=0.8$, interpreted due to periastron passage. Moreover, the power-law shaped X-ray spectrum appears harder in the high flux state \\citep{bosch-ramon:05a}.  Because inhomogeneous jets could emit $\\gamma-$rays via inverse Compton scattering, \\citet{paredes:00,paredes:02} have proposed the source to be the counterpart of the unidentified high energy EGRET source 3EG J1824-1514 \\citep{hartman:99} and also a component of the MeV emitter GRO J1823-12, which is a superposition of three EGRET sources \\citep{collmar:03}.  The source has likewise been discovered to emit at TeV energies with the High Energy Stereoscopic System (H.E.S.S.) \\citep{aharonian:05a}. This demonstrated that X-ray binaries are capable of accelerating particles to TeV energies. The H.E.S.S. measurements have revealed that the flux and energy spectrum of the source are indeed modulated with the $\\sim3.9$~d orbital period of the binary system \\citep{aharonian:06a}. The VHE $\\gamma$-ray emission is largely confined to half of the orbit and peaks around the inferior conjunction (INFC) epoch of the compact companion, where a hardening of the spectrum is observed. According to the authors these findings indicate that $\\gamma-$ray absorption with pair production occurs in the system. Interestingly, orbital variability at TeV $\\gamma-$ray energies has also been discovered in the similar system LS~I+61$^\\circ$303 with MAGIC  \\citep{albert:06} and VERITAS \\citep{acciari:08}.  The nature of the compact object in the system RX~J1826.2-1450\\,/\\,LS~5039 as well as the orgin of the TeV emission is still under debate. A microquasar scenario \\citep[see e.g.][]{paredes:00,bosch-ramon:04,torres:08} as well as a pulsar scenario  \\citep[see e.g][]{dubus:06,sierpowska:07,bosch-ramon:08e} have been discussed. In this Letter, based mainly on archival data from \\textsl{INTEGRAL} which span 4.6 years between March 2003 and October 2007, we confirm the detection of the source reported by \\citet{bird:07} in the energy range from 25~keV up to 200~keV. We report the discovery of the modulation of the hard X-ray flux with orbital phase.  \\begin{figure} \\includegraphics[width=0.33\\textwidth,angle=90.]{11397fg1.ps} \\caption{A sketch of the system. The orbital values are listed in \\citet{casares:05} and the definition of phase intervals is from \\citet{aharonian:06a}.} \\label{fig:sketch} \\end{figure} ",
        "conclusions": "In this paper we report on a $3$~Msec \\textsl{INTEGRAL} monitoring of the HMXRB RX~J1826-1450\\,/\\,LS~5039. The faintness of the source and the fact that it lies in a crowded field required a very careful, non-standard handling of the data.  Our results, obtained by applying different analysis methods, point out that the source significantly emits at hard X-rays ($25-200$~keV) and that this emission varies with the orbit in phase with the very high energy $\\gamma$-rays detected with H.E.S.S. (Fig. \\ref{fig:phasogramintehess}). The spectrum at the inferior conjunction is well described by a power law, while at the superior conjunction the hard X-ray emission, if any, is below the sensitivity of \\textsl{INTEGRAL}. Our result indicates that accretion might not be the mechanism for the production of the hard emission, since, in this case, we would expect a rather sharp flux maximum near periastron.  Moreover, the observed close correlation of hard X-ray and VHE $\\gamma$-ray emission in phase suggests that the emission would originate from the same source region, possibly from particles produced in the same acceleration process.  This casts doubt on the simple explanation that the TeV variability is mainly caused by $\\gamma$-$\\gamma$ absorption, as suggested in \\citet{aharonian:06a}.  We note that, from the correlation alone, more complicated scenarios involving $\\gamma$-$\\gamma$ absorption such as described by \\citet{dubus:08} cannot be ruled out. However the cascades initiated in the $\\gamma$-$\\gamma$ absorption process produce secondary electrons \\citep{bednarek:97} also emitting X-ray synchrotron emission. Therefore, soft and hard X-ray flux levels can be used to constrain the amount of absorption involved. The detection of TeV emission at SUPC together with the upper limit in the hard X-ray domain derived here may pose problems for models invoking only one emission region close to the compact object, as recently discussed in \\citet{bosch-ramon:08e}.  Radio observations interpreted as jets have led to a microquasar interpretation for this source. In this scenario, particles are accelerated to $\\approx$ TeV energies \\citep{perucho:08} due to the interaction between the jet and the stellar wind of the companion. Alternatively, \\citet{dubus:06} has suggested the source to be a pulsar binary system, where VHE emission arises from the pulsar wind material shocked by the interaction with the stellar companion wind and afterwards flowing away in a comet-shape tail. Unfortunately, the observations presented here do not enable us to discriminate between the two scenarios, although the predicted X-ray variability by \\citet{paredes:06} seems to be not in accordance with our observations. Finally, we remark that a hadronic scenario \\citep{romero:03,romero:08} is not ruled out by our findings. A longer and more detailed interpretation of our results is beyond the scope of this Letter, and will be the subject of further investigations."
    },
    "0812/0812.2233_arXiv.txt": {
        "abstract": "Digitized images of the drawings by J.C.~Staudacher were used to determine sunspot positions for the period of 1749--1796. From the entire set of drawings, 6285~sunspot positions were obtained for a total of 999~days. Various methods have been applied to find the orientation of the solar disk which is not given for the vast majority of the drawings by Staudacher. Heliographic latitudes and longitudes in the Carrington rotation frame were determined. The resulting butterfly diagram shows a highly populated equator during the first two cycles (Cycles~0 and~1 in the usual counting since 1749). An intermediate period is Cycle~2, whereas Cycles~3 and~4 show a typical butterfly shape. A tentative explanation may be the transient dominance of a quadrupolar magnetic field during the first two cycles. ",
        "introduction": "\\label{Introduction} Sunspots typically appear at latitudes between $10^\\circ$ and $40^\\circ$ heliographic latitude. At the beginning of the solar cycle, latitudes tend to be relatively high, while the appearance locations of sunspots shift to low latitudes as the cycle goes on. A plot of the spot appearance latitudes versus time leads to the butterfly diagram (Maunder, 1904). The diagram is typically plotted starting in 1874 with the Greenwich drawings of the solar disk. The shape of the appearance latitude as a function of time has not significantly changed since. Regular updates of the butterfly diagram are provided by Hathaway in publications and in the Internet (see e.g. Hathaway {\\it et al.}, 2003).  It is desirable  to extend the butterfly diagram into the past. Especially after the Maunder minimum, the butterfly diagram may tell us about the characteristics of the solar dynamo when it was coming back to normal after an activity lull with very few sunspots seen between 1645 and~1715.  There is a considerable set of drawings of the solar disk made by Johann Staudacher from 1749 to~1796. The total set of 848~drawings were digitized and described by Arlt (2008, hereafter Paper~I). The drawings contain information about 1031~days in the above period, with several days combined in one drawing, and notes about days when nothing was seen (the number has slightly increased since Paper~I, because of additional notes taken into account). For 999~of them, sunspots were plotted. Note that Wolf (1857) was aware of Staudacher drawings and counted the sunspots for his sunspot number time series which is still used today. But the positions have never been determined since.  In this study, we present sunspot positions for 999~of the drawings by Staudacher and present the first butterfly diagram obtained for the 18th~century. Section~\\ref{rotational_direction} evaluates the direction of the rotation in the drawings which appeared to have changed during the entire observing period of 47~years. Section~\\ref{orientation} deals with the various methods to derive the orientation of the solar equator, and Section~\\ref{positions} describes the actual position measurements once the equator is given. Section~\\ref{butterfly} shows the results of the coordinate determinations in form of a butterfly diagram and discusses its features and limitations. Finally, Section~\\ref{discussion} discusses possible implications for the theory of the solar dynamo. ",
        "conclusions": "} Positional measurements of sunspots observed by Johann Staudacher in the period of 1749--1796 are presented. The data provide us with a butterfly diagram of almost five solar cycles covered. The distribution shows a surprising behaviour. The first two cycles (Cycle~0 and~1) exhibit a different behaviour than the one we know from the ``modern'' butterfly diagram as recorded since 1874. The highly populated equator indicates the presence of a dynamo mode which is symmetric with respect to the equator. We do not have the magnetic polarities though. Nevertheless, a polarity change across such a populated equator seems very unlikely.  An interesting option is a differential rotation different from today. The assumptions made for the rotational fits and matches will then be less suitable. It is unlikely that the differential rotation changed dramatically during the 18th century, but it will be an interesting topic for future research to actually derive the rotation profile from the spot drawings.  The results presented here may also be affected by a gradual change of observing skills and knowledge of Staudacher. The precise reproductions of the solar eclipses -- with the first one drawn in 1753 -- indicate, however, that Staudacher projected the image directly on the paper he used for drawing. Such direct copies are likely to be fairly accurate at different stages of knowledge and experience. While there is certainly some development of skill in detecting small sunspots, the positional accuracy should remain fairly constant.  The sunspot cycle has not always shown sunspots as frequent as during the last 250~years. Strongly reduced activity was observed in the Maunder minimum period from about 1645 to 1715. The very few spots were then mostly present on the southern hemisphere (Ribes and Nesme-Ribes, 1993), apart from very few exceptions, until 1711. The cycle near the end of the Maunder minimum starting in about 1713 shows both hemispheres being populated again, with a considerable number of spots below $10^\\circ$ heliographic latitude. Two activity cycles are missing before the results presented here take over. It was recently suggested that the north-south asymmetry may also be relevant for the prolonged Cycle~4 (Zolotova and Ponyavin, 2007).  A number of dynamo models show the appearance of grand minima as a result of nonlinear effects between the magnetic fields and the differential rotation (Tobias, 1997; K\\\"uker, Arlt, and R\\\"udiger, 1999; Bushby, 2006). Recent mean-field models with a meridional circulation controlling the solar cycle have been successful in constructing solar-like activity cycles (Dikpati and Charbonneau, 1999). The butterfly diagram is thought to be a direct consequence of the equator-ward flow at the bottom of the solar convection zone in these models. An interesting fact is the close excitation limits of the dipolar and quadrupolar dynamo modes in these flux-dominated mean-field dynamos (Dikpati and Gilman, 2001). A transient dominance of the quadrupolar (symmetric with respect to the equator) mode can either be explained by the chaotic nature of a deterministic system (Weiss and Tobias, 2001) or by stochastic variations in the turbulence characteristics (Brandenburg and Spiegel, 2008). Such variations have led to distortions of the butterfly diagram very similar to the ones observed by Staudacher in the 18th century. The observations utilized here support the idea that the solar dynamo is modulated by nonlinear interactions between the dipolar mode -- dominant at the present time -- and a mode of quadrupolar symmetry.  Since the inspection and analysis requires a lot of man power, we have derived here only the positions of the individual spots, but have no information on the sizes of individual spots. There is very likely information still hidden in the individual areas, but we have to postpone these measuring efforts to a future project. The sunspot positions of 1749--1796 are available on request from the author.           \\begin{acks} The author is grateful to Jan Meyer for the photographic reproductions of the drawings and to Regina von Berlepsch for the support in the library of the Astrophysikalisches Institut Potsdam. The work also benefited from the support by ISSI, Switzerland, where the utilization of the drawings was discussed in an International Team. \\end{acks}"
    },
    "0812/0812.2972_arXiv.txt": {
        "abstract": "We calculate the cross-correlation function (CCF) between damped Ly-$\\alpha$ systems (DLAs) and Lyman break galaxies (LBGs) using cosmological hydrodynamic simulations at $z=3$.  We compute the CCF with two different methods.  First, we assume that there is one DLA in each dark matter halo if its DLA cross section is non-zero. In our second approach we weight the pair-count by the DLA cross section of each halo, yielding a cross-section-weighted CCF.  We also compute the angular CCF for direct comparison with observations.  Finally, we calculate the auto-correlation functions of LBGs and DLAs, and their bias against the dark matter distribution.  For these different approaches, we consistently find that there is good agreement between our simulations and observational measurements by \\citet{Cooke06a} and \\citet{Ade05f}.  Our results thus confirm that the spatial distribution of LBGs and DLAs can be well described within the framework of the concordance $\\Lam$CDM model. We find that the correlation strengths of LBGs and DLAs are consistent with the actual observations, and in the case of LBGs it is higher than would be predicted by low-mass galaxy merger models. ",
        "introduction": "\\label{section:intro}  According to the cold dark matter (CDM) model of structure formation, the spatial distribution of galaxies can be understood as a result of gravitational instability of density fluctuations in the CDM, and the dark matter halo mass function can be well described by analytic models \\citep{Sheth99}.  More precisely, hierarchical CDM models predict that the massive galaxies at high redshift (hereafter \\highz) are clustered together in high-density regions, while low-mass galaxies tend to be more evenly spaced \\citep{Kaiser84, Bardeen86}. Under the assumption that galaxies are produced from primordial density fluctuations owing to gravitational instability, one can estimate the average mass of galaxy host haloes based on clustering data.  For example, \\citet{Ade98} estimated the typical halo mass of LBGs at $z\\sim 3$ to be $10^{10.8}\\Msun \\lesssim \\Mhalo \\lesssim 10^{12}\\,\\Msun$ from observations of their auto-correlation function (ACF).  Damped Lyman-$\\alpha$ systems (DLAs), defined as quasar absorption systems with column density of $\\NHI > 2 \\times 10^{20}$ atoms cm$^{-2}$ \\citep{Wolfe86}, probe the \\HI\\ gas associated with \\highz\\ galaxies.  Since stars are hardly formed in warm ionized gas and are tightly correlated with cold neutral clouds, the amount of \\HI\\ gas is very important, being the precursor of molecular clouds \\citep{Wolfe03a}.  DLAs dominate the \\HI\\ content of the Universe at $z\\simeq 3$ and contain a sufficient amount of \\HI\\ gas mass to account for a large fraction of the present-day stellar mass \\citep{Storr00}.  The gas kinematics and chemical abundances of DLAs can be measured and are documented in detail.  However, the masses of DLA host haloes (hereafter DLA haloes) remain poorly constrained, because only about $20\\%$ of quasars exhibit DLA absorption per unit redshift \\citep{Nag07a}, and the scattered distribution of DLAs in quasar sight lines precludes the use of DLAs as tracers of dark matter halo mass.  Alternatively, the mass of DLA haloes can be probed by the cross-correlation between DLAs and a galaxy population whose clustering and halo mass are well understood.  \\citet{Cooke06a, Cooke06b} used 211 LBG spectra and 11 DLAs to measure the three dimensional (3-D) LBG ACF and DLA-LBG CCF \\citep[see also][]{Gawiser01, Bouche04, Bouche05}.  Their analysis started by counting the number of LBGs in 3-D cylindrical bins centred on each of 11 DLAs, following the method of \\citet{Ade03}.  They detected a statistically significant result of DLA-LBG CCF, and estimated an average DLA halo mass of $\\avg{M_{\\rm DLA}} \\approx 10^{11.2} \\Msun$, assuming a single galaxy per halo.  On the theoretical side, \\citet{Nag07a} calculated the average DLA halo mass using a series of cosmological hydrodynamic simulations with different box sizes, resolution and feedback strengths. They found a mean DLA halo mass of $\\avg{M_{\\rm DLA}} = 10^{12.4}\\,\\Msun$ with their Q5 run which is somewhat larger than $\\avg{M_{\\rm DLA}} = 10^{11.2}\\,\\Msun$ of \\citet{Cooke06a}. More recent work by \\citet{Pontzen08} showed that the DLA cross-section is predominantly provided by intermediate mass haloes, $10^9 < M_{\\rm vir}/\\Msun < 10^{11}$.  These results motivate us to further examine the distribution of DLAs relative to that of LBGs.  In this paper, we compute the DLA-LBG CCF in cosmological SPH simulations, using the sample of LBGs and DLAs obtained by \\citet{Nag04g, Nag04f}.  We compare our results with the observational results by \\citet{Ade05f} and \\citet{Cooke06a, Cooke06b}.  Our paper is organized as follows.  In Section~\\ref{sec:simulation}, we briefly describe the features of our cosmological SPH simulations used in this paper.  In Section~\\ref{sec:ccf} and Section~\\ref{sec:wccf}, we describe and report the methodology, binning method, and the results for `unweighted' and `weighted' DLA-LBG CCF, respectively.  We then discuss the projected angular CCF for the direct comparison with observational result by \\citet{Cooke06a, Cooke06b} in Section~\\ref{sec:angular}.  The ACFs of LBG-LBG and DLA-DLA are discussed in Subsections~\\ref{sec:LBGauto} and \\ref{sec:DLAauto}, while the bias results are reported in Section~\\ref{sec:bias}.  Finally, we discuss the implications of our work in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  Our study represents a first attempt to calculate the DLA-LBG cross-correlation function at $z=3$ using cosmological SPH simulations.  We calculated the DLA-LBG CCFs in several different approaches: 3-D, angular, unweighted, and $\\sDLA$-weighted.  We also computed the auto-CF of LBGs and DLAs, and the bias against dark matter.  In comparison to the observational data by \\citet{Ade05f, Cooke06a, Cooke06b}, we find good agreement between our simulations and observational measurements.  Our results suggest that the spatial distribution of DLAs and LBGs are strongly correlated.  Let us summarise some of the main conclusions of this work. In the first part of this paper, our results on the 3-D CCF calculated with spherical shells (Table~\\ref{table:ccfit}) are to be compared with the 3-D spherical shell result by Cooke (private communication), $\\ro = 3.39 \\pm 1.2 \\,\\himpc$ and $\\gamma = 1.61 \\pm 0.3$.  Our results are consistent with Cooke's within the error.  The shallow slope of Cooke's above estimate probably owes to the limited sample size in the spherical shell at small distances, as we discussed in Sections~\\ref{sec:ccf} and \\ref{sec:angular}.  In the second part, we have replaced the spherical shell method with the projected approach used in \\citet{Ade03} and \\citet{Cooke06b}, and calculate the best-fitting values given in Table~\\ref{table:acf}. Encouragingly, our results are within the upper and lower limits of the observational measurement by \\citet{Cooke06a,Cooke06b}. We corrected all CFs in this paper with the integral constraint.  Finally, we also analysed the auto-correlation functions of LBGs and DLAs at $z=3$ (Table~\\ref{table:acfitdladla}) found in our simulations. Our results for the best-fitting parameters of the LBG ACF agree well with \\citet{Ade05f}. Our results show that LBGs are more strongly correlated than DLAs, and have higher mean halo mass.  Figure~\\ref{fig:param} summarises the best-fitting power-law parameters for all the correlation functions that we obtained in the earlier sections.  In most cases, the slope $\\gamma$ falls into the range $\\approx 1.5-1.7$ and the variation is not very large. The correlation length $\\ro$ shows a larger variation from $2.5\\,\\himpc$ to $4\\,\\himpc$, depending on the sample and calculation method.  This trend is similar to that seen by \\citet[][Fig.\\,8]{Cooke06b}.  In general, the $\\sDLA$-weighted method gives a larger $\\ro$ than the unweighted method.   \\begin{figure} \\begin{center} \\resizebox{8.2cm}{!}{\\includegraphics{f8.eps}} \\caption{Summary of best-fitting power-law parameters for all correlation functions that we obtained in earlier sections. Long blue, red, and green dashed cross lines are for the LBG ACF, the angular CCF, and the 3-D CCF of \\citet{Cooke06a, Cooke06b}, respectively. The LBG ACF of \\citet{Ade03} is shown in a short blue dashed line and of \\citet{Ade05f} is shown in red.} \\label{fig:param} \\end{center} \\end{figure}  Finally, the LBG bias, derived from the LBG ACF in Section~\\ref{sec:bias}, has led to the upper and lower limits of the LBG dark matter halo mass of $\\log \\avg{\\Mhalo} = 11.53^{+0.22}_{-0.06}$ (see Table~\\ref{table:bias}).  This result is consistent with observational estimates of the LBG halo mass of $\\Mhalo\\sim 10^{12} \\Msun$, \\citep[e.g.,][]{Steidel98, Ade98} and within the limit of $\\Mhalo=10^{11.2} - 10^{11.8}M_{\\odot}$ \\citep{Ade05f}. Similarly, we derived the DLA biases, and obtained the mean DLA halo masses as shown in Table~\\ref{table:bias}. \\citet{Cooke06a}'s measurement showed a DLA galaxy bias of $b_{\\rm DLA}\\sim 2.4$ and an average DLA halo mass of $\\Mhalo \\sim 10^{11.2} \\Msun$.  Our average DLA bias ($\\overline{b}=1.94$ and $\\overline{b}=2.17$ for un-weighted DLA ACF and weighted DLA ACF, respectively) and halo mass estimates ($\\log\\avg{\\Mhalo^{\\rm DLA}}$=10.71 and 11.02 for un-weighted DLA ACF and weighted DLA ACF, respectively) are in good agreement with theirs. We also examined the ratio of bias values defined as $\\alpha\\equiv\\overline{b}_{\\rm CCF}/\\overline{b}_{\\rm ACF}$ \\citep{Bouche04}, and found that our value of $\\alpha = 0.727$ agrees well with the observational estimates.  This again shows that the mean halo mass of DLAs is less than that of the LBGs. The fact that $\\avg{\\Mhalo^{\\rm LBG}}$ is greater than $\\avg{\\Mhalo^{\\rm DLA}}$ is a natural outcome because the LBG sample is limited to the bright star-forming galaxies with $R_{\\rm AB} < 25$ and $M_{\\star} \\simeq 10^{10}-10^{11} \\Msun$, whereas the DLA \\HI\\ gas is present in numerous lower mass haloes below the LBG threshold.  Motivated by our earlier successes of our simulations to reproduce the physical properties of LBGs such as stellar mass, SFR, and colours \\citep{Nag04e}, in this work we examined another observational method, i.e. DLA-LBG CCF, to further check the consistency between observations and simulations. We found good agreement between our results and observations. Furthermore, there are accumulated evidence that suggest high halo masses for LBGs \\citep[e.g.,][]{ Mo96, Ade98,Baugh98, Giavalisco98, Steidel98, Kau99, Mo99, Katz99, Papovich01, Shapley01}. Therefore, the scenario that the majority of LBGs is merger-induced starburst systems associated with low-mass haloes \\citep{Lowenthal97, Sawicki98,Som01,Wea03} does not appear to be a favored model for LBGs.  In our simulations, we estimated the \\HI\\ column densities using a pixel size that is much larger than the typical quasar beam size, which is of the order of parsecs. This may have some impact on our estimates of $\\NHI$ and the corresponding statistics such as the \\HI\\ column density distribution function. For example, if the ISM is clumpy on smaller scales than our pixel size, there could be high-density neutral clouds below our resolution scale that are self-shielded and contain larger amounts of \\HI.  Unfortunately, owing to limitations in computational resources, it is not possible for us at the moment to run such a high-resolution cosmological simulation with the same box size as we have used in this paper.  In future work, we will nevertheless attempt to check the dependence of our $\\NHI$ estimates on numerical resolution, and perform more rigorous resolution tests."
    },
    "0812/0812.0370_arXiv.txt": {
        "abstract": "Control of systematic uncertainties in the use of Type Ia supernovae as standardized distance indicators can be achieved through contrasting subsets of observationally-characterized, like supernovae.  Essentially, like supernovae at different redshifts reveal the cosmology, and differing supernovae at the same redshift reveal systematics, including evolution not already corrected for by the standardization.  Here we examine the strategy for use of empirically defined subsets to minimize the cosmological parameter risk, the quadratic sum of the parameter uncertainty and systematic bias.  We investigate the optimal recognition of subsets within the sample and discuss some issues of observational requirements on accurately measuring subset properties. Neglecting like vs.\\ like comparison (i.e.\\ creating only a single Hubble diagram) can cause cosmological constraints on dark energy to be biased by $1\\sigma$ or degraded by a factor 1.6 for a total drift of 0.02 mag.  Recognition of subsets at the 0.016 mag level (relative differences) erases bias and reduces the degradation to 2\\%. ",
        "introduction": "}  Distance-redshift measurements of Type Ia supernovae (SN) provide direct mapping of the cosmic expansion history.  The peak brightness of most SN have tighter dispersion than any other cosmological object and this can be standardized with a simple light curve amplitude-width relation, first established by \\cite{phillips} in the early 1990s. This allows a SN to be calibrated to 0.15 mag or about 7\\% in distance, and provided the technique to discover the accelerated cosmic expansion in the late 1990s \\cite{perl99,riess98}. See \\cite{leibundgut} for a review as of 2001. For revealing the nature of the physics causing the acceleration, generically called dark energy, SN have continued to play a central role (e.g.\\ \\cite{knop03,riess04,astier05,riess06,essence,davis07,komatsu08,kowalski, rubin}).  The limits to the standardization for SN are not known; a second parameter to further reduce the intrinsic dispersion is actively sought among the SN observables (see, for example, \\cite{aspen}), and more detailed measurements in spectroscopy and a wider range of wavelength bands may turn up new observables and correlations. The uncertainty on cosmological parameters improves as the intrinsic scatter decreases, both more rapidly than linearly as the reduced dispersion further improves color and dust corrections, and less rapidly as measurement uncertainties remain.  Reduction in scatter can also be achieved by characterizing each supernova with a detailed array of measurements, expecting that supernovae with identical observed properties must also have identical intrinsic luminosities. These empirical observations can define subsets of SN.  Note that the converse does not necessarily hold -- SN that differ in some property, e.g.\\ position in the host galaxy or its metallicity, may not diverge in luminosity (see \\cite{howell08100031} for one recent study).  This was referred to as the mapping of subsets (empirical differences) to subclasses (intrinsic luminosity differences) in \\cite{snlow}.  Mere differences in luminosities are not sufficient to affect cosmological parameter estimation, since they will be absorbed into the ``nuisance'' fit parameter for the intrinsic luminosity (which will be impacted).  A further ingredient must be present: population drift, or evolution of the relative fraction of each subclass with redshift.  Note that SN are not per se aware of the Hubble expansion: the explosions and radiation transport take place on scales $10^{-13}$ times smaller than the Hubble length.  So SN should not evolve in a cosmic sense; rather they may be affected by their immediate environment and progenitor conditions.  Since the full diversity of environments from higher redshifts also exists at low redshifts (e.g.\\ stars and galaxies continue to form today), only the proportion of different environments changes, population drift is a more accurate description of possible changes in SN luminosity.  It is important to note here that while subsets of SN have been recognized, and the proportion of some subsets has been seen to change with redshift, {\\it current data show no definite indication that SN luminosity evolves\\/} -- other than is automatically corrected for in using a standard single parameter light curve amplitude-width relation.  That is, we know of no subsets that are subclasses.  This article looks to the future when suites of observations on large samples of SN, more detailed measurements than we have on any individual SN today, {\\it may\\/} show that indeed some subset, defined through those observational characteristics, is a subclass having a different luminosity. The basic method of comparing subsets of like SN -- likes vs.\\ likes, or SN demographics -- was explained clearly in \\cite{coping}, and we follow this approach while extending it to calculating detailed effects on cosmological parameter determination.  Systematics is emphatically the name of the game in accurate science. Understanding the level of control is essential: without an intrinsic floor, SN are only limited by cosmic variance (from the number of SN within a Hubble volume) to 0.003\\% in distance precision.  And of course a biased answer can be worse than an imprecise one.  For the reader wanting a quick conclusion, see Fig.~\\ref{fig:riskfomsimple}. In \\S\\ref{sec:pdf} we establish the formalism of subclass luminosity functions and calculate the effects of population drift on the mean and variance of the full sample luminosity function.  Using this in \\S\\ref{sec:cos}, we identify three distinct impacts on cosmology determination, and show that the bias is a major effect.  In \\S\\ref{sec:risk} we examine the interplay of bias and uncertainty as we investigate strategies for controlling systematics, such as adding fit parameters for observationally recognized subclasses. We address aspects of the observational requirements for identifying subclasses in \\S\\ref{sec:obs}, and conclude in \\S\\ref{sec:concl}. ",
        "conclusions": "}  Without any systematics, Type Ia supernovae would be statistically the most powerful tool for probing the accelerating expansion of the universe. One of the key approaches for controlling systematics is that of likes vs.\\ likes, or supernova demographics, carefully comparing sample properties through a suite of observational characterizations.  The simple concept is that like supernovae at different redshifts accurately reveal the cosmology, while supernovae at the same redshift, differing in essential ways, can define subsets giving clues to reining in systematics.  The issue is not one of evolution, but {\\it uncorrected\\/} evolution and unrecognized evolution.  This article examines techniques for evaluating the cosmological consequences of systematic control or the lack of it, and strategies for implementing such control.  The main pitfall is bias of the cosmological parameters -- this is a bad thing, not just because it degrades the effective dark energy figure of merit calculated in terms of the risk, but because physics is lost.  One may end up with an impressively precise but simply inaccurate conclusion.  Having high observational acuity and using all this information to define robust subsets is the optimal strategy.  We quantify this and demonstrate that this holds even at the price of additional subset parameters in the fit.  Analyzing the data in a single Hubble diagram can lead to biases of order a full statistical sigma and (secondarily) loss of figure of merit by a factor 1.6.  To avoid these consequences, one uses the recognized subsets to add fit parameters; this restores essentially all the cosmological leverage, as long as the subsets are sufficiently well defined by the observations that these subset mean luminosities do not evolve.  The observational requirements to define the subsets is a complex subject but we consider three categories, of separated, overlapping, and continuous, or unrecognized, confused, and unpinpointed, subsets, and quantify some requirements within simplistic models.  We further briefly consider subsets whose luminosities do in fact evolve and speculate that principal component analysis applied to supernovae spectra may prove the best path to robust control.  In the appendix we illustrate how combining separate data sets, especially from different redshift ranges, can act similarly to evolution and have significant deleterious effects.  A supernovae survey designed without the controls enabled by high observational acuity is taking a gamble on the astrophysics and supernova properties being kind.  While we do not yet know exactly what subsets to define, the capability and flexibility to do so, as measured quantitatively along the lines of the simple calculations here, are required to ensure confidence in the cosmological results."
    },
    "0812/0812.2696_arXiv.txt": {
        "abstract": "We show that any massive cosmological relic particle with small self-interactions is a super-fluid today, due to the broadening of its wave packet, and lack of any elastic scattering.  The WIMP dark matter picture is only consistent its mass $M \\gg M_{\\rm Pl}$ in order to maintain classicality.  The dynamics of a super-fluid are given by the excitation spectrum of bound state quasi-particles, rather than the center of mass motion of constituent particles.  If this relic is a fermion with a repulsive interaction mediated by a heavy boson, such as neutrinos interacting via the $Z^0$, the condensate has the same quantum numbers as the vierbein of General Relativity.  Because there exists an enhanced global symmetry $SO(3,1)_{space}\\times SO(3,1)_{spin}$ among the fermion's self-interactions broken only by it's kinetic term, the long wavelength fluctuation around this condensate is a Goldstone graviton.  A gravitational theory exists in the low energy limit of the Standard Model's Electroweak sector below the weak scale, with a strength that is parametrically similar to $G_N$. ",
        "introduction": "In the early universe, relics including photons, neutrinos and dark matter evolve out of thermal equilibrium as their interaction strength becomes small at low temperature in a process known as ``freeze-out''.  This calculation is essentially classical, assuming particles are point-like and using the Boltzmann equation~\\cite{griest_cosmic_1987,srednicki_calculations_1988}.  After freeze-out the number density of particles is fixed, and the temperature just evolves with Hubble expansion.  Their time evolution is given only by the free particle kinetic term.  It is usually assumed that the interaction strength is so weak that it can be neglected and that particles remain localized point particles forever.  The free particle Hamiltonian propagates particles and also broadens their wave packets, described by their uncertainty $\\Delta x$.  This is due to the fact that the localization of particles causes them to not be an eigenstate of the Hamiltonian if they are massive.  There are two limits of interest for the particle uncertainty $\\Delta x$ relative to the number density $n$.  The classical gas limit is $\\Delta x \\ll n^{-1/3}$.  Elastic scattering collisions and the Boltzmann equation describe this system.  The opposite limit, $\\Delta x \\gg n^{-1/3}$ is a quantum liquid.  Because particles have wave function overlap with their neighbors, one must take into account collective effects due to contact interactions.  If there exists an attractive interaction in any partial wave, then the vacuum energy can be lowered by forming bound state quasi-particles.  The system will undergo a phase transition to a super-fluid described by quasi-particles.  If the system contains global symmetries that are broken when the system becomes a super-fluid, then Goldstone bosons will emerge.  As these are massless, their dynamics are extremely important.  The idea of gravity emerging from spinors is not new and fairly obvious, as one can construct a spin-2 particle as the direct product of spinors~\\cite{ohanian_gravitons_1969,kraus_photons_2002}. However no workable theory has been yet constructed.  The first idea of this type is due to Bjorken~\\cite{bjorken_dynamical_1963}, who attempted to formulate the photon and graviton as a composite state.  The most recent attempt and the most successful is due to Hebecker and Wetterich~\\cite{hebecker_spinor_2003,wetterich_gravity_2003}. Their theory can be regarded as a reformulation of gravity in terms of spinors, but they give no dynamics for the spinors which would lead to such a theory.  This line of research was largely killed by the paper of Weinberg and Witten~\\cite{Weinberg:1980kq}, which showed that a spin-2 particle could not couple to a covariant conserved current.  Two ways out of this theorem are to quantize geometry (the approach of string theory), or to abandon diffeomorphism invariance as an exact symmetry.  Sakharov originally suggested that the graviton could be emergent, and in such theories, diffeomorphism invariance can only be approximate~\\cite{Sakharov:1967pk}. ",
        "conclusions": "We have shown that massive cosmological relics are not classical gasses. If they have attractive interactions or are fermions, they instead are a super-fluid.  This implies that WIMP dark matter scenarios are inconsistent: WIMPs cannot both be decoupled and localized for the age of the universe.  Cosmic background neutrinos must exist. They are a super-fluid, and their self-interactions are a gravitational theory.  These dynamics arise in the SM, which is a renormalizable quantum field theory.  We suggest that this may actually be the gravity that we observe."
    },
    "0812/0812.1189_arXiv.txt": {
        "abstract": "{} {To model the interaction of the solar wind with the plasma tail of a comet by means of numerical simulations, taking into account the effects of viscous-like forces.} {A 2D hydrodynamical, two species,  finite difference code has been developed for the solution  of the time dependent continuity, momentum and energy conservation equations, as applied to the problem at hand.} {We compute the evolution of the plasma of cometary origin in the tail as well as the properties of the shocked solar wind plasma around it, as it transfers momentum on its passage by the tail. Velocity, density and temperature profiles across the tail are obtained. Several models with different flow parameters are considered in order to study the relative importance of viscous-like effects and the coupling between species on the flow dynamics. Assuming a Mach number equal to 2 for the incident solar wind  as it flows past the comet's nucleus, the flow exhibits three transitions with location and properties depending on the Reynolds number for each species and on the ratio of the timescale for inter-species coupling to the crossing time of the free flowing solar wind. By comparing our results with the measurements taken $in \\ situ$ by the Giotto spacecraft during its flyby of comet Halley we constrain the flow parameters for both plasmas. } { In the context of our approximations, we find that our model is qualitatively consistent with the $in \\ situ$ measurements as long as the Reynolds number of the solar wind protons and of cometary H$_2$O+ ions is low, less than 100,  suggesting that viscous-like momentum transport processes may play an important role in the interaction of the solar wind and the plasma environment of comets} ",
        "introduction": "The nature of the interaction of the solar wind with the plasma environment of comets as they approach the Sun, has been under investigation since the early days of space physics as a discipline (Biermann 1951, Alfv\\'en 1957, see reviews by Cravens \\& Gombosi 2004, and Ip 2004). The basic elements of the interaction were developed in the 20 years following the work of Biermann (1951), who proposed that the interaction between the solar wind and the comet's plasma is responsible for the observed aberration angle of plasma tails with respect to the Sun-comet radius vector. Based on the inefficiency of Coulomb collisional processes in the coupling of the solar wind and cometary plasmas, Alfven (1957) proposed that the interplanetary magnetic field (IMF) is a fundamental ingredient in the solar wind-comet interaction; being responsible for channelling the cometary ions as it drapes into a magnetic tail. Biermann et al. (1967) suggested that as cometary ions are created and incorporated (picked-up) into the solar wind, the loading of the flow with this additional mass results in a modification of the flow properties as the solar wind approaches a comet; an idea further developed by Wallis (1973) (for a review see Szego et al. 2000). The IMF and mass loading are thus the main dynamical agents generally considered when developing models for the interaction of the solar wind with cometary ionospheres, as well as with other solar system bodies having an ionosphere and without a strong intrinsic magnetic field.  However, as has been pointed out by Perez-de-Tejada et al. (1980)  and Perez-de-Tejada (1989), several features of the flow dynamics in the cometosheath and plasma tail of comets can be attributed to the action of viscous-like forces as the solar wind interacts with cometary plasma. Such interaction processes are believed to be similar to those known to be occurring in other solar system bodies that have an ionosphere and no significant intrinsic magnetic field, particularly Venus and Mars (for a review see Perez-de-Tejada 1995, Perez-de-Tejada 2009 and references therein). {\\em In situ} measurements indicate that, as in Venus and Mars, the solar wind flow in the ionosheath of comet Halley exhibits an intermediate transition, also called the ``mystery transition'', located approximately half-way between the bow shock and the cometopause (Johnstone et al. 1986, Goldstein et al. 1986, Reme 1991, Perez-de-Tejada 1989 and references therein). Below this transition, as we approach the cometopause, the antisunward velocity of the shocked solar wind decreases in a manner consistent with a viscous boundary layer (Perez-de-Tejada 1989). Also indicative of the presence of viscous-like processes is that the temperature of the gas increases, and the density decreases, as we move from the intermediate transition to the cometopause. Taking the distance between the intermediate transition and the cometopause as the thickness of a viscous boundary layer, which depends on the effective Reynolds number of the flow ($R_{\\rm eff}$), Perez-de-Tejada (1989) estimated that $R_{\\rm eff} \\approx 300$ for the solar wind flow in the cometosheath is necessary to reproduce the flow properties measured {\\em in situ} by the Giotto spacecraft on its flyby of comet Halley.  An additional argument  suggesting the importance of viscous-like effects in the dynamics of the flow in the cometosheath and tail regions, follows from the comparison of the magnitude of the terms corresponding to momentum transport due to viscous-like forces and ${\\bf J} \\times {\\bf B}$ forces in the momentum conservation equation. Perez-de-Tejada (1999, 2000) has  argued that downstream from the terminator in the ionosheath of Venus, a scenario analogous to the one considered in this paper, the fact that the flow is superalfvenic, as found from the {\\em in situ} measurements of the Mariner 5 and Venera 10 spacecraft,  suggests that viscous-like forces may dominate over ${\\bf J} \\times {\\bf B}$ forces in the flow dynamics in the boundary layer formed in the interaction of solar wind and ionospheric plasma. If the flow is characterized by a low effective Reynolds number, $R_{\\rm eff}$, this layer extends over a significant portion of the ionosheath of the planet.  In comets, {\\em in situ} measurements obtained during the passage of the ICE spacecraft through the tail of comet Giacobinni-Zinner (Bame et al. 1986, Slavin et al. 1986, Meyer-Vernet 1986, Reme 1991) indicate that along the inbound trajectory (which lies slightly tailward of the comet nucleus) the magnetic field in the so-called transition and sheath regions, is approximately 10 nT, the number density is approximately 10 cm$^{-3}$ and the tailward flow velocity varies from $\\sim$ 400 km/s (near the bow shock) down to 100 km/s. According to Perez-de-Tejada (1999), the ratio of viscous-like to magnetic forces is essentially the square of the alfvenic Mach number, $M_A^2 = V^2 / (B^2/8\\pi\\rho)$. From the data of the ICE spacecraft cited above, we find that $M_A^2$ ranges between 4 and 40 across the cometosheath and hence, viscous-like stresses may dominate over ${\\bf J} \\times {\\bf B}$ forces by a similar amount, or more, throughout the cometosheath region tailward of the nucleus. In the vicinity of the plasma tail, the measurements of the ICE spacecraft (Bame et al. 1986, Slavin et al. 1986) indicate that the midplane density, dominated by cometary ions, reaches values of 200 cm$^{-3}$ at the point where the magnetic field is a maximum 50 nT. With flow speeds of approximately 20 km/s, the square of the alfvenic Mach number reaches a minimum value of 2-3 so that, even in the plasma tail, viscous-like forces are, at least, as important as ${\\bf J} \\times {\\bf B}$ forces following the arguments of Perez-de-Tejada (1999).  The fact that $M_A^2 >> 1$ in the cometosheath means that the magnetic energy density is much smaller than the kinetic energy associated with the inertia of the plasma. This implies that ${\\bf J} \\times {\\bf B}$ forces are not the dominant dynamical factor responsible for the large scale properties of the flow in the region. In fact, one can argue that the formation of a magnetic tail is an indication that in the cometosheath, the large-scale magnetic field does not dominate the dynamics, it is merely carried around by the superalfvenic flow. If the dynamics were controlled by the magnetic forces, field lines would not bend onto a magnetic tail and the direction of the ion tail would not be essentially in the direction of the local solar wind velocity. We believe that the magnetic field does play a crucial role in the momentum transfer between the solar wind and the cometary plasma, but it is the small scale, ``turbulent'' magnetic field component, that mediates the microscopic interaction between charged particles leading to the transfer of momentum that we are modelling as an effectively viscous process.  \\subsection{On the origin of ``viscosity''}  The precise origin of the viscous-like momentum transfer processes invoked in the viscous flow interpretation of the intermediate transition, in the ionosheath of comet Halley and in other ionospheric obstacles to the solar wind, is not yet clear.  Typical properties of solar wind and cometosheath plasma result in a ``normal'' viscosity, as it appears in the Navier-Stokes equations when derived from Boltzmann's equation, that can be considered negligible in the flow dynamics. Using for example properties of the shocked solar wind in the vicinity of the tail measured at comet Giacobini-Zinner, $n_i$ = 10 cm$^{-3}$, $|B| = 10$ nT and $T$ = 3 $\\times$ 10$^5$  K (Bame et al. 1986, Slavin et al. 1986) one calculates the viscosity coefficient resulting from particle interactions according to Spitzer (1962, eqn. 5-55) to be $\\mu \\sim 10^{17}$ g cm$^{-1}$ s$^{-1}$ . This extremely low value most likely represents a lower limit for the viscosity coefficient, since it reflects the ability to transport momentum across field lines in a plasma threaded by a strong, uniform magnetic field. A more appropriate expression for the plasma viscosity coefficient in the conditions of a cometosheath is probably given by the coefficient presented in Cravens et al. (1980), which corresponds to a plasma in a strongly fluctuating magnetic field. Perez-de-Tejada (2005) has calculated the viscosity coefficient for the ionosheath of Venus based on these results. If we use the same procedure to calculate the viscosity coefficient for the solar wind around the tail of a comet (with the conditions measured at Giacobini-Zinner) we find $\\mu \\sim 10^{-11}$ g cm$^{-1}$ s$^{-1}$.  With typical values for  the solar wind velocity and mass density in the cometosheath around the tail of comet Giacobini-Zinner, $V$ = 200 km/s and $\\rho = 1.67 \\times 10^{-23}$ gm cm$^{-3}$ respectively (Bame et al. 1986), and adopting a characteristic lengthscale of 10$^5$ km for  the variation of the flow velocity (roughly the thickness of the sheath region), we find that the corresponding Reynolds number for the flow, based on the ``normal'' viscosity coefficient estimated above, is $Re > 10^5$. This indicates that viscous effects resulting from the collisions between particles in this environment are negligible. Assuming that the Prandtl number is not very different from unity, as argued by Perez-de-Tejada (2005), we can also neglect heat conduction resulting from particle collisions.  However, as in Venus and Mars, strong turbulence has been measured in the ionosheath of comets Halley and Giacobinni-Zinner (Baker et al. 1986, Scarf et al. 1986, Klimov et al. 1986, Tsurutani and Smith 1986) and, as it generally occurs in many fluid dynamics applications, turbulence is characterized (sometimes even defined) by a dramatic increase in the efficiency of transport processes, viscosity included, in the flow (Lesieur 1990).  The likely importance of turbulent viscosity in this scenario is also expected in view of the large value of the Reynolds number estimated above. Also, as discussed by Shapiro et al. (1995) and  Dobe et al. (1999 and references therein) conditions in the ionosheath of Venus and Mars favour the development of plasma instabilities leading to effective wave-particle interactions. If this mechanism operates also in the cometosheath, it may lead, as in these planets, to increased coupling between the solar wind and cometary plasma in a viscous-like manner as suggested by Perez-de-Tejada (1989). In our opinion this justifies a detailed study of the hypothesis of viscous-like effects on the flow dynamics in solar wind-comet interactions. It is the purpose of this paper to begin these investigations.  In this paper we present results of 2D hydrodynamical, numerical simulations of the flow of solar wind and cometary H$_2$O+ ions in the tail and tailward cometosheath of a comet. This is our first attempt to model the interaction of the solar wind with the plasma environment of a comet taking into account viscous-like forces which. We review the estimation of the effective Reynolds number of Perez-de-Tejada (1989), based on the comparison of {\\em in situ} measurements at comet Halley with results from numerical simulations of the viscous-like, compressible flow of the solar wind over a dense, cold and slow velocity gas representing the plasma tail of a comet. We also study the relative importance of viscous-like forces and the coupling between the fast moving protons of the solar wind and the slow H$_2$O+ ions in the tail. We do the latter by comparing models with different values of the effective Reynolds number, the parameter controlling viscous-like effects, and the effective coupling timescale between both species.  The paper is organized as follows. In section 2 we present the formulation of the problem, the basic equations, approximations and parameters. Section 3 presents results of a series of simulations with different model parameters. A comparison of our results with {\\em in situ} measurements at comet Halley is discussed in section 4. Finally, in section 5 we summarize our main results and present our conclusions. ",
        "conclusions": "We have presented results for the numerical simulation of the interaction between the solar wind and the plasma in the tail of a comet, taking into account the effect of viscous-like stresses previously argued to be important by Perez-de-Tejada et al (1980). To our knowledge, this is the first time that viscous-like effects have been incorporated into such studies. Our results indicate the existence of 3 distinct transitions in the flow properties: outermost we find a shock front, innermost we have the cometopause and an intermediate transition which we can identify with the height of the boundary layer characterized by a fast decline in the anti-sunward flow velocity, and the onset of plasma heating and expansion due to viscous-like dissipation. The location of these transitions depends on the flow parameters, namely the effective Reynolds number of the flow for each species, $R_{\\rm eff}^{a,b}$, and the inter-species coupling parameter, $\\nu_o$.  By comparing the flow properties from our numerical simulations to the location of the shock front and intermediate transition, as measured by the Giotto spacecraft as it approached the nucleus of comet Halley, we find that, almost irrespective of the strength of the inter-species coupling, $\\nu_o$; a low value of the effective Reynolds number, approximately $R_{\\rm eff}^{a,b} \\lesssim 20$ for both species, is required to reproduce the measured transition locations. This implies, in the context of our model, that the measured flow properties cannot be explained if one does not take into account the viscous-like forces in the interaction of the solar wind and the plasma tail of a comet. Although the conclusions drawn from this study are strictly applicable only to comet Halley and solar wind conditions at the time the {\\em in situ} measurements were taken, one may speculate that viscous-like processes may be important in the solar wind-comet interaction in general.  It is important to emphasize that, this being the first attempt to include viscous-like forces in the numerical simulation of the interaction of the solar wind with a comet's plasma environment, there are many pending issues still to be addressed that could have potentially important consequences on the details of the solutions obtained under our simplified treatment. First and foremost, the precise forms we are using for the viscous like stress and effective interspecies coupling, may be questioned. As we have argued in the Introduction, plasma properties imply that ``normal'' viscosity is negligible in the region under consideration. Hence, we are invoking an effective viscosity presumably resulting from plasma turbulence and/or wave-particle interactions. However, the precise form of the terms corresponding to viscous-like momentum transfer in the equations of motion (Bousinessq hypothesis) is not formally demonstrated. Also, as we have discussed in the Formulation section of the paper, the interspecies coupling terms we are using can not be strictly derived for a plasma as the one we are modelling. In view of these arguments, one may consider that the work reported in this paper is only an academic exercise of questionable applicability to the problem of solar wind-cometary plasma interaction. In such case, a similar conclusion must be reached in regards to many other studies of fluid dynamics that use similar approaches to modelling effective viscosity and interspecies coupling.  Additional important effects still to be considered are the following: geometrical effects due to the curvature of the ionosphere are required for a more direct, quantitative comparison between {\\em in situ} measurements by the Giotto spacecraft and the results of simulations; the interaction of the charged species with neutral gas ejected from the comet which, especially in the vicinity of the nucleus, is the most abundant species; the effect of the magnetic field on the flow (particularly in the dayside and around the midplane of the near-tail region), 3D effects, incoming flow time dependence, etc. We believe that the further assessment of the relevance of these factors is beyond the present study. They are the subject of work currently in progress and will be reported in future contributions."
    },
    "0812/0812.1971_arXiv.txt": {
        "abstract": "{In the recent papers, we introduced a method utilised to measure the flow field. The method is based on the tracking of supergranular structures. We did not precisely know, whether its results represent the flow field in the photosphere or in some sub-photospheric layers. In this paper, in combination with helioseismic data, we are able to estimate the depths in the solar convection envelope, where the detected large-scale flow field is well represented by the surface measurements. We got a clear answer to question what kind of structures we track in full-disc Dopplergrams. It seems that in the quiet Sun regions the supergranular structures are tracked, while in the regions with the magnetic field the structures of the magnetic field are dominant. This observation seems obvious, because the nature of Doppler structures is different in the magnetic regions and in the quiet Sun. We show that the large-scale flow detected by our method represents the motion of plasma in layers down to $\\sim$10~Mm. The supergranules may therefore be treated as the objects carried by the underlying large-scale velocity field.} ",
        "introduction": "The supergranulation is a convection-like pattern that remains a puzzle since its discovery by \\cite{1954MNRAS.114...17H}. There are many papers published examining the parametric properties of supergranular cells. Typical values of the sizes are around 20--30~Mm, lifetime of 24~hours \\citep[see e.g.][ and references herein]{2004ApJ...616.1242D}. There is no explanation yet to give a clear statement about the origin of supergranules. The classical mechanism invoked to explain the origin of the supergranulation is the latent heat of the recombination of He$^{2+}$ into He$^+$ at roughly 10~Mm below the photosphere \\citep{1985ApJ...288..795G}. In a fluid at rest, such a heat release may trigger an instability, which can turn into motions at scales comparable to the supergranular scale. The issue in the Sun is that the plasma in the convection zone is highly turbulent and the mentioned thermal instability can be suppressed by turbulent motions. There is also a lot of works doubting the existence of the supergranulation as the convection mode. The surface properties of the supergranules may be explained as non-linear interaction between granules triggered by exploding granules \\citep[e.g.][]{2000AA...357.1063R}. The wave-like properties of the supergranulation were also discussed \\citep[e.g.][]{2003Natur.421...43G}.  The time-distance local helioseismology \\citep{1993Natur.362..430D} is a tool that allows to examine the structure of the plasma flows beneath the solar surface. The $p$-modes of the solar oscillations are excited in the upper convection zone and travel through the convective envelope. Plasma motions and variations of the sound-speed change the speed of the wave-packet propagation. From the difference between the theoretical travel-time from one point in the solar photosphere to another one and the theoretically calculated one, one can map the disturbances causing the travel-time deviations in the subphotospherical layers. Initial results of the time-distance helioseismology from SOHO/MDI \\citep{1997SoPh..170...63D} showed that a pattern of the surface flows can persist 2--3~Mm below the surface, improved results gave 8~Mm \\citep[][]{1998ESASP.418..581D}. Further time-distance analyses suggested the depth of the supergranular flow to be $\\sim15$~Mm \\citep{2003ESASP.517..417Z}. In that study, the authors calculated the correlation of the horizontal divergence at different depths with the horizontal divergence just below the surface. They found a positive correlation in the depths of 0--5~Mm. Deeper down it changes the sign. It may mean that the return flow in the supergranules starts to dominate the intercellular flow structure.  The supergranules are well visible on the entire solar disc in Dopplergrams. This is caused by the prevailing horizontality of the internal velocity field within the supergranular cells. If we assume that supergranules are objects carried by the flow field on the larger scale, the visibility of supergranules in the Dopplergrams gives an opportunity to use them to map this underlying large-scale flow field. It was first done by \\cite{1990ApJ...351..309S}. Recently, we developed a method based on the local correlation tracking algorithm \\citep[LCT; ][]{1986ApOpt..25..392N} with a similar utilisation. This method \\citep[for details see][]{2006AA...458..301S} uses the pattern of the supergranulation in the SOHO/MDI Dopplergrams \\citep{1995SoPh..162..129S} to measure the motions on larger scales. The issue of the method is that we measure displacements of the supergranular structures in the series of the processed Dopplergrams and interpret them as the large-scale velocity field in the solar photosphere. However, we cannot establish the range of depths in the solar convection zone, where the large-scale flows are well represented by the surface measurements. With the use of the data provided by the time-distance helioseismology we can determine this unknown parameter. We can also verify the assumption of our analysis: whether the supergranules are subjected to the transport by the velocity field on the larger scales.  \\begin{figure}[!b] \\centering \\resizebox{0.40\\textwidth}{!}{\\rotatebox{0}{\\includegraphics{fig5.pdf}}} \\caption{A sample of acoustic ray paths used for the time-distance helioseismology, shown in a vertical plane. The shadowed regions illustrate ranges of averaging. The vertical and horizontal lines show a grid used for inversion of acoustic travel time data. After \\cite{2006SSRv..124....1K}, who used the same dataset as in this study.} \\label{fig:fig5} \\end{figure}  There is a list of works suggesting that the supergranules do not represent the actual differential rotation or meridional circulation. \\cite{2000SoPh..193..333B} mentioned that rotation rate of the supergranules determined from MDI Dopplergrams is 5\\,\\% larger than the corresponding spectroscopic rate. Some explanations of this phenomenon appeared in the following papers. \\cite{2003Natur.421...43G} suggested the modulation of the convective pattern by travelling waves. Already \\cite{1995ApJ...443..423W} noted that the effective viscous dissipation of Rossby-waves ($r$-modes), which happen in the Sun at 0.932~$R_\\odot$ ($\\sim$50~Mm below the photosphere), could modulate the convection in the near-subphotospherical layer. \\cite{2006ApJ...644..598H}, however, used synthetic data to explain the apparent super-rotation of the supergranulation as the projection effects on the line-of-sight Doppler velocity signal. \\cite{2007AA...466..691M} re-analysed Hathaway's idea and found that the projection effects can explain the most of the supergranular super-rotation, but not all of it. \\cite{2007AA...466..691M} concluded that the supergranules are unreliable in the analysis of the motions on supergranular scales. However, in \\cite{2007SoPh..241...27S} the results showed that the large-scale apparent motion of supergranules measured by the method introduced in \\cite{2006AA...458..301S}, and used also in this study, is in almost perfect agreement with the near-surface mass flow measured by the time-distance helioseismology. It is perhaps a $k$--$\\omega$ filter used for the suppression of the noise, which makes this method more useful than the previous ones based on similar principles.  The large-scale flows are a combination of many types of motions on various spatial scales (such as the differential rotation, meridional circulation, possibly the giant cell convection, and others), which unfortunately cannot be reliably separated in components. These components may vary independently with time and depth. This implies that, in principle, when dealing with the large-scale flows, we cannot obtain unambiguous results that do not allow any alternative interpretation. Nonetheless, the realistic numerical simulation should reveal the combined large-scale plasma flow, properties of which should be comparable with the measurements. Such numerical simulation, which is not present at the time, shall allow to distinguish various components of the detected large-scale flow and to study them separately. The results in this paper therefore provide encouraging large-scale flow properties, which are to be reproduced by models. ",
        "conclusions": "Let's put two separate results obtained in the previous Section together. Using the comparison of the large-scale horizontal flows obtained by our method based on the tracking of the supergranular structures and by the local helioseismology we state that the results of the tracking method represent the horizontal dynamical behaviour in the depths down to $\\sim$10~Mm below the solar surface. We believe that this result may be important for those, who develop numerical models of the large-scale convection.  The supergranulation is coherent within layers, which show correlated large-scale horizontal flows. The interpretation of supergranules is closer to the real convective cells carried by the underlying velocity field. The depth of the loss of coherence in large-scale horizontal flows coincides with the layer, where the supergranules as a strong pattern should be formed.  The results on the vertical structure of the supergranulation we show in this paper, however, cannot be presented as the general results for the supergranulation all over the Sun. We have to keep in mind that these results were obtained only using one suitable dataset in the quiet Sun regions near the very large active region. Much more datasets are needed to describe the behaviour of the supergranulation at different stages of the magnetic activity. Perhaps the helioseismic data expected from the Solar Dynamics Observatory (SDO) will provide some new clues in this topic.  The ray approximation used in the computation of the helioseismological data is often in doubts, suggesting that better (e.g. Born approximation) approach is necessary. It is known \\citep[e.g.][]{2001ApJ...561L.229B} that the ray approximation underestimates the small-scale perturbations. The real perturbations inside the Sun are likely to be stronger than those inferred using the ray theory. Inversions, however, involve complicated averages of observed travel times and are regularised. As the result, the details of the inaccuracies of the ray approximation on ray-based inversions are not yet clear. Some comparative studies \\citep[e.g.][]{2001ApJ...553L.193J} comparing the results of the ray and Born approximation approaches showed that for the large-scale flow they reasonably match. It was argued that the ray approximation could give credible results with fewer computations. The structure of the flow on supergranular scales was tested in \\cite{2001ApJ...557..384Z} with satisfying results.  At this point we cannot exclude that our results instead of bringing the new information in the behaviour of the supergranules rather put limits on the reliability of the time-distance helioseismology. If this interpretation applies, then the time-distance measurements of the large-scale flows are reliable within upper 10~Mm in the quiet Sun regions and within upper 5~Mm in the magnetised areas. Only the improvement of the time-distance helioseismology methods can resolve this ambiguity.  \\begin{ack} M.~\\v{S}, M.~K., and M.~S. were supported by the Grant Agency of Academy of Sciences of the Czech Republic under grant IAA30030808, M.~\\v{S} additionally by ESA-PECS under grant 98030. The Astronomical Institute of ASCR is working on the Research project AV0Z10030501 (Academy of Sciences of CR), the Astronomical Institute of Charles University on the Research program MSM0021620860 (Ministry of Education of CR). SOHO is a project of international cooperation between ESA and NASA.  \\end{ack}  \\newcommand{\\SortNoop}[1]{}"
    },
    "0812/0812.3235_arXiv.txt": {
        "abstract": "{The interaction of microquasar jets with their environment can produce non-thermal radiation as is the case for extragalactic outflows impacting on their surroundings. Observational evidences of jet/medium interactions in galactic microquasars have been collected in the past years, although few theoretical work has been done regarding the resulting non-thermal emission.} {In this work we investigate the non-thermal emission produced in the interaction between microquasar jets and their environment, and the physical conditions for its production.} {We have developed an analytical model based on those successfully applied to extragalactic sources. The jet is taken to be a supersonic and mildly relativistic hydrodynamical outflow. We focus on the jet/shocked medium structure when being in its adiabatic phase, and assume that it grows in a self-similar way. We calculate the fluxes and spectra of the radiation produced via synchrotron, Inverse Compton and relativistic Bremsstrahlung processes by electrons accelerated in strong shocks. A hydrodynamical simulation is also performed to further investigate the jet interaction with the environment and check the physical parameters used in the analytical model.} {For reasonable values of the magnetic field, and using typical values for the external matter density, the non-thermal particles could produce significant amounts of radiation at different wavelengths, although they do not cool mainly via radiative channels but through adiabatic losses. The physical conditions of the analytical jet/medium interaction model are consistent with those found in the hydrodynamical simulation.} {Microquasar jet termination regions could be detectable at radio wavelengths for current instruments sensitive to $\\sim$~arcminute scales. At X-rays, the expected luminosities are moderate, although the emitter is more compact than the radio one. The source may be detectable by XMM-Newton or Chandra, with 1--10~arcseconds of angular resolution. The radiation at gamma-ray energies may be within the detection limits of the next generation of satellite and ground-based instruments.} ",
        "introduction": "\\noindent Microquasars (MQ) are radio emitting X-ray binaries (REXB) that display relativistic ejections (Mirabel \\& Rodr\\'iguez \\cite{Mirabel99}). The binary system is formed by a compact object, a neutron star or a black hole, and a normal (non-degenerate) star. The relativistic ejections may be transient or persistent (Rib\\'o \\cite{Ribo2005}), with duty cycles of jet activity that change from source to source. The steady jets are formed during the so-called Low/Hard state (Fender et al. \\cite{Fender2001}), and are expected to be only mildly relativistic (Gallo et al. \\cite{Gallo2003}). Presently, about 15 MQs are known in the Milky Way (Paredes \\cite{Paredes2005}), although some authors have proposed that many if not all REXBs could be MQs (e.g. Fender \\cite{Fender2004a}). Thus, the real number of MQs in our galaxy could be notably higher.  \\noindent MQs are considered scaled-down versions of distant quasars since they present many of the characteristics of these extragalactic sources. They serve as suitable scenarios for understanding a number of processes, such as mass accretion onto the compact object or jet formation and evolution, in timescales inaccessible in the case of quasars. In analogy with radio-loud quasars, which show jets impacting on the intracluster medium, one may expect to find radiative signatures due to strong shocks in the termination regions of MQ jets as well. Hot spots and double-lobe morphologies are common features of the powerful extragalactic FRII sources (Faranoff \\& Riley \\cite{Faranoff1974}), in which a variety of large-scale non-thermal structures are revealed at radio wavelengths. In the case of MQs, convincing evidence of jet interaction with the interstellar medium (ISM) are not numerous, partially due to the relatively small number of known sources. Nevertheless, hints or evidence of such interactions have been observed, at different wavelengths, in SS~433 (Zealey et~al. \\cite{zealey80}), 1E~1740.7$-$2942 (Mirabel et al. \\cite{mirabel92}), XTE~J1550$-$564 (Corbel et~al. \\cite{corbel02}), Cygnus~X-3 (Heindl et al. \\cite{heindl03}), Cygnus~X-1 (Mart\\'i et al. \\cite{marti96}; Gallo et~al. \\cite{gallo05}), GRS~1915+105 (Kaiser et al. \\cite{Kaiser04}), H1743$-$322 (Corbel et~al. \\cite{corbel05}), LS~I~+61~303 (Paredes et al. \\cite{paredes07}a), and Circinus~X-1 (Tudose et al. \\cite{tudose06}), although only some theoretical work has been done regarding the hydrodynamics of the interaction, or the resulting non-thermal emission (e.g. Aharonian \\& Atoyan \\cite{aharonian98}; Vel\\'azquez \\& Raga \\cite{velazquez00}; Heinz \\& Sunyaev \\cite{heinz02b}; Bosch-Ramon et al. \\cite{bosch05}; Perucho \\& Bosch-Ramon \\cite{Perucho08}).   \\noindent MQ ejections can transport large amounts of kinetic energy and momentum very far from the binary system. The age of the source times the jet kinetic power, $\\tau_{\\rm MQ}\\times Q_{\\rm jet}$, can be as high as $10^{12}$~s~($3\\times 10^{4}$~yr)~$\\times~10^{37}-10^{39}$~erg~s$^{-1} \\sim 10^{49}-10^{51}$~erg (see. e.g., Cygnus~X-1, Gallo et al. \\cite{gallo05}; SS~433, Zealey et al. \\cite{zealey80}). Even for low radiative efficiencies converting this energy into non-thermal emission, MQ jet termination regions could produce significant fluxes if they were at distances of few kpc. It could well be that MQs were associated, via interaction with their surrounding medium, to some of the extended non-thermal radio sources detected in the Galaxy of unknown origin (e.g., Paredes et al. \\cite{Paredes07}b).    \\noindent In this work, we investigate whether a typical MQ fulfills the conditions to be detectable when interacting with the surrounding external gas. We adopt some reasonable assumptions for the MQ power and age, and the density of the ISM, as well as for the non-thermal energy and magnetic field equipartition fraction in the interaction regions. We study the case of a high-mass system in order to see the role of a massive and hot companion. In the case of a low-mass system, the radiation field of the companion star and therefore the Inverse Compton (IC) contribution would be much lower.   \\noindent This paper is organized as follows. In Section~2, we present an analytical interaction model to characterize the three shocked zones; for each region, the main properties of the non-thermal emitters are studied, and the physical conditions and geometry are established. In Sect.~3, we calculate the total emission produced through synchrotron, relativistic Bremsstrahlung, and IC processes for the different set of parameters used in the model.  We carried out hydrodynamical simulations to study in more detail the jet evolution when interacting with the ambient medium. These simulations allow us also to test and enforce the validity of the assumptions of the analytical model. This is done in Sect.~4. The details of the hydrodynamical simulations can be found in Appendix A. Finally, in Sect.~5 we discuss the obtained results, giving detectability predictions at different energy bands. ",
        "conclusions": "\\noindent We have explored whether non-thermal emission can be expected from the bow shock, the cocoon and the reconfinement regions in the interaction of MQ jets with their environment. In the extragalactic framework, non-thermal radiation is supposed to come only from the cocoon as extended radio emission, as well as locally in the hot spots at the jet tips. Although the shell of shocked ambient material plays an important role in the analytical models describing the growing of FR~II galaxies (Blandford et al. \\cite{blandford1974}, Scheuer \\cite{Scheuer1974}), no radio emission is usually assumed to be produced there (see, however, Rudnick et al. \\cite{Rudnick1988}). In our model, the bow shock velocity is still large enough to accelerate particles, so it is worth to compute the expected non-thermal emission also from this region. Our results show a contribution from this zone that is comparable to that of the cocoon and higher than that coming from the reconfinement region. The geometry of the interaction structures, however, makes it difficult to disentangle the emission from the cocoon and the bow shock region since they could appear co-spatial in the plane of the sky. A clear relativistic Bremsstrahlung would favour a shell origin of the emission.   \\noindent The highest radiation output within the studied set of parameters corresponds to the case: $Q_{\\rm jet}=10^{37}$~erg~s$^{-1}$, $t_{\\rm MQ}=10^{5}$~yr and $n_{\\rm ISM}=1$~cm$^{-3}$. In case the emitting source were located at 3 kpc, this would imply a flux density of $\\sim 150$~mJy at 5~GHz. The emitting size would be of a few arcminutes, since the electron cooling timescale is longer than the source lifetime and they can fill the whole cocoon/shell structures. Considering this angular extension and taking a radio telescope beam size of $10''$, radio emission at a level of $\\sim 1$~mJy~beam$^{-1}$ could be expected. At the X-ray band, we find a bolometric flux in the range 1--10~keV of $F_{\\rm 1-10 keV}$ $\\sim 2 \\times 10^{-13}$~erg~s$^{-1}$~cm$^{-2}$. The electrons emitting at X-rays by synchrotron have very short time-scales, and the emitter size cannot be significantly larger than the accelerator itself. Although the X-rays produced in the shell through relativistic Bremsstrahlung are expected to be quite diluted, the X-rays from the cocoon would come from a relatively small region close to the reverse shock, and could be detectable by {\\it XMM-Newton} and {\\it Chandra} at scales of few arcseconds. In the gamma-ray domain, the flux between 100 MeV and 100 GeV is $F_{\\rm 100~MeV<E<100~GeV}\\sim 10^{-14}$~erg~s$^{-1}$~cm$^{-2}$, while the integrated flux above 100~GeV is $F_{\\rm E>100GeV}\\sim 10^{-15}$~erg~s$^{-1}$~cm$^{-2}$. These values are too low to be detectable by current Cherenkov telescopes. Taking into account the rough linearity between $Q_{\\rm jet}$, $n_{\\rm ISM}$, $t_{\\rm MQ}$, $\\chi$ and $d^{-2}$ with the gamma-ray fluxes obtained, sources with higher values of these quantities than the ones used here may render the MQ jet termination regions detectable by present Cherenkov telescopes. For the weakest jets, i.e, lowest ISM densities and youngest sources adopted in our model ($Q_{\\rm jet}=10^{36}$~erg~s$^{-1}$, $t_{\\rm MQ}=10^{4}$~yr and $n_{\\rm ISM}=0.1$~cm$^{-3}$), the fluxes are strongly suppressed. In the radio band, the specific flux is $F_{\\rm 5~GHz}\\sim 0.1$~mJy~beam$^{-1}$; the integrated flux at X-rays in this case is $F_{\\rm 1-10~keV}\\sim 10^{-14}$~erg~s$^{-1}$~cm$^{-2}$, and at gamma-rays $F_{\\rm 100MeV<E<100GeV}\\sim 2\\times 10^{-17}$~erg~s$^{-1}$~cm$^{-2}$ and $F_{\\rm E>100GeV}\\sim$ $2\\times 10^{-16}$~erg~s$^{-1}$~cm$^{-2}$.    \\noindent We remark that the fluxes showed above strongly depend on the non-thermal luminosity fraction $\\chi$. In the present model we use a quite conservative value of $\\chi=0.01$. In the case of a source able to accelerate particles at a higher efficiency at the interaction shock fronts, then the expected non-thermal fluxes would be enhanced by a factor ($\\chi / 0.01$).   \\noindent The non-thermal emission from the interaction regions in MQs presents characteristic features that can be distinguished from emission coming from the central binary system. First of all, the interaction structures are localized at distances up to $\\sim~10^{20}$~cm. The lifetime of electrons radiating at high and very high energies is relatively short thus the emission region should be localized near the accelerator (but far away from the central system). Regarding radio emission, synchrotron cooling times are expected to be larger than the source lifetime, $t_{\\rm MQ}=10^{4}$ - $10^{5}$~yr. The cocoon and bow shock emitting regions would be expected to form a kind of radio nebula around the central system and the flux densities at the level showed above would then come from a quite extended source.   \\noindent The detection of non-thermal emission from the interaction zones would be a proof that efficient acceleration of particles up to relativistic energies is taking place far away from the central binary system. The averaged kinetic power carried away in the jets could be better constrained, eventually showing that it can be much higher than that inferred directly from observations of the inner jet emission alone (Gallo et al. \\cite{gallo05}; Heinz \\cite{heinz06})   \\noindent Despite we focus on the non-thermal emission from the MQ jet termination regions, thermal Bremsstrahlung emission should be expected from the shell. Although the shocks considered here are still adiabatic, a non-negligible fraction of the jet kinetic luminosity of up to a few \\% may be radiated via thermal Bremsstrahlung. For bow-shock velocities of few times $10^7$~cm~s$^{-1}$, the thermal emission would peak at UV-soft X-rays, energy band that is strongly affected by absorption in the ISM. Observations of the thermal radiation in radio and optical from the interaction structures have been used to extract information of the shell physical conditions (e.g. Cygnus X-1, Gallo et al. \\cite{gallo05}; Russell et al. \\cite{Russell07}).   \\noindent The role of accelerated protons in the shocks deserves a few words, since it may be relevant in some specific situations. Given the conditions in the strong shocks we are considering, relativistic protons may reach energies of about 100 TeV; for shell densities $\\sim$~1~cm$^{-3}$, the accelerated protons have lifetimes of about 10$^{15}$~s. To reach the gamma-ray fluxes detectable for the present generation of satellite borne and ground based gamma-ray instruments above $\\sim 100$~GeV, $\\sim 10^{-13}$~erg~cm$^{-2}$~s$^{-1}$, the energy in non-thermal protons stored in the shell should be as high as $\\sim 10^{48}$~erg at few kpc distances. For a source age of $t_{\\rm MQ}=10^{5}$~yr, the required injected power in relativistic protons should be about $\\sim 3\\times 10^{35}$~erg~s$^{-1}$, thus implying that moderate levels of hadronic emission from MQ jet termination regions may be eventually detected from powerful jets, i.e. $Q_{\\rm jet} \\gtrsim 10^{37}$~erg~s$^{-1}$.    \\noindent The reason why some MQs show non-thermal emission from the jet/ISM interaction regions, whereas in other cases such emission remains undetected, is still unclear. In the context of our model, we can study the effects of varying the set of parameters defining the source and their environment, and predict some cases in which the interaction structures may or may not be detectable. First of all, the energy input injected to the medium should be high enough, and there exist strong differences in the jet kinetic power from source to source. In addition, it could be also the case that the density of the surrounding medium is so low that the shell and the cocoon get very large and their radiation too diluted to be detectable (see, for instance, Heinz \\cite{Heinz02}). Moreover, MQ jets could get disrupted at some source age, as it is found in FR~I galaxies. If this happened within times much shorter than the MQ lifetime, the probability to detect a cocoon/shell structure in the MQ surroundings would be smaller (although the detection of some other kind of structures is not discarded). On the other hand, some sources may be too far, or the non-thermal fraction too small, to detect significant emission from the interaction regions.  \\noindent The evolution of the pressure, mass density, the velocities and the Mach number predicted by the analytical model are in good agreement with those found in the hydrodynamical simulations for the shell and the cocoon regions. Otherwise, these simulations show that several conical shocks may be present within the jet as a consequence of pressure adjustments with the surrounding cocoon, instead of the one strong shock adopted in the analytical treatment. Finally, the length and width of the structures in the model and those found through the numerical simulations are also similar, with a constant ratio $R\\sim 3$ in both cases, implying that the physical assumptions used in the analytical treatment are valid to first order.  \\noindent The results of this work show that the surroundings of some MQs could be extended non-thermal emitters from radio to gamma-rays. In addition, from a comparison with observations, the magnetic field and the particle acceleration efficiency in the jet/ISM interaction regions can be constrained, giving an insight on the physics of these interaction structures. It is interesting to note that, although the adopted model is rather simple, it already accounts for cases when the sources should remain undetectable and cases in which radiation could be detected."
    },
    "0812/0812.3017_arXiv.txt": {
        "abstract": "In the framework of the braneworld models, rotating black holes can be described by the Kerr metric with a tidal charge representing the influence of the non-local gravitational (tidal) effects of the bulk space Weyl tensor onto the black hole spacetime. We study the influence of the tidal charge onto profiled spectral lines generated by radiating tori orbiting in vicinity of a rotating black hole. We show that with lowering the negative tidal charge of the black hole, the profiled line becomes to be flatter and wider keeping their standard character with flux stronger at the blue edge of the profiled line. The extension of the line grows with radius falling and inclination angle growing. With growing inclination angle a small hump appears in the profiled lines due to the strong lensing effect of photons coming from regions behind the black hole. For positive tidal charge ($b>0$) and high inclination angles two small humps appear in the profiled lines close to the red and blue edge of the lines due to the strong lensing effect. We can conclude that for all values of $b$, the strongest effect on the profiled lines shape (extension) is caused by the changes of the inclination angle. ",
        "introduction": "String theory and M-theory describing gravity as a truly higher-dimensional interaction becoming effectively 4D at low-enough energies inspired studies of the braneworld models, where the observable universe is a 3-brane (domain wall) to which the standard model (non-gravitational) matter fields are confined, while gravity field enters the extra spatial dimensions the size of which may be much larger than the Planck length scale $l_\\mathrm{P}\\sim 10^{-33}\\, \\mathrm{cm}$ \\cite{Ark-Dim-Dva:1998:}. Gravity can be localized near the brane at low energies even with a non-compact, infinite size extra dimension with the warped spacetime satisfying the 5D Einstein equations with negative cosmological constant \\cite{Ran-Sun:1999:}.  Significant deviations from the Einstein gravity occur at very high energies in vicinity of compact objects (see e.g., \\cite{Maa:2004:,Ger-Maa:2001:,Ali-Gum:2005:}). The high-energy effects produced by the gravitational collapse are disconnected from the outside space by the horizon, but they could have a signature on the brane, influencing properties of black holes \\cite{Maa:2004:}. There are high-energy effects of local character influencing pressure in collapsing matter  The non-local corrections of \"back-reaction`` character arise from the influence of the Weyl curvature of the bulk space on the brane - the matter on the brane induces  Weyl curvature in the bulk which makes influence on the structures on the brane due to the bulk gravitation stresses \\cite{Maa:2004:,Dad-etal:2000:}. The combination of high-energy (local) and bulk stress (non-local) effects alters significantly the matching problem on the brane, as compared to the 4D Einstein gravity; for spherical objects, matching no longer leads to a Schwarzchild exterior in general \\cite{Dad-etal:2000:,Ger-Maa:2001:}.  The Weyl stresses induced by bulk gravitons imply that the matching conditions do not have unique solution on the brane; in fact, knowledge of the 5D Weyl tensor is needed as a minimum condition for uniqueness \\cite{Ger-Maa:2001:}. However, recently no exact 5D solution is known. Assuming spherically symmetric metric induced on the 3-brane the effective gravitational field equations of vacuum type in both the brane and bulk can be solved, giving a Reissner-Nordstr\\\"{o}m static black hole solutions endowed with a \"tidal'' charge parameter $b$ \\cite{Dad-etal:2000:}, instead of the standard electric charge parameter $Q^2$ \\cite{MTW}. The tidal charge reflects the effects of the Weyl curvature of the bulk space, i.e., from the 5D gravitation stresses \\cite{Dad-etal:2000:,Maa:2004:} with bulk gravitation tidal effects giving the name of the charge. Note that the tidal charge can be both positive and negative, and there are some indications that the negative tidal charge should properly represent the ``back-reaction`` effects of the bulk space Weyl tensor on the \\cite{Maa:2004:,Dad-etal:2000:,Sasaki:2000:}  The stationary and axisymmetric solution of the constrained equations describing rotating black holes localized in the Randall-Sundrum braneworld were derived in \\cite{Ali-Gum:2005:}. The solutions are determined by metric tensor of the Kerr-Newman form with a tidal charge describing the 5D correction term generated by the 5D Weyl tensor stresses. The tidal charge has an ``electric'' character and arises due to the 5D gravitational coupling between the brane and the bulk, reflected on the brane through the ``electric'' part of the bulk Weyl tensor \\cite{Ali-Gum:2005:}, in analogy with the spherically symmetric black-hole  case \\cite{Dad-etal:2000:}. When the electromagnetic field is introduced, the non-vacuum solutions of the effective Einstein equations on the brane are much more complex in comparison with the standard Kerr-Newman solutions \\cite{Ali-Gum:2005:}.  Here we consider optical phenomena  in the Kerr-Newman type of solutions describing the braneworld rotating (Kerr) black holes with no electric charge, since in astrophysically relevant situations the electric charge of the black hole must be exactly zero, or very small \\cite{MTW}. Then the results obtained in analyzing the behavior of test particles and photons or test fields around the Kerr-Newman black holes could be used assuming both positive and negative values of the braneworld tidal parameter $b$ (used instead of the charge parameter $Q^2$).  The information on the properties of strong gravitational fields in vicinity of compact objects, namely of black holes, is encoded into optical phenomena of different kind that enable us to make estimates of the black hole parameters, including its tidal charge, when predictions of the theoretical models are confronted with the observed data. From this point of view, the spectral profiles of accretion discs around the black holes in galactic binaries, e.g., in microquasars, are most promising \\cite{McCli-Nar-Sha:2007:}, along with profiled spectral lines in the X-ray flux \\cite{Bao-Stu:1992:,Stu-Bao:1992:,Laor:1991:,Mat-Fab-Ros:1993:,Viergutz:1993:,Rau-Bla:1994:,Fan-Cal-Fel-Cad:1997:,Zak:2003:,Zak-Rep:2006:,Czerny_etal:2007:}. Important information could also be obtained from the quasiperiodic oscillations observed in the X-ray flux of some low-mass black hole binaries of Galactic origin \\cite{Rem:2005:,Rem-McCli:2006:ARASTRA:}, some expected intermediate black hole sources \\cite{Strohmayer:2007:}, or those observed in Galactic nuclei \\cite{Asch:2004:ASTRA:,Asch:2007:}. The most promising orbital resonance model then enables relative exact measurement of the black hole parameters \\cite{Tor-Abr-Klu-Stu:2005:,Tor:2005a:,Tor:2005b:} that should be confronted with the predictions of the optical modelling \\cite{McCli-Nar-Sha:2007:}. In the case of our Galaxy centre black hole Sgr A$^*$, we could be able to measure the detailed optical phenomena, as compared with the other sources, since it is the nearest supermassive black hole with mass estimated to be $\\sim 4\\times 10^6 M_\\odot$ \\cite{Ghe-etal:2005:}, enabling us to measure the \"silhouette`` of the black hole and other subtle GR phenomena in both weak and strong field limits \\cite{Cun-Bar:1973:,SS:a:RAGTime:2007:Proceedings,SSJ:RAGTime:2005:Proceedings}.  Here we present an introductory study on the role of the braneworld tidal charge parameter in the optical phenomena related to profiled spectral lines generated by thin radiating tori in the braneworld Kerr black-hole backgrounds. Of course, for $b>0$, the results hold as well for standard Kerr-Newman spacetimes due to the correspondence $b\\rightarrow Q^2$ with $Q^2$ being the squared charge parameter of the Kerr-Newman black hole. We use here the transfer function method \\cite{Bao-Stu:1992:} which seems to be the most appropriate in our problem. In section 2, the geometry and photon equations of motion are given. In section 3, the geodetic Al circular orbits in the equatorial plane of the Kerr spacetimes with a tidal charge are presented. In section 4, the specific energy flux of photons is given and optical phenomena related to the thin ring of radiating particles are discussed, namely the frequency shift and the focusing. In section 5, the profiled spectral lines are calculated using the transfer function method and assuming isotropically radiating sources with photon energy fixed to $E_0$ in the rest frame of the radiating source. The line profiles are determined in dependence on the viewing angle $\\theta_0$ and the ring radius $r_e$, which is assumed to be located in the inner part of a thin accretion disc. Specially, it is assumed near marginally stable orbit $r_{ms}$ and in the orbit corresponding to the resonance radius $r_{3:2}$ where the ratio of vertical and radial epicyclic frequencies is $\\sim 3/2$, corresponding to the QPOs frequency ratios commonly observed in microquasars \\cite{Tor-Abr-Klu-Stu:2005:,Tor:2005b:}. In Section 6, concluding remarks are presented. ",
        "conclusions": "We have shown that the tidal charge of the braneworld rotating black holes has a general tendency to make profiled lines of the radiating rings in the inner part of an accretion disc to be wider and flatter. We constructed some theoretically grounded quantities based on characteristics of the profiled lines that could be in principle compared with data determining profiled lines observed in microquasars and active galactic nuclei. We have shown that the influence of the inclination angle on the line profiles could be stronger than that of even high values of $b$ (Cf. Fig.\\ref{figure1} and \\ref{figure2}) that could make the situation difficult for estimating the black hole parameters when we have no limits on the inclination angle, but makes the estimates relatively convincing when we can strongly limit the inclination angle, as we can, at least, for some microquasars. The point is that the inclination angle influence gives some specific features on the profiled lines (e.g. humps for very high $\\theta_0>85^\\circ$) that could serve as an additional test of the black hole parameters. It is quite important that one such a hump appears in the case of negative braneworld parameters being located near the blue end of the line, while two humps appear for positive braneworld parameters (or equivalently for the Kerr-Newman spacetimes), one near the red end of the line, the other near the blue end of the spectral line. Clearly, all the phenomena deserve attention and bring interesting, principally new results that could be extended for development of methods enabling estimates of the black hole parameters."
    },
    "0812/0812.4771_arXiv.txt": {
        "abstract": "{A new semi-analytical model that explains  the formation and sizes of the 'great walls' - the largest structures observed in the universe is suggested. Although the basis of the model is the Zel'dovich approximation it has been used in a new way very different from the previous studies. Instead of traditional approach that evaluates the nonlinear density field it has been utilized for identification of the regions in Lagrangian space that after the mapping to real or redshift space (depending on the kind of structure is studied) end up in the regions where shell-crossing occurs. The set of these regions in Lagrangian space form the progenitor of the structure and after the mapping it determines the pattern of the structure in real or redshift space. The particle trajectories have crossed in such regions and the mapping is no longer unique there. The progenitor after mapping makes only one stream in the multi-stream flow regions therefore it does not comprise all the mass. Nevertheless, it approximately retains the shape of the structure. The progenitor of the structure in real space is determined by the linear density field along with two non-Gaussian fields derived from the initial potential.  Its shape in Eulerian space is also affected by the displacement field. The progenitor of the structure in redshift space also depends on these fields but in addition it is strongly affected by  two anisotropic fields that determine the pattern of great walls as well as their huge sizes.  All the fields used in the mappings are derived from the linear potential smoothed at the current scale of nonlinearity which is $R_{nl} = 2.7$ {\\hmpc} for the adopted parameters of the \\lcdm universe normalized to $\\sigma_8 = 0.8$.  The model predicts the existence of walls with sizes significantly greater than 500 {\\hmpc} that may be found in sufficiently large redshift surveys. } ",
        "introduction": "\\label{sec:intro} The first indications of the existence of superclusters of galaxies bridging the clusters of galaxies \\cite{gre-tho78,chi-roo79} was confirmed by the CfA redshift survey  \\cite{del-etal86}. In the following two decades the discovery of huge concentrations of galaxies spanning over 200 {\\hmpc}  were reported   ({\\eg} \\cite{gel-huc89,got-etal05}). They were dubbed  \"great walls\" due to their sizes and geometry in redshift space. The current record belongs to \"a Sloan Great Wall of galaxies 1.37 billion light years long, 80\\% longer than the Great Wall discovered by Geller and Huchra and therefore the largest observed structure in the universe\" \\cite{got-etal05}.  Taking into account that the SDSS redshift survey is much larger and deeper  than CfA one may wonder whether the Sloan Great Wall will remain the greatest wall or even greater walls will be found in the future redshift surveys and if it is so how large it may be. The answer to this question may depend to certain extent on the exact definition of the walls. For instance, some definitions may result in the percolating system of filaments and walls that would span throughout the whole volume. However even in this case the essential constituents  can be probably identified and  measured. We show that by restricting the analysis to the issues of overall geometry and scale of the largest walls one can alleviate some of the problems of this kind and make some progress. In particular, the details of the density distribution within the walls, such as the accurate positions and masses of halos become less important if the scales of interest exceeds $\\sim$10 {\\hmpc}.  We begin with a brief  discussion of the structure in real space. The scale separating the nonlinear regime of the gravitational growth from  linear or mildly nonlinear regimes  is roughly  around 5  {\\hmpc} which is close to the galaxy correlation scale. It seems  to be by far too small to be directly relevant to the scale of filaments some of them are a least 70 -- 100 {\\hmpc} long. Bond, Kofman and Pogosyan \\cite{bon-etal96}  (hereafter BKP) suggested that the filamentary network in real space with a typical scale of $\\sim$30 {\\hmpc} can be explained  by invoking the correlation bridges between relatively high and therefore rare peaks in the linear density field filtered on scales greater than $R_{\\rm b}$ such that $\\sigma_{\\rho}(R_{\\rm b}) \\lesssim 1$. Unfortunately,  the BKP model has not provided a framework for a quantitative evaluation  of the density contrast in the filaments bridging the clusters of galaxies from the enhancement of the conditional correlation function between two or more peaks. Besides, the gravitational growth of the structure in the universe is a deterministic process if the linear perturbation field is specified. Therefore, one may seek a deterministic  relation between the initial field and the final structure at least in  cosmological simulations where the full information is available. In addition, the cluster analysis of the nonlinear dark matter density field in real space obtained by N-body simulation in the \\lcdm model   \\cite{virgo98} revealed the filaments with lengths up to $\\sim$100 {\\hmpc} \\cite{sh-she-sah04} {\\ie} significantly greater than 30 {\\hmpc}. A similar analysis of the mock galaxy catalogs  \\cite{col-etal98} demonstrated the presence of even longer filaments in redshift space, up to $\\sim$ 150 {\\hmpc}  \\cite{she04}. Both numbers are almost certainly  affected by the simulation box sizes (240 {\\hmpc} in the former and 346 {\\hmpc} in the latter case) and could be even longer in larger boxes. The simulated dark matter density fields and especially mock galaxy catalogs look similar to the observed distribution of galaxies as well as they have similar statistics of various kinds that makes the model viable. However, the emergence of the nonlinear structures (great walls) $\\sim$ 50 -- 100 times greater than the scale of nonlinearity remans unexplained even in the cosmological N-body simulations where the full information about the both initial conditions and formed structure is available.  It has been realized for a long time that the structures in real space and observed in redshift space have different pdfs, power spectra, correlation functions, and higher order moments ({\\eg} \\cite{pee80, kai87,hiv-etal95,sza-etal98,hui-kof-sh00,pea-etal01}). The redshift-space correlation function has an oblate shape on scales greater than a few {\\hmpc} indicating the presence of anisotropic structures in redshift space flattened in the radial direction. However, the measurements do not go beyond $\\sim$ 30 {\\hmpc} where the correlation function drops to $\\sim 0.1$ and the signal drowns in  noise \\cite{pea-etal01}. Theory has not made any specific predictions concerning  the structure in the universe on scales greater than 100 {\\hmpc} therefore the conventional wisdom is that  the perturbations on such scales must be in the linear regime.  The 'finger of God' is the most conspicuous and best understood effect in redshift space. However, it is a relatively small scale effect and cannot explain the coherence of the great walls over a few hundred {\\hmpc}. It has been also demonstrated in \\cite{pra-etal97,mel-etal98} that the size of the structures in the redshift space is considerably greater than that of the parent structures in real space.  In addition, the authors emphasized a characteristic circular pattern in redshift space and the increased spacing between the structures in the redshift direction.  However, they neither  mentioned the enlargement of sizes of the structures in the transverse direction nor provided explanation to their observations.  The presence of walls up to 150 {\\hmpc} in length in redshift space in three-dimensional simulations based on the adhesion approximation was emphasized  in \\cite{wei-gun90}. The authors pointed out that the correlation function was weak or even negative at these scales. They also observed that $\\sim$ 100 {\\hmpc} walls grow from several favorably aligned high-density peaks in the linear density field, each coherent over $\\sim$ 20 {\\hmpc}. This seems to be a factual description of the process, however similarly to the BKP model it requires a dynamical model different from the linear theory of gravitational instability. The major problem with the structures spanning over $\\sim$ 300 {\\hmpc} arises from a lack of a natural scale of this magnitude in the linear density field. If the field is smoothed with say 100 {\\hmpc} or greater scale than its amplitude becomes so low that such a field would remain almost perfectly linear at the present time.  We view the formation of the structure as a continuous mapping from the initial practically uniform state in Lagrangian space to the final highly inhomogeneous and anisotropic distribution of mass on scales up to $\\sim$ 100 {\\hmpc} in Eulerian space followed by another mapping to redshift space where the largest structures become greater than 300 {\\hmpc}. These mappings involve the transport of mass (although not physical in the case of the mapping to redshift space), however the rms displacement of mass elements is only about 15 {\\hmpc} and therefore it cannot produce the structures of needed sizes by itself. Thus, the certain initial fields {\\ie} the fields present in the linear stage and playing an important role in the nonlinear dynamics must have scales in the range of 100 {\\hmpc} or even greater. The major goal of this paper is to show that the fields having needed properties exist in Lagrangian space. The outcome of their combined effect is the definition of two parent structures that are the progenitors of the large-scale structures: one in real and the other in redshift space. Being closely related two progenitors are still quite distinct. The properties of the progenitors are determined by a  combination of Gaussian and non-Gaussian fields derived from the initial gravitational potential smoothed with a Gaussian filter at the scale of nonlinearity $R_{\\rm nl}$ defined by the equation  $<\\delta^2_{\\rm lin}(R_{\\rm nl})> =1$, where $\\delta \\equiv \\delta \\rho/\\bar{\\rho}$. In  the standard \\lcdm model normalized to $\\sigma_8 \\approx 0.8$ this scale is $R_{\\rm nl}$ = 2.7 {\\hmpc}. The choice of the filter and filtering scale for 2D illustrations approximately corresponds to an optimal choice of the filter and scale in 3D simulations that studied the accuracy of the Truncated Zel'dovich Approximation (TZA) in a set of power law models in the Einstinen-de Sitter model \\cite{mel-pel-sh94}. Some adjustment will be probably needed for the \\lcdm model in 3D. Based on that study I would conservatively guess that the change in $R_{\\rm nl}$ would not be greater than 30\\%. No fields smoothed at greater scales are involved in the mapping. The structure resulting from the mapping of the progenitors cannot and is not supposed to reproduce accurately the mass or galaxy number densities, however the overall shape and sizes of the structure in real space and great walls in redshift space must be depicted quite well.  We use TZA as a mapping device for both mapping \\cite{col-mel-sh93,mel-pel-sh94,mel-sh-wei94,hui-kof-sh00}. In this paper we will concentrate on the lengths  of the progenitors in the Lagrangian space and their transformation caused by the mapping to Eulerian and redshift space.  We illustrate the model by two-dimensional plots. Although no two-dimensional model can reproduce all the properties of a three-dimensional system it may provide a reasonably good description of major ideas. In order to make the appearance of illustrations more realistic we use the linear power spectrum $P^{2D}(k) = k\\,P^{3D}(k)$, where $P^{3D}(k)$ is the power spectrum of the linear density perturbations in the \\lcdm model with parameters  $h=0.7$, $\\Omega_m=0.3$, $\\Omega_b = 0.047$, $n=1$. We use the BBKS transfer function \\cite{bar-etal86} with the $\\Gamma$-parameter given by \\cite{sug95}. This choice of  two-dimensional power spectrum retains all the moments $<k^j>$ of the three-dimensional spectrum and therefore all the scales of the two-dimensional field determined by the ratio of two moments ({\\eg}  the scale of peaks $R_*^2 \\sim <k^2>/<k^4>$) remain similar to that of  the three-dimensional field except a small factor $\\sqrt{3/2}=1.2$. The two-dimensional fields were generated in 1024 {\\hmpc} box on 1024$^2$ mesh.  The rest of the paper is organized as follows. Section 2 outlines the dynamical model and fields involved in both mappings. Section 3 provides two dimensional illustrations that elucidate the main concepts of the model, the axes in the figures showing various fields are given in {\\hmpc}. Section 4 discuss the differences between 2D and 3D cases, Sec. 5 summarizes the results and outline the further problems, finally the appendix gives the joint density probability function of three invariants of the deformation tensor that are used in the mappings. ",
        "conclusions": "We propose a new theoretical model that describes the formation of the large-scale structure in redshift space. It naturally explains the sizes of the observed great walls and predicts that even greater walls will be discovered in future deeper and larger redshift surveys. Although based on the Zel'dovich approximation \\cite{z70} the suggested model uses it in a new and very different way. Instead of traditional approaches that focused on the evaluation of density or positions and velocities of the particles we introduce the notion of the progenitor of the large-scale structure.  It is defined as a set of regions in Lagrangian space that after mapping to Eulerian or redshift space (depending on the kind of mapping we study) end up in the regions where the particle trajectories have crossed and the mapping is no longer unique. Neglecting the thermal velocities and assuming smooth initial condition the boundaries of these regions in Eulerian or redshift space are caustics. In reality due to hierarchical clustering process the caustics do not form but the density of both matter and galaxies is still highest in these regions ({eg} \\cite{lit-wei-par91}). Even in the case of smooth initial conditions the evaluation of the density within the shell-crossing regions is not easy and we do not attempt to estimate it in this paper. Instead we utilize the Zel'dovich approximation for identification of the progenitor of the structure in Lagrangian space and then for  mapping it to Eulerian or redshift space. In particular we  demonstrate the role of various fields some of which are non-Gaussian or/and anisotropic. This is the most important factor in explanation of the lengths of filaments in Eulerian space. Non-Gaussian $I_2(\\vecq)$ and $I_3(\\vecq)$ fields fill the gaps between peeks of Gaussian $\\delta(\\vecq) = I_1(\\vecq)$ field significantly boosting the length of the progenitor of the filaments in Lagrangian space. It is worth stressing that the mapping of the progenitor to Eulerian space does not increase its lengths noticeably although it significantly affects its appearance.  The mapping to redshift space has  both the progenitor in Lagrangian space and the structure in redshift space significantly different than the mapping to Eulerian space. In both cases the thickness of the filaments in Eulerian space and walls in redshift space is probably the least accurate in this approximation. The thickness of filaments in Eulerian space is probably exaggerated \\cite{sh-z89}. The thickness of the walls in redshift space may be affected by nonlinear effects that change the velocities in the shell-crossing regions. However, the estimate of the extent of the walls in the transverse direction is expected to be quite robust.  Although both the suggested model and one suggested by \\cite{bon-etal96} (BKP) are based on the Zel'dovich approximation (eq. \\ref{eq:zappr}) there are a few radical differences  between them.  The current model can explain the sizes of both the structures in Eulerian and redshift space emphasizing that the structure in redshift space must have significantly greater sizes in transverse directions than in the line of sight direction and therefore the wall kind structures must be more conspicuous in redshift space. The BKP model discusses only the structure in Eulerian space. The BKP model uses the smoothing scale $R_b$ defined by an inequality $\\sigma_{\\delta}(R_b ) \\lesssim 1$ such that it corresponds to \"the conceptual picture of a smooth stately large-scale single-stream flow\" \\cite{bon-etal96} while the current model uses the scale of nonlinearity and therefore assumes the crossing of trajectories and presence of multi-stream flows. Although it is not excluded that the further comparison with cosmological N-body simulations may slightly change the scale for better fit to N-body simulations it will not change the basic construction used by the model - the progenitor of the structure.  The model defines the progenitor of the structure  as a set of regions in Lagrangian space where the Jacobian of the mapping becomes negative indicating that these regions end up in multi-stream flow regions after mapping to Eulerian or redshift space. There is also one more principal difference. The BKP model uses only linear {\\ie} Gaussian fields and conditional correlations between them. This model utilizes both  non-Gaussian fields ($I_2, I_3, M_{qq}$) one of which is anisotropic as well as anisotropic Gaussian fields $d_{qq}, s_q$ along with linear density contrast field. The displacement fields $s_i$ and $s_q$  also play important roles in the appearances  of the structure in Eulerian and redshift spaces. The suggested approach allows easily to isolate the role of every field and thus helps to better understand the complex nonlinear dynamics of the structure formation.  Further studies including the more detailed studies of the statistical properties of the fields involved in the mappings as well as a direct comparison with cosmological N-body simulations are obviously needed. Some work has already started and the results will be reported in the following papers."
    },
    "0812/0812.0364_arXiv.txt": {
        "abstract": "We present an analysis of 23 L dwarfs whose optical spectra display unusual features. Twenty-one were uncovered during our search for nearby, late-type objects using the Two Micron All-Sky Survey while two were identified in the literature. The unusual spectral features, notably weak FeH molecular absorption and weak \\ion{Na}{1} and \\ion{K}{1} doublets, are attributable to low-gravity and indicate that these L dwarfs are young, low-mass brown dwarfs. We use these data to expand the spectral classification scheme for L0 to L5-type dwarfs to include three gravity classes. Most of the low-gravity L dwarfs have southerly declinations and distance estimates within 60~pc. Their implied youth, on-sky distribution, and distances suggest that they are members of nearby, intermediate-age ($\\sim$10--100 Myr), loose associations such as the $\\beta$~Pictoris moving group, the Tucana/Horologium association, and the AB Doradus moving group. At an age of 30~Myr and with effective temperatures from 1500 to 2400~K, evolutionary models predict masses of 11--30~$M_{Jupiter}$ for these objects. One object, 2M~0355+11, with $J-K_s=2.52\\pm0.03$, is the reddest L dwarf found in the field and its late spectral type and spectral features indicative of a very low gravity suggest it might also be the lowest-mass field L dwarf. However, before ages and masses can be confidently adopted for any of these low-gravity L dwarfs, additional kinematic observations are needed to confirm cluster membership. ",
        "introduction": "Brown dwarfs are star-like objects with insufficient mass to sustain hydrogen burning ($M\\la75~M_{Jupiter}$). As a result, they never reach the main sequence and instead, continually cool with time and evolve through the MLT spectral sequence \\cite[and references therein]{Kirkpatrick05}. Unlike main sequence stars, the spectral type of a field brown dwarf does not provide a unique constraint on its mass. For example, an early-L dwarf could be an old low-mass star, a young low-mass brown dwarf, or an even younger planetary-mass brown dwarf.  Observations of spectral features that are sensitive to surface-gravity effects ($g \\propto M/R^2$) have long been used to distinguish compact dwarf stars from massive, extended giants with similar effective temperatures. The recognition of similar gravity-sensitive features in the spectra of brown dwarfs has begun to help break the age-mass degeneracy that has complicated their study. Brown dwarfs younger than $\\sim$100~Myr have low gravities because they have both larger radii (since they are still contracting) and lower masses than their same spectral type counterparts in the field \\citep{Burrows01}. Spectral features indicative of low-gravity (weak CaH, \\ion{K}{1}, \\ion{Na}{1}, and strong VO) have been identified in the optical spectra of late-M dwarf members of young 1--5~Myr clusters \\citetext{e.g., IC 348, \\citealp{Luhman99_IC348}; Taurus, \\citealp{Luhman03_taurus}}, the juvenile $\\sim$10~Myr TW Hydrae association \\citep{Gizis02}, and the adolescent 100~Myr Pleiades open cluster \\citep[e.g.,][]{Martin96}.  While optical (6000--9000~\\AA) spectra of very young M dwarfs are relatively common, similar data for L dwarfs, which have extremely faint optical magnitudes, are sparse. That said, optical spectroscopy has revealed several candidate members of $\\sigma$~Orionis, Lupus, Chamaeleon I, and Chamaeleon II to have early-L spectral types \\citep{BarradoyNavascues01,Jayawardhana06,Luhman08a,Luhman08b}. However, the low resolution and low signal-to-noise of those data preclude any examination of the low-gravity features that might be present. Near-infrared (1--2.5~\\micron) spectra of the faintest young cluster brown dwarfs are more feasible to obtain and have revealed candidate L dwarf members of Orion, Chameleon II, Ophiuchus, and Upper Scorpius \\citep{Lucas01, McGovern04, Allers07,lodieu08_spectra}. Unfortunately, near-infrared data of young objects systematically yield later spectral types than optical spectra \\citep{Luhman03_IC348}. For example, \\citet{Herczeg08_lris} find that several of the latest L dwarf candidates identified in Upper Sco with near-infrared spectra by \\citet{lodieu08_spectra} are actually late-M dwarfs in the optical. In short, unlike for M dwarfs, there are presently no optical spectra of 1--10~Myr-old L dwarfs with sufficient resolution or signal-to-noise ratios to serve as age benchmarks for L dwarfs found in the field.  There are two well-studied L dwarfs outside young clusters that display unambiguous low-gravity features in their optical and near-infrared spectra: G196-3B and \\objectname[2MASS J01415823-4633574]{2MASS J01415823$-$4633574} (hereafter 2M~0141$-$46 \\footnote{We use abbreviated notation for 2MASS sources throughout the text; e.g., 2M~hhmm$\\pm$dd, where the suffix is the J2000 sexagesimal right ascension (hours and minutes) and declination (degrees). Full source names and coordinates are provided in Table~\\ref{tab:lowg}.}). G196-3B was identified as an L2-type companion to the young ($\\sim$20--300~Myr) early-M dwarf G196-3A \\citep{Rebolo98, K01_gstar} while 2M~0141$-$46 was discovered as a single L0-type dwarf in the field \\citep{Kirkpatrick06}. \\citet{McGovern04} performed a thorough optical and near-infrared spectral analysis of G196-3B while this was done for 2M~0141$-$46 in its discovery paper. These authors find that G196-3B and 2M~0141$-$46 resemble early-type L dwarfs but have peculiar optical and near-infrared spectral features indicative of low surface-gravity ($log(g) \\sim~4.0 \\pm 0.5$). In particular, the spectra display weak \\ion{Na}{1} doublets (8183, 8195~\\AA; 1.13, 1.14~\\micron), weak and sharp \\ion{K}{1} doublets (7665, 7699~\\AA; 1.17, 1.24~\\micron), strong VO absorption bands (7300--7550, 7850--8000~\\AA; 1.05~\\micron), and weak FeH absorption bands (0.98, 1.19~\\micron). Based on the age estimate for G196-3A, \\citeauthor{Rebolo98} estimate the mass of G196-3B to be $25_{-10}^{+15} M_{Jupiter}$. Comparing the spectra of 2M~0141$-$46 to both models and late-type dwarfs with well-constrained ages, \\citeauthor{Kirkpatrick06} estimate 2M~0141$-$46 to have an age of 1--50~Myr and a mass of 6--25~$M_{Jupiter}$. Given its estimated age and location deep in the southern hemisphere, \\citeauthor{Kirkpatrick06} also hypothesize that 2M~0141$-$46 might be a member of the 12~Myr-old $\\beta$~Pictoris moving group or the 30~Myr-old Tucana/Horologium association.  In addition to these two benchmark objects, other L dwarfs with low-gravity features in their optical spectra have been identified in the field. Three were presented by \\citet{Cruz07} and fourteen were listed by \\citet{Paper10}. Combining these discoveries with new ones and improved data, \\citet{Kirkpatrick08} examined in detail the optical spectra of twenty low-gravity objects, including eight late-M dwarfs and eleven L dwarfs. In this paper, we focus only on L dwarfs, complementing the efforts of \\citeauthor{Kirkpatrick08}. Here we identify and examine twenty-three low-gravity L dwarfs found in the field and use extant data to compile a spectral sequence that spans from L0 to L5 and includes three gravity classes. These data include eight very low-gravity L0-type objects with spectra nearly identical to that of 2M~0141$-$46; two additional L0 dwarfs with spectral features suggesting intermediate gravities; and thirteen objects, including G196-3B, that have low gravities and even later spectral types (L1--L5).  In \\S~\\ref{sec:id} we briefly outline how the low-gravity objects were identified during the course of the NStars census of late-type dwarfs within 20~pc of the Sun. In \\S~\\ref{sec:types} we introduce a spectral typing scheme for low-gravity L dwarfs. In \\S~\\ref{sec:spectra}, the unusual spectra are described and compared to spectral standards via both overplotting and spectral indices. We use empirical arguments to justify low-gravity as the explanation for the observed spectral peculiarities. In \\S~\\ref{sec:discussion}, we discuss age, distance, and mass estimates for the low-gravity objects and, based on their distribution on the sky, suggest that they are likely members of nearby young moving groups. We summarize our results in \\S~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Ages} \\label{sec:ages}  Low gravity in brown dwarfs implies both low mass and youth. However, it is difficult to estimate the ages of our low-gravity objects since the optical spectra of young L dwarfs in clusters are limited to early spectral types (L0--L1) and are not of sufficient signal-to-noise nor resolution to use as fiducials. Despite the lack of an age-gravity calibration, we can still put rough age estimates on our low-gravity objects using their gravity-sensitive spectral features and their distribution on the sky.  As discussed by \\citet[see their Fig.~7]{Kirkpatrick08}, only objects Pleiades age and younger ($\\le$100~Myr) display low-gravity features discernible at the modest resolution of our data. At these young ages, brown dwarfs are still contracting and have larger radii, and thus lower gravities, than older objects of the same temperature \\citep{Burrows01}. Since we can clearly see low-gravity effects in the spectra of our objects, we adopt 100~Myr as the upper limit on their age. To estimate a lower age limit, we use the observation that none of our low-gravity objects are near young, dense star forming regions. Since the clusters these objects originated from appear to have dissipated, the objects are most likely $\\sim$10~Myr or older. Thus, we adopt an age range of $\\sim$10--100~Myr for our low-gravity L dwarfs.   As illustrated in the spectra shown in Figure~\\ref{fig:sequence1} and~\\ref{fig:sequence2} and quantitatively with the indices in Figures~\\ref{fig:lines} and~\\ref{fig:bands}, the low-gravity features in objects of the same spectral subtype have a range of strengths. The spectra are shown from top-to-bottom within each spectral subtype in decreasing order of their spectral peculiarity with the normal-gravity standard shown last. On the spectral index plots, objects with spectra indicative of very low-gravity are plotted as open circles while objects with intermediate-gravity features are plotted as shaded circles. The weakness of the \\ion{Na}{1} doublet and the \\ion{K}{1} doublet's cores and wings are the primary features used to classify objects as either very low-gravity or intermediate-gravity. As weaker alkali-metal features indicate lower gravities, the very low-gravity objects are also likely younger than the intermediate-gravity objects. Thus we hypothesize that within each subtype, ordering the spectra based on their implied gravity also creates an age sequence. Additionally, as long as the the effective temperatures of the younger objects are not significantly hotter than those of the older objects of the same spectral type, this sequence in gravity and age is also a mass sequence, with the youngest objects being the least massive.  \\citet{Kirkpatrick08} use the behavior of low-gravity features in late-M dwarfs of different ages (see their Fig.~7) to demonstrate that gravity-sensitive features distinguish ages to $\\approx$1~dex and we expect this remains true for L dwarfs. For example, among the L0 objects, the low-gravity features suggest that 2M~0141$-$46 and its clones (L0$\\gamma$) are the youngest with ages $\\approx$10--30~Myr; 2M~1552+29 and DENIS~0357$-$44 (L0$\\beta$) have intermediate ages of $\\approx$100~Myr; and the L0 standard 2M~0345+25 is the oldest with an age $>$200~Myr. Generalizing this idea implies that the objects we have classified as very low-gravity ($\\gamma$) are closer to $\\approx$10~Myr while the intermediate-gravity objects ($\\beta$) are more likely $\\approx$100~Myr. The validity of this hypothesis for L dwarfs will likely not be determined until cluster membership is confirmed for our objects. On the other hand, an age-gravity calibration for M dwarfs is possible and is underway and will likely shed light on the properties of low-gravity early-L dwarfs (Kirkpatrick et al., in preparation).  \\subsection{Distances/Temperatures} \\label{sec:distances}  We have estimated crude distances for the low-gravity objects with the $M_J$-spectral type relation derived by \\citet{Cruz03}. The uncertainties on the distances were estimated by propagating through the uncertainty on $M_J$ found for a spectral type uncertainty of $\\pm$~1 and using the covariance matrix of the best fit parameters of the $M_J$-spectral type relation. Field dwarfs with measured parallaxes listed by \\cite{Dahn02} were used to derive this relation and it is not strictly valid for low-gravity objects. The $M_J$-spectral type relation has been updated to include new results and revised to include T dwarfs \\citep[most recently by ][]{Looper08_T}. However, we continue to use the \\citeauthor{Cruz03} relation to maintain consistency within our sample of objects found by the NStars 20-pc survey. The difference between the $M_J$ predicted by the two relations is at most 0.17~mag for L6 but is less than 0.1~mag for L4 and earlier.  Evaluating if the field $M_J$-spectral type relation under- or over-estimates the true distances to low-gravity objects is not straightforward. On one hand, the low-gravity objects could be brighter than normal objects of the same spectral type and the $M_J$-spectral type relation would underestimate luminosities and distances for low-gravity objects. This is indeed the case for higher-mass stars where the low-gravity objects (i.e., giants) have the similar effective temperatures to their normal gravity counterparts (i.e., dwarfs), but much larger radii. This scenario would also be valid if the temperature scale for low-gravity objects is hotter than that of normal dwarfs as derived by \\cite{Luhman99_IC348} and \\cite{Luhman03_IC348}.  There are four objects with independently derived distances that might shed light on the validity of the $M_J$-spectral type relation for young objects. \\citet{Kirkpatrick06} studied the low-gravity early-L dwarf 2M~0141$-$46 in detail and compared its observed optical and near-infrared spectra to model spectra; based on this thorough analysis, they estimate a distance of $\\sim$35~pc. Using a spectral type of L0$\\pm1$ for 2M~0141$-$46, the $M_J$-spectral type relation yields a distance of 41$\\pm5$~pc. Similarly, \\cite{Rebolo98} estimate a distance to G196-3A of $21\\pm6$ and the $M_J$-spectral type relation gives $27\\pm5$ using L3$\\pm1$ for the low-gravity secondary. In both cases, the two values agree within the uncertainties. On the other hand, \\citet{Gizis07} and \\cite{Teixeira08} find trigonometric parallax distances of $54.0^{+3.2}_{-2.8}$~pc and 55~pc for 2M~1207$-$39 and SSSPM~1102$-$34, respectively. Both of these objects are late-type M dwarfs in the TW Hydrae Association. The $M_J$-spectral type relation yields a distance of 25$\\pm$2 and 22$\\pm$2 to the two systems respectively, an underestimate by a factor of two. Each of these four systems are not ideal for evaluating the $M_J$-spectral type relation for young objects: the distance estimate for 2M~0141$-$46 is based on un-tested models and is not robust, the G196-3AB system might be older than our typical low-gravity L dwarf, and both SSSPM~1102$-$34 and 2M~1207$-$39 harbor active accretion disks \\citep{Morrow08,Riaz08}. Clearly, until trigonometric parallaxes are obtained for a relatively large sample of objects, both the temperature and distance scales will remain uncertain for low-gravity L dwarfs.  Regardless of the complications, the $M_J$-spectral type relation gives a rough estimate of the distance and at least provides a sense of relative distances to the low-gravity objects. For the lack of a better method, we adopt the distances predicted by the $M_J$-spectral type relation (with an uncertainty of $\\pm$1 on the spectral type) for the low-gravity objects and list them in Table~\\ref{tab:lowg}.  \\subsection{Distribution on the Sky and Coincidence with Moving Groups}  The location of our candidate low-gravity brown dwarfs on the sky is shown in Figure~\\ref{fig:map} (\\textit{open and shaded circles}). Also shown are members of the AB Doradus moving group ($\\sim$100~Myr, \\textit{blue plusses}), the Tucana/Horologium Association (Tuc/Hor, $\\sim$30~Myr, \\textit{green crosses}), the $\\beta$ Pictoris moving group ($\\sim$10~Myr, \\textit{red five-point asterisks}), and the TW Hydrae association (TWA, $\\sim$10~Myr, \\textit{purple six-point asterisks}) as identified by \\citet{Zuckerman04}. The positional coincidence and agreement between the estimated distances and ages all strongly suggest that our low-gravity L dwarfs are members of these groups. Six of the nine very low-gravity L0$\\gamma$ objects (\\textit{open circles}) overlap the core of the Tuc/Hor association. The intermediate-gravity dwarfs (\\textit{shaded circles}) are more widely distributed on the sky and appear to have a distribution most similar the members of the AB Dor moving group. G196-3B and 2M~1022+58 are the outliers on the figure at 10$^{\\mbox{h}}$, +50$\\degr$ and do not not appear to be connected to any of the young associations considered here. This figure is similar to one presented in \\citet{Kirkpatrick08} but here we show only L dwarfs and do not consider late-M dwarfs. (As mentioned previously, the sample of twenty-three L dwarfs considered here and the twenty objects studied by \\citet{Kirkpatrick08} have seven L dwarfs in common.)  The majority of the candidate low-gravity objects were identified via a color-magnitude query of the 2MASS All-Sky survey and thus their spatial distribution is not a selection effect. By far, the largest position-based cuts in our 2MASS search were to avoid the crowded and reddened regions of the Galactic plane ($|b|>15\\degr$) and known star formation regions that extend beyond the plane (e.g., Upper Sco, Taurus). Other small regions were also excluded such as M31, M33, and the Large and Small Magellanic Clouds. As shown in Figure~1 of \\cite{Paper10}, the distribution of the sources surviving these cuts, which is the parent population of our low-gravity objects, is uniform. Again, the selection criteria and sky coverage of our 2MASS search are discussed in detail by \\cite{Cruz03, Cruz07} and \\citeauthor{Paper10}.  In addition to the coincidence of the projected location on the sky of our low-gravity objects with the moving groups, their distance estimates are also broadly consistent with each other. The M$_J$-spectral type relation yields distances to the low-gravity objects all within 85~pc; the majority are estimated to be between 13~and 50~pc. (As described in the previous section, these distances are very uncertain.) The mean distances to the moving groups are also between 20 and 50~pc; however, the spread of distance estimates within each group is quite large. For example, as listed by \\citet{Zuckerman04}, the distances to $\\beta$~Pic members range from 15--60~pc and 37--66~pc for Tuc/Hor members. As such, unlike most of their young cluster counterparts, these nearby moving groups overlap with each other and have neither a unique distance nor unique location on the sky. Thus, accurate distances alone for the low-gravity objects would not determine which group they belong to. For these reasons, we are currently targeting all of our low-gravity objects for proper motion, radial velocity, and parallax measurements in order to derive their UVW space motions and to confirm or refute cluster membership \\citep{Mamajek05}.  None of the low-gravity L dwarfs appear to be members of TWA. Indeed, the only TWA member uncovered by the 20-pc census is \\object[2MASS J11395113-3159214]{2MASS J11395113$-$3159214} (M9). The other late-M dwarf members were excluded by the color cuts: \\object[2MASS J12073346-3932539]{2MASS J12073346$-$3932539} failed the $J-K_s$ cut while \\object[2MASS J11020983-3430355]{2MASS J11020983$-$3430355} failed the $R-J$ cut \\citep{Cruz03}. With a mean distance of $\\sim$54~pc, TWA is the most distant of the nearby associations and any L dwarf members were probably also excluded by the color-magnitude search criteria designed to constrain the sample to only the nearest/brightest and latest-type/reddest dwarfs. While recent work using the Deep Near-Infrared Southern Sky Survey (DENIS) has identified one new late-type M dwarf member of TWA \\citep{Looper07_1245}, no L dwarfs were found by this effort (Looper 2008, private communication). Despite its large spatial extent, a deep survey of TWA is warranted to uncover its later-type members.  \\subsection{Masses} Evolutionary models, combined with reasonable assumptions about ages and temperatures, can be used to estimate the masses of the low-gravity L dwarfs. The temperature scale for field L dwarfs has been found to range from 2400~K at L0 down to 1500~K for L5--L8 \\citep{Golimowski04}. Even though the temperature scale for low-gravity L dwarfs is likely to be different than that of normal-gravity field dwarfs, it is sufficient to make the rough mass estimates we are interested in here. For field-age (1--5 Gyr) objects within that temperature range, the evolutionary models of \\cite{Burrows01} and \\cite{Baraffe03} predict masses of 48--85~$M_{Jupiter}$. If, however, the low-gravity L dwarfs are indeed young as we suspect, the models predict very-low masses: at 30~Myr, the age of Tuc/Hor, 1500--2400~K corresponds to 11--30~$M_{Jupiter}$.  One individual object, 2M~0355+11, merits comment. We characterize this object's peculiar spectrum as a very low-gravity L5$\\gamma$ (Figure~\\ref{fig:sequence2}). This interpretation puts this object at the lowest end of both our age and temperature estimates and we speculate it has on effective temperature of 1500--1700~K (based on its similarity to an L5) at an age of 10--30~Myr (based on its very  low-gravity spectral features). In this scenario, the evolutionary models predict a mass of 8--13~$M_{Jupiter}$. If indeed this object is quite young and very low-mass, it would be the lowest-mass L dwarf identified in the field and an analog to 2M~1207$-$39B, the very cool L dwarf companion found in TWA \\citep{Chauvin05}. Both of these objects have very red $J-K$ colors and are the two reddest L dwarfs known: 2M~0355+11 and 2M~1207$-$39B have $J-K_s$ colors of 2.52$\\pm$0.03 and 3.07$\\pm$0.23 respectively \\citep{Mohanty07}. However, further age and effective temperature constraints are necessary before a robust mass estimate for 2M~0355+11 can be made.     We have identified 23 L dwarfs with very peculiar spectral features in their optical spectra. The peculiar features that we observe include weak alkali doublets (\\ion{Na}{1}, \\ion{K}{1}, \\ion{Rb}{1}, \\ion{Cs}{1}); weak \\ion{K}{1} wings; weak absorption bands of FeH, CrH, and TiO; and strong VO absorption bands at early types. All of these spectral peculiarities, except weak TiO, are consistent with the spectral features characterizing low-gravity objects (i.e., giants and young M dwarfs). After considering and ruling out other possible explanations for the observed spectral peculiarities, we conclude low-gravity is the most probable underlying cause. We use these spectra to expand the \\cite{K99} scheme to include three gravity classes spanning L0 to L5. The gravity of the spectrum is indicated by a greek suffix: $\\alpha$ (or lack of suffix) for normal gravity, $\\beta$ for intermediate-gravity, and $\\gamma$ for very-low gravity. We do not propose any modification to the \\cite{K99} method for assigning spectral types to normal-gravity L dwarfs. Compared to normal dwarfs with the same spectral type, the low-gravity objects have red $J-K_s$ colors. Contrary to conventional wisdom, only 9/23 display \\ion{Li}{1} in absorption and only 3/23 objects display H$\\alpha$ in emission. However, higher signal-to-noise observations are required to robustly measure the detection fraction of these lines in low-gravity L dwarfs.  Based on the low-gravity features present in the spectra of our objects and their lack of association with young, dense star-forming regions, we estimate their ages to be approximately 10--100~Myr. Further, given the behavior of low-gravity features in the spectra of late-M dwarfs of different ages shown by \\citet{Kirkpatrick08}, we assert that the objects we have classified as very low gravity are closer to $\\approx$10~Myr while the intermediate-gravity objects are more likely $\\approx$100~Myr. While the temperature and distance scales for low-gravity L dwarfs are still uncertain, we estimate rough distances to our objects using the \\citet{Cruz03} $M_J$-spectral type relation for field dwarfs. There is positional, age, and distance agreement between the low-gravity L dwarfs and southern, nearby loose associations, most notably Tuc/Hor. However, before membership in any association or moving group can be assigned and ages adopted, astrometric and radial velocity measurements are required and those observations are underway. Future statistically robust studies of these nearby associations could yield an initial mass function that probes the planetary-mass regime and possibly identify the true lowest-mass objects formed via star formation."
    },
    "0812/0812.0478_arXiv.txt": {
        "abstract": "{A pulsar searching system has been operating at 18-cm band at the Urumqi Observatory 25-m radio telescope since mid-2005. Test observations for known pulsars show the system can perform pulsar searching. Perspect of using this system to observe 3EG sources and other target searching are prepared and discussed. \\keywords {pulsar, search, radio} }  \\authorrunning{Jiang Dong}            % \\titlerunning{Finding Radio Pulsar in 3EG Sources at Urumqi Observatory}  % ",
        "introduction": "Among the more than 1700 known pulsars, seven are seen at $\\gamma$-ray frequency (Lorimer \\& Kramer, 2005). It is worth mentioning that six of the seven have detected radio emission (McLaughlin 2001). This might indicate the possible relationship between $\\gamma$-ray sources and radio pulsars (McLaughlin 2001, Lorimer 2003, Gonthier et al. 2002, Cheng et al. 2004, Qiao et al. 2004). Encouraged by such phenomenon, we started the program of searching for radio pulsar in the error boxes of 3EG sources.  The Compton Gamma Ray Observatory (CGRO) was the second of NASA's Great Observatories which was operated from April 5, 1991 to June 4, 2000. The Energetic Gamma Ray Experiment Telescope (EGRET) provides the highest energy gamma-ray window for the Compton Observatory. Its energy range is from 20~MeV to 30~GeV. EGRET is 10 to 20 times larger and more sensitive than previous detectors operating at these high energies and has made detailed observations of high-energy processes associated with diffuse gamma-ray emission, gamma-ray bursts, cosmic rays, pulsars, and active galaxies known as blazars.  In next section we will introduce the pulsar searching system at Nanshan, Urumqi Observatory, and discuss using of small radio telescolpe (SRT) to do pulsar search and possible transient radio source search (Cordes, J.M., et al., 2004). We report the prime results and discuss the perspect of future observations in Section 3. ",
        "conclusions": ""
    },
    "0812/0812.0014_arXiv.txt": {
        "abstract": "After nearly a decade of quiescence, the soft gamma-ray repeater \\src\\ reactivated on 2008 May 28 with a bursting episode followed by a slowly decaying enhancement of its persistent emission. To search for the still unknown spin period of this SGR taking advantage of its high flux state, we performed on 2008 September 27--28 a 120 ks long X-ray observation with the \\xmm\\ satellite. Pulsations with $P=2.594578(6)$ s were detected at a higher than 6$\\sigma$ confidence level, with a double-peaked pulse profile. The pulsed fraction in the 2--12 keV range is $19\\%\\pm3\\%$ and $24\\%\\pm3\\%$ for the fundamental and the second harmonic, respectively. The observed 2--10 keV flux is $3.4\\times10^{-13}$ \\flux, still a factor of $\\sim$ 5 above the quiescent pre-burst-activation level, and the spectrum is well fitted by an absorbed power law plus blackbody model (photon index $\\Gamma\\simeq0.6$, blackbody temperature $kT\\simeq0.5$ keV, and absorption $N_{\\rm{H}}\\approx1.2\\times10^{23}$ cm$^{-2}$). We also detected a shell of diffuse soft X-ray emission which is likely associated with the young supernova remnant G337.0--0.1. ",
        "introduction": "Two classes of young isolated neutron stars, the anomalous X-ray pulsars (AXPs) and the soft gamma-ray repeaters (SGRs) are currently believed to host magnetars, i.e. isolated neutron stars endowed with ultra-strong magnetic fields, $B\\sim 10^{14}$--$10^{15}$ G on the surface and one (or more) order of magnitude higher in the interior \\citep{duncan92,thompson93,thompson95}. SGRs and AXPs share a number of properties, including rotation periods in the 2--12 s range, large period derivatives ($\\dot P\\sim 10^{-13}$--$10^{-9}$ s s$^{-1}$), inferred surface magnetic dipole field strength $B\\gtrsim10^{14}$ G, large and variable X-ray luminosities (exceeding that available from the braking of their rotation), and the emission of short bursts (see \\citealt{woods06} and \\citealt{mereghetti08} for recent reviews).\\\\ \\indent The five known Local Group SGRs were discovered as sources of short ($<$1 s) and intense ($10^{41}$--$10^{43}$ \\lum) bursts of gamma rays that they emit during sporadic periods of activity interrupted by long stretches of quiescence.  Once localized through their bursts and flares, all of them were found to be persistent X-ray emitters, with luminosities of about $10^{33}$--$10^{36}$ \\lum. In four SGRs, \\src\\ being the exception,  pulsations were detected with periods from 5 to 8 s and period derivatives of $10^{-11}$--$10^{-9}$ s s$^{-1}$.\\\\ \\indent \\src\\ was discovered in 1998 June by the \\emph{Compton Gamma Ray Observatory}  thanks to the intense bursts it emitted at that time, about 100 in six weeks \\citep{woods99}. Soon after the discovery of the bursts, its persistent X-ray emission was detected by \\sax\\ at  a flux level of $\\sim$ $7\\times10^{-12}$ \\flux\\ (unabsorbed, 2--10 keV), corresponding to a luminosity of $\\sim$ $10^{35}$ \\lum\\ for a source distance $d= 11$ kpc \\citep{corbel99}. This led to the possible association of the SGR with the young supernova remnant (SNR) G337.0--0.1 in the CTB\\,33 complex \\citep{woods99,hurley99_1627}. In the following 10 years no further bursting activity was reported while the persistent emission steadily decayed to about $10^{33}$ \\lum\\ (for $d= 11$ kpc; \\citealt{corbel99}), the lowest value ever observed from an SGR. At the same time the spectrum softened \\citep{kouveliotou03,mereghetti06}.\\\\ \\indent The long-term fading of \\src\\ was suddenly interrupted by its burst reactivation in 2008 May. This episode was associated with a temporary enhancement of the persistent X-ray flux and a marked spectral hardening \\citep{eiz08}. \\\\ \\indent The detection of SGR-like bursts from \\src\\ made it a bona fide member of the SGR class. However, two strong pieces of evidence in favor of this identification, the measure of the source period and period derivative, were still missing.\\footnote[9]{\\citet{woods99} reported a candidate periodicity of marginal significance at 6.41 s, but the signal was not confirmed by subsequent observations.} In order to search for pulsations taking advantage of the enhanced flux after the reactivation, we requested a Target of Opportunity \\xmm\\ observation, that was performed as soon as the satellite visibility constraints allowed it. ",
        "conclusions": "We have discovered pulsations with period of 2.6 s in \\src, which was the only known magnetar with no spin period yet measured. The pulse shape is double peaked, with pulsed fractions of $\\sim$ $19\\%$ and $24\\%$ in the two harmonics. No dependence of the pulse profile on the energy band has been found. Among the known magnetars, only 1E\\,1547.0--5408 has a shorter period (2.07 s; \\citealt{camilo07}) than that of \\src. We searched for a similar periodicity in all available archival data with sufficient time resolution, but no pulsation was found. Relying only on the last \\xmm\\ observation we could set an upper limit on the period derivative,  $|\\dot{P}|<6\\times 10^{-10}$ s s$^{-1}$ which is not particularly constraining for a magnetar, and yields $B\\lesssim1.2\\times10^{15}$ G. Based on the peak luminosities of bursts observed from \\src, its dipole magnetic field was estimated to be $B\\gtrsim1.8\\times10^{14}$ G \\citep{woods99,eiz08}. This limit can be used to infer a period derivative $\\dot{P}\\gtrsim1.2\\times10^{-11}$ s s$^{-1}$, a spin-down luminosity $\\dot{E}\\gtrsim8\\times10^{33}$ \\lum, and a characteristic age $\\tau_c= P/(2\\dot{P})\\lesssim3.4$ kyr. Overall, and lacking any information on the pulsation period and profile in the past, we can only conclude that the timing properties of \\src\\ reported here appear quite typical of a member of the SGR class.\\\\ \\indent Our observation  caught \\src\\ at a flux of $3.4\\times10^{-13}$ \\flux, the highest seen with \\xmm\\ and a factor of $\\sim$ 5 above the level preceding its 2008 May burst activation (see Figure~1 of \\citealt{eiz08} for a complete flux history of the source). Its spectrum was well fitted in earlier data by a single-component model, either a heavily absorbed power law or a blackbody \\citep{mereghetti06}. The higher-quality spectrum of the new data cannot be adequately fitted by these simple models, while it is consistent with a power law plus blackbody model, and it is significantly harder than in any observation of \\src\\ taken before 2008 May \\citep{kouveliotou03,mereghetti06}. The presence of a soft thermal component and the spectral hardening as the source moves from a quiescent to a (burst) active phase makes \\src\\ akin to other SGRs and AXPs (e.g. \\citealt{mte05,rea05}). This latter behavior can be interpreted in the framework of the ``twisted-magnetosphere model'' as due to a progressive growth of the shear in the magnetosphere as magnetic helicity is transferred from the internal to the external field \\citep{tlk02,fernandez07,nobili08}.\\\\ \\indent \\src\\ has been proposed to be associated (\\citealt{woods99,hurley99_1627}, but see also \\citealt{gaensler01}) to the radio SNR G337.0--0.1 \\citep{frail96,whiteoak96}. The shell of diffuse soft radiation (Section~\\ref{imaging}), consistent with emission from a young SNR at the distance of \\src, is likely to be the first detection of this SNR in X-rays. Due to the complexity of the region, the extension of the SNR at radio frequencies is debated: \\citet{sarma97} claimed a diameter of $\\sim$ 95 arcsec with \\src\\ outside the shell, while \\citet{brogan00} proposed a different morphology extending the SNR size towards the southwest and encompassing the SGR. The position and  extension of the diffuse X-ray emission are consistent with the latter hypothesis, supporting the possible association between \\src\\ and G337.0--0.1.\\\\"
    },
    "0812/0812.0825_arXiv.txt": {
        "abstract": "A population of  very light primordial black holes  which  evaporate before nucleosynthesis begins is unconstrained unless   the decaying black holes leave stable relics. We show that gravitons Hawking radiated from these black holes would source a substantial stochastic background of high frequency gravititational waves ($10^{12}$ Hz or more) in the present universe. These black holes may lead to a transient period of matter dominated expansion. In this case the primordial universe could be temporarily dominated by large clusters of ``Hawking stars\" and the resulting gravitational wave spectrum is independent of the initial number density of primordial black holes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.3401_arXiv.txt": {
        "abstract": "Jet physics is again flourishing as a result of \\chandra's ability to resolve high-energy emission from the radio-emitting structures of active galaxies and separate it from the X-ray-emitting thermal environments of the jets.  These enhanced capabilities have coincided with an increasing interest in the link between the growth of super-massive black holes and galaxies, and an appreciation of the likely importance of jets in feedback processes.  I review the progress that has been made using \\chandra\\ and \\xmm\\ observations of jets and the medium in which they propagate, addressing several important questions, including:   Are the radio structures in a state of minimum energy?  Do powerful large-scale jets have fast spinal speeds?  What keeps jets collimated? Where and how does particle acceleration occur? What is jet plasma made of? What does X-ray emission tell us about the dynamics and energetics of radio plasma/gas interactions? Is a jet's fate determined by the central engine? ",
        "introduction": "\\label{sec:stage}  \\subsection{Historical perspective} \\label{sec:historical}  In the 1970s and 1980s the powerful capabilities of radio interferometry gave birth to the study of extragalactic radio jets. It became clear that radio jets are plasma outflows originating in the centres of active galaxies, seen through their synchrotron emission. After much debate, properties such as the relative one-sidedness of the jets, and the measurement of apparent superluminal expansion, by Very Long Baseline Interferometry (VLBI), were accepted as due to the outflows having relativistic bulk speeds.  Early attempts at unifying source populations based on special relativity and apparent source properties \\citep[e.g.,][]{scheuer-readhead} have developed over the years into comprehensive unified schemes \\citep[e.g.,][]{barthel} whereby quasars are explained as radio galaxies whose jets are at small angles to the line of sight and so are boosted by relativistic effects.  By the mid 1990s, the study of radio jets had reached something of a hiatus, and major groups around the world turned their attention to other pursuits such as gravitational lensing and the study of the Cosmic Microwave Background (CMB) radiation.  A turning point was the sensitivity and high-fidelity mirrors of the \\chandra\\ X-ray Observatory \\cite{weisskopf}, which resulted in the detection of resolved X-ray emission from many tens of well-known extragalactic radio sources (see \\cite{harris-kraw} for a source compilation as of 2006: the number continues to increase).  When combined with X-ray measurements of the ambient gas made with \\chandra\\ and \\xmm, and multiwavelength data, many important questions related to the physics of jets can be addressed.  Progress towards answering those questions is the substance of this review.  The enhanced capabilities for the X-ray study of jets have coincided with strong interest from the wider astronomical community in the growth of supermassive black holes (SMBHs), following the links that have been made between SMBH and galaxy growth \\citep[e.g.,][]{richstone, gebhardt}.  SMBHs (and indeed compact objects of stellar mass) commonly produce jets, as an outcome of accretion processes responsible also for black-hole growth.  It is also clear that extragalactic jets are capable of transferring large amounts of energy to baryonic matter in the host galaxies and surrounding clusters at large distances from the SMBH.  The way in which heating during the jet mode of AGN activity might overcome the problem of fast radiative cooling in the centre of clusters is now intensely studied in nearby objects (see \\S \\ref{sec:dynamics}), and heating from `radio mode' activity is included in simulations of hierarchical structure formation \\citep[e.g.,][]{croton}.  We need therefore to understand what regulates the production of jets and how much energy they carry.  X-ray measurements of nuclear emission probe the fueling and accretion processes, and those of resolved jet emission and the surrounding gaseous medium probe jet composition, speed, dynamical processes, energy deposition, and feedback.  \\subsection{Radiation processes} \\label{sec:processes}  The two main jet radiation processes are synchrotron radiation and inverse-Compton scattering. Their relative importance depends on observing frequency, location within the jet, and the speed of the jet. The thermally X-ray-emitting medium into which the jets propagate plays a major r\\^ole in the properties of the flow and the appearance of the jets.  The physics of the relevant radiation processes are well described in published work \\citep[e.g.,][]{ginzburg, blum-gould, pachol, spitzer, rlightman, sarazin, longair}, and most key equations for the topics in this review, in a form that is independent of the system of units, can be found in \\cite{worr-santiago}.  It is particularly in the X-ray band that synchrotron radiation and inverse-Compton emission are both important.  X-ray synchrotron emission depends on the number of high-energy electrons and the strength and filling factor of the magnetic field in the rest frame of the jet. Inverse Compton X-ray emission depends on the number of low-energy electrons, the strength of an appropriate population of seed photons (such as the CMB, low-energy jet synchrotron radiation, or emission from the central engine), and the geometry of scattering in the rest frame of the jet.  In an ideal world, observations would be sufficient to determine the emission process, and this in turn would lead to measurements of physical parameters.  In reality, X-ray imaging spectroscopy, even accompanied by good measurements of the multiwavelength spectral energy distribution (SED), often leaves ambiguities in the dominant emission process.  Knowledge is furthered through intensive study of individual sources or source populations.  \\subsection{Generic classes of jets} \\label{sec:classes}    In discussing jets, it is useful to refer to the Fanaroff and Riley \\cite{fr} classification that divides radio sources broadly into two morphological types, FRI and FRII.  A relatively sharp division between FRIs and FRIIs has been seen when sources are mapped onto a plane of radio luminosity and galaxy optical luminosity \\cite{ledlow-owen} -- the so called Ledlow-Owen relation. FRIIs are of higher radio luminosity, with the separation between the classes moving to larger radio luminosity in galaxies that are optically more massive and luminous. The distinct morphologies \\citep[e.g.,][]{muxlow} are believed to be a reflection of different flow dynamics \\citep[e.g.,][]{leahy}.  \\begin{figure}[t] \\centering \\includegraphics[height=5.0cm]{n315radio.eps} \\includegraphics[height=5.0cm]{n315xray.eps} \\caption{Roughly 6.6~kpc (projected) of the inner jet of the $z=0.0165$ FRI radio galaxy NGC~315. Left: 5~GHz \\vla\\ radio map showing a knotty filamentary structure in diffuse emission. Right: Smoothed \\chandra\\ X-ray image of $\\sim 52.3$~ks livetime also showing knotty structure embedded in diffuse emission. The ridge-line defined by the radio structure is shown in white, and indicates a level of correspondence between the radio and X-ray knots. Figure adapted from \\protect\\cite{worr-N315}. } \\label{fig:n315jet} \\end{figure}   FRI sources (of lower isotropic radio power, with BL Lac objects as the beamed counterpart in unified schemes) have broadening jets feeding diffuse lobes or plumes that can show significant gradual bending, usually thought to be due to ram-pressure as the source moves relative to the external medium.  The jet emission is of high contrast against diffuse radio structures, implying that the jet plasma is an efficient radiator. kpc-scale jets are usually brightest at a flaring point some distance from the active galactic nucleus, and then fade gradually in brightness at larger distances from the core, although this pattern is often interrupted by bright knots seen when the jet is viewed in the radio or the X-ray.  Such an example is shown in Figure~\\ref{fig:n315jet}\\footnote{Values for the cosmological parameters of $H_0 = 70$~km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm {m0}} = 0.3$, and $\\Omega_{\\Lambda 0} = 0.7$ are adopted throughout this review.}. The jets are believed to slow from highly-relativistic to sub-relativistic flow on kpc-scales from entrainment of the external interstellar medium (ISM), perhaps enhanced by stellar mass loss within the jet.  The strong velocity shear between the jet flow and the almost stationary external medium must generate instabilities at the interface \\cite{birkinshaw-kh}, and drive the flow into a turbulent state.  The physics of the resulting flow is far from clear, although it can be investigated with simplifying assumptions \\citep[e.g.,][and see \\S \\ref{sec:jetcollimate}]{bick-31, bick-315}.  \\begin{figure}[t] \\centering \\includegraphics[height=6.0cm]{3c200-smo3-xray-reduced.ps} \\caption{The $z=0.458$ FRII radio galaxy 3C~200. A smoothed 0.3-5~keV \\chandra\\ X-ray image of $\\sim 14.7$~ks livetime  is shown with radio contours from a 4.86~GHz \\vla\\ radio map \\cite{leahy-atlas} (beam size $0.33''\\times 0.33''$).  Both nuclear \\protect\\cite{belsole-cores} and extended X-ray emission are detected. A rough correspondence of some of the extended X-ray emission with the radio lobes has resulted in the claim for inverse-Compton scattering of the CMB by electrons in the lobes \\protect\\cite{croston-lobes}, but most of the extended emission over larger scales is now attributed to cluster gas \\protect\\cite{belsole-env}. } \\label{fig:3c200} \\end{figure}   FRII sources (of higher isotropic radio power, with quasars as the beamed counterpart in unified schemes) have narrower jets that are sometimes faint with respect to surrounding lobe plasma and that terminate at bright hotspots (Fig.~\\ref{fig:3c200}).  The jets are often knotty when observed with high resolution, and the jets can bend abruptly without losing significant collimation (see \\S\\ref{sec:0637} and \\S\\ref{sec:FRIpolariz} for examples). The bending is often large in quasar jets, supporting the conjecture that quasars are viewed at small angle to the line of sight and that bends are amplified through projection.  In contrast to FRI jets which are in contact with the external medium, the standard model for FRII jets is that they are light, embedded in lobe plasma, and remain supersonic with respect to the external gas out to the hotspots.  The energy and momentum fluxes in the flow are normally expected to be sufficient to drive a bow shock into the ambient medium.  The ambient gas, heated as it crosses the shock, forces old jet material that has passed through the hotspots into edge-brightened cocoons.  FRII jets are thus low-efficiency radiators but efficient conveyors of energy to large distances.  They are often hundreds of kpc in length (particularly when deprojected for their angles to the line of sight), crossing many scale heights of the external medium from relatively dense gas in a galaxy core to outer group or cluster regions where the external density and pressure are orders of magnitude lower.  State-of-the-art three-dimensional magneto-hydrodynamical simulations that incorporate particle transport and shock acceleration do well at reproducing the essential characteristics of synchrotron emission from such a source, and suggest that the shock and magnetic-field structures of the hotspots and lobes are extraordinarily complex and unsteady \\cite{tregillis, tregillis2}.  \\subsection{Lifetimes and duty cycles} \\label{sec:dutycycles}  Individual FRI and FRII radio galaxies are thought to live for at most some tens of millions of years \\citep[e.g.,][]{mack, kaiser}.  Age estimates are based on measuring curvature in the radio spectra caused by radiative energy losses of the higher-energy electrons over the lifetime of the sources \\citep[e.g.,][]{alexanderleahy}.  In contrast to the relative youth of observed radio structures, present-day clusters were already forming in the young Universe.  Ideas that radio sources have an important r\\^ole in heating cluster gas (see \\S\\ref{sec:dynamics}) then require a correct balance between the duty-cycle of repeated radio activity and heating efficiency as a function of jet luminosity.  The duty cycle can be probed by searching for evidence of repeated activity from individual sources.  Radio sources classified as GHz-Peaked Spectrum (GPS) or Compact Steep Spectrum (CSS) are small and believed to be either young or have their growth stunted by the external medium \\cite{odea-GPS}, and source statistics suggest that if they evolve to kpc-scale sizes they must dim while so doing \\cite{readhead}.  VLBI kinematic studies provide convincing evidence that sources in the Compact Symmetric Object (CSO) subset, at least, are young, with current ages less than $10^4$ years \\cite{conway}.  The fact that it is relatively uncommon to see GPS sources with extended radio emission that may be a relic of previous activity has been used to argue that periods between sustained activity are generally at least ten times longer than the radiative lifetime of the radio emission from the earlier activity \\cite{stanghellini}.  This is consistent with a time between episodes of activity in FRIIs of between about $5 \\times 10^8$ and $10^9$ years that is estimated using optical- and radio-catalog cross correlations coupled with an average source lifetime of about $1.5 \\times 10^7$ years from modelling projected source lengths \\cite{bird}. Of course, within the lifetime of an individual radio source there might be shorter-term interruptions or variations of activity (see \\S\\ref{sec:evolution}). ",
        "conclusions": "\\label{sec:conclusions}   The last decade has seen massive progress in our understanding of the X-ray properties of extragalactic radio jets and their environments.  \\chandra's sub-arcsec spatial resolution has been of paramount importance in measuring resolved X-ray emission from kpc-scale jet structures, and in extending studies of X-ray nuclei to sources other than beamed quasars and BL~Lac objects by separating the emission of weaker nuclei from that of the jets and X-ray emitting environments.  The assumption that radio structures roughly lie in a state of minimum energy between their relativistic particles and magnetic fields is broadly verified in a few tens of sources through combining X-ray inverse Compton with radio synchrotron data (\\S\\ref{sec:minenergytest}). This is the assumption commonly adopted in the absence of other information, and so its verification is reassuring, although much sub-structure is likely to occur and there is no reason to expect minimum energy to hold in dynamical structures.  The increase in numbers of known resolved kpc-scale X-ray jets has been remarkable, from a handful to the several tens of sources that \\chandra\\ has mapped in detail.  There are grounds to believe that there are X-rays from synchrotron radiation in sources both of FRI and FRII types (\\S\\ref{sec:Xsyn}), requiring in-situ particle acceleration to TeV energies.  The steepening in spectral slope which most commonly occurs at infra-red energies may be related more to acceleration processes than energy losses, but more multiwavelength observational work is required to characterize the acceleration sites and support a theoretical understanding.  The fact that X-ray synchrotron emission with an X-ray to radio flux-density ratio, $S_{\\rm 1~keV}/S_{\\rm 5~GHz}$, between about $10^{-8}$ and $10^{-7}$ is so common in jets where the bulk flow is inferred to be relativistic implies that there will be many more X-ray jet detections with current instrumentation in sufficient exposure time.  The dominant X-ray emission mechanism in resolved quasar jets remains uncertain, but it is likely that beamed emission from scattering of CMB photons is dominant in jets at small angles to the line of sight.  This requires that highly relativistic bulk flows exist far from the cores, contradicting earlier radio studies but possibly understandable in the context of transverse velocity profiles. The knotty appearance of these jets is then possibly a result of variable output from the nuclei. Much of the knotty X-ray appearance of FRI jets, on the other hand, likely arises from spatial variations in the strength of particle acceleration (\\S\\ref{sec:CenAn315}).  Jet theory has had some pleasing successes, such as the agreement between X-ray pressure profiles and predictions from hydrodynamical models for low-power jets in the regions where they are believed to be slowed by entrainment of the external medium or stellar mass loss (\\S\\ref{sec:jetcollimate}).  We are still largely ignorant of jet composition, and this is a difficult problem to solve since jet dynamics are governed by the energy of the constituent particles rather than their mass.  There is generally growing support for a strong presence of relativistic protons (\\S\\ref{sec:jetcomposition}).  The observation of bubbles and cavities in cluster gas produced dynamically by radio structures has renewed interest in the mechanisms by which active galaxies introduce heat into gaseous atmospheres.  A few nearby bright systems have been the subject of intense study with \\chandra\\ (\\S\\ref{sec:FRIdynamicscluster}).  Although the way in which energy is deposited on the large scale is still far from clear, information on morphology and temperature has been used to infer the underlying energetics of the structures.  An area where work is still in its infancy is that of understanding the triggering of radio sources, and the possible r\\^ole played here by galaxy and cluster mergers in promoting or inhibiting radio-source development (\\S\\ref{sec:mergers}).  The emerging picture shows that very different accretion structures can host radio jets, with a tendency for quasar-type nuclei to be associated with more powerful jets.  How jets are powered by these different accretion structures and gas infall, and the duration of a given mode relative to typical lifetimes of radio sources, remain to be better understood.  The future is bright.  \\chandra\\ and \\xmm\\ are now mature observatories.  Operational experience is enabling both more ambitious and more speculative programs to be undertaken.  For example,  \\chandra\\ is completing sensitive exposures of all 3CRR radio sources within a redshift of 0.1, and a large shallow survey of quasar jets to study the X-ray-emission mechanism in a statistical sense and seek out more sources for deep, detailed study.  Observations of a somewhat more speculative nature are also being made, such as observing radio sources of different inferred ages, and studying how galaxy and cluster mergers are impacting the radio-source structures and their influence on the surrounding atmospheres.  These are just examples.  At the same time, \\suzaku\\ is making spectral measurements of active-galaxy nuclei, and testing the spin characteristics of black holes hosting radio sources through searching for relativistic broadening in Fe lines.   We can expect fantastic results from continuing X-ray work, and many surprises.  New facilities coming on line will enrich the X-ray results. \\spitzer\\ has measured dust, stars, and non-thermal cores in the centres of radio galaxies, placing constraints on the central structures.  It has also detected a number of kpc-scale jets, helping to tie down the all-important breaks in the spectral distributions of the synchrotron radiation that are likely to be connected to the process of particle acceleration.  \\herschel\\ will continue such work.  The characteristics of the non-thermal emission at energies higher than the X-ray provide a sensitive test of emission mechanisms and a probe of jet composition.  The Fermi Gamma-ray Space Telescope is providing such data, particularly for the embedded small-scale jets of highly-beamed quasars and BL~Lac objects, as are ground-based Cerenkov telescopes sensitive to TeV emission.  \\alma\\ will probe the cool component of gas in active galaxies, and provide information on one possible component of accretion power. Radio measurements with \\emerlin\\ and \\evla\\ will probe spatial scales intermediate between pc and kpc, important in the launching and collimation of jets.  They will also provide improved information on transverse jet structure.  Extending polarimetry to the X-ray, as is under study in the community, will provide key tests of jet emission and acceleration mechanisms, just as such work with \\hst\\ is starting to do in the optical.  Most importantly, a future X-ray observatory that has the sensitivity and spectral resolution to probe gas dynamics associated with radio sources is crucial for confirming and extending source modelling that is currently in its infancy.  Such capabilities will come with the launch of a new facility such as the International X-ray Observatory currently under study by ESA and NASA."
    },
    "0812/0812.4198_arXiv.txt": {
        "abstract": "The technique of \\gr\\ astronomy at very high energies (VHE:$>100$~GeV) with ground-based imaging atmospheric Cherenkov telescopes is described, the \\hess\\ array in Namibia serving as example. Mainly a discussion of the physical principles of the atmospheric Cherenkov technique is given, emphasizing its rapid development during the last decade. The present status is illustrated by two examples: the spectral and morphological characterization in VHE \\grs\\ of a shell-type supernova remnant together with its theoretical interpretation, and the results of a survey of the Galactic Plane that shows a large variety of non-thermal sources. The final part is devoted to an overview of the ongoing and future instrumental developments. ",
        "introduction": "\\label{intro}  At \\gr\\ energies $E_{\\gamma}~ \\gsim~ 30$~MeV electron-positron pair production in the Coulomb field of an atomic nucleus dominates the interactions of photons with normal matter. Cosmic \\grs\\ of such energies can therefore be most effectively measured in either gas-filled or solid state satellite detectors above the atmosphere registering the tracks of the resulting $e^{\\pm}$-pairs, or using the Earth's atmosphere itself as detector in observing secondary Cherenkov radiation with ground-based telescopes. Satellite detectors work from $\\sim 30$~MeV up to some tens of GeV, where they become statistics-limited as a result of the limited size of detectors that can be transported by rockets into space. In the Very High Energy range (VHE; $E_{\\gamma} > 100$~GeV) large optical telescopes are used, often with diameters of 10~m or more. \\begin{figure*} \\includegraphics[width=\\textwidth]{HESS_array.eps} \\caption[The \\hess\\ telescope array]{The \\hess\\ telescope array \\cite{hesspic} in the Khomas Highland of Namibia at an elevation of 1800 m a.s.l. The four 13~m telescopes stand at the corners of a square with a side length of 120 m.} \\label{fig:1}       % \\end{figure*}  The history, development and results of satellite \\gr\\ detectors have been described in this volume by K. Pinkau. Addressing the complementary ground-based technique, we shall in the following summarize the physics of UV/Optical imaging atmospheric Cherenkov telescopes which register the Cherenkov radiation from the atmospheric secondary particles. Subsequently we shall discuss two important astrophysical results that demonstrate the power of this observation method. Given our personal involvement in the \\hess\\ (High Energy Stereoscopic System) array in Namibia we shall in most instances use this experiment as example (see Fig.~\\ref{fig:1}), although we shall also emphasize the other large arrays existing worldwide. The scientific results are astonishing, given the slow development of the field over the three decades before the mid-1990s. We conclude with a glimpse at the next generation European project that envisages a further sensitivity increase by an order of magnitude. ",
        "conclusions": ""
    },
    "0812/0812.4151_arXiv.txt": {
        "abstract": "{ Over the last decade, X-ray observations unveiled the existence of several classes of isolated neutron stars (INSs) which are radio-quiet or exhibit radio emission with properties much at variance with those of ordinary radio pulsars. The identification of new sources is crucial in order to understand the relations among the different classes and to compare observational constraints with theoretical expectations. A recent analysis of the 2XMMp catalogue provided less than 30 new thermally emitting INS candidates. Among these, the source \\zersixfiv\\ appears particularly interesting because of the softness of its X-ray spectrum, $kT=117\\pm14$\\,eV and $\\nh=(3.5\\pm1.1)\\times10^{21}$\\,cm$^{-2}$ (3\\,$\\sigma$), and of the present upper limits in the optical, $m_{\\rm B}\\ga26$, $m_{\\rm V}\\ga25.5$ and $m_{\\rm R}\\ga25$ (98.76\\% confidence level), which imply a logarithmic X-ray-to-optical flux ratio $\\log(F_{\\rm X}/F_{\\rm V})\\ga3.1$, corrected for absorption. We present the X-ray and optical properties of \\zersixfiv\\ and discuss its nature in the light of two possible scenarios invoked to explain the X-ray thermal emission from INSs: the release of residual heat in a cooling neutron star, as in the seven radio-quiet \\ros-discovered INSs, and accretion from the interstellar medium. We find that the present observational picture of \\zersixfiv\\ is consistent with a distant cooling INS with properties in agreement with the most up-to-date expectations of population synthesis models: it is fainter, hotter and more absorbed than the seven \\ros\\ sources and possibly located in the Carina Nebula, a region likely to harbour unidentified cooling neutron stars. The accretion scenario, although not entirely ruled out by observations, would require a very slow ($\\sim$\\,10\\,km\\,s$^{-1}$) INS accreting at the Bondi-Hoyle rate. } ",
        "introduction": "A major outcome of the \\ros\\ mission has been the discovery of a group of radio-quiet isolated neutron stars (INSs) with properties clearly at variance from the bulk of the population of rotation-powered pulsars. At present, the group comprises seven objects (whence the nickname ``Magnificent Seven'', or \\msev) sharing similar properties, which include purely thermal X-ray spectra with very soft temperatures ($kT\\sim40$--100\\,eV) and low absorption column densities ($\\nh\\sim$ few $10^{20}$\\,cm$^{-2}$), long spin periods ($P\\sim3$--10\\,s), intense magnetic fields ($B\\sim10^{13}$--10$^{14}$\\,G), very faint optical counterparts ($m_{\\rm B}\\ga25$ when detected, implying X-ray-to-optical flux ratios $F_{\\rm X}/F_{\\rm{opt}}\\ga10^4$) and no associations with supernova remnants or stellar companions \\citep[see e.g.][for a recent review]{hab07}.  For some time, following the pioneering work by \\cite{ors70} and \\cite{shv71}, it was thought that these sources might belong to the long sought population of old INSs powered by accretion of the interstellar medium (ISM), which was expected to be detected by \\ros\\ \\citep[see][for a review]{tre00}. However, the measurement of large proper motions \\cite[e.g.][and references therein]{mot07,mot08} strongly supports the idea that the \\msev\\ represent a local population ($d\\la500$\\,pc, \\citealt{pos07,ker07a}) of middle-aged neutron stars ($\\sim$\\,10$^5$--10$^6$\\,yr), which give off thermal radiation as they cool down. The similar spin periods and their position on the $P-\\dot{P}$ diagram raised the possibility that they could be related to other populations of INSs -- like magnetars, high magnetic field radio pulsars (HBPSRs) and rotating radio transients (RRATs; \\citealt{lau06}) -- which is an intriguing and still debated issue \\citep[e.g.][]{pop06,kap08,kea08}.  \\begin{table*} \\begin{center} \\begin{tabular}{l c c c c c c c l c} \\hline\\hline OBSID & Observation Date & Detector & Filter & $t_{exp}$ & $\\theta$ & $t_{eff}$ & Counts & Rate & Near Gap?\\\\ &                  &          &        & (ks)      & (arcmin) & (ks)      &        & (10$^{-2}$\\,s$^{-1}$) & \\\\ \\hline 112580601 & 2000-07-26 & pn & thick  & 27.7 & 9.7  &  15.0 & 558  & $3.73\\pm0.18$ & yes\\\\ &            & M1 & thick  & 33.2 & 8.7  &  17.8 & 244  & $1.37\\pm0.10$ & no \\\\ &            & M2 & thick  & 30.1 & 9.1  &  15.7 & 198  & $1.26\\pm0.10$ & yes\\\\ 112580701 & 2000-07-27 & pn & thick  & 8.0  & 9.7  &   4.3 & 161  & $3.7\\pm0.3  $ & yes\\\\ &            & M1 & thick  & 10.9 & 8.7  &   5.9 &  88  & $1.50\\pm0.19$ & no \\\\ &            & M2 & thick  & 7.9  & 9.1  &   4.1 &  41  & $1.00\\pm0.18$ & yes\\\\ 112560101 & 2001-06-25 & pn & thick  & 21.5 & 15.8 &   6.7 & 291  & $4.4\\pm0.3  $ & yes\\\\ 112560201 & 2001-06-28 & pn & thick  & 20.5 & 15.9 &   6.3 & 200  & $3.2\\pm0.3  $ & yes\\\\ 145740101 & 2003-01-25 & M1 & thick  & 6.9  & 8.3  &   4.6 &  78  & $1.68\\pm0.22$ & yes\\\\ 145740201 & 2003-01-27 & M1 & thick  & 6.9  & 8.3  &   4.6 &  73  & $1.59\\pm0.21$ & yes\\\\ &            & M2 & thick  & 6.9  & 8.0  &   4.8 &  69  & $1.44\\pm0.20$ & yes\\\\ 145740301 & 2003-01-27 & M1 & thick  & 6.8  & 8.3  &   4.5 &  59  & $1.30\\pm0.20$ & yes\\\\ &            & M2 & thick  & 6.8  & 8.0  &   4.7 &  68  & $1.44\\pm0.19$ & yes\\\\ 145740501 & 2003-01-29 & M2 & thick  & 6.9  & 7.9  &   4.8 &  67  & $1.41\\pm0.19$ & yes\\\\ 160160101 & 2003-06-08 & M1 & thick  & 16.5 & 8.5  &   8.8 & 127  & $1.44\\pm0.15$ & yes\\\\ &            & M2 & thick  & 15.5 & 8.2  &   8.6 & 114  & $1.33\\pm0.15$ & no \\\\ 160160901 & 2003-06-13 & M1 & thick  & 31.1 & 8.5  &  16.6 & 207  & $1.25\\pm0.10$ & no \\\\ &            & M2 & thick  & 31.1 & 8.3  &  17.1 & 217  & $1.27\\pm0.10$ & no \\\\ 145780101 & 2003-07-22 & M1 & thick  & 8.4  & 9.4  &   4.6 &  76  & $1.66\\pm0.21$ & no \\\\ 160560101 & 2003-08-02 & M2 & medium & 11.8 & 9.1  &   6.3 & 121  & $1.93\\pm0.19$ & no \\\\ 160560201 & 2003-08-09 & M1 & thick  & 12.2 & 8.6  &   6.8 & 105  & $1.54\\pm0.17$ & no \\\\ &            & M2 & medium & 12.1 & 9.1  &   6.4 &  96  & $1.50\\pm0.18$ & no \\\\ 160560301 & 2003-08-18 & M1 & thick  & 18.5 & 8.6  &  10.8 & 155  & $1.42\\pm0.13$ & no \\\\ &            & M2 & medium & 18.5 & 9.4  &  10.0 & 160  & $1.60\\pm0.14$ & no \\\\ 206010101 & 2004-12-07 & pn & medium & 19.3 & 17.3 &   5.8 & 366  & $6.3\\pm0.4  $ & no \\\\ 311990101 & 2006-01-31 & pn & thick  & 24.3 & 7.7  &  15.9 & 837  & $5.27\\pm0.26$ & no \\\\ &            & M2 & thick  & 65.3 & 8.2  &  44.2 & 731  & $1.65\\pm0.15$ & no \\\\ 9488      & 2008-09-05 & ACIS-I & OBF & 60.0 & 6.7 & 57.0  & 567  & $0.95\\pm0.04$ & no \\\\ \\hline \\end{tabular} \\caption{Description of the set of \\xmm\\ and \\chan\\ observations that serendipitously detected source \\candzsf. The exposure times ($t_{exp}$), filtered for background flares, and the effective exposures ($t_{eff}$) on \\candzsf, accounting for vignetting, are reported; $\\theta$ is the source off-axis angle. Net source photons and the on-axis count rates are in the 0.15--3\\,keV (\\xmm) and 0.5--3\\,keV (\\chan) energy bands.\\label{tab_Xdata}} \\end{center} \\end{table*} Considering that within 1\\,kpc the \\msev\\ appear in comparable numbers as young ($\\la$ few Myr) radio and $\\gamma$-ray pulsars \\citep{pop03}, they may represent the only identified members of a large, yet undetected, elusive population of radio-quiet and thermally emitting INSs. The discovery of new candidates is then mandatory in order to make any progress towards an understanding of their properties as a population and of their relations with other classes of Galactic INSs.  Despite the relatively small field of view and low sky coverage of the \\xmm\\ Observatory\\footnote{In a near future, \\eROS\\ will become an excellent tool to search for INSs (\\texttt{http://www.mpe.mpg.de/projects.html\\#erosita}).}, its large effective area and good positional accuracy at soft X-ray energies make it ideal to look for faint INS candidates. Recently, \\cite{pir08} reported on the preliminary results of a programme aimed at identifying new thermally emitting INSs in the 2XMMp catalogue\\footnote{\\texttt{http://xmmssc-www.star.le.ac.uk/Catalogue/xcat\\_\\\\public\\_2XMMp.html}}. This version of the catalogue, released on July 2006, contains more than 120 thousand sources, covering a total sky area of $\\sim$\\,285 deg$^2$. Out of the considered $\\sim$\\,7.2$\\times10^4$ EPIC pn sources with count rates above 0.01\\,s$^{-1}$, less than 30 good candidates met all the selection criteria -- i.e. well detected point-like sources showing soft thermal spectrum ($kT\\le200$\\,eV and $\\nh=10^{19}-10^{22}$\\,cm$^{-2}$), with no optical candidates and no identifications in over 170 astronomical catalogues. The brightest and most promising INS candidate is source \\zersixfiv\\ (hereafter \\candzsf). In this work we discuss the properties of \\candzsf\\ in connection with the two main options concerning its nature, a cooling or an accreting INS. The description of the selection procedure and the results of follow-up optical observations on a sample of the X-ray brightest INS candidates will be presented elsewhere (Pires et al., in preparation). ",
        "conclusions": "Overall, the present observational picture of \\candzsf\\ suggests that it has the right properties to be identified with a younger and more distant thermally emitting INS, possibly the eighth member of the \\ros-discovered \\msev. The analysis of its X-ray emission, although based on archival data obtained with non-optimal configurations, reveals an intrinsically soft energy distribution, possibly variable on long time scales. However, uncertainties on the calibration accuracy of the large off-axis angles at which the source was observed may account for the variations of the derived spectral parameters. Its mean temperature, column density and 0.15--3\\,keV observed flux are $kT=117\\pm14$\\,eV, $\\nh=(3.5\\pm1.1)\\times10^{21}$\\,cm$^{-2}$ and $f_{\\rm X}=(1.03\\pm0.06)\\times10^{-13}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$, where we adopted the weighted means of the subsample of the 10 best \\xmm\\ observations. The present optical limits imply a logarithmic X-ray-to-optical flux ratio greater than $\\sim$\\,3.1, already high enough to safely exclude standard classes of X-ray emitters (late-type stars, AGN and CVs). The column density is consistent with the source being located within the Carina Nebula, roughly at a distance of $\\sim$\\,2\\,kpc. Assuming that the X-ray emission is well described by the mean parameters given above, \\candzsf\\ may be more luminous and youngish than the \\msev\\ and thus perhaps still close to its birth place. \\candzsf\\ is unique in the sense that it is hotter than the seven nearby sources and may represent an evolutionary missing link between the different classes of magnetars, radio-transient and radio-quiet INSs. The accretion scenario would require an unlikely low-velocity neutron star, combined with high accretion efficiency, to account for the X-ray luminosity of \\candzsf. A systematic search for pulsations is crucial in order to establish the true nature of this X-ray source."
    },
    "0812/0812.1774_arXiv.txt": {
        "abstract": "{ We present $UBVI$ photometry of the old open cluster NGC 1193. Color-magnitude diagrams (CMDs) of this cluster show a well defined main sequence and a sparse red giant branch. For the inner region of $r<50\\arcsec$, three blue straggler candidates are newly found in addition to the objects Kaluzny (1988) already found. The color-color diagrams show that the reddening value toward NGC 1193 is $E(B-V) =0.19 \\pm 0.04$. From the ultraviolet excess measurement, we derived the metallicity to be [Fe/H]$=-0.45 \\pm 0.12$. A distance modulus of $(m-M)_0 =13.3 \\pm 0.15$ is obtained from zero age main sequence fitting with the empirically calibrated Hyades isochrone of Pinsonneault et al. (2004). CMD comparison with the Padova isochrones by Bertelli et al. (1994) gives an age of log $t =9.7 \\pm 0.1$. } ",
        "introduction": "Since the open clusters are composed of stars with almost the same age and metallicity, they can be used in testing the theoretical models for stellar evolution. The relatively close open clusters can be good targets for even small telescopes to get a good enough quality photometric data (e.g.: Kyeong, Byun, \\& Sung 2001; Kim \\& Sung 2003). In particular, the old open clusters are very useful in testing the stellar isochrones because they show clear evolutionary sequences.  There have been a few observational studies on NGC 1193, which has a condensed central region.  One is the $BV$ photometry of Kaluzny (1988) which is used to derive the basic parameters.  The other is Friel \\& Janes (1993)'s spectroscopic study of several stars of NGC 1193 to determine its metallicity.  In this paper, we describe the $UBVI$ CCD observations of the old open cluster NGC 1193, which is located at $\\alpha$=3$^h$ 05$^m$, $\\delta$=$+$44$^\\circ$ 23$'$ (J2000). Although the bright star in the western side hamper the wide field observation, we present accurate $UBVI$ photometry for this cluster and derive the basic parameters such as reddening, metallicity, distance modulus and age from the photometric data.  The robust determination of reddening is significant for the estimation of distance, and subsequently age and metallicity. So, we estimate the reddening through two-color diagram method ($U-B, B-V$) and other cluster parameters are then derived by various conventional methods.  Section II describes the observations and data reduction processes. The comparison with the previous photometric data is presented in Section III. Color-magnitude diagrams (CMDs) and blue straggler candidates are explained in Section IV. Cluster parameters such as the reddening, metallicity, distance modulus and age are derived in Section V.  Section VI summarizes our results.  \\begin{table*} \\begin{center} {\\bf Table 1.}~~Journal of Observations for NGC 1193 \\\\ \\begin{tabular}{ccccc} \\hline \\hline \\noalign{\\smallskip} Date &UT(Start) & Filter & ${\\rm T_{exp}}$(sec) &Airmass \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} & 11:06 &$U$ &600 $\\times$ 3 &1.06 \\\\ 2006 August 30 & 10:50 &$B$ &300 $\\times$ 3 &1.07 \\\\ & 11:38 &$V$ &200 $\\times$ 3 &1.04 \\\\ & 10:37 &$I$ &60 $\\times$ 3 &1.09 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\end{table*}  \\begin{figure*}% \\vskip 5mm \\epsfxsize=8.9cm \\epsfysize=8.9cm \\centerline{\\epsffile{fig1.ps}} {\\small {\\bf ~~~Fig. 1.}---Synthetic field (4$\\farcm$3$\\times$4$\\farcm$3) of NGC 1193 as observed in $V$ frame.  North is at the top and east is to the left. All the detected stars are plotted.  The X- and Y-axes are in pixels and 1 pixel corresponds to $0\\farcs25$. Circle and square sizes are proportional to the brightnesses. The large cross means the cluster center determined by Kaluzny (1988). The filled squares represent the blue straggler candidates (See the text), while their sizes are also proportional to their brightnesses. } \\end{figure*} ",
        "conclusions": "The goal of this paper is to present the $UBVI$ photometric data and derive the basic physical parameters of the old open cluster NGC 1193. The CMDs of NGC 1193 show the characteristics of the old open clusters, such as the sub giant branch and red giant branch.  The blue straggler candidates inside of $r=50\\arcsec$ from the center is identified and their magnitudes are also measured.  Using color-color diagrams, the reddening value is determined to be $E(B-V) =0.19 \\pm 0.04$, which is somewhat larger than the Kaluzny (1988)'s value 0.12. The metallicity of NGC 1193 determined by using the UV excess method is [Fe/H]$= -0.45 \\pm 0.12$, which is in the middle of the two previously values in the literature, $-0.50$ (spectroscopic data of four stars, Friel \\& Janes 1993), $-0.29$ ($BV$ isochrone fitting, Kaluzny 1998). Our derived distance modulus is $(m-M)_0 =13.3 \\pm 0.15$. This value is well matched with the Kaluzny's value of 13.4 within its error. Finally, from the Padova isochrone fitting to the observed CMDs, the age is estimated to be log $t =9.7 \\pm 0.1$. The summary of the fundamental parameters for NGC 1193 is presented in Table 4.  \\vspace{4mm}  We would like to thank the staff of Observatorio Astron\\'omico Nacional in the Sierra San Pedro M\\'atir (OAN-SPM) in Mexico for the use of observing facilities and their warm support during our observations. This research has made use of the SIMBAD database, operated by CDS, Strasbourg, France. This research has also made use of the WEBDA database, operated at the Institute for Astronomy of the University of Vienna."
    },
    "0812/0812.0291_arXiv.txt": {
        "abstract": "{}{In the approximation of linear dissipative magnetohydrodynamics (MHD) it can be shown that driven MHD waves in magnetic plasmas with high Reynolds number exhibit a near resonant behaviour if the frequency of the wave becomes equal to the local \\alf (or slow) frequency of a magnetic surface. This near resonant behaviour is confined to a thin region, known as the dissipative layer, which embraces the resonant magnetic surface. Although driven MHD waves have small dimensionless amplitude far away from the resonant surface, this near-resonant behaviour in the dissipative layer may cause a breakdown of linear theory. Our aim is to study the nonlinear effects in Alfv\\'en dissipative layer}{In the present paper, the method of simplified matched asymptotic expansions developed for nonlinear slow resonant waves is used to describe nonlinear effects inside the \\alf dissipative layer.} {The nonlinear corrections to resonant waves in the \\alf dissipative layer are derived and it is proved that at the \\alf resonance (with isotropic/anisotropic dissipation) wave dynamics can be described by the linear theory with great accuracy.}{} ",
        "introduction": "Magnetic fields are ubiquitously present in solar and space plasmas. For regions where plasma-beta (the ratio of the kinetic and magnetic pressures) is less than one, magnetism controls the dynamics, topology and thermal state of the plasma. The magnetic field in the solar atmosphere is not dispersed, but it tends to accumulate in thinner or thicker entities often approximated as magnetic flux tubes. These magnetic flux tubes serve as an ideal medium for guided wave propagation.  One particular aspect of the solar physics that has attracted much attention since the 1940s is the very high temperature of the solar corona compared with the much cooler lower regions of the solar atmosphere requesting the existence of some mechanism(s) which keeps the solar corona hot against the radiative cooling. One of the possible theories proposed is the transfer of omnipresent waves' energy into thermal energy by resonant absorption or resonant coupling of waves \\citep[see e.g.][]{poedts1990, sakurai1, goossens1995b}.  Waves which were initially observed sporadically mainly in radio wavelengths \\citep[see e.g.,][]{kai73, asch1992} are now observed in abundance in all wavelengths, especially in (extreme) ultraviolet \\citep[see e.g.,][]{deforest98, asch1999, nak99, robb01, king03, erd08, mar08}. Since the plasma is non-ideal, waves can lose their energy through transport processes, however, the time over which the waves dissipate their energy is far too long. In order to have an effective and localized energy conversion, the plasma must exhibit transversal inhomogeneities relative to the direction of the ambient magnetic field. It was recognized a long time ago that solar and space plasmas are inhomogeneous, with physical properties varying over length scales much smaller than the scales determined by the gravitational stratification. Homogenous plasmas have a spectrum of linear eigenmodes which can be divided into slow, fast and Alfv\\'{e}n subspectra. The slow and fast subspectra have discrete eigenmodes whereas the Alfv\\'{e}n subspectrum is infinitely degenerated. When an inhomogeneity is introduced the three subspectra are changed. The infinite degeneracy of the Alfv\\'{e}n point spectrum is lifted and replaced by the Alfv\\'{e}n continuum along with the possibility of discrete Alfv\\'{e}n modes occurring, the accumulation point of the slow magnetoacoustic eigenvalues is spread out into the slow continuum and a number of discrete slow modes may occur, and the fast magnetoacoustic point spectrum accumulates at infinity \\citep[see e.g.,][]{goedbloed1975, goedbloed1984}.  According to the accepted wave theories, effective energy transfer between an energy carrying wave and the plasma occurs if the frequency of the wave matches one of the frequencies in the slow or \\alf continua, i.e. at the slow or \\alf resonances. The \\alf resonance has been more frequently associated with heating of coronal structures given the low-$\\beta$ regime of the solar corona. Nevertheless, slow resonance cannot be ruled out as an additional source of energy transfer. From a mathematical point of view, a resonance is equivalent to regular singular points in the equations describing the dynamics of waves, but these singularities can be removed by, e.g., dissipation. Recently resonant absorption has acquired a new applicability when the observed damping of waves and oscillations in coronal loops has been attributed to resonant absorption. Hence, resonant absorption has become a fundamental constituent block of one of the newest branches of solar physics, called \\textit{coronal seismology} \\citep[see e.g.,][]{nak99, rudermanconf2002, goossens2002, arregui2007, ballai08, goossens08, terradas2008} when applied to corona and solar magneto-seismology when applied to the entire coupled solar atmosphere \\citep[see e.g.,][]{erdelyi07, verth07a, verth07b}.  Given the complexity of the mathematical approach, most theories describing resonant waves are limited to the linear regime. Perturbations, in these theories, are considered to be just small deviations from an equilibrium despite the highly nonlinear character of MHD equations describing the dynamics of waves and the complicated interaction between waves and plasmas. Initial numerical investigations of resonant waves in a nonlinear limit \\citep[see e.g.,][]{ofman1995} unveiled that the account of nonlinearity introduces new physical effects which cannot be described in the linear framework.  The first attempts to describe the nonlinear resonant waves analytically appeared after the papers by \\citet{ruderman1997b, ruderman3} which were followed by further analysis by, e.g., \\citet{Ballai1998a, Ballai1999}; \\citet{Ruderman2000}; \\citet{Clack2008}; however, all these papers focused on the slow resonant waves only. These studies revealed that nonlinearity does affect the absorption of waves. In addition, the absorption of wave momentum generates a mean shear flow which can influence the stability of resonant systems.  The present paper is the first analytical study on the nonlinear resonant \\alf wave, where we obtain governing equations using techniques made familiar from previous studies on nonlinear slow resonant MHD waves. Before embarking on the actual derivation, let us carry out a qualitative discussion. First of all we should point out that in plasmas with high Reynolds numbers (as in the solar corona) efficient dissipation only operates in a thin layer embracing the resonant surface. This layer is called \\emph{the dissipation layer}. This restriction on the effect of dissipation makes the problem more tractable from a mathematical point of view, as outside the dissipative layer the dynamics of waves is described by the ideal MHD. Dissipation is a key ingredient of the problem of resonance. As it was mentioned earlier dissipation removes singularities in mathematical solutions. From a physical point of view dissipation is important as it is the mechanism which relaxes the accumulation of energy at the resonant surface and eventually contributes to the global process of heating.  It is important to stress that the choice of dissipation has to be related to the very physics which is described as different waves are sensitive to different dissipative mechanisms. Due to the dominant role of the magnetic field in the solar corona, transport processes are highly anisotropic. Possible dissipation mechanisms acting in coronal structures can be described within the framework of Braginskii's theory \\citep{braginskii} as it was shown in applications by, e.g. \\citet{erdelyi1995, ofman1995, mocanu2008}. \\alf waves are incompressible and transversal (in polarization), therefore, it is sensible to adopt shear viscosity and magnetic resistivity. Despite both transport processes being described by rather small coefficients, the net effect of dissipation can be increased considerably when the dissipative coefficients are multiplied by large transversal gradients.  The paper is organized as follows. In the next section we introduce the fundamental equations and discuss the main assumptions. In Sect. III, we derive the governing equation for wave dynamics inside the \\alf dissipative layer. Section IV is devoted to calculating the nonlinear corrections at \\alf resonance. Finally, in Sect. V we summarize and draw our conclusions, pointing out a few applications and further studies to be carried out in the future. ",
        "conclusions": "In the present paper we have investigated the nonlinear behaviour of resonant \\alf waves in the dissipative layer in one-dimensional planar geometry in plasmas with anisotropic dissipative coefficients, a situation applicable to solar coronal conditions. The plasma motion outside the dissipative layer is described by the set of linear, ideal MHD equations. The wave motion inside the dissipative layer is governed by Eq. (\\ref{eq:clacky2}). This equation is linear, despite taking into consideration (quadratic and cubic) nonlinearity. The Hall terms of the induction equation in the perpendicular direction relative to the ambient magnetic field cancel each other out.  The nonlinear corrections were calculated to explain why Eq. (\\ref{eq:clacky2}), describing the nonlinear behaviour of wave dynamics, is always linear. We found that, in the second order of approximation, magnetoacoustic modes are excited by the perturbations of the linear order of approximation. These secondary waves act to counteract the small pressure and density variations created by the first order terms. In addition, these waves are not resonant in the \\alf dissipative layer. In the third order approximation the perturbations become Alfv\\'{e}nic, however, these perturbations are much smaller than those in the linear order of approximation. Equations (\\ref{eq:finalexpansionlarge}) and (\\ref{eq:finalexpansionsmall}) describe the expansion of large and small variables, respectively, and demonstrate that all higher order approximations of both large and small variables at the Alfv\\'{e}n resonance are smaller than the linear order approximation, provided condition (\\ref{eq:BASICASSUMPTION}) is satisfied. This condition ensures that the oscillation amplitude remains small inside the dissipative layer. From a naive point of view the linear theory is applicable as soon as the oscillation amplitude is small. The example of slow resonant waves clearly shows that this is not the case. The nonlinear effects become important in the slow dissipative layer as soon as $\\epsilon\\sim R^{-2/3}$, i.e. as soon as the oscillation amplitude in the dissipative layer, which is of the order of $\\epsilon R^{1/3}$, is of the order of $R^{-1/3}\\ll1$. We also found that any dispersive effect due to the consideration of ions' inertial length (Hall effect) is absent from the governing equation.  This calculation of nonlinear corrections to resonant \\alf waves in dissipative layers allows us to apply the already well-known linear theory for studying resonant \\alf waves in the solar corona with great accuracy, where the governing equation, jump conditions and the absorption of wave energy are already derived \\citep[see e.g.,][]{sakurai1, goossens1995b, erdelyi1998}.  It is interesting to note that this work can be transferred to isotropic plasma rather easily. Shear viscosity, supplied by Braginskii's viscosity tensor (see Appendix A), acts exactly as isotropic viscosity. Therefore, replacing $\\eta$ by $\\rho_{0_a}\\nu$ in Eq. (\\ref{eq:clacky2}) provides the required governing equation for resonant \\alf waves in isotropic plasmas. Moreover, the work on the nonlinear corrections presented in this paper is also unaltered by anisotropy. This implies that we can consider resonant \\alf waves in dissipative layers throughout the solar atmosphere and still use linear theory if condition (\\ref{eq:BASICASSUMPTION}) is satisfied."
    },
    "0812/0812.2294_arXiv.txt": {
        "abstract": " ",
        "introduction": "The sun's 22-year cycle of activity is most clearly seen in solar surface magnetic fields (see for example the review by \\inlinecite{2003A&ARv..11..287O} for a good discussion of the solar dynamo and solar magnetic fields). Although the Sunspot Number (SSN) is the most readily available indicator of the state of the solar cycle, the magnetic fields over the whole surface provide a more complete measure of the dynamo process. A variety of questions arise in using the surface magnetic fields to study the solar cycle -- which fields are most important, the weak general field or the strong field and associated sunspots, how do the field strengths change over various time scales from instabilities with changes in seconds or less to trends lasting many decades. The connection between photospheric magnetic fields and magnetic fields near earth or at interplanetary spacecraft requires knowledge of the field strength at the solar surface. Is the overall strength of the magnetic field stationary or does it have any multi-decade trends (cf: \\inlinecite{2002JGRA.107jSSH16A})? These surface fields are also needed to calculate the strength measured by interplanetary spacecraft such as {\\it Ulysses} \\cite{1999ApJ...520..871G}. The topology of the field in the solar surface regions introduces additional uncertainty since the field of one polarity lower in the atmosphere may be canceled by the opposite polarity before it can emerge into the corona and heliosphere \\cite{2007ApJ...671..936A}. Questions such as these involving magnetic field strength can only be addressed when we are confident in the quantitative interpretation of the measured quantity.  The 150-foot solar tower telescope on Mt.\\ Wilson has been dedicated to the synoptic measurement of solar surface magnetic fields using a Babcock magnetograph observing the Fe I spectral line at $\\lambda5250$\\AA.  Although other spectral lines and observing techniques offer advantages over these choices, the uniformity of the data set extending back to 1967 is a strong reason to leave the observing system as nearly unaltered as possible.  The application of these data to modeling the heliospheric magnetic field then depends on our ability to understand how to interpret the magnetograms.  An early step toward this goal was carried out by \\inlinecite{1972SoPh...22..402H} and subsequently extended by \\inlinecite{1992ASPC...26..265U} and \\inlinecite{2002ApJS..139..259U} using a cross-correlation of fields observed simultaneously from two or more pairs of spectral samples.  A compelling feature of the scatter diagrams from these simultaneous measurements is the linearity between the magnetic fields obtained from the different spectral samplings.  Each spectral sample measures a different altitude in the solar atmosphere and is influenced by different thermal properties of the atmosphere at that altitude.  If the altitude dependent part of the relationship were also to depend on the field strength at each point on the solar surface, the scatter diagram would not be confined to a straight-line band passing through the zero-zero point.  The observed straight line distributions found on the scatter diagram establish that variable filling factor is the only cause for varying field strength, all flux tubes being essentially identical.  The fact that the slope is not unity can be used to study the altitude dependence of the flux tubes.  Throughout this paper we will use the term ``magnetic field strength'' and usually mean the apparent field strength which is the product of a flux tube field strength and a filling factor.  This quantity is actually the magnetic flux per unit area and is the quantity we wish to measure.  The extensive literature devoted to the determination of the flux tube fields and the filling factors is of interest to us and is referenced where appropriate below.  However, questions of flux tube physics are of secondary importance to our objectives in this paper. Where the flux tube field strength is discussed, it is made explicit.  Otherwise, the term magnetic field strength should be read to mean the magnetic flux per unit area as measured in the observed pixels.  The essential step in our study of magnetic field strength is the choice of a comparison line.  Our objective is to learn the relationship between measured circular polarization of a relatively large pixel (12 arcseconds squared or 20 arcseconds squared) and the magnetic flux in that area.  We are less concerned with questions of the structure of magnetic flux tubes or the strength of the field within the flux tubes -- we only seek to learn the best way to estimate the product of the filling factor and the magnetic field in the flux tubes.  In addition, we need to be able to observe the line simultaneously with $\\lambda5250$\\AA. For this purpose the line at $\\lambda5233$\\AA\\ is ideally suited due to its large width and depth and due to the good linearity of its line wings.  For a related objective of determining the flux tube magnetic field strength the comparison line chosen by \\inlinecite{1973SoPh...32...41S} was $\\lambda5247$\\AA\\ to take advantage of the fact that the two lines have differing $g_{\\rm eff}$ values but very similar thermal response and line formation properties.  However in both this paper and the subsequent study by \\inlinecite{1978A&A....70..789F} the reference to $\\lambda5233$\\AA\\ was retained with the assumption that magnetic fluxes from this line are obtained directly from the polarization factors and the line profile.  An interesting citation to a private communication by Livingston in this last paper suggests that the magnetic fields derived from $\\lambda5233$\\AA\\ depend on position within this line.  We verify below that the deduced magnetic field strength for any particular area on the sun's surface that depends on the spectral sampling of $\\lambda5233$. Our previous results based on $\\lambda5233$\\AA\\ yielded a correction factor of 4.5 to be applied as a multiplier for the magnetic fields obtained from $\\lambda5250$\\AA\\ whereas \\inlinecite{1973SoPh...32...41S} found only a factor of two.  Our correction factor is larger due partly to the fact that the shifted Zeeman components at $\\lambda5250$\\AA\\ are shifted beyond the sampling pass-band of the MWO Babcock magnetograph.  We refer to the reduced apparent field due to the large line shift as the ``classical saturation effect'' since it depends only on the actual line profile, the strength of the flux tube magnetic fields and the spectral resolution of the observing system. ",
        "conclusions": "\\label{disc}  Our principal result is given above and provides a small correction to the procedure previously applied to the Mt.\\ Wilson Observatory magnetogram observations.  The previous description of this correction factor as a saturation correction is false.  The various interpretations on the polarimetry of the spectral lines used in our observations differ from one another due to processes involved with transfer of radiation through an inhomogeneous, magnetized atmosphere.  A full analysis of this problem is beyond the scope of the present investigation but is obviously needed to fully understand the changes in shape of the polarized spectral lines.  We base our recommended correction factor on the fact that the magnetic field deduced from the spectral line bisector is less variable than the field deduced from the monochromatic Stokes $V$ parameter.  We further adopt the magnetic field implied by the properties of the spectral line near its core on the grounds that this part of the line is closest to the outer edge of the solar atmosphere.  We find that the line profile properties are most stable in regions where the field strength is significantly larger than that found in the completely quiet sun.  The line profiles for the weakly magnetized regions are well determined by our system but they show substantial deviations from symmetry.  We suspect that this is a consequence of the reduced effect of adjacent flux tubes as a constraint on the flux tube structure in cases where the field is so weak that there are no adjacent flux tubes.  The unconstrained flux tubes can then adopt a wider range of structure including perhaps cases where the field lines return locally to the solar interior rather than extending out of the solar atmosphere.  The fact that the error of measurement of the fields is well below 1 gauss for our line profile system means that the measured profiles can be used to constrain models of flux tubes in regions where the field strength is very low.    Our determination that the differing deduced magnetic field strength for differing spectral sampling configurations has a general potential impact on the measurement of magnetic fields with other observing systems.  All programs provide maps which give a single magnetic field strength for each observed point.  Had the adjustment of the Mt.\\ Wilson fields been a consequence of a saturation effect unique to this system, the potential implications would also have been limited to the Mt.\\ Wilson system.  However, our determination that the effect comes from a spectral transfer effect means that other observing programs need to determine how to relate their field measurements to measurements made by other systems.  Even within a single observing programs, multiple deduced magnetic fields could occur depending on the exact specification of the algorithm bringing the spectropolarimetric data to a deduced magnetic field strength.  Examples where a direct comparison has been made on a pixel by pixel basis include \\inlinecite{2005ApJS..156..295T} and \\inlinecite{2008SoPh..tmp..125D}.  We are also in a position to update discussion from \\ref{profandmag} which compares our scale factor to the deductions of \\inlinecite{2003SoPh..213..213B}.  Those authors concluded from Advanced Stokes Polarimeter data that the MDI magnetic fields are $0.64\\pm0.013$ times the correct value.  After combining the factor from \\inlinecite{2005ApJS..156..295T} to the new corrected $\\lambda5233$\\AA\\ scale we find that the MDI magnetic fields are $0.619\\pm0.018$ times the recommended value at the center-to-limb distance used by \\inlinecite{2003SoPh..213..213B}. This agreement within errors of measurement gives support to our approach in the calibration of the $\\lambda5250$\\AA\\ magnetograms.    \\begin{acks} The authors thank J.\\ Harvey and J.\\ Todd Hoeksema for their comments on drafts of this paper.  This work has been supported by NASA, NSF and ONR through a series of grants.  Most recently support has come from the project NASA/HMI 16165880-26967-G at Stanford University and from NSF through grant NSF ATM-0517729. \\end{acks}"
    },
    "0812/0812.2893_arXiv.txt": {
        "abstract": "The proto-planetary Red Rectangle nebula is powered by HD~44179, a spectroscopic binary (P = 318 d), in which a luminous post-AGB component is the primary source of both luminosity and current mass loss.  Here, we present the results of a seven-year, eight-orbit spectroscopic monitoring program of HD~44179, designed to uncover new information about the source of the Lyman/far-ultraviolet continuum in the system as well as the driving mechanism for the bipolar outflow producing the current nebula.  Our observations of the H$\\alpha$ line profile around the orbital phase of superior conjunction reveal the secondary component to be the origin of the fast (max. v $\\sim560$ km s$^{-1}$) bipolar outflow in the Red Rectangle. The outflow was previously inferred on the basis of a single, broad H$\\alpha$ emission line profile. The variation of total H$\\alpha$ flux from the central H II region with orbital phase also identifies the secondary or its surroundings as the source of the far-ultraviolet ionizing radiation in the system. The estimated mass of the secondary ($\\sim0.94$ M$\\sun$) and the speed of the outflow suggest that this component is a main sequence star and not a white dwarf, as previously suggested. We identify the source of the Lyman/far-ultraviolet continuum in the system as the hot, inner region (T$_{max} \\ge 17,000$ K) of an accretion disk surrounding the secondary, fed by Roche lobe overflow from the post-AGB primary at a rate of about $2 - 5\\times10^{-5}$ M$\\sun$ yr$^{-1}$. The speed of the accretion-driven, bipolar outflow was found to be strongly modulated by the inter-component separation in the highly eccentric ($e = 0.37$) orbit, with maximum speeds occurring near periastron and minimum speeds found around apastron. The total luminosity of the accretion disk around the secondary is estimated to be at least 300~L$\\sun$, about 5\\% of the luminosity of the entire system. ",
        "introduction": "The Red Rectangle (RR) and its associated star HD~44179 are an extensively studied protoplanetary nebula system that is frequently discussed as a prototype for intermediate-mass stars traversing the brief evolutionary phase between the top of the asymptotic giant branch (AGB) and the onset of the planetary nebula stage, when the star's hot core begins to ionize the expelled circumstellar matter (van Winckel 2003, 2007; Siodmiak et al. 2008; Szczerba et al. 2007). The RR was discovered on the basis of its strong near- and mid-infrared excess (Price \\& Walker 1976) and it was subsequently identified with a compact bipolar nebula associated with the post-AGB star HD~44179,  earning  its name on the basis of its rectangular appearance on existing red photographic plates (Cohen et al. 1975). Subsequent high-resolution imaging at optical (Osterbart et al. 1997; Cohen et al. 2004; Vijh et al. 2006) and near-infrared wavelengths (Tuthill et al. 2002) as well as spectroscopic mapping (Schmidt \\& Witt 1991) revealed the RR as an X-shaped biconical nebula, emerging from an optically thick disk seen nearly edge-on. We will refer to this disk, resolved by HST, as the circumbinary disk, to distinguish it from the much smaller accretion disk around the photometric secondary in HD~44179, to be discussed later. The circumbinary disk was explored in detail by Men'shchikov et al. (2002); it is estimated to have inner and outer radii of 14 AU and 43,000 AU, respectively, with a thickness of 90 AU. The distance to the RR is uncertain, but recent detailed modeling of the stellar and nebular emission by Men'shchikov et al. (2002) places the object at a distance of 710 pc, which we adopt for the purpose of the present study.  Soon after the discovery of the RR and other similar objects, two different models emerged for the explanation of the bipolar morphology of such nebulae. Icke (1981, 2003) proposed a biconical outflow pattern above a luminous disk as the shaping mechanism, with radiation pressure as the driving force. On the other hand, Morris (1981, 1987) envisioned binary star systems at the centers of bipolar nebulae, with bipolar jets launched by an accretion disk surrounding the secondary as both the driving and shaping force for the nebula. We note that in the case of HD~44179 its nature as a spectroscopic binary was not established until 1995 by van Winckel et al. (1995). In recent years, the binary star model as the cause of bipolar nebulae was further developed by Mastrodemos \\& Morris (1998, 1999) and, independently, by Soker (2000, 2005). Two recent papers by Icke (2003) and Soker (2005) continue the long-standing debate about the shaping mechanism for the special case of the RR. Nordhaus \\& Blackman (2006) discuss three possible modes of the evolutionary outcome of the common-envelope phase of a close binary including an AGB star (Livio et al. 1979), one of which may describe an earlier stage of evolution of the RR (e.g. Akashi \\& Soker 2008), during which the outer circumbinary disk was formed.  Another controversy concerns the nature of the secondary star in the HD~44179 system. Following the initial identification of HD~44179 as a spectroscopic binary by van Winckel et al. (1995),  Waelkens et al. (1996) confirmed the orbital period of $318\\pm3$ days and reported the photometric variability of the system with the same period. The phase dependence of the brightness variation was attributed to the variable effective scattering angle under which the primary component of HD~44179 is being viewed by a process of scattering off dust in the upper layers of the optically thick circumbinary disk. This viewing geometry was subsequently confirmed by high-resolution imaging at optical (Osterbart et al. 1997; Cohen et al. 2004) and infrared wavelengths (Tuthill et al. 2002). An important consequence of this indirect view of HD~44179 is that its ``effective inclination'' is considerably smaller than otherwise implied by the near-edge-on orientation of the system, based on the appearance of the circumbinary disk. Waelkens et al. (1996) estimated an ``effective inclination'' of 35\\degr\\  for HD~44179. When coupled with the mass function of the single-line spectroscopic binary,  \\begin{equation} f(m_1,m_2)=\\frac{m_2^3\\ sin^3\\ i}{(m_1 +m_2)^2}=\\frac{(a_1\\ sin\\ i)^3}{P^2}=0.049 M\\sun , \\end{equation}  \\noindent where $P$ is the period  of the binary, $i$ is the inclination of the orbit, $m_1$ and $m_2$ are the masses of the primary and secondary stars, respectively, and $a_1$ is the semi-major axis of the observed primary, the mass of the secondary is larger than the current mass of the photometrically dominant post-AGB primary of the system. Waelkens et al. (1996) suggested, therefore, that the secondary is most likely a low-luminosity main sequence star of less than 1 M$\\sun$. However, in such a configuration neither the photometrically dominant post-AGB star (T$_{\\mathrm{eff}} < 8000$ K) nor the even cooler main-sequence companion star can explain the radio detection of a small, compact HII region at the center of the RR by Jura et al. (1997). Consequently, Men'shchikov et al. (2002) developed a detailed model for the RR system assuming that the secondary is a hot He white dwarf star, which could supply the required Lyman continuum flux for the observed photo-ionization of hydrogen (Jura et al. 1997) and of helium (Kelly \\& Latter 1995), as well as supply the far-UV continuum required for the excitation of the extended red emission (ERE; Witt et al. 2006) and the high-excitation CO emission (Sitko 1983; Sitko et al. 2008) observed in the RR. This scenario, however, implies a greatly more complicated evolutionary history of mass loss and exchange, and it is based upon an ``effective inclination'' of 79\\degr\\ (Men'shchikov et al. 2002), inconsistent with the results of Waelkens et al. (1996).  In this paper we present the results of a program of high-resolution spectroscopic observations, which provide complete phase coverage over eight orbits of the HD~44179 binary from 2001 to 2008. These observations provide the first direct confirmation of the Morris (1987) model for the shaping of the  RR through a jet emanating from the vicinity of the secondary star. The observed jet speed is consistent with the expected outflow from an accretion disk centered on a low-mass main sequence star, as proposed by Waelkens et al. (1996). If such an accretion disk is fed by the mass loss from the post-AGB component of HD~44179 through Roche lobe overflow, the resulting disk temperatures and disk luminosity can account fully for the Lyman continuum luminosity required for the production of the compact HII region at the center of the RR, observed by Jura et al. (1997).  We are reporting the details of our observations in \\S2 of this paper. In \\S3 we present the evidence supporting our claim that the source of the bipolar outflow as well as the source of the Lyman contininuum are located in the vicinity of the secondary, as it moves through its orbit. In \\S4 we discuss the mass of the secondary as well as the process of mass accretion from the post-AGB component. In \\S5 we examine the predictions of steady mass-accretion disk models, including the expected disk spectra and luminosities as a function of the mass transfer rate. This is followed by a discussion and conclusions in \\S6. ",
        "conclusions": "\\subsection{Binaries in Proto-Planetary Nebulae} Ongoing research into the shapes and shaping mechanisms of bipolar planetary nebulae and proto-planetary nebulae (Balick \\& Frank 2002; Soker and Livio 1994; Nordhaus \\& Blackman 2006; Akashi 2007) suggests that there may be a number of distinctly different mechanisms at work in different objects. It is also possible that different mechanisms may dominate the morphological evolution of a single object during different phases of its evolution (e.g. Waters et al. 1998). Our current result concerning the presence of a bipolar jet associated with the secondary star in HD~44179, which appears to be energized by mass accretion into a circumstellar disk through Roche lobe overflow on part of the post-AGB primary is illustrated as an artistic representation in Figure 11. We believe HD~44179 to be a system where a binary has survived an earlier common-envelope phase, which produced the outer circumbinary disk, within which the remaining semi-detached binary is now engaged in mass transfer. This particular configuration may be the result of a specific set of initial conditions and it may also be short-lived. This picture should, therefore, not be generalized too far and should not be applied indiscriminately to all bipolar nebulae (Allen et al. 1980).  Our data support the model proposed  by Morris (1987) for the RR and developed later by Soker (2000, 2005). This, however, may only explain the cause of the \\emph{present} outflow from the central cavity of the RR. We emphasize that the accretion disk around the secondary discussed here is much smaller than and physically distinct from the much larger circumbinary disk, which has been resolved in high-resolution near-infrared (Tuthill et al. 2002) and optical images (Osterbart et al. 1997; Cohen et al. 2004; Vijh et al. 2006). The presence of this outer disk, which is in Keplerian revolution around the central binary (Bujarrabal et al. 2005), very likely assists in restraining any current mass outflow into bipolar directions. This outer disk may have had its origin during an earlier common-envelope phase, one of the mechanisms discussed by Nordhaus \\& Blackman (2006). Both the current mass ejection process in the RR as well as a possible earlier common-envelope phase depend critically upon the binarity of HD~44179, and this particular aspect of the RR appears to be rather common among bipolar nebulae (e.g. Yamamura et al. 2000; de Ruyter et al. 2006; van Winckel et al. 2006). A recent \\emph{GALEX} search for ultraviolet excesses among AGB stars by Sahai et al. (2008) revealed detectable excesses in about 43\\% of their sample, suggesting the presence of either hot companions or accretion upon low-mass secondaries. This is about three times the confirmed binary frequency among central stars in planetary nebulae (De Marco 2006), and goes quite far in supporting the contention that binarity is required for the production of planetary nebulae.  \\subsection{Excitation of Extended Red Emission } The RR appears to be unique among known bipolar proto-planetary nebulae for the remarkably high intensity of extended red emission (ERE) (Schmidt et al. 1980; Witt \\& Boroson 1990), which is responsible for the X-shape morphology of the nebula at red wavelengths. Detailed structural studies of the RR (Schmidt \\& Witt 1991;  Cohen et al. 2004;  Vijh et al. 2006) have found the ERE in the RR to be most intense on nebular surfaces which are illuminated directly by the central source, consistent with an ERE excitation mechanism involving photo-excitation by far-ultraviolet photons with energies E $>$ 10.5 eV (Witt et al. 2006). A potentially embarrassing problem for this scenario is the fact that neither the post-AGB primary nor the main-sequence secondary in HD~44179 emits nearly enough radiation with E $>$ 10.5 eV to account for the total number of ERE photons observed from the RR. The suggestion that the secondary might be a hot He white dwarf advanced by Men'shchikov et al. (2002) appeared to be a welcome solution to this problem, but our current observations do not support this proposal. Now, the hot, innermost region of the proposed accretion disk with a spectrum as shown in Figure~\\ref{diskspectra} replaces the role of the white dwarf companion in providing the needed far-ultraviolet photons for the excitation of the ERE. Witt et al. (2006) estimate the total ERE luminosity in the RR, corrected for extinction, as L$_{\\mathrm{ERE}} \\sim 10$ L$\\sun$, and they proposed a two-stage process of creation of the ERE carrier through ionization of a precursor, followed by optical pumping with near-UV/optical photons to explain the production of ERE. The luminosity of far-UV photons (E $>$ 10.5 eV) obtained from the disk for a mass accretion rate of $2\\times10^{-5}$ M$\\sun$\\ yr$^{-1}$ shown in Figure~\\ref{diskchars}, with a spectrum as shown in Figure~\\ref{diskspectra}, would be higher than the luminosity of Lyman continuum photons of E $>$ 13.6 eV  and sufficient to accomplish the critical first stage of this process, leaving the main task of optical pumping of the ionized ERE carrier to the post-AGB primary with its comparatively large near-UV/optical luminosity. In order to reproduce the conditions leading to the high ERE luminosity observed in the RR, apparently three conditions must be met: a powerful source of far-ultraviolet photons for the production of the ERE carriers, requiring a high mass-accretion rate onto an accretion disk surrounding the secondary; a still more powerful source of near-UV/optical photons for the optical pumping of the ERE carriers, in our case the HD~44179 post-AGB primary; and a carbon-rich chemistry with C/O $>$ 1. There is new, strong evidence (Lagadec \\& Zijlstra 2008; Woitke 2006) showing that among AGB stars of luminosity L $< 10^4$ L$\\sun$ and low metallicity only C-rich stars achieve mass-loss rates of $10^{-5}$ M$\\sun$\\ yr$^{-1}$ or above. These are all conditions that apply to HD~44179 and may point to the reason for the apparently unique prominence of ERE in the nebula associated with this star.  \\subsection{Conclusions} We summarize the findings of this investigation as follows: \\begin{enumerate} \\item RV measurements spanning 8 orbits during the years between 2001 and 2008 confirmed the orbital period of 318 days originally published by Waelkens et al. (1996) and constrain the uncertainty to $\\pm$1 day. They do not support the period of 322 days reported later by Men'shchikov et al. (2002). \\item The emitting gas at the center of the RR characterized by narrow emission lines, e.g. by the narrow H$\\alpha$ core and the KI 7699~\\AA\\ line, does not participate in the orbital motions of the binary components in HD~44179. Instead, the narrow-line gas is nearly at rest with respect to the system. This agrees with earlier results by Hobbs et al. (2004). \\item The complex H$\\alpha$ emission line profile undergoes systematic changes that are periodic with the orbital period of 318 days. Five separate components are found to contribute to the H$\\alpha$ profile: (a) a narrow (18 km s$^{-1}$ FWHM) emission spike at rest with respect to the RR system; (b) a broad emission plateau (FW $\\sim 200$ km s$^{-1}$) moving with the orbital motion of the secondary while simultaneously changing its width from a maximum near periastron to a minimum near apastron; (c) weak, extended emission wings indicating maximum outflow velocities of $\\sim560$ km s$^{-1}$; (d) a modulation of (a) and (b) by an underlying H$\\alpha$ absorption line arising in the atmosphere of the post-AGB primary, with the orbital period of 318 days; and (e) a strongly blue-shifted H$\\alpha$ absorption appearing near the phase of superior conjunction, when the lines of sight from the observer to the post-AGB primary pass over the secondary. \\item We interpreted these spectroscopic observations as indicating that the fast bipolar outflow in the RR originates in the vicinity of the secondary star. The speed of the outflow as well as a re-assessment of the likely mass of the secondary in light of the viewing geometry of HD~44179 proposed by Waelkens et al. (1996) suggest that the secondary most likely is a main-sequence star of 0.94 M$\\sun$. The energy and momentum seen in the fast outflow appear to be generated by mass accretion from the mass-losing post-AGB primary onto the secondary. Our interpretation of the outflow supports the early model of Morris (1987). Our conclusions concerning the nature of the secondary support an earlier suggestion by Waelkens et al. (1996) but conflict with Men'shchikov et al. (2002), who proposed a secondary in the form of a low-mass helium white dwarf. \\item We found the H$\\alpha$ spike emission flux to vary with the orbital period. The highest flux was observed near the phase of inferior conjunction for the primary, when the secondary is on the far side of its orbit. This suggests that the source of the ionizing radiation in HD~44179 travels with the secondary. The large amplitude of the H$\\alpha$ spike emission flux as a function of phase can be understood in view of the large shadowing effect for ionizing radiation caused by the large disk of the primary. \\item We estimated the mass accretion rate needed to produce the disk luminosity and temperatures to generate the Lyman-continuum luminosity required to maintain the central compact HII region in the RR, detected at radio-continuum frequencies by Jura et al. (1997). We employed a steady accretion disk model and found that a mass accretion rate of $\\sim2\\times10^{-5}$ M$\\sun$\\ yr$^{-1}$ is sufficient to produce the needed photon flux, although more realistic modeling of the Lyman continuum portion of the disk spectrum may push the mass accretion rate higher, to $\\sim5\\times10^{-5}$ M$\\sun$\\ yr$^{-1}$. Such a disk is expected to have a luminosity of about 300 L$\\sun$, spread out over a broad spectrum ranging from the Lyman continuum through the optical and the near infrared, as illustrated in Figure 9. The predicted disk luminosity is about 5\\% of the luminosity of the post-AGB primary. \\item The estimated mass accretion rate for the secondary is in excellent agreement with observationally determined mass-loss rates of post-AGB stars in the literature. This suggests that the mass transfer in the HD~44179 system occurs with very high efficiency through Roche lobe overflow. This mechanism is supported also by the estimated size of the primary, suggesting that it fills its Roche lobe; the high orbital eccentricity, which can be maintained by enhanced mass overflow at the phase of periastron; and by the observation that shows that the speed of the bipolar outflow originating near the secondary decreases by close to a factor two between the time of periastron and apastron. This suggests that the liberated accretion energy varies with the orbital period in response to a mass overflow rate, which is highest near periastron. \\item We show that the far-UV luminosity of the accretion disk is expected to be unaffected by the modulation in the mass accretion rate with orbital phase, in view of a long radial drift time scale. The momentum carried by the jets in HD~44179 greatly exceeds the radiative momentum available from the disk's luminosity. We conclude that the jets must be driven and collimated by magneto-hydrodynamical processes occurring over a large extent of the accretion disk. \\item A unique feature of the RR is the intense extended red emission (ERE), which requires far-UV radiation for its initiation, in addition to the presence of large carbonaceous molecules (Witt et al. 2006). The proposed accretion disk around the secondary generates enough far-UV radiation to power the ERE observed in the RR. Other systems at a similar evolutionary stage may have a lower mass transfer rate, generating insufficient UV radiation (Figure~\\ref{diskchars}) or lack the required carbon-rich chemistry and thus lack the preconditions for ERE, making the RR relatively unique in this regard. \\end{enumerate}"
    },
    "0812/0812.3187_arXiv.txt": {
        "abstract": "We model the evolution of the abundances of light elements in carbon-enhanced metal-poor (CEMP) stars, under the assumption that such stars are formed by mass transfer in a binary system. We have modelled the accretion of material ejected by an asymptotic giant branch star on to the surface of a companion star. We then examine three different scenarios: one in which the material is mixed only by convective processes, one in which thermohaline mixing is present and a third in which both thermohaline mixing and gravitational settling are taken in to account. The results of these runs are compared to light element abundance measurements in CEMP stars (primarily CEMP-$s$ stars, which are rich in $s$-processes elements and likely to have formed by mass transfer from an AGB star), focusing on the elements Li, F, Na and Mg. None of the elements is able to provide a conclusive picture of the extent of mixing of accreted material. We confirm that lithium can only be preserved if little mixing takes place. The bulk of the sodium observations suggest that accreted material is effectively mixed but there are also several highly Na and Mg-rich objects that can only be explained if the accreted material is unmixed. We suggest that the available sodium data may hint that extra mixing is taking place on the giant branch, though we caution that the data is sparse. ",
        "introduction": "Carbon-enhanced metal-poor (CEMP) stars are defined as stars with [C/Fe]\\footnote{[A/B] = $\\log (N_\\mathrm{A}/N_\\mathrm{B}) - \\log (N_\\mathrm{A}/N_\\mathrm{B})_\\odot$}$>$+1.0 \\citep{2005ARA&A..43..531B}, with [Fe/H]$\\ <-2$ in most cases. These objects appear with increasing frequency at low metallicity \\citep{2006ApJ...652L..37L}. The study of CEMP stars is being used to probe conditions in the early universe. For example, CEMP stars have been used to infer the initial mass function in the early Galaxy \\citep[e.g.][]{2005ApJ...625..833L}. Chemical abundance studies have revealed that the majority of the CEMP stars are rich in $s$-process elements like barium \\citep{2003IAUJD..15E..19A}, forming the so-called CEMP-$s$ group. Recent survey work has detected radial velocity variation in around 68\\% of these CEMP-$s$ stars and this is consistent with them all being in binary systems \\citep{2005ApJ...625..825L}.  Binary systems provide a natural explanation for these objects, which are of too low a luminosity to have been able to produce their own carbon. The primary\\footnote{This is the initially more massive star in the system; the secondary is the initially less massive star.} of the system was an asymptotic giant branch (AGB) star which became carbon-rich through the action of third dredge-up (the deepening of the convective envelope into regions of the star where material has undergone nuclear burning, see e.g. \\citealt{1983ARA&A..21..271I}) and transferred material on to the low-mass secondary (most likely via a stellar wind). The primary became a white dwarf and has long since faded from view, with the carbon-rich secondary now being the only visible component of the system.  It has commonly been assumed that accreted material remains unmixed on the surface of the recipient star during the main sequence and only becomes mixed with the stellar interior during first dredge-up. However, \\citet{2007A&A...464L..57S} pointed out that this picture neglects thermohaline mixing. Thermohaline mixing is the process that occurs when the mean molecular weight of the stellar gas increases toward the surface. A gas element displaced downwards and compressed will be hotter than its surroundings. It will therefore lose heat, become denser and continue to sink. This leads to mixing on thermal timescales until the molecular weight difference is eliminated \\citep{1972ApJ...172..165U,1980A&A....91..175K}.  Accreted material has undergone nuclear processing in the interior of the companion AGB star and hence has a higher mean molecular weight than the pristine material of the companion. The accreted AGB material should therefore become mixed with the pristine material of the secondary because of the action of thermohaline mixing. In their models, \\citet{2007A&A...464L..57S} showed that thermohaline mixing could, under optimal circumstances, lead to the accreted material being mixed with nearly 90 per cent of the secondary.  The extent to which accreted material is mixed into the accretor during the main sequence has been questioned. Using a sample of barium-rich CEMP stars (i.e. a set of stars where we expect the AGB mass transfer scenario to hold), \\citet{2008ApJ...678.1351A} showed that the distribution of [C/H] values in turn-off stars (i.e. those stars that have reached the end of their main-sequence lives, are still of low-luminosity and have yet to become giants) was different from that in giants suggesting that significant mixing only happened at first dredge-up. A similar point was made by \\citet{2008ApJ...679.1541D} using the data of \\citet{2006ApJ...652L..37L}. These authors showed that the abundance patterns of turn-off stars and giants were consistent with coming from the same distribution if first dredge-up resulted in the [C/H] value (and also the [N/H] value) being reduced by around 0.4 dex. This drop in both C and N is the result of the accreted material being mixed with pristine stellar material during the deepening of the stellar envelope at first dredge-up. They find that this result is consistent with having an accreted layer of material mixed to an average depth of about 0.2\\ms\\ (or alternatively having an accreted layer of 0.2\\ms\\ that remains unmixed until first dredge-up). \\citet{2008ApJ...677..556T} also questioned the efficiency of mixing in the binary system CS~22964-161. This system consists of two turn-off stars which are both carbon-rich. The system is also found to be lithium-rich, with $\\epsilon(\\mathrm{Li})=+2.09$\\footnote{$\\epsilon(\\mathrm{Li})=\\log_{10}\\left({N_\\mathrm{Li}\\over N_\\mathrm{H}}\\right)$+12 where $N$ is the number abundance of the element.}. As lithium is such a fragile element and is easily destroyed at temperatures above around $10^6$\\,K, any extensive mixing of accreted material will result in lithium being depleted. \\citet{2008ApJ...677..556T} suggest that gravitational settling could in principle inhibit the action of thermohaline mixing as helium diffusing away from the stellar surface would create a stabilising mean molecular weight gradient (a so-called `$\\mu$-barrier').  This idea was tested by \\citet{2008MNRAS.389.1828S}, who included the physics of both thermohaline mixing and gravitational settling in their models of stars accreting material from a putative AGB donor. Their grid of models over a range of accreted masses  (0.001-0.1\\ms) from donors of different initial mass (1-3.5\\ms) showed that gravitational settling could only seriously inhibit thermohaline mixing if a small quantity of material had been accreted. \\citet{2008MNRAS.389.1828S} compared their predicted surface abundances of carbon and nitrogen to those of observed CEMP stars, concluding that none of their model sets did a good job of reproducing the observed abundances patterns. The aim of this work is to extend the secondary models of SG08 to a greater range of isotopes (rather than just the C and N looked at in that paper) in the hope that these additional isotopes may help to illuminate the processes that happen in the AGB mass transfer formation scenario.  Recently, there has been considerable progress in measuring the abundances of light elements in carbon-enhanced metal-poor stars. \\citet{2008ApJ...678.1351A} have provided a sample of over 70 barium-rich CEMP stars that have other abundance determinations, including for sodium and magnesium. The element fluorine has also been measured for the first time in a CEMP star \\citep{2007ApJ...667L..81S}. Each element that we have measurements for potentially gives us another way to explore both the nature of the donor AGB stars and what happens to the material that is accreted on to the secondary star. This paper looks at what the light elements in CEMP stars can tell us about these topics. The AGB mass transfer scenario is likely to apply to the majority of CEMP stars (specifically those with $s$-process enrichments which are classified as CEMP-$s$ stars), but not all CEMP stars form this way. In addition, the elements discussed in this paper do not necessarily come from AGB stars (for example, Na and Mg in the Galaxy mostly come from supernovae). The reader should bear these points in mind throughout. ",
        "conclusions": "The measurement of a high Li abundance in an unevolved CEMP star precludes the possibility of {\\it extensive} mixing of accreted material. This could either be because there is no thermohaline mixing or because only a small quantity of material is accreted (likely a few thousandths of a solar mass). The inclusion of gravitational settling does not substantially increase the very stringent limit imposed on the amount of mass that can be accreted. It is possible that the Li abundance in accreted AGB material is substantially higher than computed at present owing to the occurrence of some additional mixing mechanism during the thermally pulsing asymptotic giant branch. However, the Li enhancements required to make the thermohaline mixing models fit the observations  would have to be at least two orders of magnitude above the current yield. This seems an unlikely level of enhancement.  Comparing measured sodium abundances to the models suggests that the majority, {\\it but not all}, stars undergo mixing during the main sequence. The abundance in the AGB ejecta is significantly higher than that observed in most CEMP stars. Unless sodium is produced to a much lower degree in AGB stars than the models predict (or the stars we are observing all had companions of much lower mass than we have assumed), it is difficult to reconcile models without main-sequence mixing with the observed abundances. However, it is difficult to explain the extremely sodium-rich objects with models that {\\it do} include thermohaline mixing as this would require the accreted material to have a sodium abundance considerably higher than current model predictions. It therefore seems prudent to argue that some stars mix their accreted material while others do not or that there is a spread in the efficiency of the mixing of accreted material. The cause for this is unclear and merits investigating. One possible suggestion is that rotation may play a role. \\citet{2008ApJ...684..626D} have pointed out that rotationally-driven horizontal turbulence may suppress thermohaline mixing. If the secondary stars are rotating at different rates (or the angular momentum content of the accreted material varies from star-to-star) then some stars may experience thermohaline mixing whilst others do not.  The flatness of the observed magnesium abundances also does not support models without thermohaline mixing on the main sequence. However, this may be because the models used are not $\\alpha$-enhanced and consequently are likely to be less rich in \\el{24}{Mg} than the observed stars. AGB and secondary models with $\\alpha$-enhanced compositions and abundances should certainly be investigated. As the magnesium from AGB stars is predominantly \\el{25,26}{Mg}, additional \\el{24}{Mg} in the stellar material could easily mask the dilution of the AGB material and hence we would not be able to see any changes due to mixing events. However, if the CEMP stars have received a substantial quantity of magnesium from an AGB companion, it is likely to be in the form of the neutron-rich isotopes \\el{25}{Mg} and \\el{26}{Mg}. Hence the signature of accretion from such a companion may be detectable in the ratios of \\el{25,26}{Mg} to \\el{24}{Mg}. Such measurements have been made in metal-poor but carbon-normal stars but the results are inconclusive \\citep{2003ApJ...599.1357Y,2007ApJ...659L..25M}.  The cases of extreme enhancement of F, Na and Mg on the giant branch argue against the accretion of only small quantities of material (i.e. thousandths of a solar mass) in these cases. The required abundances in the AGB ejecta would be beyond the current range of model predictions. The accretion of much greater amounts of material would be expected to show more extensive mixing on the main sequence if thermohaline mixing is efficient.  No clear picture arises from the elements studied here. The detection of lithium in turn-off objects seems to rule out extensive mixing for these cases. Some highly sodium and magnesium enriched turn-off objects can only be explained if accreted material remains unmixed during the main sequence. However, measurements of [Na/Fe] seems to suggest that most CEMP stars do efficiently mix their accreted material. A more comprehensive study involving a greater number of elements may help to elucidate what processes go on in these objects."
    },
    "0812/0812.3652_arXiv.txt": {
        "abstract": " ",
        "introduction": "In the last few years there has been great progress in understanding the non-Gaussianity of the primordial spectrum of density fluctuations. Starting from the first full computation of the non-Gaussian features in single field slow roll inflation \\cite{Maldacena:2002vr,Acquaviva:2002ud}, several alternative models have been proposed that produce a large and in principle detectable level of non-Gaussianities through different mechanisms for generating density fluctuations in the quasi de Sitter inflationary phase, both in the case of single field inflation \\cite{Arkani-Hamed:2003uz,Alishahiha:2004eh,Shandera:2006ax,Chen:2006nt,Cheung:2007st}, and in the case of multi-field inflation \\cite{Lyth:2002my,Zaldarriaga:2003my}. Further, it has been at last developed a consistent model of bouncing universe \\cite{Lehners:2007ac,Buchbinder:2007ad,Creminelli:2007aq} which, though clearly less compelling than inflation, predicts a possibly detectable amount of non-Gaussianity \\cite{Creminelli:2007aq}. If large non-Gaussianities will be detected, we would have to abandon the standard slow-roll inflation picture, but would be left with new information to understand the dynamics of the inflaton \\cite{Babich:2004gb,Cheung:2007st}.  At the same time, from the experimental side, the WMAP satellite has allowed for a huge improvement in our measurement of the properties of the CMB. Observations seem to confirm the generic predictions of standard slow roll inflation \\cite{Komatsu:2008hk}. Limits on the primordial non-Gaussianity of the CMB have been significantly improved \\cite{Creminelli:2005hu,Komatsu:2008hk}, but for the moment the data are consistent with a non-Gaussian signal. Recently it has been realized that even large scale structure surveys are relevantly sensitive to at least some particular kind of non-Gaussianities \\cite{Dalal:2007cu,Slosar:2008hx}, giving constraints comparable to those from WMAP \\cite{Slosar:2008hx}.  Oversimplifying, the limits on non-Gaussianities are set on the parameter $f_{\\rm NL}$ such that the curvature $\\zeta$ that we observe in the CMB is given by: \\bee\\label{eq:non-gauss-naive} \\zeta(\\vec x)=\\zeta_g(\\vec x)-\\frac{3}{5}f_{\\rm NL}\\left( \\zeta_g(\\vec x)^2-\\langle\\zeta_g^2\\rangle\\right)\\ , \\eee where $\\zeta_g(\\vec x)$ is a gaussian random variable~\\footnote{This oversimplifies the structure of the possible non-Gaussianities \\cite{Babich:2004gb,Cheung:2007st}, but this is enough for our introductory purposes.}. So far we have a 1-$\\sigma$ error on $f_{\\rm NL}$ of order $\\Delta f_{\\rm NL}\\simeq 30$ from WMAP \\cite{Komatsu:2008hk}, and soon the Planck satellite is expected to improve this limit to $\\Delta f_{\\rm NL}\\lesssim 4$ \\cite{Babich:2004yc}. However, it is clear by looking at eq.~(\\ref{eq:non-gauss-naive}) that the corrections due to the non-linearities of general relativity and of the Boltzmann equations that regulate the fluid dynamics are very naively expected to give rise to $f_{\\rm NL}$ of ${\\cal{O}}(1)$. This is clearly very close to what Planck will achieve, and it means that in order to be able to fully exploit the next new experimental results, the calculation of the non-Gaussianity induced by the non-linearities in general relativity and in the plasma physics needs to be done.  Of course we are not the first ones to realize the importance of this issue. In particular, an effort of finding and solving the full set of second order equations is on-going by some groups \\cite{bartolo, bartolo2, astro-ph/0703496,Pitrou:2008ak}. For example, the three point function of temperature fluctuations in the particular limit where one of the wavelengths is much longer than the horizon at recomination and than the other two wavelengths has been consistently computed, finding an effect of order $f_{\\rm NL}\\sim1$~\\cite{Creminelli:2004pv}. This is still far from the required full calculation. The derivation of the second order equation requires to deal with some subtleties such as the formulation of the Boltzmann equation in curved spacetime that were neglected in former derivations such as the one of \\cite{bartolo}, and that could affect the result at the order we care about. One of the two purposes of the present paper is to derive in a consistent way this set of equations. Unfortunately this will take us through a rather technical path, though we will try to make always clear the physical interpretation of what we are doing. We can already synthesize the main points of the derivation. We will work in the all-orders generalization of the Newtonian gauge, which is called Poisson Gauge. Writing Einstein equations in this gauge is tedious, but not particularly difficult. Writing down the Boltzmann equations in this frame is instead a little more subtle. The Boltzmann equation consists of a free-streeming term (so called Liouville term) which takes into account of the free evolution of the species one-particle distribution, and of a collision term, which takes into account how interactions among the particles affect the distribution. The collision term is expressed in terms of cross-sections that are measured and well expressible in Minkowski space. This means that, in order to write down the Boltzmann equation in Poisson gauge, one needs at each point  to find the coordinate transformation to the local inertial frame, write down the Boltzmann equation there, and then transform back to the original frame. This will give us the recipe to write down the Boltzmann equations at second order in the perturbations.  If schematically for the moment we write each perturbation $\\delta$ as $\\delta=\\delta^{(1)}+\\delta^{(2)}$ where $\\delta^{(2)}$ is of order $(\\delta^{(1)}){}^2$, the resulting first and second order set of Einstein and Boltzmann equations will take the form of a system of equations of the form: \\be &&D[\\delta^{(1)}]=0\\ , \\\\ \\nonumber &&D[\\delta^{(2)}]=S[\\delta^{(1)}{}^2]\\ . \\ee Here $D$ is a differential operator that is the same both for the first and the second order equations. The second order equation, unlike the first order one, has a source term $S$ proportional to the square of the first order perturbations \\footnote{The first order perturbations are non zero only once one assigns them some primordial non-zero initial conditions.}.  Among the many terms that contribute to the source $S$, there are some (coming from the Compton scattering collision term) in which one of the two first order perturbations is the perturbation in the number density $n_e$ of free electrons: $\\delta_e\\equiv \\delta n_e/n_e$. If one concentrates, as we will do, on the CMB, then $\\delta_e$ affects the CMB anisotropy only at second order. This can be seen in the following way. The free electron density affects the CMB temperature through the Compton collision term, which determines the visibility function and the diffusion scale. The former gives the probability for a photon to originate from a given distance along the line of sight to the observer, while the latter accounts for the Silk damping at higher multipoles $l$'s. In a homogeneous universe, before, during and after recombination the radiation temperature decreases as $a^{-1}$, $a$ being the scale factor, irrespective of the electron density. Thus, a perturbation to the electron density changes the position of the last scaterring surface and the mean free path of the photons before decoupling, but not the observed radiation temperature. Therefore, since at first order we would perturb $n_e$ and keep the other quantities unperturbed, the CMB anisotropies are not affected by $\\delta_e$ at this order. However, they clearly are so at second order. This explains why this quantity had never been computed before in detail. The first part of the paper will in fact be devoted to the computation of this quantity which is necessary to solve for the second order perturbations. As we will see, this task will not be completely straightforward because of the way recombination occurs in our universe.  In particular, even in the homogeneous universe the collision term which controls the population of free electrons does not vanish (of course, otherwise there would be no recombination in the homogeneous universe). This means that the same subtleties that we mentioned above about writing down the Boltzmann equation at second order will now apply also at first order. Further, even in the unperturbed case, recombination is not a very straightforward process, involving atomic transition between many different states. The treatment is simplified by the fact that matter and radiation are very close to equilibrium during recombination, with the most notable exception being the ground and first excited states of Hydrogen which are not mutually in equilibrium. This results in a series of approximations, which allow one to treat Hydrogen as an effective 3-level atom (ground state, first excited state and continuum) as was done in the classic work of Peebles \\cite{peebles}. In the end, one can write a single ordinary differential equation (ODE) governing the evolution of the electron density. Working in the same approximations, we shall derive the analogous equation for the perturbation of the electron density~eq.~(\\ref{eq:delta_e_first}). We also derive the equation for the perturbations of the matter temperature~eq.~(\\ref{deltaTM1})  \\footnote{One can try to directly perturb Peebles' ODE and write an equation for $\\delta_e$ which then has to be complemented with an equation for the perturbations to the matter temperature, $\\delta_{T_M}$. The calculation of $\\delta_e$ has been already attempted in \\cite{Novosyadlyj} without accounting for the effects of metric perturbations and of the perturbation to the escape probability, and in \\cite{arXiv:0707.2727v1,astro-ph/0702600v2}, where only a non rigorous  derivation of the perturbations to escape probability was given. We will show that the formulas given in \\cite{arXiv:0707.2727v1,astro-ph/0702600v2} are indeed correct apart for irrelevant corrections.}.   At least on large scales, one can consider the primordial perturbation as representing a time delay $\\delta t$ of the unperturbed solution in each point in space (of the order of $10^{-5} H^{-1}$). In this case, $\\delta n_e\\sim \\dot n_e \\delta t$, and since the time-scale of recombination is rather quick, $\\dot n_e/n_e\\sim 15\\, H$ (that is 5 times the expansion rate), we expect $\\delta_e\\sim 5 \\times 10^{-5}$, an enhanced perturbation with respect to the standard ones. This is in fact what will be the result of a detailed calculation.  Because of this enhancement, it is natural to wonder if this might affect the CMB, for example by generating an $f_{\\rm NL}$ of order 5, which could be potentially detectable by Planck.  We find the way $\\delta_e$ changes the CMB is mainly by delaying the time of recombination, by modifying the shape of the visibility function, and by changing the diffusion scale. The result is that $\\delta_e$  gives a potentially detectable effect equivalent to $f_{\\rm NL}\\sim 5$. Since the calculation and the interpretation of this result is not completely straightforward, it will be  the subject of a companion paper \\cite{inprep}.  The paper is organized as follows. In sec.~\\ref{section:review} we review recombination in the unperturbed universe. In sec.~\\ref{section:firstorder} we derive the first order equations for recombination in the perturbed universe and discuss the results of their numerical integration. In sec.~\\ref{section:secondorder} we derive the full set of second order equations for cosmological perturbations. In sec.~\\ref{section:summary} we summarize our results. In App.~\\ref{section:A} we review the formulation of kinetic theory in curved spacetime. In App.~\\ref{section:B} we derive in a rigorous way the perturbations to the escape probability. In App.~\\ref{section:C} we give an erratum to the Compton collision term at second order which is already present in the literature. ",
        "conclusions": "The evolution of $\\delta_e$ and $\\delta T_M$ is given by eq.s~(\\ref{num_cons_1}), (\\ref{N_e_1}), (\\ref{delta_K_P}) and (\\ref{deltaTM1}), where we also include the contribution from He recombination as explained at the end of sec.~\\ref{sec:first_order_eq}. The linear approximation is valid up to the reionization epoch, when the electron density perturbation becomes non-linear.  The cosmological parameters we use are $$(\\Omega_b,\\,\\Omega_\\Lambda,\\,h,\\,T_{\\rm cmb}\\,,Y_p,\\,n_s,\\tau)=(0.041,\\,0.76,\\,0.73,\\,2.726,\\,0.24,\\,0.958,\\,0.092)\\ .$$ We remind that the peak of the visibility function is at $\\eta=288.42\\,$Mpc and the present conformal time is $\\eta_0=14554.28\\,$Mpc. We find it helpful to give the following correspondence $$\\eta=(70, 288, 500, 1000, 5000)\\,{\\rm Mpc}\\quad \\rightarrow\\quad z=(5936, 1090, 507, 174, 9)\\ .$$  In the Introduction we argued that the amplitude of $\\delta_e$ at large scales should be enhanced due to the small timescale of recombination. This is valid for gauges, as Newtonian gauge for example, where the perturbations do not vanish for $k\\to 0$, and  are completely described as time delayed homogeneous solutions. Therefore, for $k\\to 0$ in the Newtonian gauge we expect \\be\\label{dedb} \\delta_e/\\delta_b=d \\ln n_e/d\\ln n_b=1-\\partial_\\eta (\\ln x_e)/(3\\mathcal{H})\\ , \\ee and analogously $\\delta_{T_M}/\\delta_{T_R}=d \\ln T_M/d\\ln T_R$. In Fig. \\ref{fig:con1} we show a plot of the amplitudes of several perturbations including $\\delta_e$ and $\\delta_{T_M}$ in the Newtonian gauge for $k=10^{-4}\\,$Mpc$^{-1}$. The electron density and matter temperature perturbations indeed follow precisely an evolution equivalent to a time delayed homogeneous universe solution as in eq.~(\\ref{dedb}) (in fact in the plot the difference is in practice indistinguishable). The electron density perturbation departs from $\\delta_b$ during HeIII$\\to$ HeII recombination at $\\eta\\sim70\\,$Mpc, and then during Hydrogen recombination. The HeII$\\to$ HeI recombination is delayed \\cite{sasselov}, and its contribution to $\\delta_e$ overlaps with that of Hydrogen recombination.  As expected,  the electron density perturbation is enhanced with respect to the other perturbations at recombination.  The plotted first order variables are (except for $\\delta_{T_M}$) the matter and radiation fluctuations which source the second order CMB anisotropies. Later we will see that this enhancement persists also for modes well inside the horizon. This  suggests that $\\delta_e$ will play an important role in sourcing the bispectrum of the CMB.  The calculation of these effects will be the subject of a companion paper  \\cite{inprep}.  Still in Fig.~\\ref{fig:con1}, before $z\\sim500$, we notice that $\\delta_{T_M}\\simeq\\delta_{T_R}$ as Compton scattering keeps baryons and photons in thermal contact.  $\\delta_{T_M}/\\delta_{T_R}$ becomes equal to 2 after $z\\sim500$ when the matter temperature decouples. At this points matter and radiation redshift differently even in an unperturbed universe, which explains the factor of 2.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=15cm]{newtonian0_0001n2n.pdf} \\end{center} \\caption{The amplitude of the perturbations to the electron density in the Newtonian gauge compared to other first order perturbations for a mode well outside the horizon. The peak of the CMB visibility function is at $\\eta=288\\,$Mpc. The normalization of the perturbations is the one in CMBFAST -- superhorizon $\\zeta=1$. The perturbations in the Newtonian gauge for superhorizon modes are equivalent to time shifted zeroth order evolution. This results in two regions of enhancement of $\\delta_e$ corresponding to HeIII$\\to$HeII, and the combined HeII$\\to$HeI and HII$\\to$HI recombination. Similarly, $\\delta_{T_M}$ coincides with $\\delta_{T_R}$ while Compton scattering is effective in keeping the photons and electrons in thermal contant. As $T_M$ decouples, $\\delta_{T_M}$ converges to $2\\delta_{T_R}$.} \\label{fig:con1} \\end{figure}  Even more interestingly, we find that the enhancement to $\\delta_e$ persists even for modes inside the horizon where the effect is clearly not gauge dependent. In Fig. \\ref{fig:con4} we show $\\delta_e$ and $\\delta_{T_M}$ for a mode corresponding to the second acoustic peak. The Kompaneets timescale $n_t T_M/\\Lambda_C$ is still smaller than one comoving Mpc for $\\eta<400\\,$Mpc (see Fig. \\ref{fig:thermal}), therefore, for the modes we consider, even after recombination, the full (zeroth+first order) $T_M$ and $T_R$ are kept in contact by Compton scattering, as can be seen in Fig. \\ref{fig:con4}. At later times, the two temperature fluctuations deviate and $\\delta_{T_M}$ becomes greatly enhanced with respect to $\\delta_{T_R}$ since overdense regions become hotter due to adiabatic contraction.  For $z<200$ ($\\eta>1000\\,$Mpc) one can neglect $\\d_\\eta \\log( x_e)/(3\\mathcal{H})\\lesssim 0.1$ in (\\ref{dedb}), and from the time shift perspective one would expect $\\delta_e/\\delta_b\\approx 1$. However, from Fig. \\ref{fig:con4} we clearly see that $\\delta_e$ is suppressed compared to $\\delta_b$ during the dark ages. As argued in \\cite{arXiv:0707.2727v1}, during the dark ages, overdense regions have enhanced $T_M$, and this is associated, as we will soon explain, to a more efficient recombination  and a suppressed $\\delta_e$. This effect induces a decrease of about $10\\%$ in $\\delta_{T_M}$ during the dark ages, which in turn influences the 21$\\,$cm power spectrum at a $2\\%$ level at $z\\sim 50$ \\cite{arXiv:0707.2727v1}. For a detailed discussion of the effect of the matter temperature perturbations on the 21cm angular-power spectrum we refer the reader to \\cite{astro-ph/0702600v2}.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=15cm]{newtonian0_04n2n.pdf} \\end{center} \\caption{The amplitude of the perturbations to the electron density in the Newtonian gauge compared to other first order perturbations for a mode corresponding approximately to the second acoustic peak. The enhancement of $\\delta_e$ is still well approximated by eq.~(\\ref{dedb}). This mode is well inside the horizon, and therefore, this enhancement is not a gauge artefact, and will leave an imprint on the CMB bispectrum.} \\label{fig:con4} \\end{figure}  From the numerical results  (see Figures \\ref{fig:con4} and \\ref{fig:viscon}), we see that eq.~(\\ref{dedb}) is still a good approximation to the ratio of $\\delta_e$ and $\\delta_b$ even for modes well inside the horizon. This does not mean anymore that $\\delta_e$ is just a time delay of the homogeneous solution. Instead, it tells us that $n_e$ is determined by the local value of $n_b$ (or equivalently of the temperature) with the same relationship as in an unperturbed universe: $\\delta n_e\\simeq \\delta n_b\\, \\dot n_e/\\dot n_b$. This can be explained by the fact that the timescale for recombination is still much faster than the mode we are considering.  In Fig. \\ref{fig:viscon} we plot $\\delta_e$ and $\\delta_b$ (which has a size comparable to the other perturbations that source the second order CMB anisotropies) as a function of $k$ at the peak of the visibility function. As one can expect, (\\ref{dedb}) well approximates $\\delta_e$ down to scales comparable to the recombination timescale ($k\\sim0.1\\,$Mpc$^{-1}$). The enhancement of $\\delta_e$ is prominent for modes with scale larger than approximately $10^{-1}\\,$ Mpc$^{-1}$ (which is comparable to the photon diffusion scale $k_D\\sim0.15\\,$Mpc$^{-1}$). The suppression at high-$k$ with respect to what implied by (\\ref{dedb}) is due to the fact that for these modes the timescale of the oscillation becomes faster than the rate of recombination (see Fig.~\\ref{fig:recomb}), inducing an averaging out of the enhancement \\footnote{This suppression at large $k$ will lead to an important simplification when we derive the CMB bispectrum from recombination \\cite{inprep}.}.   Notice that in Fig.~\\ref{fig:con4}, we can see that at late times, even for a relatively low-$k$ mode, $|\\delta_e|$ is a bit suppressed with respect to what is given by eq.~(\\ref{dedb}). Let us briefly explain why this happens. For $\\eta\\gtrsim 400$ Mpc, the timescale of the mode with $k=0.04$ Mpc$^{-1}$ becomes comparable to the recombination timescale (see again Fig. \\ref{fig:recomb}). At this time,  the contribution from Ly$\\alpha$ escape, two-photon absorption, and photoionization is negligible, and we can approximate the  homogeneous integrated Hydrogen recombination collision term as   $Q\\simeq-\\alpha_B n_{\\mathrm{HII}}n_e$. This clearly  increases in magnitude for higher $n_e$ or $n_{\\rm HII}$.  This dependence of $Q\\propto -n_{\\mathrm{HII}}$ implies that overdense regions will recombine faster, explaining the suppression of $|\\delta_e|$ relative to (\\ref{dedb}) at later times in Fig.~\\ref{fig:con4}.    \\begin{figure}[h!] \\begin{center} \\includegraphics[width=15cm]{con_vis1n.pdf} \\end{center} \\caption{A plot of some first order quantities  in the Newtonian gauge at the peak of the visibility function showing the enhancement of $\\delta_e$. The first order $\\delta_b$ shows the typical magnitude of the variables sourcing the CMB anisotropies at second order, and one can see that $\\delta_e$ dominates. We also plot the expected electron density perturbation coming from eq.~(\\ref{dedb}) (dotted line) which is a good approximation to the true $\\delta_e$ down to the recombination scale ($k\\sim0.1\\,$Mpc$^{-1}$). The enhancement for superhorizon modes is gauge dependent, but persists for modes inside the horizon, where the effect is observable, approximately down to the photon diffusion scale ($k_D\\approx0.15\\,$Mpc$^{-1}$).} \\label{fig:viscon} \\end{figure}  Let us comment on the validity of our calculation at high $k$'s. In deriving the perturbed Peebles equation, we pointed out that for $k>1\\,$Mpc$^{-1}$, the Ly$\\alpha$ line cannot be treated as quasi-static, and the perturbations to the escape probability will be modified. However, those perturbations have a comparably small effect on $\\delta_e$, since recombination is dominated by two-photon decay. The most important effect which may influence the 3-level atom calculation is the effect of nonequilibrium excited levels. That effect in the homogeneous universe leads to a correction of about $10\\%$ to $n_e$ for $\\eta\\gtrsim360$ Mpc, when photoionization becomes negligible (Fig. \\ref{fig:recomb}) \\cite{sasselov}. At this point recombination to excited states becomes enhanced, and the bound-bound transitions between $n>2$ levels is no more in thermal equilibrium (for a discussion of the effect see \\cite{1999ApJ...523L...1S}). This tells us that the timescale over which the excited states equilibrate becomes comparable to the Hubble time  at these redshifts (see again Fig.~\\ref{fig:recomb}). Since this effect occurs quite late in the recombination history, it gives only a 10\\% effect on $n_e$. In the presence of perturbations with $k$ lower than $\\sim 10^{-3}\\,$ Mpc$^{-1}$, the effect will be of the same order. For modes with higher $k$, the thermalization timescale becomes more important even earlier, and so potentially the effect could be larger. However, since the photoionization rate grows quite steeply as we move to higher redshifts (see again Fig.~\\ref{fig:recomb}), this effect becomes important only for $k\\sim 1 \\,$ Mpc$^{-1}$. This is the upper limit on the $k$'s for which our calculation is valid.                 \\subsubsection{Perturbations in the synchronous gauge}  In \\cite{inprep} we will be working in synchronous gauge to calculate the CMB bispectrum generated from recombination. In this gauge, the evolution of $\\delta_e$ and $\\delta T_M$ is given again by (\\ref{num_cons_1}), (\\ref{N_e_1}), (\\ref{delta_K_P}) and (\\ref{deltaTM1}), but with the substitutions $\\Psi\\to 0$, and $\\dot \\Phi\\to -\\dot h/6$, where $h$ is the trace of the perturbations to the spatial part of the metric as defined in \\cite{maed}.  In Fig.~\\ref{fig:vissyn} we show the results for $\\delta_e$ in the synchronous gauge. As expected, for modes well inside the horizon, the density perturbations in the synchronous and Newtonian gauges are nearly identical (cf. Fig.~\\ref{fig:viscon}).  Perturbations vanish for superhorizon modes since they correspond to time-shifted homogeneous solutions and the synchronous gauge time coordinate equals the proper time of comoving observers  which makes the time-delay vanish outside of the horizon. As we will show in \\cite{inprep}, this property of the synchronous gauge will be very useful in estimating the CMB bispectrum generated from recombination.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=15cm]{syn_vis2n.pdf} \\end{center} \\caption{The same as in Fig.~\\ref{fig:viscon} repeated in the synchronous gauge. The only important difference between the results in the Newtonian and the synchronous gauge is for superhorizon modes for which the amplitude in the synchronous gauge goes to zero.} \\label{fig:vissyn} \\end{figure}                       }  The purpose of this paper was to provide a detailed analysis of the perturbations to the recombination history of the universe and give the full corrected second order Boltzmann and Einstein equations for the photons, baryons and CDM for a flat FRW background. Those are necessary for the correct treatment of CMB secondaries coming from recombination.  We analyzed the different timescales in recombination and argued that a perturbed version of the recombination equation for the Peebles 3-level atom is adequate for obtaining the evolution of the perturbations to the electron density. We derived rigorously the perturbations to the escape probability by angle averaging the photon distribution function, and showed that it is sufficient to treat the escaping of Ly$\\alpha$ photons as only due to the local velocity divergence.  We showed that for modes longer than the timescale of recombination (corresponding to $k\\lesssim0.1\\,$Mpc$^{-1}$), it is a good approximation to consider $n_e^{(0+1)}$ as determined by the local value of $n_b^{(0+1)}$ (or equivalently of the temperature), with the same relationship as in an unperturbed universe: $\\delta n_e\\simeq \\delta n_b\\, \\dot n_e/\\dot n_b$. This is due to the fact that the timescale of recombination is much faster than the considered modes. For the same reason,  this also implies that the amplitude of the electron density perturbations is enhanced by a factor of roughly five compared to the baryon density perturbations. This enhancement is physical since it persists for modes well inside the horizon, down to scales comparable with the photon diffusion scale (corresponding to $k\\lesssim0.2\\,$Mpc$^{-1}$). The enhancement is most prominent around the peak of the photon visibility function and therefore has the potential to have an impact on the CMB bispectrum, as we will investigate in a companion paper \\cite{inprep}.  Solving for the density of free electrons in the universe, as well as writing the full set of second order Boltmann and Einstein equations requires to be careful with the definition of coordinates and collision terms. For this reason, we reviewed the formulation of kinetic theory in curved spacetime and applied it to derive the Boltzmann and Einstein equations at second order in the perturbed FRW universe. We compared our results with former pre-existing literature. We find that  the former derivations are not completely correct, and we give the necessary corrections.  Solving the full set of second order equations is a very hard task, which however is motivated by the current experimental and theoretical status. In a companion paper \\cite{inprep}, we solve approximately these second order equations restricting ourself to modes well inside the horizon and to second order sources that are proportional to $\\delta_e$, and we compute the induced bispectrum. There we find that the non-Guassian signal is equivalent to about $f_{\\rm NL}\\sim-5$, which is potentially detectable by future experiments such as Planck."
    },
    "0812/0812.4302_arXiv.txt": {
        "abstract": "Gravitational-wave signals from inspirals of binary compact objects (black holes and neutron stars) are primary targets of the ongoing searches by ground-based gravitational-wave interferometers (LIGO, Virgo, and GEO-600). We present parameter-estimation simulations for inspirals of black-hole--neutron-star binaries using Markov-chain Monte-Carlo methods. As a specific example of the power of these methods, we consider source localisation in the sky and analyse the degeneracy in it when data from only two detectors are used. We focus on the effect that the black-hole spin has on the localisation estimation. We also report on a comparative Markov-chain Monte-Carlo analysis with two different waveform families, at 1.5 and 3.5 post-Newtonian order. ",
        "introduction": "\\label{sec:intro}  Binary systems with compact objects --- neutron stars (NSs) and black holes (BHs) --- 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 \\cite{2002gr.qc.....4090C}: LIGO \\cite{1999PhT....52j..44B}, Virgo \\cite{2004CQGra..21..385A} and GEO-600 \\cite{2004CQGra..21S.417W}. Merger-rate estimates are quite uncertain and for BH-NS binaries current detection-rate estimates range from 0.0003 to 0.1\\,yr$^{-1}$ for first-generation instruments \\citeaffixed{2008ApJ...672..479O}{\\emph{e.g.}}. Upgrades to Enhanced LIGO/Virgo (2008--2009) and Advanced LIGO/Virgo (2011--2014) are expected to increase detection rates by factors of about $\\sim 10$ and $\\sim 10^3$, respectively.  The measurement of source properties holds major promise for improving our astrophysical understanding and requires reliable methods for parameter estimation. This is a challenging problem, however, because of the large number of parameters ($9$ for circular non-spinning binaries, and more for spinning systems) and the significant amount of structure in the parameter space. In the case of low-mass-ratio binaries (\\emph{e.g.}\\ BH-NS), these issues are amplified for significant spin magnitudes and large misalignments between the BH spin and the orbital angular momentum \\cite{1994PhRvD..49.6274A,2003PhRvD..67d2003G,2003PhRvD..67j4025B}. However, the presence of spins improves parameter estimation through the signal modulations, although still presenting us with a considerable computational challenge. This was highlighted in the context of LISA observations \\citeaffixed{2004PhRvD..70d2001V,2006PhRvD..74l2001L}{see} and in our first study devoted to ground-based observations \\cite{2008ApJ...688L..61V}.   In this paper we examine the potential for parameter estimation of spinning binary inspirals with ground-based interferometers. \\citename{2006CQGra..23.4895R} \\citeyear{2006CQGra..23.4895R,2007PhRvD..75f2004R} explored parameter estimation for non-spinning binaries, which requires 9 parameters. We focus on BH-NS binaries, which can exhibit significant coupling between the orbital angular momentum and the BH spin, mainly because of the high mass ratio \\cite{1994PhRvD..49.6274A}, while at the same time we are justified to ignore the NS spin, leading to a 12-dimensional parameter space. We apply a newly developed Markov-chain Monte-Carlo (MCMC) algorithm \\cite{2008CQGra..25r4011V} to spinning inspiral signals injected into synthetic noise and we derive posterior probability-density functions (PDFs) of all twelve signal parameters. In our previous study \\cite{2008ApJ...688L..61V}, we showed the accuracy obtained in sky-position determination using data from a two-detector network, where a degeneracy in the sky position exists, and from a three-detector network, where the degeneracy is broken. Following this work, we here analyze in further detail the degeneracy which is present when data from only two detectors are used. In section~\\ref{sec:skyring} we show that the degeneracy in the sky position is reduced but not lifted when a significant spin is present ($a_\\mathrm{spin}\\ge0.5$), and that a sufficient angle between spin and orbital angular momentum can break such a degeneracy ($\\theta_{SL}\\ge55^{\\circ}$). In this study, we demonstrate that these degeneracies are due to a high degree of similarity between signals from sources with significantly different parameter sets, while we made sure the observed effects are real and not in fact artifacts due to potential errors in our analysis methodology.  In section~\\,\\ref{sec:comparison}, we also demonstrate that the inclusion of higher post-Newtonian orders in the waveform can improve the accuracy of intrinsic-parameter estimation.  Meanwhile, we find that the additional structure in the parameter space of higher-order waveforms lowers the sampling efficiency of the MCMC and requires improvements to the sampling scheme. ",
        "conclusions": "\\label{sec:concl}  We have explored the degeneracies in the sky position for a gravitational-wave observation of an inspiral of a BH-NS binary with two non-colocated ground-based interferometers. Whereas simple triangulation based on time delays alone would result in a homogeneous sky ring, our MCMC runs show (in experiment 1) an incomplete ring in the sky consisting of arcs separated by gaps. While the arcs make up most of the circumference of the sky ring for an inspiral with a non-spinning BH, these arcs become smaller when spin is present and may be reduced to a single arc for the case of moderate spin and a sufficient misalignment between the BH spin and the orbital angular momentum. We demonstrated that in the spinning case, the maximum likelihood values are in fact lower in the gaps than in the arcs. In the non-spinning case the likelihood in the gaps can actually be as high as in the arcs; however the gaps can be explained by a smaller volume of support in the other extrinsic parameters, such as the binary orientation and distance (experiment 2). It is then less likely for the chains to sample this smaller volume in parameter space, resulting in gaps in the two-dimensional sky probability density function. In the non-spinning case, if a source is located in a gap, the posterior PDF still has a gap at the true source sky location (experiment 3).  The subject of estimating the position of a gravitational-wave source on the sky has been explored by many researchers, in the context of binary-inspiral and burst signals; the assumed baseline for these studies corresponds to three detectors located at the three LIGO-Virgo observatory positions, and operating at their design sensitivities. For detectable burst sources (SNR$>$5), arrival-time techniques should allow a precision of about $1^\\circ$ for the source direction~\\cite{2006PhRvD..74h2004C}. A method that takes into account burst-signal arrival time and amplitude, plus arrival-time uncertainties equivalent to what is observed with real LIGO and Virgo data, gives source uncertainties of a few degrees~\\cite{2008PhRvD..78l2003M}. Coherent techniques also exist for burst detection and sky-location determination, but by their own admission these methods are computationally costly~\\cite{2008arXiv0809.2809S}. Many techniques have also been developed for sky localization for binary-inspiral events. It has been shown, using time and mass parameter estimates, that the LIGO-Virgo network could localize the position of a binary-neutron-star inspiral to an accuracy of $4^\\circ$ if it were located in M87 (16 Mpc), or $2^\\circ$ for a source in NGC 6744 (10 Mpc)~\\cite{2008CQGra..25d5001B}; coherent methods, which also depend on estimating the time and mass parameters, give similar accuracies~\\cite{2009arXiv0901.4936A}.   MCMC parameter-estimation methods, like those used in this study, are capable of estimating the sky parameters, along with all of the other signal parameters; this is one reason why MCMC methods are computationally intensive. Coherent MCMC methods applied to signals observed by the LIGO-Virgo network will be able to resolve the sky location to ~$2^\\circ$ for signals with an SNR of 15~\\cite{2007CQGra..24..607R}. When the compact objects have spin, and the search templates account for this parameter, sky localization becomes relatively more accurate for higher values of spin~\\cite{2008ApJ...688L..61V}. It is important to remember that, in principle, the MCMC results show the best constraints one could hope to place on signal parameters (including the sky location) by displaying the true posterior PDFs. While comparisons with other particular sky-localization results may be cumbersome because different waveforms were assumed at different SNRs in different detectors, MCMC produced posterior PDFs display the statistically correct and most precise localization. Bayesian methods achieve better parameter-estimation accuracy when the template model describes the functional form of the actual signal more accurately. When a gravitational-wave detection occurs, it is likely that all possible sky localization algorithms will be used, and the methods should be considered to be complementary.   In section~\\ref{sec:comparison}, we compared MCMC results on software injections using waveform families of ${1.5}$-pN and ${3.5}$-pN order for both the injections and the MCMC parameter estimation. We have shown that the higher-order templates have the potential for more accurate parameter estimation, but that sampling the parameter space with these templates is more computationally difficult. However, a number of upcoming improvements should improve the sampling efficiency of our MCMC runs.  The analysis presented here is the second 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. We are in the process of updating our MCMC code to include the spin of the second binary member, increasing the dimensionality of the parameter space from 12 to 15. Finally, we intend to further develop our Bayesian approach into a standard tool that can be included in the follow-up analysis pipeline used for the processing of the `science data' collected by ground-based laser interferometers.      \\ack  The authors would like to thank Gareth Jones for his help in the understanding of the LAL parameter system. This work is partially supported by the National Science Foundation grant NSF-0838740, a NSF Gravitational Physics grant (PHY-0653321) to VK; NSF Gravitational Physics grant (PHY-0553422) to NC and the Max-Planck-Society (CR). Computations were performed on the Fugu computer cluster funded by NSF MRI grant PHY-0619274 to VK."
    },
    "0812/0812.0521_arXiv.txt": {
        "abstract": "We employ an analytical approach to investigate the signatures of Baryon Acoustic Oscillations(BAOs) on the convergence power spectrum of weak lensing by large scale structure. It is shown that the BAOs wiggles can be found in both of the linear and nonlinear convergence power spectra of weak lensing at about $40\\le l\\le600$, but they are weaker than that of matter power spectrum. Although the statistical error for LSST are greatly smaller than that of CFHT and SNAP survey especially at about $30<l<300$, they are still larger than their maximum variations of BAOs wiggles. Thus, the detection of BAOs with the ongoing and upcoming surveys such as LSST, CFHT and SNAP survey confront a technical challenge. ",
        "introduction": "In the early universe prior to recombination, the free electrons couple the baryons to the photons through Coulomb and Compton interactions, so these three species move together as a single fluid. The primordial cosmological perturbations on small scales excite sound waves in this relativistic plasma, which results in the pressure-induced oscillations and acoustic peak \\citep{1984ApJ...285L..45B,1998ApJ...496..605E}. The memory of these baryon acoustic oscillations (BAOs) still remain after the epoch of recombination. The BAOs leave their imprints through the propagating of photons on the last scattering surface and produce a harmonic series of maxima and minima in the anisotropy power spectrum of the cosmic microwave background (CMB) at $z\\approx1000$. In addition, due to the significant fraction of baryons in the universe, BAOs can also be imprinted onto the late-time power spectrum of the non-relativistic matter \\citep{1984ApJ...285L..45B,1996ApJ...471..542H,1998ApJ...496..605E}, which have been detected in the large-scale correlation function of Sloan Digital Sky Survey (SDSS) luminous red galaxies \\citep{2005ApJ...633..560E}, and the power spectrum of $21$ cm emission generated from the neutral hydrogen from the epoch of reionization through the underlying density perturbation \\citep{2008ApJ...673L.107M,2008PhRvL.100i1303C}. Essentially, the BAOs can give rise to the wiggles in the matter power spectrum of large scale structure during the evolution of the universe. Gravitational lensing can directly reveal the strenth of gravitational clustering\\citep{2003ApJ...592..664P,2005ApJ...629...23C}, and weak gravitational lensing is the direct measurement of the projected mass distribution of the large-scale structure\\citep{1999ARA&A..37..127M, 2001PhR...340..291B,2003ARA&A..41..645R,2006PhR...429....1L,2008PhR...462...67M}. Therefore, the BAOs prior to recombination should also be imprinted onto weak lensing power spectrum. Recently, the influence of baryons on the weak lensing power spectrum are investigated by many works \\citep{2004APh....22..211W,2004ApJ...616L..75Z,2006ApJ...640L.119J}. \\cite{2008arXiv0802.2416Z} study self calibration of galaxy bias in spectroscopic redshift surveys of baryon acoustic oscillations to show that SKA is able to detect BAO in the velocity power spectrum, and the precision measurement of cosmic magnification are also demonstrated\\citep{2005PhRvL..95x1302Z,2006MNRAS.367..169Z}. The ongoing and upcoming surveys such as the Canada-France-Hawaii-Telescope (CFHT) Legacy Survey\\footnote{{\\tt http://www.cfht.hawaii.edu/Science/CFHLS/}}, the SuperNova Acceleration Probe\\footnote{{\\tt http://snap.lbl.gov}: SNAP is being proposed as part of the Joint Dark Energy Mission (JDEM)} (SNAP) and the Large Synoptic Survey Telescope\\footnote{{\\tt http://www.lsst.org}} (LSST) will significantly reduce the statistical errors to a few percent level in the measurement of weak lensing power spectrum. Therefore, it is necessary to explore the feasibility of detecting the BAOs in weak lensing surveys. In this paper, we concentrate on the wiggles of BAOs on the power spectrum of weak lensing by large scale structure and their detectability for current weak lensing survey. ",
        "conclusions": "In this paper, we use an analytical approach to present signatures of BAOs on the convergence power spectrum of weak lensing. We show that, in both of the linear and nonlinear matter power spectra of mixed baryons+CDM, the BAOs wiggles can be clearly seen at about $k\\sim 0.1\\,h$Mpc$^{-1}$. With the decrease of redshift, the clustering of structures are enhanced for both of the pure CDM and mixed baryons+CDM models. Compared with pure CDM case, the mixed baryons+CDM case can suppress the linear and nonlinear matter power spectrum by an order of a few percents or more at about $k\\ge 0.1\\,h$Mpc$^{-1}$. The nonlinear evolution of structure can suppress the amplitude of BAOs wiggles and shift them to small scales with decrease of redshift. For the convergence power spectrum of weak lensing, we show that the BAOs wiggles can be visible in both of the linear and nonlinear power spectra for the mixed baryons+CDM model at about $40\\le l\\le600$, but they are weaker than that of matter power spectrum. With the increase of $z_s$ from 0.5 to 2, the BAOs wiggles are shifted to small scales for both of linear and nonlinear power spectra, i.e., from the range of about $40\\le l\\le250$ to $100\\le l\\le600$. We also study the detectability of BAOs's wiggles for weak lensing survey. Although the statistical error for LSST are greatly smaller than that of CFHT and SNAP survey especially at about $30<l<300$, they are still larger than the their maximum variations of BAOs wiggles. Thus, the detection of BAOs with the ongoing and upcoming surveys such as LSST, CFHT and SNAP survey i confront a technical challenge. Therefore, we expect future weak lensing survey with more lower statistical errors to capture BAOs wiggles signal from the convergence power spectrum. In addition, we have shown that the BAOs signatures can be imprinted onto the entire history of cosmic structure evolution since the epoch of recombination. The future observation of BAOs on weak lensing will provide a complementary knowledge of current BAOs measurement. On the other hand, the BAOs provide a ``standard ruler'' for the determination of cosmological parameters especially the probes of dark energy, so the combination of BAOs measurement by weak lensing together with other detection of BAOs signature at different stage: the last scattering surface, reionization and large scale distribution of galaxies in local universe, will supply more robust constraint on cosmological models in the future.  \\noindent{\\bf Acknowledgments.} We are very grateful to the anonymous referee for his valuable comments and suggestions that greatly improve this paper. We also thank Dr.Pengjie Zhang for valuable discussion and useful comment. This work was supported by the National Science Foundation of China (Grants No.10473002, 10533010), the Ministry of Science and Technology National Basic Science program (project 973) under grant No.2009CB24901 and the Scientific Research Foundation for the Returned Overseas Chinese Scholars, State Education Ministry."
    },
    "0812/0812.2181_arXiv.txt": {
        "abstract": "Differential rotation plays a crucial role in the alpha-omega dynamo, and thus also in creation of magnetic fields in stars with convective outer envelopes. Still, measuring the radial differential rotation on stars is impossible with the current techniques, and even the measurement of surface differential rotation is difficult. In this work we investigate the surface differential rotation obtained from dynamo models using similar techniques as are used on observations, and compare the results with the known radial differential rotation used when creating the Dynamo model. ",
        "introduction": "The model used in this work was originally used to investigate the flip-flop phenomenon (\\cite{elst_kor}, \\cite{kor_elst}). In the current work is is used to study the correlations between surface and radial differential rotation. The dynamo is modelled with a turbulent fluid in a spherical shell. The internal rotation is similar to the solar one with equatorial regions rotating faster than the polar regions, but here a smaller difference between core and surface rotation is used. The inner boundary of the convective zone is at the radius 0.4R$\\star$. The mean electromotive force contains an anisotropic alpha-effect and a turbulent diffusivity. The nonlinear feedback of the magnetic field acts on the turbulence only. The boundary conditions describe a perfect conducting fluid at the bottom of the convection zone and at the stellar surface the magnetic field matches the vacuum field.  We use standard cross-correlation methods to study the changes in the magnetic pressure maps obtained from the dynamo calculations. As these maps can be treated the same way as the temperature maps obtained from Doppler imaging, we can use the same techniques as in the case of the real observations to analyse the snapshots from the dynamo calculations. We have taken magnetic pressure maps from 36 time points over the activity cycle. The chosen time points, which are separated by about 50 days, are shown in Fig.~\\ref{fig:maps}.  \\begin{figure} \\begin{center} \\includegraphics[width=12.5cm]{Korhonen_Elstner_f1.eps} \\caption{Six examples of the 36 snapshots of the Dynamo model.} \\label{fig:maps} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0812/0812.1846_arXiv.txt": {
        "abstract": "Using numerical simulations of galactic disks that resolve scales from $\\sim 1$ to several hundred pc, we investigate dynamical properties of the multiphase interstellar medium in which turbulence is driven by feedback from star formation.  We focus on effects of HII regions by implementing a recipe for intense heating confined within dense, self-gravitating regions.  Our models are two-dimensional, representing radial-vertical slices through the disk, and include sheared background rotation of the gas, vertical stratification, heating and cooling to yield temperatures $T\\sim 10-10^4\\K$, and conduction that resolves thermal instabilities on our numerical grid.  Each simulation evolves to reach a quasi-steady state, for which we analyze the time-averaged properties of the gas. In our suite of models, three parameters (the gas surface density $\\Sigma$, the stellar volume density $\\rho_*$, and the local angular rotation rate $\\Omega$) are separately controlled in order to explore environmental dependencies. Among other statistical measures, we evaluate turbulent amplitudes, virial ratios, Toomre $Q$ parameters including turbulence, and the mass fractions at different densities.  We find that the dense gas ($n>100\\ \\cm^{-3}$) has turbulence levels similar to those observed in giant molecular clouds and virial ratios $\\sim 1-2$.  Our models show that the Toomre $Q$ parameter in the dense gas evolves to values near unity; this demonstrates self-regulation via turbulent feedback. We also test how the star formation rate $\\Sigma_{\\rm SFR}$ depends on $\\Sigma$, $\\rho_*$, and $\\Omega$.  Under the assumption that the star formation rate is proportional to the amount of gas at densities above a threshold $n_{\\rm th}$ divided by the free-fall time at that threshold, we find that $\\Sigma_{\\rm SFR} \\propto \\Sigma ^{1+p}$ with $1+p\\sim 1.2-1.4$ when $n_{\\rm th}=10^2$ or $10^3 \\cm^{-3}$, consistent with observed Kennicutt-Schmidt relations.  Estimates of star formation rates based on large-scale properties (the orbital time, the Jeans time, or the free-fall time at the mean density within a scale height), however, depart from rates computed using the measured amount of dense gas, indicating that resolving the ISM structure in galactic disks is necessary to obtain accurate predictions of the star formation rate. ",
        "introduction": "The interstellar medium (ISM) is commonly envisioned as a self-regulating system in statistical quasi-equilibrium.  Multiple components of gas with varying densities and temperatures coexist \\citep{1969ApJ...155L.149F,1974ApJ...189L.105C,1977ApJ...218..148M}, animated by turbulence that pervades the whole volume \\citep{2004ARA&A..42..211E}. Different components of gas play different roles in the ISM ecosystem, with the coldest and densest portions responsible for star formation. Massive stars, when they are born, energize the ISM through the HII regions and supernova blasts they create \\citep{1978ppim.book.....S}; this energy input is important in replenishing continual losses through turbulent dissipation.  UV radiation from young massive stars is also crucial in heating the gas.  The rate of star formation is determined by the available supply of dense gas, which in turn is regulated by the interplay between dynamics and thermodynamics in the ISM, and is affected by the galactic environment in which the ISM is contained \\citep{2004RvMP...76..125M,2007ARA&A..45..565M}.  While this overall framework is generally accepted and is supported by existing theory and observations, much work remains on both fronts to quantify the dependence of statistical properties on the global system parameters, and to establish when and how self-regulated quasi-steady states are achieved.  Given the importance of time-dependent processes and interdependencies in the ISM, complex theoretical models are needed in order to address even rather basic questions. For example, what sets the relative proportions of the different gas components?  In an idealized classical picture such as that of \\citet{1969ApJ...155L.149F} for the atomic medium, given a pressure and a mean density $\\bar n$, thermal equilibrium defines a density for each of two stable phases, $n_{\\rm warm}$ and $n_{\\rm cold}$, and the ratio of cold to warm gas is given by a simple algebraic relation: $M_{\\rm cold}/M_{\\rm warm}= (n_{\\rm warm}^{-1} - \\bar n^{-1})/ (\\bar n^{-1} - n_{\\rm cold}^{-1})$.  In the real ISM, however, which is a time-dependent system, thermal equilibrium only holds to the extent that the radiative times are short compared to dynamical times for compressions and rarefactions. Furthermore, the value of the mean density $\\bar n$ and pressure (averaged over large scales) are not even known {\\it a priori} for a given ISM surface density, since the vertical distribution of gas is sensitive to its dynamical state. This dynamical state itself depends on the (unknown) dense gas fraction, since more dense gas produces more feedback from star formation, and hence more turbulence that inflates the disk vertically to reduce $\\bar n$ (and also produces local variations in density and pressure through compressions and rarefactions).  Multidimensional effects (the ISM is not simply stratified perpendicular to the galactic plane, but is composed of filamentary clouds) and self-gravity additionally complicate the situation.  In recent years, a number of groups have begun developing models of the turbulent, multiphase ISM using time-dependent computational hydrodynamics simulations that include feedback from star formation \\citep[e.g.,][]{ 1999ApJ...514L..99K, 2004A&A...425..899D, 2005A&A...436..585D, 2007ApJ...665L..35D, 2005ApJ...626..864M, 2006ApJ...653.1266J, 2006ApJ...638..797D}, self-gravity of the gas \\citep[e.g.,][]{ 1999ApJ...516L..13W, 2002ApJ...577..197W}, and both of these effects \\citep[e.g.,][]{ 2001ApJ...547..172W, 2007ApJ...660..276W, 2005MNRAS.356..737S, 2006ApJ...641..878T, 2008ApJ...673..810T, 2008ApJ...680.1083R}. The treatment of feedback in these simulations is to inject thermal energy in regions identified as sites of star formation; most models focus on the energy input from supernovae. In very large scale simulations that have minimum resolution of only 50-100 pc, feedback implemented via thermal energy deposition is not in practice very effective, because the input energy is easily radiated away.  With finer numerical resolution, feedback regions expand adiabatically at first to make hot diffuse bubbles, driving shocks that sweep up surrounding gas and ultimately generate turbulence throughout the computational domain. A number of different issues have been addressed by these recent simulations, including investigating departures from thermal equilibrium in density and temperature PDFs, measuring the relative velocity dispersions of various gas components, and testing whether relationships between star formation and gas surface density emerge that are similar to empirical Kennicutt-Schmidt laws.   Even before the advent of supernovae, massive stars photoionize their surroundings, creating HII regions within molecular clouds that are highly overpressured and expand.  HII regions may in fact be the most important dynamical agents affecting the properties of dense gas in giant molecular clouds (GMCs), since the original turbulence inherited from the diffuse ISM is believed to dissipate within a flow crossing time over the cloud \\citep{1998ApJ...508L..99S,1998PhRvL..80.2754M}, while GMCs are thought to live for at least a few crossing times \\citep{2007prpl.conf...81B}. Analytic and semi-analytic treatments find that GMCs with realistic sizes, masses, and  star formation rates can indeed be maintained by the energy input from HII regions for a few crossing times, ultimately being destroyed through a combination of photoevaporation and kinetic energy inputs that unbind the remaining mass \\citep{1979MNRAS.186...59W, 1994ApJ...436..795F, 1997ApJ...476..166W, 2002ApJ...566..302M,2006ApJ...653..361K}. Recent three-dimensional numerical studies have begun to address this process in detail \\citep[e.g.][]{2006ApJ...647..397M,2007ApJ...668..980M,2007ApJ...671..518K}, focusing on regions within GMCs.   In the present work, we consider how the large-scale dynamical state of the ISM is affected by star formation feedback in the form of expanding HII regions.  Our main interests are in exploring how the turbulence driven by HII regions affects the properties of dense gas (we measure statistics of density, temperature, and velocity), in testing ideas of global self-regulation by feedback (we evaluate Toomre $Q$ parameters and virial ratios), and in exploring how galactic environment systematically affects the character of the ISM, including its ability to form stars. Complete ISM models should of course include feedback from supernovae as well as those from HII regions, and it is our intention to do this in future work.  However, we consider it useful to adopt a sequential approach, independently testing the effects of HII region feedback to provide a baseline for more comprehensive simulations. In addition to developing a physical understanding of the ways in which feedback affects the ISM, another goal of our work is to investigate the sensitivity of numerical results to prescriptions that are a necessary -- but not always fully tested -- aspect of galactic-scale studies of star formation.  In particular, we examine how the choice of density threshold in commonly-adopted recipes for star formation affects the resulting dependence of the star formation rate on ISM surface density.   Our approach to exploring the effects of galactic environment is to conduct a large suite of local simulations that cover a range of values for three basic parameters: the total surface density of gas in the disk ($\\Sigma$), the local midplane stellar density ($\\rho_*$), and the local rate of galactic rotation ($\\Omega$).  The parameter range covers a factor of six in gas surface density and galactic angular rotation rate, and a factor of 30 in stellar density. Our suite is divided into four series, each of which has one independent parameter that is systematically varied. We also include comparisons with hydrostatic models that are identical in terms of their input parameters to the fully-dynamic models, but do not include feedback and hence are not turbulent.  For this first set of pilot studies, we have not implemented full radiative transfer to evaluate the extent of HII regions (we intend to do so in the future), but instead introduce a simple prescription in which the boundaries of HII regions are determined by the gravitational potential.  Using this approach (rather than, for example, adopting a single fixed outer radius) has the advantage that the volume of the heated region expands as the density surrounding the source drops.  Since our treatment of HII regions does not attempt to be exact, we do not consider our specific results for e.g. velocity dispersions to be more than approximate (although in fact we find similar values for velocity dispersions in dense gas to those that are observed in GMCs). Instead, we shall emphasize the general properties of a multiphase ISM system in which turbulence is driven from within the dense phase.   This paper is organized as follows: In \\S 2 we describe our numerical methods, and in particular the recipe for star formation feedback. The control parameters for our disk models, and the properties of each model series, are presented in \\S 3.  Section 4 gives an overview of evolution based on our fiducial model. In \\S 5 we present the statistical properties of the gas in each model, and test environmental influences by intercomparing the model series. The implications of our results for star formation, both in real galaxies and in numerical simulations, is analyzed in \\S 6. We conclude with a summary and discussion in \\S \\ref{Summary}. ",
        "conclusions": "}  We have developed a numerical hydrodynamic code to study the life-cycle of multiphase, turbulent interstellar gas in disk galaxies; our model includes gas self-gravity, the vertical gravity associated with a fixed stellar disk, radiative cooling and heating in the temperature range of $10 ~\\mbox{K} \\le T \\le 10,000 ~\\mbox{K}$, sheared rotation in the background galactic potential (we adopt a flat galactic rotation curve), and a prescription for feedback in the form of HII regions that originate within massive, dense clouds.  Our simulation domains represent slices in the radial-vertical plane of a galactic disk.  We focus on scales large enough to include vertical stratification ($L_z$ up to 410 pc) and significant shear of the disk angular velocity ($L_x$ up to 1.6 kpc), but small enough to resolve sub-structure within dense, self-gravitating clouds that form (typical zone resolution is $\\sim 1$pc).  Our models are 2.5-dimensional, in the sense that all three components of the velocity are time-dependent functions.  For feedback to occur, we impose thresholds on both the local volume density and on the gravitational potential, so that HII regions only occur within massive clouds (consistent with observations).  The expansion of HII regions drives turbulence in all the components of the gas.  We have performed a large suite of simulations, covering a factor of six in gas surface density $\\Sigma$. In order to explore the dependence of ISM properties on galactic environment (in particular, the stellar vertical gravity and the angular momentum content of the gas), we have considered four different model series.  In Series Q, we vary $\\Sigma$, stellar volume density $\\rho_*$, and the disk rotation rate $\\Omega$ in tandem. In Series K, we vary $\\Sigma$ and $\\rho_*$ together while holding $\\Omega$ fixed.  In Series R, we vary $\\Sigma$ and $\\Omega$ together while holding $\\rho_*$ fixed.  Finally, in Series S the values of $\\Sigma$ and $\\rho_*$ are held constant while $\\Omega$ is varied.  Our main conclusions, and their relation to other recent work, are as follows:  1. {\\it Density, temperature, and pressure distributions}  We find that in spite of time-dependent effects, the density and temperature distributions of the gas retain bimodal profiles reminiscent of the classical \\citet{1969ApJ...155L.149F} two-phase model of the ISM. Although large-amplitude turbulence heats and cools via $P dV$ work and entropy production in shocks, most of the gas (by mass) remains near the curve in the pressure-density phase plane that is defined by radiative equilibrium: $n^2 \\Lambda - n \\Gamma=0$.  This is possible because the cooling time is generally shorter than the turbulent dynamical times, for our models. If turbulent compressions or expansions were more extreme in magnitude and also rapid in time compared to radiative times, then these adiabatic changes would lead to density-temperature pairs that more strongly departed from thermal equilibrium (see e.g. \\cite{2002ApJ...577..768S,2005A&A...433....1A,2008ApJ...683..786H} for a discussion of the dependence on various physical timescales involved). \\citet{2005A&A...433....1A}, \\citet{2005ApJ...629..849P}, and \\citet{2007ApJ...663..183P} similarly found that even with large-amplitude turbulence -- and no direct heating of the gas -- mass-weighted density and temperature PDFs have two (broadened) peaks, although for very high amplitude turbulence the trough tends to fill in (e.g. \\citealt{2007A&A...465..431H}). \\citet{2008ApJ...681.1148K} have emphasized the importance of maintaining sufficient numerical resolution, as numerical diffusion associated with flow across the grid broadens warm/cold interfaces, populating the thermally-unstable regime in the phase diagram. We note that bimodal character is most easily seen in mass-weighted rather than volume-weighted PDFs, although many of the results in the literature show only volume-weighted PDFs. Our results on the bimodal thermal distribution of gas are consistent with observations of atomic HI in the midplane of the Milky Way \\citep{2003ApJ...586.1067H}, which at the same time show interesting evidence of out-of-equilibrium gas, particularly at high latitudes.  2. {\\it Turbulence}  We find that appreciable turbulence can be excited in all components of the gas.  The values of the velocity dispersion in dense gas ($n>10^2 \\cm^{-3}$) of $\\sim 2-4\\ \\kms$ are similar to, though slightly lower than, those observed in Milky Way and other local-group GMCs (e.g. \\citealt{1987ApJ...319..730S, 2008ApJ...675..330S, 2008arXiv0807.0009B, 2008arXiv0809.1397H}).  For lower-density gas in our models, velocity dispersions are slightly higher, but still lower by a factor two compared to observed velocity dispersions of $\\sim 7-10~\\kms$ of Solar-neighborhood warm and cold HI seen in 21 cm emission and absorption \\citep[e.g.,][]{2003ApJ...586.1067H,2004JApA...25..185M}.  It is not surprising that the turbulence levels in our simulations are only moderate, given that we have included only one of the many sources of turbulence that is present in the ISM.  Turbulent driving in the ISM has been reviewed by e.g. \\citet{2004RvMP...76..125M,2004ARA&A..42..211E}. Supernova are widely-believed to be the most important source of turbulence for diffuse gas, and this has been demonstrated by numerical simulations \\citep[e.g.,][]{ 1999ApJ...514L..99K, 2005A&A...436..585D}. In the outer galaxy where star formation rates are low, driving by the magnetorotational instability may, however, dominate \\citep{1999ApJ...511..660S,2005ApJ...629..849P,2007ApJ...663..183P}. Spiral shocks are also effective in driving turbulence in the warm ISM, especially at high latitudes \\citep{2008ApJ...681.1148K}. The interaction between large-scale self-gravity, rotation, and shear can drive near-sonic turbulence at large scales \\citep{2002ApJ...577..197W,2007ApJ...660.1232K}, although the amplitude of this at scales less than the disk scale height may be modest. It is not known how effective these other mechanisms are for driving turbulence within GMCs, however, which are very dense and therefore present a small effective crossection to the diffuse ISM.   3. {\\it Feedback and the Toomre $Q$ parameter}  We have measured the Toomre parameter for each of our models, considering both the entire medium and just the dense gas, and comparing ``thermal-only'' with ``turbulent+thermal'' values. For all models, we find that the ``turbulent+thermal'' $Q$-value for dense gas is in the range $1-2$, and is much greater than the thermal-only value. A further interesting point is that the turbulence level evidently adjusts with surface density in order to reach a marginally-stable state. In Series Q and R, which have $\\Omega/\\Sigma=const.$, the velocity dispersions are relatively independent of $\\Sigma$, yielding marginally-stable $Q$ in the cold, dense gas. In Series K, which has constant $\\Omega$ and therefore is highly unstable at large $\\Sigma$ in the absence of turbulence, the velocity dispersions strongly increase with $\\Sigma$, as a consequence of much higher levels of feedback activity.  These higher velocity dispersions lift $Q$ to near unity. Thus, our simulations give direct evidence of feedback leading to a self-regulated quasi-steady state.  We note that $Q$ values vary for different components; \\citet{2002ApJ...577..197W} similarly found a large range of $Q$ when measured in local patches within their disk simulations.   Depending on what exactly is included in a model, the threshold for gravitational instability in previous non-turbulent simulations is measured to be at $Q\\sim 1.5$ (see \\citealt{2007ARA&A..45..565M}), which is similar to the values we find here when turbulence is included in $Q$. In the actively-star-forming regions of galaxies, measured values of the Toomre parameter are not constant, but show a limited range \\citep{2001ApJ...555..301M}. Evidence for star formation thresholds tied to $Q$ are more mixed \\citep{2001ApJ...555..301M, 2003MNRAS.346.1215B, 2007ApJS..173..524B}, possibly because star formation in outer disks primarily takes place in spiral arms \\citep{1998ApJ...506L..19F,2007ApJS..173..538T,2008ApJ...683L..13B} which strongly compress the gas above ambient densities.     4. {\\it The virial ratio}  We measure the virial ratio $\\cal R$ (eq. \\ref{eq:vir}) for all our models, separating into different density regimes.  We find that dense gas ($n>100\\ \\cm^{-3}$) generally has $\\cal R$ between $1 - 2$, whereas lower-density gas has large values of $\\cal R$. $\\cal R$ does not vary strongly with $\\Sigma$ in any of the series. In particular, we note that in spite of the large range of velocity dispersions in the dense gas in Series K at different $\\Sigma$, $\\cal R$ varies only weakly with $\\Sigma$.   This indicates that feedback can effectively regulate the dynamics within massive, dense clouds, independent of the larger-scale galactic environment. This is consistent with both older studies based on $^{12}$CO observations \\citep{1987ApJ...319..730S}, and recent studies based on $^{13}$CO observations \\citep{2008arXiv0809.1397H}, both of which find $\\cal R$ near $1-2$ for Milky Way GMCs. Although masses based on CO are less certain in external galaxies, virial ratios also likely near unity for GMCs observed in the Local Group \\citep{2008arXiv0807.0009B}  5. {\\it Cloud surface densities}  We estimate the surface density of typical clouds by measuring the mass-weighted vertically-integrated column of gas; we define this as $\\Sigma_{\\rm cloud}$.  For the values of the feedback threshold that we adopt, we find that $\\Sigma_{\\rm cloud}$ is in the range $70-150 ~\\msun \\pc^{-2}$ for most models.  This is comparable to the typical GMC surface densities that are observed in the Milky Way and in Local Group galaxies \\citep{1987ApJ...319..730S,2008ApJ...675..330S, 2008arXiv0807.0009B, 2008arXiv0809.1397H}.  Whether observational selection effects or physical processes impose a limited range of column densities for GMCs is an open question.  For example, magnetic fields may impose a minimum surface density for formation of gravitationally-bound structures in the ISM, at a value $\\Sigma = B/(2 \\pi\\sqrt{G})= 30 ~ \\msun \\pc^2 (B/10 \\mu {\\rm G})$ \\citep{2007ARA&A..45..565M}.  Here, we find that altering the volume heated in our HII region prescription does not appreciably change $\\Sigma_{\\rm cloud}$, but changing the gravitational potential threshold for star formation feedback does:  when potential thresholds are low, $\\Sigma_{\\rm cloud}$ is also low.   Comparison of cloud properties with observations may turn out to be a much more critical test of whether an ISM model is realistic than some other measures, such as the star formation rate.    6. {\\it Dependence of $\\Sigma_{\\rm SFR}$ on $\\Sigma$}  We obtain estimates of the dependence of surface star formation rate $\\Sigma_{\\rm SFR}$ on gas surface density $\\Sigma$ in our models by assuming that the timescale for star formation is proportional to the free-fall time in gas above some density threshold $n_{\\rm th}=\\rho_{\\rm th}/\\mu$.  We compare results for two different threshold densities, $n_{\\rm th} = 100\\ \\cm^{-3}$ and $n_{\\rm th} = 10^3\\ \\cm^{-3}$.  For Series Q and R (which most resemble real galaxies), we find (for either choice of threshold) relations that are well-described by power laws: $\\Sigma_{\\rm SFR}\\propto \\Sigma^{1+p}$ with $1+p=1.2-1.4$.  In Series K, a power law is a less-good fit; the slope is also shallower (closer to unity).   These results are consistent with empirical Kennicutt-Schmidt laws \\citep{1959ApJ...129..243S,1963ApJ...137..758S,1998ARA&A..36..189K}, which show similar values of $1+p$ when all the gas is included in $\\Sigma$ (e.g. \\citealt{1989ApJ...344..685K,1998ApJ...498..541K,2002ApJ...569..157W, 2007A&A...461..143S,2007ApJ...671..333K}). Recent work has suggested that $1+p$ is close to unity if just CO-emitting molecular gas is included \\citep{Bigiel2008} ; this implies all CO-emitting gas (most of which is at $n=10^2-10^3 \\cm^{-3}$) has the same star formation rate independent of galactic environment. Our prescription that the star formation rate per unit mass of dense gas is constant is equivalent to empirically finding $1+p=1$ if only molecular gas is included.   It is encouraging that the results we find for Kennicutt-Schmidt relations in our disk models with feedback are compatible with observations. We also find, however, that star formation rates predicted from hydrostatic models are in fact similar to those predicted from the hydrodynamic models, with similar slopes at large $\\Sigma$.  This is true even though the hydrostatic models are not at all like real galaxies. Thus, one must be cautious in considering a numerical model successful if it yields reasonable star formation rates, since this can simply be a consequence of choosing reasonable initial conditions in a simulation.  Indeed, a number of recent numerical studies have found results similar to observed Kennicutt-Schmidt laws, regardless of the detailed physics that they included in the models (e.g. \\citealt{2006ApJ...639..879L,2006ApJ...641..878T, 2008ApJ...673..810T,2008ApJ...680.1083R}). \\citet{2008MNRAS.383.1210S} have also recently emphasized that reproducing empirical star formation scaling laws is not by itself a critical test of an ISM model.    7. {\\it Density-dependence of star formation efficiency}  The star formation efficiency per free-fall time can be defined locally as a function of threshold density by $\\epsilon_{\\rm ff}(\\rho_{\\rm th}) = t_{\\rm ff}(\\rho_{\\rm th}) \\Sigma_{\\rm SFR}/\\Sigma(\\rho> \\rho_{\\rm th})$; corresponding global measures can also be obtained.  For a given true star formation rate, $\\epsilon_{\\rm ff}(\\rho_{\\rm th}) \\propto \\tau_{\\rm SF}$ where the scaled star formation time $\\tau_{\\rm SF}$ is shown for two different threshold densities in Fig. \\ref{fig:tSF}. We find that $\\tau_{\\rm SF}$ (or $\\epsilon_{\\rm ff}(\\rho_{\\rm th})$) decreases with increasing $\\rho_{\\rm th}$.   This is not a strong effect, however: it is less than a factor 2 for an order of magnitude difference in $\\rho_{\\rm th}$.  In Series Q and R, the ratio of efficiencies at different threshold densities are also independent of $\\Sigma$ (although this is not true for Series K).  \\citet{2007ApJ...654..304K} recently compiled a range of observations of $\\epsilon_{\\rm ff}$ (which they refer to as SFR$_{\\rm ff}$).  They point out that for threshold densities above $\\sim 100 \\cm^{-3}$, the value of $\\epsilon_{\\rm ff}$ does not vary strongly with density.  They also find a smaller value for gas traced by HCN than for gas traced by CO, which since HCN has a higher critical density than CO is consistent with our finding that $\\epsilon_{\\rm ff}$ decreases with increasing $\\rho_{\\rm th}$.   Even though scaled star formation times at high density vary together independent of $\\Sigma$, the free-fall time at the vertically-averaged density is {\\it not} proportional to $\\tau_{\\rm SF}$.  Instead, we find that $\\tau_{\\rm SF}/t_{\\rm ff}(\\rho_{\\rm ave})$ increases at increasing $\\Sigma$.  This implies that a prescription for star formation based on the mean density within one ``average'' scale height in a disk (or from a simulation that does not resolve high-density gas) would increasingly overestimate the star formation rate at high $\\Sigma$. The same is true for the orbital time and the Jeans time:  an assumption that the star formation rate varies $\\propto \\Sigma/t_{\\rm orb}$ or $\\propto \\Sigma/t_{\\rm J}$ would increasingly overestimate $\\Sigma_{\\rm SFR}$ at high $\\Sigma$ compared to the value obtained from measuring the mass of gas in dense, gravitationally-bound regions.  Thus, ISM models must resolve self-gravitating structures at scales less than the disk thickness in order to make accurate predictions of the star formation rate."
    },
    "0812/0812.3809.txt": {
        "abstract": "This paper presents a review of the history, motivation and current status of high energy neutrino telescopes. Many years after these detectors were first conceived, the operation of kilometer-cubed scale detectors is finally on the horizon at both the South Pole and in the Mediterranean Sea. These new detectors will perhaps provide us the first view of high energy astrophysical objects with a new messenger particle and provide us with our first real glimpse of the distant universe at energies above those accessible by gamma-ray instruments. Some of the topics that can be addressed by these new instruments include the origin of cosmic rays, the nature of dark matter, and the mechanisms at work in high energy astrophysical objects such as gamma-ray bursts, active galactic nuclei, pulsar wind nebula and supernova remnants. ",
        "introduction": "Neutrino astronomy was born with the first detection of low energy neutrinos from an astrophysical source, the Sun, by the pioneering work of Ray Davis and John Bahcall. Using a  100,000 gallon tank of perchloroethylene in the Homestake mine, Ray Davis and his group painstakingly extracted and counted Argon molecules created through inverse beta decay by solar neutrinos captured on Chlorine atoms as they traversed the tank.   The flux of neutrinos inferred from this careful measurement was a mere 1/3 of that calculated by Bahcall for models of solar fusion~\\cite{davis, bahcall}. The disagreement between the measurements of Davis and the predictions of Bahcall led to what was known as the \u00d2solar neutrino problem\u00d3.  In the years since, ongoing inquiries into the deficit of solar neutrinos by the Super-Kamiokande and SNO observatories,  along with radiochemical measurements using gallium, have yielded surprising revelations about the very nature of the messengers themselves~\\cite{SK, SNO, sage, gallex}.  These collaborations have established that solar neutrinos undergo flavor oscillations, explaining the discrepancy first suggested by Davis' observations decades before, and showing that neutrinos are not the massless  particles once thought~\\cite{sk-atm}.  %The Super-Kamiokande and SNO experiments relied on the detection of secondary leptons, produced when neutrinos elastically scattered off target electrons or capture on a target nucleus, traverse an optically transparent medium at high energies, producing Cherenkov radiation that could be sensed by photomultiplier tubes installed within the medium. This enabled Super-Kamiokande in particular, which relied on elastic scattering for solar neutrinos, to make a high statistics measurement of the solar neutrino flux with directional information, establishing that the neutrinos were coming from the Sun. The water Cherenkov technique, which had also previously been employed in the search for proton decay and measurement of the atmospheric neutrino flux, has become the mainstay of modern neutrino astronomy. In fact, the directional capability of the water Cherenkov technique had enabled the Super-Kamiokande collaboration to use atmospheric neutrinos to show the first definitive evidence for neutrino flavor oscillations and hence nonzero neutrino mass in 1995.~\\cite{sk-atm}  The first observation of neutrinos from a source outside of our solar system came quite fortuitously on February 23, 1987 at 7:35 UT when a burst of neutrinos from the collapse of a blue supergiant in the Large Magellanic Cloud, was simultaneously detected by the IMB~\\cite{IMB} and Kamiokande-II~\\cite{ka1987} detectors.  A total of 19 neutrinos were observed within a time interval of about 6 seconds. %in excellent agreement with models of neutron star formation that predicted a sudden, intense release of energy via the reaction $e^- + p \\rightarrow \\nu_e + n$  at the time of the core collapse. It was sighted in the optical three hours later. This event, which became known as Supernova 1987A, was the nearest observed supernova explosion to earth in nearly 400 years and therefore provided a rare opportunity, not only to confirm models of supernova formation, but to constrain basic properties of neutrinos, including their masses and lifetimes.  This again established the unique potential of neutrino telescopes, both as astrophysical observatories and as long baseline neutrino detectors.   %It has been over thirty years since the pioneering ideas for using neutrinos as messengers to study the universe at the highest energies were conceived. By the time Supernova 1987A was observed, the idea of using neutrinos as messengers to study the universe at the highest energies had already been conceived.  In the ensuing years, the technology and methods required to build an array sensitive enough for extragalactic observations were developed through the construction and operation of a number of arrays on an intermediate scale. The first detectors on the scale of a cubic kilometer, which are large enough to be sensitive to the expected flux of astrophysical neutrinos, are now nearing completion and are coming into operation. These new neutrino telescopes are truly discovery instruments. They hold the promise of opening a new window on the high energy universe by extending the observable energy range for distant and optically opaque objects and by using a new particle to complement the many observations using gamma rays.  This paper begins by explaining the motivation to use neutrinos as astrophysical messengers at high energies, the connection of the expected neutrino flux to the known cosmic ray flux, and why detectors on the scale of a kilometer are necessary to reach a sensitivity of astrophysical interest. In section two, the development of the techniques used in neutrino astronomy and the performance requirements common to all detectors of this type are described followed by the status of the major neutrino telescopes now in operation or under construction. Section three presents the results from current experiments along with the expected sensitivity for the new detectors that are coming on line. %add comment about future prospects here     \\subsection{Why neutrinos and why a kilometer cubed?}  One of the most compelling arguments for building large scale neutrino observatories is to shed light on the mystery of cosmic rays. Since their discovery almost 100 years ago, the origin of cosmic rays has presented a puzzle. The energy spectrum and flux of charged particle cosmic rays bombarding the Earth has now been extensively measured, and the range of the observed energies is truly astounding, with a handful of events recorded with an energy in excess of $10^{20}$ eV (Figure 1 right). These measurements have established the fact that there are cosmic charged particle accelerators. Yet, to this day, no unambiguous source of hadronic acceleration has been found.  \\begin{figure}[htbp] %  figure placement: here, top, bottom, or \\includegraphics[width=3in]{figure1a.png} \\includegraphics[width=3.5in]{figure1b.png} \\caption{ Left:  Messenger particles energy vs distance they propagate. Photons (blue) are attenuated by interaction with the CMB and IR background radiation. Protons (red) interact with the CMB while undergoing magnetic deflection at lower energies~\\cite{nu_prop}.  Right: The spectrum of cosmic rays.  The flux is shown as a function of energy~\\cite{cr_spec}. } \\label{fig:messengers} \\end{figure}   What information we do have on the astrophysical sources themselves has come from astronomical observations using photons which, because they are neutral, traverse space without bending and point back to their source. These measurements, which extend from the radio and infrared through high energy gamma-rays, have provided invaluable information in our attempt to assemble a comprehensive theory on the origin of cosmic rays and the acceleration mechanisms at work in high energy astrophysical objects. However, using gamma ray measurements alone, it is difficult to determine if the gamma rays are produced only by relativistic electrons or from hadronic processes that could also be the source of the galactic cosmic rays, since high energy electrons will produce synchrotron radiation and can transfer their energy to ambient photon fields through inverse Compton scattering.   %The observation of neutrinos from these galactic sources, such as those observed by Milagro~\\cite{Milagro} and HESS~\\cite{HESS}, would be unambiguous evidence for hadronic acceleration. The observation of neutrinos from a source would be unambiguous evidence for hadronic acceleration, since cosmic proton accelerators are expected to produce cosmic rays, gamma rays and neutrinos with similar luminosities.  This prediction arises from the conditions required for acceleration; a massive bulk flow of relativistic charged particles and a magnetic field to confine the charged particles within the acceleration region.  Candidate acceleration sites might be powered by large gravitational forces such as those found near black holes. Electrons accelerated in these regions would loose energy through synchrotron radiation thus, if protons were also being accelerated at the site, they would collide with the resulting radiation field,  producing mesons in processes such as \\begin{eqnarray*} p+ \\gamma &\\rightarrow \\Delta^+ \\rightarrow \\pi^0 + p \\\\ &\\rightarrow \\Delta^+ \\rightarrow \\pi^+ + n.\\\\ \\end{eqnarray*} If neutral pions are produced, they will decay to high energy gamma rays whose energy will be shifted downward through pair production as they escape the source. Charged pions will decay mainly to a muon and a muon neutrino.  Although the protons may remain confined in the magnetic field, the neutrons and neutrinos  may escape, with the neutron subsequently decaying into a proton, an electron and a neutrino.  It is also possible to imagine a ``hidden source'' which acts as a beam dump from which only neutrinos could escape, however, such sources would not contribute to the observed cosmic ray flux.  In recent years experiments such as Whipple, Milagro~\\cite{Milagro} and HESS~\\cite{HESS} have discovered a number of sources of very high energy gamma rays, both galactic sources and distant sources such as Active Galactic Nuclei.  Many of the gamma-ray emitters in our galaxy are associated with pulsar wind nebula and supernova remnants, which are generally believed to be the source of galactic cosmic rays. %Using gamma ray measurements alone, it is difficult to determine if the gamma rays are produced only by relativistic electrons or from hadronic processes that could also be the source of the galactic cosmic rays. Estimates for neutrino fluxes for these galactic sources indicate they may be detectable by cubic-kilometer scale detectors~\\cite{beacom, halzen2}.  For astrophysical sources outside our galaxy the pervasive cosmic microwave background (CMB) and infrared (IR) radiation limits the propagation of gamma-rays and protons as messenger particles (Figure~\\ref{fig:messengers} left).  Photons of a high enough energy will pair produce on the CMB and IR radiation limiting the horizon for TeV gamma rays to tens of Megaparsec.  Although protons can penetrate larger distances, because they carry an electric charge, the trajectories of protons are deflected by the magnetic field of our galaxy and do not point back to their sources.  However, at the high end of the proton energy spectrum  (above $\\approx 10^{18}$ eV) this deflection becomes small enough that proton astronomy becomes possible. Currently operating extensive air shower arrays such as HiRes and Auger have exploited this fact in looking for the sources of ultra high energy cosmic rays.  Unfortunately, at these energies proton interaction with the CMB (the GZK effect~\\cite{GZK1,GZK2}) again results in pion production and limits the propagation distance to tens of Megaparsec.  A steep sharpening of the cosmic ray energy spectrum has recently been observed by both the HiRes~\\cite{hires} and Auger~\\cite{auger} observatories at around $5 \\times 10^{19}$ eV, which corresponds to the expected energy of the GZK cutoff (red line in Fig~\\ref{fig:messengers}).  This could indicate that some of  the sources of the high energy cosmic rays have reached their maximum acceleration energy, or it could indicate that the highest energy cosmic rays are indeed  extragalactic in origin and the energy of the observed proton spectrum is being shifted downward as they propagate over long distances.  This raises several confounding questions.  Could the spectrum of energies produced at the source extend even higher?  Could the high energies cosmic rays we see be the product of decaying pions produced on the CMB? Since neutrinos interact only weakly, they can propagate indefinitely and still point back to their source.  In addition, as for the galactic sources, the presence of neutrinos is the hallmark of hadronic acceleration, providing insight into what fuels these most powerful astrophysical engines. %Measuring the flux of neutrinos at high energies could determine whether or not the GZK cutoff has been observed. Whatever the origin of cosmic rays, the acceleration of charged particles to the observed energies requires enormous gravitational forces such as those found in accreting black holes at the center of active galactic nuclei (AGN) or at the core of supernova remnants (SNR).  Another candidate for cosmic ray acceleration sites are gamma-ray bursts (GRB).  Conventional models describe a rapidly expanding fireball of unknown origins where gamma rays in the fireball collide with high energy protons, creating a neutrino afterglow through similar mechanisms to those described above.  In fact, it is a tantalizing coincidence that the energy injection rate of photons integrated over all gamma-ray bursts is very similar to the energy injection rate of ultra high energy cosmic rays, (UHECRs) leading many to speculate that gamma-ray bursts and the source of UHECRs may be one and the same~\\cite{waxman, halzen1}.   \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=3.7in]{figure2.jpg} \\end{center} \\caption{ The Waxman-Bahcall upper limit of the cosmological neutrino flux (dot dashed lines) based on the observed cosmic ray spectrum.  The solid lines labeled ``Jet'' show the predicted flux based on gamma-ray emission from blazars, while the GRB line is based upon models of GRB emission.  The dashed lines show predictions for neutrinos at the core of AGNs from which the parent protons cannot escape.~\\cite{wb}} \\label{fig:bound} \\end{figure}  %If these objects are indeed the source of high energy cosmic rays, it is hard to imagine the particles could escape such a dense environment without interacting.  These cosmic beam dumps would therefore be the source of both high energy gamma rays and neutrinos, as the products of decaying pions exit the source.  %John Bahcall and collaborator Eli Waxman made a series of calculations connecting the expected diffuse flux of neutrinos to the observed high energy cosmic ray spectrum\\cite{wb, wb2}.  Earlier models of neutrino production in the jets emanating from active galactic nuclei (blazars) assumed that protons in the jet underwent Fermi shock acceleration, producing a proton spectrum of   $dN_p/dE_p \\propto E^{-2}_{p}$  extending up to an energy of $10^{19}$  eV and that pion production from protons interacting on the photon spectrum, which would follow a similar power law structure, results in a harder neutrino spectrum which falls with energy as $E^{-1}$. Since roughly half of the pions produced will be neutral and decay into gamma rays, and the other half will be charged producing muon neutrinos, assumed to carry 5\\% of the energy of the parent proton, an estimation of the neutrino flux, shown by the lines labeled ``Jet'' in Figure~\\ref{fig:bound}, can be deduced from the observed gamma-ray emission from blazars.  Waxman and Bahcall assumed that some fraction of the energy of the protons accelerated in the source lost their energy to pion photoproduction, and calculated the upper limit of the diffuse neutrino flux by working backward from the observed cosmic ray spectrum.  This robust upper limit is shown by the dash-dotted line in Figure~\\ref{fig:bound}, and differs from the jet models by up to two orders of magnitude.  This calculation does not take into account the possibility of neutrinos produced at the core of the AGN near the central black hole, where the parent protons could not escape (Figure~\\ref{fig:bound} dashed line).  The line labeled ``GRB'' is their calculation of neutrino flux based on a model of gamma-ray bursts (GRB) and is consistent with their upper bound. The calculations of Waxman and Bahcall set a benchmark for the flux sensitivity of the current generation of neutrino telescopes.  Even if fluxes are not sufficient to pinpoint individual sources of neutrinos, a study of the flux and energy spectrum of the diffuse glow of neutrinos integrated over the sky will provide constraints on models of high energy cosmic ray production.  A neutrino observatory with an instrumented area of 1 square kilometer will be sensitive to fluxes well below Waxman and Bahcall's upper limit.  If the same acceleration mechanisms are responsible for the observed spectrum of cosmic rays and gamma rays, the connection between the expected luminosities and neutrinos can be used to predict the neutrino flux, and ultimately, constrain proposed models of high energy phenomena such as AGNs and GRBs.  John Bahcall and collaborator Eli Waxman made a series of calculations connecting the expected diffuse flux of neutrinos to the observed high energy cosmic ray spectrum\\cite{wb, wb2}. They claimed that the cosmic ray spectrum placed the most stringent limits on the flux of cosmic neutrinos from ``transparent sources\", those from which neutrons may escape.   Note that the cosmic ray spectrum does not limit the flux from hidden sources.  Waxman and Bahcall  assumed a proton spectrum of   $dN_p/dE_p \\propto E^{-2}$ at the source, since this spectrum is typical of Fermi shock acceleration.  They then calculated the energy production rate in protons at the source assuming that the same sources produced the observed spectrum of cosmic rays.  From particle physics, they were then able to deduce the expected energy density of muon neutrinos.  Their upper bound is shown by the dash-dotted line in Figure~\\ref{fig:bound}.  This limit contradicted several AGN models labeled ``Jet\" in the same figure. Noting that most of the extragalactic gamma ray energy is contained in a diffuse glow rather rather than from discrete sources suggests another approach; assuming a particular spectrum at the source, and normalizing it to the observed diffuse flux of gamma rays. More recently, a  calculation of the expected neutrino flux based on the cosmic ray spectrum has been made by Mannheim, Protheroe, and Rachen~\\cite{mpr} which did not assume a particular injection spectrum but instead assumed a specific spectrum for the observable cosmic ray flux, taking into account  production, source evolution, and propogation effects, that yielded a more stringent limit than those of Waxman and Bahcall.  In any case, a neutrino observatory with an effective area of 1 ${\\rm km^3}$ will be sensitive to diffuse glows of neutrinos that are well below the upper bounds calculated by any of these methods.   In summary, the motivation for pursuing high energy neutrino astrophysics is the unique and complementary role that neutrinos could play in our observation of the high energy universe. The observation of neutrino emission from astrophysical sources would provide unambiguous evidence for hadronic acceleration and the source of cosmic rays. The neutrino flux estimates discussed above indicate that neutrino detectors of cubic kilometer scale are necessary to approach the sensitivity needed. Furthermore, neutrinos propagate through the universe almost undisturbed to the highest energies without deflecting in  magnetic fields and are therefore usable as astronomical messenger particles at energies where gamma-rays or protons are not, and therefore hold the promise of opening a new energy window on the high energy universe. ",
        "conclusions": "It is truly an exciting time in particle astrophysics. Building on the long history of cosmic-ray and gamma-ray measurements, new experiments are coming on line that significantly extend our ability to observe the highest energy astrophysical objects, giving us new insight into the mechanisms at work in these objects and perhaps into the old question of the source of the cosmic rays. These new instruments include high energy cosmic ray detectors, gamma-ray detectors on the surface and in space, and neutrino telescopes with sensitivities that are of astrophysical interest.  Neutrino telescopes have a unique and complementary role to play in this new era of particle astronomy with their ability to look for distant sources at energies not accessible to gamma-ray instruments, and because neutrinos, as new messengers, can provide unique information on the acceleration mechanisms at work in these high energy astrophysical sources. The next generation neutrino observatories will give us an unprecedented sensitivity for sources in both the northern and southern skies. With detectors on the scale of a cubic kilometer operating at South Pole in the near future, and planned for the Mediterranean soon after, we are optimistic that more than thirty years since they were first conceived, we will finally open this new window on the high energy universe.     \\ack The author wises to thank John Carr for his valuable input on the status of the ANTARES array, and Greg Sullivan and Jordan Goodman for their helpful comments on this manuscript. %We also wish to thank our collaborators Gary Hill and Francis Halzen for helpful discussions. This work was supported by the National Science Foundation  through grant No. PHY-0502709."
    },
    "0812/0812.4839_arXiv.txt": {
        "abstract": "We present a new multi--fluid, grid MHD code PIERNIK, which is based on the Relaxing TVD scheme~\\cite{Jin95}. The original scheme (see Trac \\& Pen~\\cite*{trac} and Pen~\\etal~\\cite*{pen}) has been extended by an addition of dynamically independent, but interacting fluids: dust and a diffusive cosmic ray gas, described within the fluid approximation, with an option to add other fluids in an easy way.  The code has been equipped with shearing--box boundary conditions, and a selfgravity module, Ohmic resistivity module, as well as other facilities which are useful in astrophysical fluid--dynamical simulations. The code is parallelized by means of the MPI library. In this paper we present Ohmic resistivity extension of the original Relaxing TVD MHD scheme, and show examples of magnetic reconnection in cases of uniform and current--dependent resistivity prescriptions. ",
        "introduction": "\\enlargethispage{\\baselineskip} Piernik MHD code is capable of dealing with resistive MHD equations \\linebreak \\cite{pawlaszek-08}, which involve the resistive dissipation term in the induction equation \\begin{equation} \\label{induction_equation} \\partial_t \\mathbf{B} = \\nabla\\times \\left(\\mathbf{v}\\times\\mathbf{B} \\right)-\\nabla\\times(\\eta\\mathbf{J}), \\end{equation} where $\\mathbf{J}$ is the electric current density. Due to the conservative nature of the basic MHD algorithm, any decrement of magnetic energy during the update of magnetic field results in a corresponding increment of thermal energy.  Therefore, no additional source term related to resistivity is needed in the energy equation. \\par The resistivity algorithm of PIERNIK code relies on additional terms of resistive origin in the electromotive force, supplemented to the original Constraint Transport (CT) scheme implemented by Pen~\\etal~\\cite*{pen} within the Relaxing TVD scheme. In the present version of PIERNIK we have implemented a current--dependent resistivity, according to Ugai~\\cite*{ugai} \\begin{equation} \\eta=\\eta_{1}+\\eta_{2}(J^2-J^2_{\\indx{crit}})\\Theta(|\\mathbf{J}|-|\\mathbf{J}_{\\indx{crit}}|), \\end{equation} where $\\eta_{1}$ is a constant corresponding to uniform resistivity, $\\eta_{2}$ represents the anomalous resistivity, which switches on when the current density exceeds a certain critical value $J_{\\indx{crit}}$,  and $\\Theta$ is the Heaviside step function. ",
        "conclusions": ""
    },
    "0812/0812.0247_arXiv.txt": {
        "abstract": "We study different stages of the neutron star cooling by computing neutron star properties at various temperatures and entropies using an effective chiral model including hadronic and quark degrees of freedom. Macroscopic properties of the star such as mass and radius are calculated and compared with observations. It can be seen that the effects of chiral restoration and deconfinement to quark matter in the core of the neutron star at different stages of the evolution can be significant for the evolution of the star and allow insight into the behavior of matter at extreme densities. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.2074_arXiv.txt": {
        "abstract": "A new diagnostic method, $Om$   is applied to generalized Chaplygin gas (GCG) model as the unification of dark matter and dark energy. On the basis of the recently observed data: the Union supernovae, the observational Hubble data, the SDSS baryon acoustic peak and the five-year WMAP shift parameter, we show the discriminations between GCG and $\\Lambda$CDM model. Furthermore, it is calculated that the current equation of state  of dark energy  $w_{0de}=-0.964$ according to GCG model. ",
        "introduction": "$}   {\\small {~~~}}~The type Ia supernovae (SNe Ia) investigations \\cite{SNe}, the cosmic microwave background(CMB) results from WMAP \\cite{CMB} observations, and surveys of galaxies \\cite{LSS} all suggest that the expansion of present universe is speeding up rather than slowing down. The accelerated expansion of the present universe is usually attributed to the fact that dark energy (DE) is an exotic component with negative pressure. Many kinds of DE models have already been constructed such as $\\Lambda$CDM \\cite{LCDM}, quintessence \\cite{quintessence}, phantom \\cite{phantom}, quintom \\cite{quintom}, generalized Chaplygin gas (GCG) \\cite{GCG}, modified Chaplygin gas \\cite{MCG}, holographic dark energy \\cite{holographic}, agegraphic dark energy\\cite{agegraphic}, and so forth. In addition, model-independent method\\footnote{For example, using mathematical fundament  one expands  equation of state  of DE $w_{de}$ or deceleration parameter $q$ with respect to scale factor $a$ or redshit $z$, such as  $w_{de}(z)=w_{0}$=const \\cite{independent1}, $w_{de}(z)=w_{0}+w_{1}z$\\cite {independent2}, $w_{de}(z)=w_{0}+w_{1}\\ln(1+z)$ \\cite{independent3}, $w_{de}(z)=w_{0}+\\frac{w_{1}z}{1+z}$ \\cite{independent4}, $q(z)=q_{0}+q_{1}z$ \\cite{independent1}, $q(z)=q_{0}+ \\frac{q_{1}z}{1+z}$ \\cite{independent5}, and so forth. Where $w_{0}$, $w_{1}$, or $q_{0}$, $q_{1}$ are model parameters. For more information about model-independent method, please see review paper \\cite{review}. } and modified gravity theories (such as scalar-tensor cosmology \\cite{Scalar}, braneworld models \\cite{braneworld})  to interpret accelerating universe have also been discussed. So, a general and model-independent manner to distinguish these models introduced by different theories or methods is necessary. Statefinder diagnostic method is presented in Refs. \\cite{rs}, and it has been applied to a large number of DE models \\cite{rspapers}. Recently, another geometrical diagnostic $Om$ is also introduced in Ref. \\cite{Om} to differentiate $\\Lambda$CDM with other models. An important property for $Om$ diagnostic is that it  can be used to distinguish DE models with small influence from density parameter $\\Omega_{0m}$, though the current observations suggest an uncertainties of at least 25\\% in the value of current matter density $\\Omega_{0m}$ \\cite{5yWMAPSNeBAO}.  In this paper, we apply $Om$  diagnostic to GCG model.  The paper is organized as follows. In section 2, the GCG model  as the unification of dark matter and dark energy is introduced briefly. Based on the recently observed data: the Union SNe Ia \\cite{307Union}, the observational Hubble data (OHD) \\cite{OHD}, the baryon acoustic oscillation (BAO) peak from Sloan Digital Sky Survey (SDSS) \\cite{SDSS} and the five-year WMAP CMB shift parameter \\cite{5yWMAP},  $Om$ diagnostic  is used to GCG model  in section 3. Section 4 is the conclusions. ",
        "conclusions": "$} \\label{section4}  On the basis of the recently observed data: the Union SNe Ia data, the nine observational Hubble data, the SDSS baryon acoustic peak and the five-year WMAP result, we apply a geometrical diagnostic $Om$ to distinguish GCG model and $\\Lambda$CDM model. From Fig. \\ref{Omwdefigure}, it is shown that the larger error for  the evolution of $w_{de}(z)$ may be produced by the erroneous estimation of matter density $\\Omega_{0m}$. And the $Om(z)$ is a better quantity than $w_{de}(z)$ to truly distinguish DE models and to show the properties of DE. According to the $Om$ diagram, it is easy to see that for the constraint from the single standard candle data, the difference between GCG model and $\\Lambda$CDM model is obvious, while for the combined constraint, the best fit evolutions of $Om(z)$ for them are similar and the difference between these two models is not clear at $1\\sigma$ confidence level. In addition, we also calculate the value of current EOS of DE, $w_{0de}=-0.964$, by using the value of $Om(0)$ for GCG model. Here the $Om(z)$ diagram is not sensitive to the variation of  density parameter.   \\textbf{\\ Acknowledgments } The research work is supported by NSF (10703001) of PR China."
    },
    "0812/0812.0592_arXiv.txt": {
        "abstract": "We present the results of a stellar population analysis of 30 Lyman alpha emitting galaxies (LAEs) at z $\\sim$ 0.3, previously discovered with the {\\it Galaxy Evolution Explorer} ({\\it GALEX}). With a few exceptions, we can accurately fit model spectral energy distributions to these objects, representing the first time this has been done for a large sample of LAEs at z $<$ 3, a gap of $\\sim$ 8 Gyr in the history of the Universe.  From the 26/30 LAEs which we can fit, we find an age and stellar mass range of 200 Myr -- 10 Gyr and 10$^{9}$ -- 10$^{11}$ $M$\\sol, respectively.  These objects thus appear to be significantly older and more massive than LAEs at high-redshift.  We also find that these LAEs show a mild trend towards higher metallicity than those at high redshift, as well as a tighter range of dust attenuation and interstellar medium geometry.  These results suggest that low-redshift LAEs have evolved significantly from those at high redshift. ",
        "introduction": "Over the past decade thousands of Lyman alpha emitting galaxies (LAEs) have been discovered from 2.4 $<$ z $<$ 6.96 (e.g., Cowie \\& Hu 1998; Rhoads et al. 2000; Rhoads et al. 2001; Gawiser et al. 2006a; Kashikawa et al. 2006; Ouchi et al. 2008).  It was predicted that strong Ly$\\alpha$ emission would be a signpost of primitive galaxies in formation, and thus these objects could represent some of the first galaxies in the Universe (Partridge \\& Peebles 1967).  However, with the new results coming from stellar population modeling, the identification of LAEs as primitive galaxies has come into doubt.  Numerous studies have begun to detect dust in LAEs, which is an indicator that a prior generation of stars has lived and died, thus making these objects not primitive (Chary et al. 2005; Lai et al. 2007; Pirzkal et al. 2007; Nilsson et al. 2007; Kuiper et al. 2009; Yuma et al. 2009; Finkelstein et al. 2008, 2009).  While the presence of dust was originally proposed as the reason why LAEs were not detected 20 years ago, it is also possible that the dust can help enhance the value of the \\lya~EW by selectively attenuating the continuum (Neufeld et al. 1991; Haiman \\& Spaans 1999; Hansen \\& Oh 2006).  In our previous work (Finkelstein et al. 2008, 2009), we demonstrated that dust enhancement from a clumpy interstellar medium (ISM) can help explain the spectral energy distributions (SEDs) of a sample of z $\\sim$ 4.5 LAEs, with clumpy dust appearing to exist in 10 out of a sample of 14 LAEs, even though 12/14 of these LAEs were young ($<$ 15 Myr).  We also found plausible evidence for dust in every object in our sample, although their derived metallicities were rather low (with half of the sample having Z $<$ 0.2 $Z_{\\odot}$).  This suggests that these objects are not primordial, as a previous generation of stars was likely responsible for the creation of the dust.  While we have estimates for these fundamental physical parameters of distant LAEs, it becomes interesting to study them at lower redshift where we can observe them in greater detail.  However, large samples of LAEs have only been studied down to z $\\sim$ 3.1 (Gawiser et al. 2006a; McLinden et al. 2009; Nakamura et al. 2009), with a few recent samples currently under analysis at z $\\sim$ 2 (Nilsson et al. 2008; Reddy et al. 2008).  Given the loss of CCD quantum efficiency at $\\sim$ 3600 \\AA~(and the atmospheric ultraviolet (UV) cut-off at $\\sim$ 3100 \\AA), detection of LAEs at z $\\lesssim$ 2 is not possible from the ground.  However, LAEs could be detected using the space-based UV telescope, {\\it Galaxy Evolution Explorer} ({\\it GALEX}), at z $\\lesssim$ 1.  Deharveng et al. (2008) recently published the discovery of 96 LAEs at 0.2 $<$ z $<$ 0.35 via {\\it GALEX} spectroscopy, using data from the publicly available {\\it GALEX} GR2 data release.  We present the results from a stellar population modeling analysis of 30 of these objects, with an emphasis on their differences from high-redshift LAEs.  Where applicable, we assume cosmology with $\\Omega_{m}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7 and H$_{0}$ = 70 km s$^{-1}$ Mpc$^{-1}$ (c.f. Spergel et al. 2007).  All magnitudes in this paper are listed in AB magnitudes (Oke \\& Gunn 1983). ",
        "conclusions": "We have performed stellar population model fitting on a sample of 30 GALEX discovered z $\\sim$ 0.3 \\lya~emitting galaxies.  We derived the stellar population ages and masses, as well as  metallicities, dust extinction, and ISM geometries.  The age and mass results imply that low-redshift LAEs are both older and more massive than their high-redshift counterparts, with the youngest low-redshift LAE having a best-fit age of $\\sim$ 200 Myr, much older than the average age at high redshift.  We also found that the metallicity in LAEs is higher at low redshift, showing evolution of the mean metallicity of the stars inside these galaxies, as would be expected over a period of $\\sim$ 9 Gyr of cosmic time.  Perhaps most interesting, we found that the ISM at low redshift appears to be very stable, in a quasi-clumpy state.  At high redshift, we found that the ISM in LAEs ran the whole range from extremely clumpy to very homogeneous.  At z $\\sim$ 0.3, we found that while a few galaxies were fairly clumpy, most resided in the regime of q $\\sim$ 1, where dust will not affect the \\lya~EW.  This result is confirmed by the study of local LAE analogs by Atek et al (2008).  Our results show that LAEs at z $\\sim$ 0.3 show significant differences from their high-redshift cousins, although further study of a larger sample will help confirm this.  Although low-redshift LAE studies are difficult as they need to be done from space, help should be on the way with the near-UV sensitive Wide Field Camera 3 (WFC3) due to be installed on the {\\it Hubble Space Telescope}."
    },
    "0812/0812.0586_arXiv.txt": {
        "abstract": "We describe 2D hydrodynamic simulations of the migration of low-mass planets ($\\leq 30 M_{\\oplus}$) in nearly laminar disks (viscosity parameter $\\alpha < 10^{-3}$) over timescales of several thousand orbit periods.  We consider disk masses of $1$, $2$, and $5$ times the minimum mass solar nebula, disk thickness parameters of $H/r = 0.035$ and $0.05$, and a variety of $\\alpha$ values and planet masses.  Disk self-gravity is fully included.  Previous analytic work has suggested that Type~I planet migration can be halted in disks of sufficiently low turbulent viscosity, for $\\alpha \\sim 10^{-4}$.  The halting is due to a feedback effect of breaking density waves that results in a slight mass redistribution and consequently an increased outward torque contribution.  The simulations confirm the existence of a critical mass ($M_{cr} \\sim 10 M_{\\oplus}$) beyond which migration halts in nearly laminar disks.  For $\\alpha \\ga 10^{-3}$, density feedback effects are washed out and Type~I migration persists.  The critical masses are in good agreement with the analytic model of Rafikov (2002).  In addition, for $\\alpha \\la 10^{-4}$ steep density gradients produce a vortex instability, resulting in a small time-varying eccentricity in the planet's orbit and a slight outward migration.  Migration in nearly laminar disks may be sufficiently slow to reconcile the timescales of migration theory with those of giant planet formation in the core accretion model. ",
        "introduction": "\\label{sec:intro}  The standard theory of Type~I (low planet mass) migration presents a challenge for understanding planet formation. According to the core accretion model, the growth time from planetesimals to gas giant planets is dominated by a phase that occurs when a newly formed solid core with mass $\\sim 10 M_{\\oplus}$ accretes gas (Pollack et al 1996, Hubickyj et al 2007).  The duration of this slow phase, $\\sim 10^6 y$, is determined by a thermal bottleneck that prevents the gaseous envelope from contracting, until it achieves sufficient mass. On the other hand, the standard theory of Type~I migration (Tanaka et al 2002) predicts that planets in the slow growth phase would migrate into the disk center in $\\sim10^5 y$ for the minimum solar mass nebula.  These migration timescales have been confirmed by multidimensional hydrodynamical simulations (Bate et al 2003, D'Angelo \\& Lubow 2008).  Recently, Ida \\& Lin (2008) and Schlaufman, Lin, \\& Ida (2008) have studied the effect of the ice line on the disk surface density profile and consequently on the Type I migration. They obtained much better agreement with the observed extrasolar planet mass-semimajor axis distribution if the Type I migration is reduced by an order-of-magnitude from the linear theory values.  The shortness of the standard migration timescale has motivated investigations of possible effects to slow or even reverse migration. These include magnetic fields (Terquem 2003), magneto-rotational instability (MRI) turbulent fluctuations (Nelson \\& Papaloizou 2004), and density traps (Menou \\& Goodman 2004).  Recently, protoplanet migration in the non-isothermal disks has been investigated (Paardekooper \\& Mellema 2006; Baruteau \\& Masset 2008; Paardekooper \\& Papaloizou 2008; Kley \\& Crida 2008). These simulations show indications of slowing migration due to coorbital torques for certain ranges of gas diffusivity and turbulent viscosity.  In this Letter we discuss another mechanism that naturally occurs when the disk turbulent viscosity is sufficiently small.  The Tanaka et al (2002) Type~I migration rates were derived under the assumption that the disk density distribution is unaffected by the presence of the planet.  Numerical simulations commonly adopt turbulent viscosity parameter values $\\alpha \\ga 10^{-3}$. Such values are suggested by considering the observationally inferred disk masses and accretion rates for T Tauri stars (e.g., Hartmann et al 1998). With such $\\alpha$ values, turbulent diffusion suppresses disk disturbances for planets of mass less than $0.1 M_J$.  The numerical simulations then satisfy the Tanaka et al (2002) assumptions and yield migration rates that are in close agreement.  The various models of planet formation by core accretion typically involve lower $\\alpha$ values, $\\alpha < 10^{-3}$ (see Cuzzi \\& Weidenschilling 2006), which we refer to as nearly laminar values.  To form planetesimals via gravitational instability from small dust particles (Safronov 1969, Goldreich \\& Ward 1973) requires the dust layer to be very thin $\\sim 10^{-4} H$, suggesting $\\alpha \\ll 10^{-4}$. For the solids to be dynamically decoupled from the gas requires a dust layer disk of thickness $\\la 10^{-2} H$, again suggesting nearly laminar conditions. Cuzzi \\& Weidenschilling (2006) estimate that $\\alpha \\la 2 \\times 10^{-4}$ in order that meter size solids avoid destructive effects of collisions due to turbulent motions.  In the planet formation regions of disks, considerations of the MRI (Balbus \\& Hawley 1991) suggest that the disk maybe unstable only in surface layers, due to the low levels of ionization below these layers (Gammie 1996).  A major uncertainty is the abundance of small grains that can suppress the instability.  The disk may be nearly laminar for the purposes of planet formation.  However, surface layer turbulent fluctuations may propagate disturbances to the disk midplane. They may provide some effective turbulence in that region as well (Fleming \\& Stone 2003; Turner \\& Sano 2008).   In nearly laminar disks, density waves launched by a planet at various Lindblad resonances can redistribute disk mass as they damp. The disk turbulent viscosity needs to be sufficiently small for the density perturbation to not diffuse away.  The redistributed gas slightly enhances (reduces) the disk density interior (exterior) to the orbit of an inwardly migrating planet.  Some analytic studies have shown that such density feedback effects could slow and even halt the migration (Hourigan \\& Ward 1984; Ward \\& Hourigan 1989; Ward 1997; Rafikov 2002).  The critical planet mass at which the feedback becomes important depends on the efficiency of the damping of waves excited by the planet.  Wave damping may be due to the shock dissipation and will generally occur nonlocally, at some distance from the radius where the wave is launched.  In this {\\em Letter}, we explore the consequences of nearly laminar disks on planet migration by means of nonlinear 2D numerical simulations. ",
        "conclusions": "\\label{sec:diss}   The local wave damping model of Ward \\& Hourigan (1989) suggests critical planet masses of $M_{\\rm cr} \\sim \\Sigma r_p^2 h^3$, which evaluates to $0.006 f M_{\\oplus}$ and $0.02 f M_{\\oplus}$ for $h = 0.035$ and $0.5$, respectively.  These values differ from Table 1 by a factor of more than 100. However the scaling of the critical mass with disk thickness is close to $h^3$ as given by this theory. The scaling of $M_{\\rm cr}$ with surface density in Table 1 is, however, weaker than that suggested by this theory.  In another local wave damping model, Ward (1997) suggests that $M_{\\rm cr} \\sim 0.2 \\Sigma r_p^2 h$, which evaluates to $1.1 f M_{\\oplus}$ and $1.6 f M_{\\oplus}$ for $h = 0.035$ and $0.05$, respectively. Although these values are numerically closer to the simulation values in Table 1, the predicted linear scaling in both $f$ and $h$ does not agree with the trends in the Table 1. The analytic model of Rafikov (2002) includes the effects of nonlocal damping by means of shocks.  In that case, the critical masses are given by \\begin{equation} M_{cr} = \\frac{2 c_s^3}{3 \\Omega G}~ {\\rm min}\\left[5.2 Q^{-5/7}, \\, 3.8 ( Q/h )^{-5/13} \\right], \\end{equation} where $Q = \\Omega c_s/(\\pi G \\Sigma)$. All the simulated cases correspond to the strong feedback branch of $M_{cr}$ that is given by the second argument of the $min$ function. Values for these critical masses are also given in Table 1.  We see that the agreement between the simulation and theory is quite good.  The critical masses are in the range of the core masses during the slow phase of gas accretion in the core accretion model of planet formation. This result suggests that planet migration might not be a limiting factor in planet formation. The phase of run-away mass accretion follows the slow evolution phase.  Previous studies suggest that run-away mass accretion to $1 M_J$ in a disk with $\\alpha \\simeq 0.004$ occurs in about $10^{5} y$ (e.g., D'Angelo \\& Lubow 2008). During this phase, we speculate that more modest levels of turbulent viscosity $ 10^{-4} < \\alpha < 10^{-3}$ may provide sufficient accretion to form a $1 M_J$ planet within $\\sim 10^{6} y$, while remaining in this nearly laminar disk regime of slow Type I migration.  Although this picture is suggestive of a resolution of the migration problem, the longer term evolution of nearly laminar disk-planet systems requires further exploration.  The slowly migrating planet will continue to create a deeper gap over time. The steepening density gradients should lead to the vortex instability (Koller et al.  2003; Li et al. 2005).  The consequences of the vortex instability should be explored."
    },
    "0812/0812.2060_arXiv.txt": {
        "abstract": "{The recent experimental evaluation of the $^{18}$F($\\alpha$,$p$)$^{21}$Ne reaction rate, when considering its associated uncertainties, presented significant differences compared to the theoretical Hauser-Feshbach rate.  This was most apparent at the low temperatures relevant for He-shell burning in asymptotic giant branch (AGB) stars. Investigations into the effect on AGB nucleosynthesis revealed that the upper limit resulted in an enhanced production of $^{19}$F and $^{21}$Ne in carbon-rich AGB models, but the recommended and lower limits presented no differences from using the theoretical rate. This was the case for models spanning a range in metallicity from solar to [Fe/H] $\\sim -2.3$. The results of this study are relevant for observations of F and C-enriched AGB stars in the Galaxy, and to the Ne composition of mainstream silicon carbide grains, that supposedly formed in the outflows of cool, carbon-rich giant stars. We discuss the mechanism that produces the extra F and summarize our main findings.}  \\FullConference{10th Symposium on Nuclei in the Cosmos\\\\ July 27 - August 1 2008\\\\ Mackinac Island, Michigan, USA}  \\begin{document}  \\bibliographystyle{unsrt} ",
        "introduction": "Until 2006 the only rate for the $^{18}$F($\\alpha$,p)$^{21}$Ne reaction was the theoretical estimate available in the Brussels nuclear reaction-rate library \\cite{aikawa05}.  The first experiment aimed at determining the $^{18}$F($\\alpha$,p)$^{21}$Ne rate over a large range of stellar temperatures was carried out by Lee \\cite{LeeThesis}. This experimental evaluation, when considering its associated uncertainties, presented significant differences compared to the theoretical rate, especially at the low temperatures relevant for He-shell burning in asymptotic giant branch (AGB) stars ($T \\approx 0.3$~GK). We investigate the effect of such differences on the nucleosynthesis occurring in AGB models of various initial mass and composition.  The theoretical estimate of the $^{18}$F($\\alpha$,p)$^{21}$Ne rate was not present in previous studies e.g., \\cite{karakas06a}, although we had included the species \\iso{18}F because of its important role in the reaction chain \\iso{14}N($\\alpha,\\gamma$)\\iso{18}F($\\beta^+\\nu$)\\iso{18}O, leading to the production of \\iso{18}O in the He shell. It was found that the inclusion of the $^{18}$F($\\alpha$,p) reaction resulted in an increase in the production of the stable \\iso{19}F; see \\S\\ref{results} for more details on the production mechanism. This is of interest because AGB models do not synthesize enough $^{19}$F to match the [F/O] abundances observed in AGB stars \\cite{jorissen92}. Also, the cosmic origin of fluorine is still uncertain, with massive stars \\cite{woosley95} playing a significant role in producing fluorine alongside AGB stars. However, AGB stars and their progeny (e.g., post-AGB stars, planetary nebulae) are still the only confirmed site of fluorine production thus far \\cite{federman05, werner05}. Observations of an enhanced F abundance ([F/Fe] = 2.90) in a carbon-enhanced metal-poor halo star \\cite{schuler07} is further motivation to understand the details of F production in AGB stars \\cite{lugaro08}.  In this proceedings, we summarize the results of calculations published in Karakas et al. \\cite{karakas08a}. ",
        "conclusions": "\\begin{figure} \\begin{center} \\includegraphics[width=0.7\\textwidth]{f1.eps} \\caption{Comparison of fluorine abundances observed by \\cite{jorissen92} and model predictions for the 3$M_{\\odot}$, $Z=0.008$ model. The predictions are normalized such that the initial $^{19}$F abundance corresponds to the average F abundance observed in K and M stars. Each symbol on the prediction lines represents a TDU episode. Solid lines represent calculations performed using no $^{18}$F($\\alpha,p$)$^{21}$Ne reaction, which are equivalent to using the current lower limit, recommended value and Brussels library rate. Dotted lines are calculations performed using the current upper limit of the rate.} \\end{center} \\label{fig1} \\end{figure}  From Fig.~\\ref{fig1}, we see that the surface [$^{19}$F/$^{16}$O] ratios from the 3$M_{\\odot}$, $Z=0.008$ model are a factor of $\\sim 2.2$ times higher when employing the new upper limit of the $^{18}$F($\\alpha,p$)$^{21}$Ne reaction. We chose to show this model because it produces the largest F abundances, although the metallicity of this model is probably at the lower end of the distribution of Galactic carbon stars. Table~\\ref{tab:yields} and Fig.~\\ref{fig1} shows that a match between the stellar models and the stars with the highest observed $^{19}$F abundances is possible, but only for very high C/O ratios of $\\sim 4-5$ found in the 3$M_{\\odot}$, $Z = 0.008$ model.   These high C/O ratios are likely not realistic, and the inclusion of carbon-rich, low-temperature opacities into the stellar models would cause the TP-AGB evolution to end before the model star reached such C/O ratios. A solution to the mystery of the high F abundances at modest C/O ratios is still missing, but a re-evaluation of the F and C abundances in the sample of AGB stars considered by \\cite{jorissen92} may help, along with a detailed examination of model uncertainties.  The $^{18}$F($\\alpha$,p)$^{21}$Ne reaction also affects the abundance of $^{21}$Ne in the He-shell of AGB stars. There is a long-standing puzzle concerning the isotopic composition of Ne measured in stellar silicon carbide (SiC) grains extracted from meteorites, which formed in the envelopes of carbon-rich AGB stars e.g., \\cite{lewis90,zinner06}. Models computed with the upper limit of the $^{18}$F($\\alpha, p$)$^{21}$Ne reaction rate show an increase in the $^{21}$Ne abundance, and hence in the $^{21}$Ne/$^{22}$Ne ratio in the intershell of up to a factor of 6; see \\cite{karakas08a} for further details. We conclude that the \\iso{18}F($\\alpha,p$)\\iso{21}Ne reaction rate being closer to its upper limit may be a promising explanation for the $^{21}$Ne/$^{22}$Ne ratios in SiC grains. The modeling uncertainties related to convection and mass loss do not affect the intershell compositions and thus do not apply to the discussion of the Ne composition of stellar SiC grains. For this reason, the results for Ne are also a more reliable hint that the $^{18}$F($\\alpha$,p)$^{21}$N reaction is indeed closer to its upper limit than the comparison of F in AGB stars.  However, more experimental data for this reaction at temperatures below 0.4~GK are required to help verify this result."
    },
    "0812/0812.0854_arXiv.txt": {
        "abstract": "{Open clusters offer a unique possibility to study the time evolution of the radial metallicity  gradients of several elements in our Galaxy, because they span large intervals in age and Galactocentric distance, and both quantities can be more accurately derived than for field stars.} {We re-address the issue of the Galactic metallicity gradient and its time evolution by comparing the empirical gradients traced by a sample of 45 open clusters with a chemical evolution model of the Galaxy.} {At variance with previous similar studies, we have collected from the literature only abundances derived from high--resolution spectra. The clusters have Galactocentric distances $7 \\lesssim R_{\\rm GC} \\lesssim 22$~kpc and ages from $\\sim 30$~Myr to 11~Gyr. We also consider the $\\alpha$-elements Si, Ca, Ti, and the iron-peak elements Cr and Ni. Cepheids trace instead the present-day Fe gradient in the inner parts of the disk.} {The data for iron-peak and $\\alpha$-elements indicate a steep metallicity gradient for $R_{\\rm GC}\\lesssim 12$~kpc and a plateau at larger radii. The time evolution of the metallicity distribution is characterized by a uniform increase of the metallicity at all radii, preserving the shape of the gradient, with  marginal evidence for a flattening of the gradient with time in the radial range 7-12~kpc. Our model is able to reproduce the main features of the metallicity  gradient and its evolution with an infall law exponentially decreasing with radius and with a collapse time scale of the order of 8~Gyr at the solar radius. This results in a rapid collapse  in the inner regions, i.e. $R_{\\rm GC}\\lesssim 12$~kpc (that we associate with an early phase of disk formation from the collapse of the halo) and in a slow inflow of material per unit area in the outer regions at a constant rate with time (that we associate with accretion from the intergalactic medium). An additional uniform inflow per unit disk area would help to better reproduce the metallicity plateau at large Galactocentric radii, but it is difficult to reconcile with the present-day radial behaviour of the star formation rate.} {Our results favour a scenario where the Galactic disk is formed inside-out by the rapid collapse of the halo and by a subsequent continuous accretion of intergalactic gas} ",
        "introduction": "\\label{intro}  The radial metallicity gradient and its time evolution provide strong constraints on our understanding of the formation and the evolution of galaxies.  Detailed models of Galactic chemical evolution (GCE) have been continuously improved over the last decades (e.g., Lacey \\& Fall \\cite{lacey}; Tosi \\cite{tosi88}; Ferrini et al.~\\cite{ferrini92}, \\cite{ferrini94}; Giovagnoli \\& Tosi \\cite{giova}; Chiappini, Matteucci \\& Gratton \\cite{chiappini97}; Boissier \\& Prantzos \\cite{boissier99}; Hou et al.~\\cite{hou00}; Chiappini et al.~\\cite{chiappini01}). However, the issue of the time evolution of the radial abundance gradients is far from being settled unequivocally, both from a theoretical (e.g., Goetz \\& Koeppen~\\cite{goetz92}; Koeppen~\\cite{koeppen94}; Moll{\\`a} et al.~\\cite{molla97}; Henry \\& Worthey~\\cite{henry99}; Chiappini et al.~\\cite{chiappini01}) and an observational (e.g., Friel \\cite{friel95}; Friel et al.~\\cite{friel02}; Maciel~\\cite{maciel00}; Maciel et al.~\\cite{maciel03}; Stanghellini et al.~\\cite{stanghellini06}) point of view.  While most GCE models are able to reproduce the present-day radial distribution of several chemical elements derived from sample objects representative of the present-day composition of the interstellar medium (ISM), such as, e.g., \\hii\\, regions, Cepheids, super-giant stars, they generally disagree on the predicted behaviour of its time evolution.  In this respect, GCE models can be divided in two groups: models which predict a {\\em steepening} of the metallicity gradient with time (e.g., Tosi \\cite{tosi88}; Chiappini et al.~\\cite{chiappini97}; Chiappini et al.~\\cite{chiappini01}) and those which predict a {\\em flattening} with time (e.g., Moll\\`a et al.~\\cite{molla97}; Portinari \\& Chiosi~\\cite{portinari99}; Hou et al.~\\cite{hou00}).  The main differences between the two groups of models are the efficiency of the enrichment processes in the inner andf outer regions of the Galactic disk (the star formation rate, SFR), and the nature of the material (primordial or pre-enriched) falling from the halo onto the disk (the infall).  In models of the former group the outer disk is pre-enriched by the previous evolution of the halo, and during the first Gyr its metallicity is affected very little by the relatively lower SFR in the disk. Since the infall decreases rapidly with Galactocentric radius, reflecting the strongly concentrated mass distribution in the halo, a larger amount of metal poor material falls in the central regions rather than in the outer Galactic disk. Thus, a positive metallicity gradient is initially established. When the halo collapse phase is over, the metallicity increases more rapidly in the inner disk than in the outer Galaxy, because of a combination of higher SFR at small radii and continuous dilution by infall at large radii, reversing the sign of the gradient.  In contrast, in models of the latter group, the initially vigorous star formation activity in the inner disk and the metal-poor composition of the infalling gas produce a rapid increase of the metallicity near the Galactic center. In the inner regions the metallicity increases rapidly, reaching its final value in the first 2--3 Gyr of disk evolution, whereas the enrichment of the outer disk is slower, resulting in a progressive flattening of the radial gradient.  In this group of models, various infall rates have been assumed: a very rapid one simulating the formation of the disk from a monolithic collapse of the halo, or one much diluted with time, simulating a hierarchical formation of the disk by continuous infall of gas from the intergalactic medium (IGM). These rates produce different time evolution of the metallicity gradient: a rapid flattening or a uniform increase of the metallicity with essentially no change of slope, respectively (e.g., Moll\\`a et al.\\cite{molla96}).  In principle, it should be possible to discriminate between the different evolutionary scenario, by comparing the metallicity gradient of the ``old'' stellar population  e.g., red giant branch (RGB) stars, planetary nebulae (PNe), and old open clusters (OCs), with the present-day gradient outlined, e.g., by \\hii\\,  regions, young stars (including Cepheids and giant stars) and young OCs.  In external galaxies, the comparison of PN and \\hii\\, region abundances, investigated with similar observational and analysis techniques, has given encouraging results (e.g., Magrini et al.~\\cite{magrini07} for the spiral galaxy M33), since PNe can be assumed at the same distance as the host galaxy.  In our Galaxy, however, the huge uncertainty on the distances to Galactic PNe restrict their use as reliable tracers of the past evolution of metallicity gradients.  On the contrary, the determination of metallicities of OCs represents a reliable approach to the study of the time evolution of the metallicity gradient. OCs represent, in fact, a unique opportunity to study abundances of a large number of elements in stars of known ages and located at a wide range of well-determined Galactocentric distances. Thus, they can fully address the issue of the time evolution of the metallicity gradient, provided that a homogeneous analysis is performed.  Several spectroscopic studies of chemical abundances in OCs have been carried out during the last years, both at low- (e.g., Friel \\cite{friel95}; Friel et al.~\\cite{friel02}) and high-resolution (e.g., Carretta et al.~\\cite{carretta04}; Carraro et al.~\\cite{carraro04}, \\cite{carraro06}, \\cite{carraro07}; Yong, Carney \\& de Almeida \\cite{yong05}; Sestito et al~\\cite{sestito06}, \\cite{sestito08}; Bragaglia et al.~\\cite{bragaglia07}). Different authors find different results, depending on the quality of the data, and in particular on the spectral resolution, and on the method of analysis adopted. These discrepancies especially affect metallicity gradients. Also, abundances for the farthest clusters, which are critical to determine the gradient up to very large Galactocentric distances, were not available until very recently. For example, Friel (\\cite{friel95}; \\cite{friel06}), Carraro, Ng \\& Portinari (\\cite{carraro98}) and Friel et al.~(\\cite{friel02}) suggest the presence of a negative [Fe/H] gradient, while Twarog, Ashman \\& Anthony-Twarog (\\cite{twarog}) and Corder \\& Twarog (\\cite{corder}) favor a step-like distribution of the Fe content with Galactocentric distance. Furthermore, recent results (e.g., Yong et al. \\cite{yong05}, Carraro et al.~\\cite{carraro07}; Sestito et al.~\\cite{sestito07}) suggest that the gradient becomes flat for Galactocentric radii $R_{\\rm GC}$ larger than $\\sim 11-12$~kpc.  For the time evolution, evidence for variation so far  has been not conclusive.  The majority of studies seem to favor a slight evolution from steeper gradients in the past to shallower one at the present time. For example, Friel et al.~(\\cite{friel02}), based on metallicities derived from low resolution spectra, find a flattening of the {\\it overall} gradient from $-0.096$~dex\\,kpc$^{-1}$ for ages greater than 6~Gyr, to $-0.023$~dex\\,kpc$^{-1}$ for clusters younger than 2~Gyr; similarly, Chen et al.~(\\cite{chen03}) derive a shallower slope for clusters younger than 0.8~Gyr than for older ones. On the contrary Salaris et al. (\\cite{salaris04}) find a steeper slope for younger clusters. These discrepant results are likely due to: {\\em i)} use of different and/or inhomogeneous abundance datasets; {\\em ii)} the different radial regions of our Galaxy which are analyzed.    In this paper we re-address the issue of the radial metallicity gradient in our Galaxy observationally and theoretically, exploiting the availability of abundances based on high resolution spectra for several open clusters. Specifically, we collected and analyzed the most recent metallicity determinations of a large sample of OCs, considering {\\em only measurements based on high resolution spectroscopy}, which by far are more reliable than those based on photometry or low resolution spectroscopy. These observations, described in Sect.~\\ref{sect_data}, are compared with the results of our chemical evolution model, able to reproduce the main observational features of our Galaxy.  In addition to Fe, the most recent studies (e.g., Friel et al.~\\cite{friel03}; Carretta et al.~\\cite{carretta04}; Friel, Jacobson, \\& Pilachowski \\cite{friel05}; Yong et al. \\cite{yong05}; Sestito et al.~\\cite{sestito06}; Randich et al.~\\cite{randich06}; Bragaglia et al.~\\cite{bragaglia07}; Sestito et al.~\\cite{sestito07}, \\cite{sestito08}) also include the analysis of other elements. These studies allow us to investigate for the first time the radial gradient of Fe-peak and $\\alpha$-elements based on the comparison of OC abundances, rather than on field stars, and model predictions. Finally, we consider the empirical gradient traced by Cepheids for the inner parts of the disk (R$_{\\rm GC} \\leq 7$~kpc).  The rest of the paper is organized as follows: in Sect.~\\ref{sect_model} we describe the model; in Sect.~\\ref{sect_confronto} we analyze the time evolution of the radial gradient of several chemical elements, including Fe and several $\\alpha$ and iron-peak elements; the abundance ratio radial behaviours are discussed in Sect.~\\ref{sect_other}, while the results of this study are discussed in Sect.~\\ref{sect_discussion}; in Sect.~\\ref{sect_summary} we summarize our conclusions. In the Appendix we describe the details of the Galactic chemical evolution model. ",
        "conclusions": "\\setcounter{table}{1} \\begin{table} \\caption{Average [Fe/H] from OC data of  different epochs and radial regions (see the text for the definition of radial regions).} \\label{tab_average} \\centering \\begin{tabular}{rrr} age (Gyr)  & Region II            & Region III \\\\ \\hline $<0.8$       & $0.016 \\pm 0.14$         & $-0.17 \\pm 0.12$   \\\\ 0.8--4       & $0.05 \\pm 0.11$          & $-0.29 \\pm 0.16$   \\\\ 4--10         & $-0.04 \\pm 0.29$         & $-0.33 \\pm 0.13$   \\\\ \\hline \\end{tabular} \\end{table}   In their paper, Maciel et al. (2006) showed the results on the time evolution of the [Fe/H] gradient obtained by several authors, both from models (Hou et al.~\\cite{hou00}, Chiappini et al.~\\cite{chiappini01}) and from observations (PNe by Maciel et al.~2006, OCs by Friel et al.~2002, Chen et al.~2003, and \\hii\\ regions, Cepheids, OB stars from several authors, see references therein).  The metallicity gradient was computed over the whole disk. However, as already remarked, the gradient does not have a constant slope over the whole disk, and, thus, a direct comparison with our data is difficult to obtain.  Therefore we compared our results with the samples of OCs by Friel et al.~(2002) and Chen et al.~(2003) in the same radial regions.  Their samples were discussed in the previous sections. In summary, they differ from our sample in data quality (high-resolution spectroscopy vs. low-resolution spectroscopy and/or photometry) and radial range (they extend up to $\\sim$16~kpc, while our sample extends up to 22~kpc).  The OCs located at high Galactocentric radius are fundamental in determining the slope in the older and intermediate age bin.  In order to compare with the results of Friel et al.~(2002) and Chen et al.~(2003) we have re-computed the slope of the older and intermediate age bins for $R_{GC}\\leq16$, obtaining a slope for the cumulative sample of -0.067~dex kpc$^{-1}$, and for  only the old bin of -0.079~dex kpc$^{-1}$, in closer agreement with their results: -0.075~dex kpc$^{-1}$ for the sample of Chen et al.~(2003) with ages $>$0.8~Gyr and -0.075~dex kpc$^{-1}$ for the sample of Friel et al.~(2002) with ages $>$4~Gyr.  This points out the importance of OCs at larger Galactocentric distances for a correct determination of the outer gradient.  The comparison with PNe is more tricky.  The accuracy of the Galactic metallicity gradient derived for PNe has  been debated for a long time because of the large uncertainties on PN distances.  Recent results (Stanghellini et al.~\\cite{stanghellini06}, Perinotto \\& Morbidelli~\\cite{perinotto06}) have shown very different gradients for the Galactic PN population than previous results.  At variance with Maciel et al.'s results, Stanghellini et al.~(\\cite{stanghellini06}) and  Perinotto \\& Morbidelli~(\\cite{perinotto06}) found an almost flat gradient at all epochs. On the other hand, some other recent works find steeper gradients (e.g., Pottasch \\& Bernard-Salas~\\cite{pottasch06}; Maciel \\& Costa~\\cite{maciel08}) pointing to the need for an accurate determination of the age in order to measure gradients. The main differences among these works are the adopted distance scales, demonstrating how this choice can affect the estimate of the metallicity gradient,  in conjunction with different age estimates for the PN progenitors.  Thus a faithful determination of the Galactic metallicity gradient from PNe is far from  being obtained.   A more appropriate way to compare the chemical evolution model with observations is to divide the Galaxy in several radial regions.  In Tables~\\ref{tab_average} and \\ref{table:grad} we list the the average abundances of iron in each age bin and radial region, and the best-fit slopes in the three radial regions. The latter were computed through a weighted fit procedure for the OCs and taken from Andrievsky et al. (\\cite{andri04}) and Lemasle et al. (\\cite{lemasle07}) for the Cepheids.  The slopes of the models are also shown.  As anticipated above, three age bins were considered: age less than or equal to 0.8~Gyr, in the first panel; $0.8<\\mbox{age}\\leq 4$~Gyr, in the second panel; and $4<\\mbox{age}\\leq 11$~Gyr in the third panel.  For OCs with ages between 4~Gyr and 11~Gyr, we excluded from the calculation of gradient the metal-rich cluster NGC~6791 (marked with an empty circle in Figs.~\\ref{iron1}, \\ref{fepeak}, \\ref{alpha1},\\ref{alpha2}), because it was probably born in a different place in the disk or could even have been captured by our Galaxy during a merger event (Carraro et al.~\\cite{carraro06}, but see also Bedin et al.~\\cite{bedin06} for a contrasting view). Thus the age range for our further analysis is between 4~Gyr and 10~Gyr.  First, we note that, although it is well known that a strict age-metallicity relationship does not hold in the Galactic disk, a tendency is present for older OC to have on average lower metallicities. Second, each radial region shows a different behaviour due to the distinct ratios of SFR and infall. To better analyze the time evolution of the metallicity gradients, we plot in Fig.~\\ref{slopes} the slopes of the gradient for the specific case of iron, the best studied element, in the three radial regions defined above. Similar behaviours are observed also for the other elements analyzed in this study.  As there are no observations of OCs in region~I, it is not possible to obtain the time evolution of the gradient in this region.  In region~II, our sample of OCs does not show any conclusive evidence for flattening with time of the slope of the gradient, although the present time gradient seems slightly flatter then the gradient in the two older age bins.  Finally, in region~III, the OC metallicities produce a gradient consistent with a zero slope at all epochs.  Our data have allowed us to consider the time evolution in different regions of the disk, showing that in the $\\sim 7$--22~kpc interval no marked evolution is present.    The observations of Cepheids by Andrievsky et al.~(\\cite{andri04}) (see Table~\\ref{tab_grad_iron} and filled squares in Fig.~\\ref{slopes}) show a good agreement with the absolute [Fe/H] predicted by the model in regions~I, II, and III at the present time (see Table~\\ref{tab_grad_iron}).  In region~I, the model predicts a flattening of the metallicity gradient from the past epochs to the present time that cannot be tested because only Cepheid data are available in this region; in regions~II and III it predicts an overall increase of metallicity at all radii but no significant change in the slope of the gradient (see Fig.~\\ref{slopes}).   This can be understood as follows. In our model, the inner part of the disk is built by the rapid collapse of the halo, whereas the outer parts by the continuous slow accretion of intergalactic gas.  In the inner disk, a steep chemical gradient is initially established by the strong concentration of infalling gas (see eq.~\\ref{eq_infall}), on a timescale probably of the order of a few Gyr or less (OCs are not able to constrain the chemical evolution of the gradient for the first $\\sim 5$~Gyr of evolution of the Galaxy). The  model predictions indicate that this initial gradient is preserved in the subsequent evolution of the disk.  Turbulent mixing of the interstellar gas seems to be inefficient in smoothing the initial metallicity gradient over the lifetime of the Galaxy, in agreement with theoretical estimates; for example, according to Talbot \\& Arnett~(\\cite{talbot73}), the timescale for mixing a region of a radial extent of 5 kpc like Region~II is about $75$~Gyr. Also the role of large-scale radial flows, often invoked in the chemical evolution of disk galaxies, seems quite marginal, since their effect is to steepen the metallicity gradient (Lacey \\& Fall~\\cite{lacey}; Tosi~\\cite{tosi88}; Portinari \\& Chiosi~\\cite{portinari00}), contrary to observational evidence. However Friedli (\\cite{friedli98}) found that flows induced by bars can flatten the gradient.  In the outer disk, on the other hand, accretion of intergalactic gas and disk evolution occur almost independently of radius, and the metallicity gradient remains flat.   We note that the slope of older OCs is significantly affected by the presence of a number of old clusters located at low Galactocentric distances with solar or super-solar metallicities. Old clusters have had enough time to move along their orbits, and the possibility exists that they were actually born in inner regions far from their present location. Although it seems unlikely that all old clusters with (over)-solar metallicities have different birthplaces, we stress that knowledge of the orbits of these older OCs, and thus their previous location, would help in definitively settling the issue. Also, a few clusters located at distances of 11-12 kpc are present with [Fe/H] somewhat below the plateau value observed at larger Galactocentric distances and characterized by large error bars. New observations for these clusters should be obtained in order to better constrain the slope of the gradient.  Finally, metallicity determinations should be performed for both young and old clusters located at Galactocentric distances lower than $\\sim 7$~kpc.  \\begin{figure} \\psfig{figure=0634fig13.eps,width=11cm, angle=-90} \\caption{Time evolution of the slope of the iron gradient in the three radial region analyzed: $4$~kpc $\\lesssim R_{\\rm{GC}}\\lesssim7$~kpc (region~I), $7$~kpc $\\lesssim R_{\\rm{GC}}\\lesssim12$~kpc (region~II), $12$~kpc $\\lesssim R_{\\rm{GC}}\\lesssim22$~kpc (region~III).  Continuous lines are the slopes of the model gradients approximated with a linear fit.  Filled squares are the slopes of the gradients of  Cepheids. Filled circles represent the slopes of OC gradients with their uncertainties in the three age bins, as in Table~\\ref{tab_grad_iron}.} \\label{slopes} \\end{figure}      s We have assembled a homogeneous sample of abundances in OCs from high-resolution spectroscopic data and compared it  to the results of a GCE model.  We examined in particular the radial gradients of Fe/H, Si/H, Ca/H, Ti/H, Ni/H, and Cr/H, and also their [X/Fe] gradients at different epochs during the lifetime of the Galaxy. We have divided our data into three age bins and three radial regions. Our main results are the following:   \\begin{itemize}  \\item[({\\em i}\\/)] From the observations of OCs and Cepheids (see Table~\\ref{tab_grad_iron}), we find that the metallicity gradients of all the elements analyzed are very steep at low Galactocentric radii, in particular for $R_{\\rm GC}\\lesssim 7$~kpc.  The gradients show a change of slope around 7~kpc, and further flattening at about 11--12~kpc, resulting in a metallicity plateau in the outer regions.  \\item[({\\em ii}\\/)] The combination of the observations of Cepheids and OCs allow us to follow the time evolution of the metallicity gradient and to compare it with the results of a GCE model. The time evolution of the gradient can be followed where observations of OCs with different ages are available for $R_{\\rm GC}> 7$~kpc.  At Galactocentric radii $7$~kpc $\\lesssim R_{\\rm GC}\\lesssim 12$~kpc the combined observations of Cepheids and OCs suggest that the slope of the metallicity gradient has not changed significantly over the past $\\sim 10$~Gyr; if anything, the gradient has flattened with time, as also suggested by the model for $R_{\\rm GC} < 7$~kpc. For $12~\\mbox{kpc} \\lesssim R_{\\rm GC} \\lesssim 22$~kpc, OCs belonging to different epochs and Cepheids all suggest a flat metallicity distribution.  \\item[({\\em iii}\\/)] The GCE model that better reproduces the data is characterized by an inside-out formation of the disk. The infall of gas is represented by an exponential law which, combined with the distribution of gas in the halo, produces a rapid collapse in the regions with R$\\lesssim$12~kpc and a uniform accretion in the outer regions.  The inner regions are rapidly evolving due to the higher infall and SFR, while the outer parts evolve more slowly.  The metallicity in the outer disk increases uniformly at all radii without a marked change of slope of the gradient.  Peculiar episodes of mergers in the past history of the Galaxy would explain the outer metallicity plateau indicated by the older clusters. \\end{itemize}  \\begin{landscape} \\setcounter{table}{0} \\begin{table*} \\scriptsize \\caption{Open clusters included in our sample. We list in Col.~1 the cluster name, in Col.~2 the cluster age, in Cols.~3 and 4 the Galactocentric radius and its reference (an asterisk indicate that we computed R$_{\\rm GC}$  assuming the cluster distance given in the quoted reference). In Cols.~5, 6, 7, 8, 9, and 10 we list the main properties of the observations and of the analysis: the signal to noise ratio, the resolving power, the number and type of stars analyzed, the adopted method (Equivalent widths, EW, or spectral synthesis), the iron abundance with its error, and the reference.} \\begin{tabular}{llllllllll} \\hline \\hline Cluster &Age            & $R_{\\rm{GC}}$ &  Ref.              & S/N      & R     &N.Stars&Method   &[Fe/H]  & Ref.\\\\ & (Gyr)         &(kpc)         & ($R_{\\rm{GC}}$)&          &       &      &         &        & ([Fe/H])     \\\\ \\hline IC4665   & 0.025        &8.18         & Manzi et al.~(\\cite{manzi08})$^*$        &   30-150    & 60,000       &18,dwa     &Synt       & $-0.03\\pm0.04$ & Shen et al..~(\\cite{shen})\\\\ IC2602   & 0.035        &8.49         & WEBDA$^*$\\footnote{http://www.univie.ac.at/webda/} & 30-100 & 18,000-44,000&9,dwa     &EW       & $-0.05\\pm0.05$ &Randich et al. (2001)              \\\\ IC2391   & 0.053        &8.50         & WEBDA$^*$                                &   30-100 & 18,000-44,000&4,dwarfs  &EW       & $-0.03\\pm0.07$ &Randich et al. (2001)                \\\\ Blanco 1 & 0.1-0.15     &8.50         &   Friel~(\\cite{friel06})                 &   70-100 & 50,000       &8,dwa     &EW       & $+0.04\\pm0.04$ & Ford et al.~(\\cite{ford05})   \\\\ Pleiades & 0.125        &8.60         &   Friel~(\\cite{friel06})                 &   70     & 45,000       &15,dwa    &EW       & $-0.03\\pm0.02$ & Boesgaard \\& Friel~(\\cite{boesgaard90})       \\\\ NGC6475  & 0.220        &8.22         & Robichon et al.~(\\cite{robichon99})$^*$  &   50-150 & 29,000       &13,dwa    &EW       & $+0.14\\pm0.06$ &Sestito et al. (\\cite{sestito03})                 \\\\ M11      &  0.250       &6.86         & Gonzalez \\& Wallerstein~(\\cite{gonzalez00})$^*$&  85,130 & 38,000 &10,gia     &EW/synthesis& $+0.17\\pm0.13$ &Gonzalez \\& Wallerstein (2000)\\\\ M34      & 0.250        &8.90         &   Friel~(\\cite{friel06})                 &   70     & 45,000       &4,dwa     &EW       & $+0.07\\pm0.03$ & Schuler et al.~(\\cite{schuler03})       \\\\ NGC3960  & 0.7          &7.96         & Friel et al.~(\\cite{friel02})            &  100-150 &45,000        &6,giants  &EW       &$+0.02\\pm0.04$ & Sestito et al.~(\\cite{sestito06})\\\\ Hyades   & 0.7          &8.50         &   Friel~(\\cite{friel06})                 &   100-200& 60,000       &55,dwa    &EW       & $+0.13\\pm0.05$ & Paulson et al.~(\\cite{paulson03})       \\\\ Praesepe & 0.7          &8.62         &   Friel~(\\cite{friel95})                 &   130     & 100000       \t   &7\t      &EW\t& $+0.27\\pm0.05$ & Pace et al. (\\cite{pace08})   \\\\ NGC2324  &0.7           & 11.65       & Friel et al.~(\\cite{friel02})            &   90-150 &45,000        &7,gia     &EW       &$-0.17\\pm0.05$ & Bragaglia et al.~(\\cite{bragaglia07})\\\\ NGC2660  & 0.8          &9.18         & Friel~(\\cite{friel95})                   &   50-80  &45,000        &5,gia     &EW       &$+0.04\\pm0.04$ & Sestito et al.~(\\cite{sestito06})\\\\ Mel71    & 0.8          &10.0         &   Friel~(\\cite{friel06})\t         &   100    & 16,000,34,000&2,gia     &EW       & $-0.32\\pm0.16$ & Brown et al.(1996)               \\\\ Ru 4     &0.8           & 12.40       &  Carraro et al.~(\\cite{carraro07})$^*$   &25--40    & 40,000       &3,gia     &EW       &$-0.09\\pm0.05$ &Carraro et al.~(\\cite{carraro07})  \\\\ Ru7      & 0.8          &13.50        &  Carraro et al.~(\\cite{carraro07})$^*$   &25--40    & 40,000       &5,gia     &EW       &$-0.26\\pm0.05$&Carraro et al.~(\\cite{carraro07})  \\\\ NGC2360  & 0.9          &9.28         &   Friel et al.~(\\cite{friel02})          &   50-100 & 28,000       &4,gia     &EW       & $+0.07\\pm0.07$ & Hamdani et al.~(\\cite{hamdani00}) \\\\ NGC2112  & 1.0          &9.21         &   Friel et al.~(\\cite{friel02})          &   70-80  & 16,000,34,000&2,gia     &EW       & $+0.16\\pm0.03$ & Carraro et al.~(\\cite{carraro08})           \\\\ NGC2477  & 1.0          &8.94         & Friel et al.~(\\cite{friel02})            &  100-150 &45,000        &6,gia     &EW       &$+0.07\\pm0.03$ & Bragaglia et al.~(\\cite{bragaglia07})\\\\ NGC7789  &1.1           &9.44         &   Friel et al.~(\\cite{friel02})          &   $>$50  & 30,000       &9,gia     &EW/synthesis& $-0.04\\pm0.05$ &Tautvaisiene et al.~(\\cite{taut05})   \\\\ NGC3680  & 1.4          &8.27         &   Friel et al.~(\\cite{friel02})          &    80    & 100,000      &2,dwa     & EW&     $-0.04\\pm 0.03$  &   Pace et al.~(\\cite{pace08})\\\\ NGC6134  & 1.6          &7.68         & WEBDA$^*$                                &   80-200 & 43,000       &3,gia     &EW       & $+0.15\\pm0.07$ & Carretta et al.~(\\cite{carretta04})\\\\ IC4651   & 1.7          &7.70         &    Friel~(\\cite{friel06})                &   100    & 48,000       &4,gia     &EW       & $+0.11\\pm0.01$ & Carretta et al..~(\\cite{carretta04})\\\\ Be73     &1.9           &17.40        &   Carraro et al.~(\\cite{carraro07})$^*$  &25--40    & 40,000       &2,gia     &EW       &$-0.22\\pm0.10$&Carraro et al.~(\\cite{carraro07})  \\\\ NGC2506  &2.2           &10.88        &   Friel et al.~(\\cite{friel02})          &   80-110 & 48,000       &3,gia     &EW       & $-0.20\\pm0.02$ & Carretta et al.~(\\cite{carretta04})   \\\\ NGC2141  &2.5           &12.42        &   Friel et al.~(\\cite{friel02})          &   100    & 28,000       &8,gia     &EW       & $-0.18\\pm0.15$ & Yong et al.~(\\cite{yong05})   \\\\ NGC6819  & 2.9          &8.18         &   Friel et al.~(\\cite{friel02})          &   130    & 40,000       &3,gia     &EW       & $+0.09\\pm0.03$ &Bragaglia et al. (2001)            \\\\ Be66     &3.7           &12.9         &   Phelps \\& Janes~(\\cite{phelpsjanes96}$^*$)&   5,15& 34,000       &2,gia     &EW       & $-0.48\\pm0.24$ & Villanova et al.~(\\cite{villanova05})   \\\\ Be22     &4.2           &11.92        &   Friel~(\\cite{friel95})                 &   20-25  & 34,000       &2,gia     &EW       & $-0.32\\pm0.19$ & Villanova et al.~(\\cite{villanova05})   \\\\ M67      & 4.3          &9.05         &   Friel et al.~(\\cite{friel02})          &   90-180 & 45,000       &10,dwa    &EW       & $+0.03\\pm0.01$ & Randich et al.~(\\cite{randich06})\\\\ Be20     &4.3           & 16.37       & Friel et al..~(\\cite{friel02})           &   40-80  &45,000        &2,gia     &EW       &$-0.30\\pm0.02$ & Sestito et al.~(\\cite{sestito08})\\\\ Be29     &4.3           & 22.0        & Friel~(\\cite{friel06})                   &   25-150 &45,000        &6,gia     &EW       &$-0.31\\pm0.03$ & Sestito et al.~(\\cite{sestito08})\\\\ NGC7142  &4.4           &  9.70       & Jacobson et al.~(\\cite{jacobson08})      &   100    &30,000        &4,gia     &EW       & $0.14\\pm0.01$ & Jacobson et al.~(\\cite{jacobson08})  \\\\ NGC6253  &  4.5         &7.10         & Bragaglia \\& Tosi~(\\cite{bragagliatosi06})&   60-140 &45,000       &5,gia     &EW       &$+0.36\\pm0.07$ & Sestito et al.~(\\cite{sestito07})\\\\ NGC2243  &4.7           &10.76        &   Friel et al.~(\\cite{friel02})          &   100    & 30,000       &2,gia     &EW       & $-0.48\\pm0.15$ & Gratton \\& Contarini~(\\cite{gratton94})  \\\\ Be75     & 5.1          &15.50        &  Carraro et al.~(\\cite{carraro07})$^*$   &25--40    & 40,000       &1,gia     &EW       &$-0.22\\pm0.20$&Carraro et al.~(\\cite{carraro07})  \\\\ Be31     &5.2           &12.02        &   Friel et al.~(\\cite{friel02})          &   100    & 28,000       &5,gia     &EW       & $-0.57\\pm0.23$ & Yong et al.~(\\cite{yong05})     \\\\ Mel66    & 5.4          &10.21        &   Friel et al.~(\\cite{friel02})          &   80-130 &45,000        &6,gia     &EW       &$-0.33\\pm0.03$ & Sestito et al.~(\\cite{sestito08})\\\\ Be32     &6.1           &11.35        & Friel et al.~(\\cite{friel02})            &  50-100  &45,000        &9,gia     &EW       &$-0.29\\pm0.04$ & Sestito et al.~(\\cite{sestito06})\\\\ Be25     &6.1           &18.20        &  Carraro et al.~(\\cite{carraro07})$^*$   &25--40    & 40,000       &4,gia     &EW       &$-0.20\\pm0.05$&Carraro et al.~(\\cite{carraro07})  \\\\ NGC188   & 6.3          &9.35         &   Friel et al.~(\\cite{friel02})          &   20-35  & 35,000-57,000&5,dwa     &EW       & $+0.01\\pm0.08$ &Randich et al. (2003)                 \\\\ Saurer1  &6.6           &19.3         &   Friel~(\\cite{friel06})                 &   80     & 34,000       &2,gia     &EW       & $-0.38\\pm0.14$ & Carraro et al.~(\\cite{carraro04})   \\\\ Cr 261   & 8.4          &7.52         & Friel et al.~(\\cite{friel02})            &   70-130 & 45,000       &7,gia     &EW       &$+0.13\\pm0.05$ & Sestito et al.~(\\cite{sestito08})\\\\ Be17     &  10.07       &11.2         &   Friel et al.~(\\cite{friel05})          &   15-100 & 25,000       &4,gia     &EW       & $-0.10\\pm0.09$ & Friel et al.~(\\cite{friel05})\\\\ NGC6791  & 10.9         &8.12         &   Friel et al.~(\\cite{friel02})          &   10-60  & 20,000       &15,gia    &synthesis& $+0.39\\pm0.08$ & Carraro et al.~(\\cite{carraro06}) \\\\  \\hline \\hline \\end{tabular} \\label{tab_ocs} \\end{table*} \\end{landscape}    \\setcounter{table}{2} \\begin{landscape} \\begin{table*} \\scriptsize \\label{table:grad} \\caption{ [Fe/H], [Ca/H], [Si/H], [Ti/H], [Ni/H], [Cr/H] gradients (dex~kpc$^{-1}$) in different radial ranges.  The observed gradients are computed from the data discussed in Sect.~\\ref{sect_data} using a weighted linear fit. Model gradients are approximated by linear fits. $^\\dag$ For OCs with $4~\\mbox{Gyr}<\\mbox{age}<11~\\mbox{Gyr}$, the presence of the two metal rich OCs discussed in the text leads to a steep gradient ($-0.16$). Excluding them we still obtain flatter gradients ($-0.13$), but still steeper than model predictions. $^{(a)}$ Andrievsky et al.~(\\cite{andri04}) in the radial ranges: 4.0--6.6~kpc, 6.6--10.6~kpc, 10.6--14.6~kpc.  $^{(b)}$ Lemasle et al.~(\\cite{lemasle07}) in the radial range 8--12~kpc.  $^{(c)}$ OCs from Table~\\ref{tab_ocs} and references therein, $^{(d)}$ in the radial range 7-12~kpc} \\centering \\hspace{-5cm} \\begin{tabular}{lllllllll} \\hline \\smallskip &                       &  model      &           &           & observations        &               &                &                 \\\\ \\hline gradient & range                 & present time & 4~Gyr ago & 11~Gyr ago & Cepheids             & OCs$^{(c)}$            &OCs$^{(c)}$             & OCs$^{(c)}$               \\\\ &(kpc)                        &             &           &           & present time         & $\\mbox{age}\\leq 0.8$~Gyr   & $0.8~\\mbox{Gyr}<\\mbox{age}\\leq 4~\\mbox{Gyr}$ & $4~\\mbox{Gyr}<\\mbox{age}\\leq 11~\\mbox{Gyr}$  \\\\ \\hline \\smallskip $\\Delta$[Fe/H]/$\\Delta R$ &4--7   & $-0.134$    & $-0.162$  & $-0.176$  & $-0.128\\pm0.029^{(a)}$  & --               & --               & --                \\\\ $\\Delta$[Fe/H]/$\\Delta R$ & 7--12 & $-0.075$    & $-0.071$  & $-0.076$  & $-0.044\\pm 0.011^{(a)}$ & $-0.053\\pm 0.029$  & $-0.094\\pm 0.008$  & $-0.091\\pm0.06$\\\\ &                     &           &              & $-0.051\\pm 0.004^{(b)}$  &                  &                  &          \\\\ $\\Delta$[Fe/H]/$\\Delta R$ &12--22 & $-0.019$    & $-0.023$  & $-0.028$  & $+0.004\\pm 0.011^{(a)}$  & --                 & $+0.005\\pm 0.050$ &  $-0.001\\pm 0.008$              \\\\ \\hline \\smallskip $\\Delta$[Ca/H]/$\\Delta R$ &4--7    & $-0.137$    & $-0.172$  & $-0.156$  &                       & --               & --               & --                \\\\ $\\Delta$[Ca/H]/$\\Delta R$ & 7--12  & $-0.073$    & $-0.067$  & $-0.076$ &                          &  --              & $0.02\\pm0.03$    & $-0.13\\pm0.03$        \\\\ $\\Delta$[Ca/H]/$\\Delta R$ &12--22  & $-0.018$   & $-0.024$  & $-0.029$  &                          &  --              & $-0.03\\pm 0.09$  &  $0.009\\pm 0.009$              \\\\ \\hline \\smallskip $\\Delta$[Si/H]/$\\Delta R$ &4--7    & $-0.124$    & $-0.153$  & $-0.126$  &                       & --               & --                    & --                \\\\ $\\Delta$[Si/H]/$\\Delta R$ & 7--12  & $-0.063$    & $-0.056$  & $-0.076$ &                          & $-0.05\\pm 0.03$  & $-0.07\\pm 0.10$       &  $-0.12\\pm0.03$           \\\\ $\\Delta$[Si/H]/$\\Delta R$ &12--22  & $-0.017$   & $-0.024$  & $-0.029$  &                          &  --              &   --                  &  $0.003\\pm 0.013$          \\\\ \\hline \\smallskip $\\Delta$[Ti/H]/$\\Delta R$ &4--7    & $-0.114$    & $-0.135$  & $-0.138$  &                       & --               & --                    & --                \\\\ $\\Delta$[Ti/H]/$\\Delta R$ & 7--12  & $-0.039$    & $-0.039$  & $-0.046$  &                         & $-0.08\\pm 0.03$  & $-0.07\\pm 0.06$        & $-0.16\\pm0.04$            \\\\ $\\Delta$[Ti/H]/$\\Delta R$ &12--22  & $-0.007$    & $-0.009$  & $-0.018$  &                         &  --              & $-0.04\\pm 0.02$       &  $0.02\\pm 0.01$              \\\\ \\hline \\smallskip $\\Delta$[Ni/H]/$\\Delta R$ &4--7    & $-0.126$    & $-0.157$  & $-0.156$     &                       & --                & --                   & --                \\\\ $\\Delta$[Ni/H]/$\\Delta R$ & 7--12   & $-0.072$    & $-0.071$  & $-0.076$    &                         & --                & $-0.07\\pm 0.06$      & $-0.13\\pm0.03$             \\\\ $\\Delta$[Ni/H]/$\\Delta R$ &12--22  & $-0.020$   & $-0.024$  & $-0.029$      &                         & --                & --                   &  $0.00\\pm 0.02$         \\\\ \\hline \\smallskip $\\Delta$[Cr/H]/$\\Delta R$ &4--7      & $-0.130$    & $-0.162$  & $-0.153$     &                       &    &        & --                \\\\ $\\Delta$[Cr/H]/$\\Delta R$ & 7--12    & $-0.072$    & $-0.068$  & $-0.075$&                         & $-0.08\\pm 0.03$              & $0.04\\pm 0.11$                   & $-0.18\\pm0.03$             \\\\ $\\Delta$[Cr/H]/$\\Delta R$ &12--22    & $-0.020$   & $-0.024$  & $-0.029$   &                         & --                & --                    &  $-0.006\\pm 0.03$              \\\\  \\hline \\end{tabular} \\hspace{-5cm} \\label{tab_grad_iron} \\end{table*} \\end{landscape}  \\appendix"
    },
    "0812/0812.2917_arXiv.txt": {
        "abstract": "We study the stochastic background of gravitational waves produced from preheating in hybrid inflation models.  We investigate different dynamical regimes of preheating in these models and we compute the resulting gravity wave spectra using analytical estimates and numerical simulations. We discuss the dependence of the gravity wave frequencies and amplitudes on the various potential parameters. We find that large regions of the parameter space leads to gravity waves that may be observable in upcoming interferometric experiments, including Advanced LIGO, but this generally requires very small coupling constants. ",
        "introduction": "Gravity waves (GW) from the early universe can carry information about inflation, (p)reheating after inflation, and even about pre-inflation.  During inflation, tensor modes of the classical large scale cosmological perturbations are produced from the quantum fluctuations of gravitons \\cite{starob79}.  Their amplitude is proportional to the energy scale of inflation, and their spectrum extends over a wide range of wavelengths.  Gravitational waves at the cosmological, long-wavelength part of the spectrum lead to B-mode polarization of the CMB anisotropy fluctuations.  In high energy models of inflation (like chaotic inflation), where the ratio of the amplitudes of the tensor to scalar modes is about $ r \\sim 0.1$, the B-mode of anisotropies should be detectable by forthcoming CMB polarization experiments.  However, in anticipation of a null signal observation of GW from inflation, one might still be able to use GW to constrain inflationary models in ways other than just constraining the overall energy scale. There are models of inflation where the total number of inflationary e-folds $N$ exceeds the minimum required to homogenize the observable universe only by a small margin. For such models with anisotropic pre-inflationary expansion, gravity waves from pre-inflation are amplified and may contribute to the large scale CMB temperature anisotropy \\cite{GKP}.  Observational limits on gravity waves from CMB polarization experiments also result in constraints on $N$.  If inflation occurs at lower energies (as in many hybrid inflation models), the amplitude of the resulting gravity waves would be too weak to be observed with CMB anistropies. However, GW from inflation at the short-wavelength part of the spectrum might fall in the amplitude-frequency range which accessible to future gravity wave astronomy projects such as BBO and DECIGO.  These experiments could probe $r$ down to the level $r \\sim 10^{-6}$.  Gravity waves may also carry unique information about the post-inflationary dynamics, in particular the inflaton decay and the subsequent evolution of its decay products towards thermal equilibrium.  Often  this process starts with preheating, a violent non-perturbative restructuring of the field configurations shortly after inflation. Preheating leads to large, non-linear field inhomogeneities which necessarily generate a classical GW background. This mechanism of GW production is complementary to the production of GW from vacuum fluctuations during inflation. In our previous paper \\cite{DBFKU} we developed a machinery for analytic and numerical calculations of GW production in preheating models. There are several other papers on the subject that we will discuss below. It is also worth mentioning that the techniques used to study GW from preheating have parallels in calculations of GW from phase transitions / hydrodynamical turbulence, see e.g. \\cite{transitions} and references therein. In \\cite{DBFKU} we developed applications of our methods to the example of preheating after chaotic inflation. The parameters of these large-field models are fixed by CMB normalization, which ensures that inflation must occur at high energy scales.  As a result, the gravity waves produced from preheating in these models have typical frequencies of order $f \\sim 10^7 - 10^8$ Hz, which is far beyond the observable range (see Section \\ref{Discussion} for observational constraints). As conjectured in \\cite{frag}, preheating after low-energy inflation may generate GW with lower frequencies $f$. One of the most popular and most studied models of inflation and preheating is hybrid inflation \\cite{hybrid} with tachyonic preheating \\cite{tachpre}. There are many hybrid inflation models, notably in the context of supergravity (see \\cite{lythriotto} for a review) and string theory (see \\cite{stringinf} for reviews). In these models, the energy scale of inflation is not fixed by CMB normalization, so it was optimistically thought that for some parameters the frequency of GW from preheating may fall into an observable range. The subject of this paper is GW production from tachyonic preheating, in particular, assessing how realistic is the range of parameters that may lead to an observable signal.  Let us briefly review the literature on the subject.  The production of gravitational waves from preheating was originally studied in \\cite{pregw1} for chaotic inflation, and recently by several groups in \\cite{pregw2, frag, pregw3, juan1, juan2, DBFKU, pregw4, pregw5}.  The first numerical methods \\cite{pregw1, pregw2} were based on the Weinberg formalism \\cite{weinberg}, which, strictly speaking is applicable only for isolated sources in a Minkowski background.  This may lead to significant differences in the resulting gravity wave spectra, as explained in \\cite{DBFKU}. In \\cite{pregw3}, the gravity wave equations in an expanding universe were solved in Fourier space. Another method was used in \\cite{juan1, juan2}, where the evolution equation for metric perturbations were solved in configuration space. In \\cite{juan1} and the earlier version of \\cite{juan2} the transverse-traceless radiative part of GW was not properly extracted\\footnote{Incorrect extraction of TT part, in particular, may lead to the incorrect conclusion that a significant amount of gravity waves is produced from the stage of scalar field ``turbulence'' after preheating, see \\cite{DBFKU} for details.}  but after discussion of the problem in \\cite{DBFKU} this was corrected in the second version of \\cite{juan2}.  Finally, in \\cite{DBFKU}, we developed a method based on the Green's function solution in momentum space, to calculate, numerically and analytically, the production of gravity waves from a stochastic medium of scalar fields in an expanding universe. At present, the numerical methods of different groups \\cite{pregw3}, \\cite{juan2} and \\cite{DBFKU} seem to agree well with each other, see \\cite{pregw4,pregw5} for a comparison of the results in a model of chaotic inflation.  Most of these papers focused on models of preheating based on parametric resonance after chaotic inflation, which leads to gravity wave signals at frequencies which are irrelevant observationally. Models of preheating after hyrbid inflation are observationally much more interesting, because they involve extra parameters and they occur at lower energy scales, and thus they may generate GW with lower frequencies.  Gravity wave production after hybrid inflation was first considered in \\cite{juangw}, extrapolating the results of \\cite{pregw1}, assuming that preheating occurs through parametric resonance. However, we know since \\cite{tachpre} that preheating after hybrid inflation occurs in the qualitatively different regime of tachyonic amplification due to the dynamical symmetry breaking, when the fields roll towards the minimum through the region where their effective mass is negative -- a process called tachyonic preheating.  It completes in one or very few oscillations of the fields and is very different from parametric resonance\\footnote{Hybrid inflation was also invoked in \\cite{pregw3} to motivate a $m^2\\,\\phi^2 + g^2 \\phi^2 \\chi^2$ model of parametric resonance at low energy scales.  However, as noted above, this model is not relevant for preheating after hybrid inflation.  The gravity wave spectra shown in \\cite{pregw3} fall into an observable range for an inflaton mass of order $10$ GeV, instead of $10^{13}$ GeV in chaotic inflation (and for a coupling constant of order $10^{-30}$).  One can also consider GW production from the non-perturbative decay of a scalar field $\\phi$ in this model not necessarily related to the inflaton, as for example in the scenario considered in \\cite{curvaton}. The present-day peak frequency $f_*$ may then be estimated, in the same way as we do below, from the typical momentum $k_* \\sim \\sqrt{g m \\Phi_0} a^{-3/4}$ amplified in this model. If $\\phi$ dominates the energy density before decaying and the decay is followed by the radiation dominated era, one finds $f_* > \\sqrt{g}\\,10^{10}\\,\\mathrm{Hz}$, independently of the mass $m$ and the initial amplitude $\\Phi_0$ of the field $\\phi$ (as long as $g \\Phi >> m$ for preheating to occur). This bound can be relaxed if $\\phi$ does not dominate the energy density, but a very small coupling $g^2$ is still required to achieve $f_* < 10^3\\,\\mathrm{Hz}$, a necessary condition for these GW to be observable.}.  Before further discussion on GW, we have to comment about specific challenges for numerical simulations of tachyonic preheating in hybrid inflation models with several parameters \\footnote{Contrary to the case of resonant preheating after single-parameter chaotic inflation}. These parameters include the self-coupling $\\lambda$ of the symmetry breaking fields and their coupling $g^2$ to the inflaton. The scales of the preheating dynamics can be significantly  different depending on the combination $g^2/\\lambda$ (roughly speaking, the ratio of effective masses around the local maximum and minimum of the potential). Therefore,  care should be taken to incorporate both scales in the simulations. An additional subtlety is related to the initial conditions for the fields around the bifurcation point, namely, relatively high or low inflaton velocity, which results in qualitatively  different initial regimes -- quantum diffusion or classical fast or slow rolls around the bifurcation point, see \\cite{tachpre} for details. The case with higher initial velocity is easier to model numerically.  The production of gravitational waves properly  from tachyonic preheating was first investigated in \\cite{juan1, juan2}, for a specific region of the parameter space ($g^2 \\sim \\lambda \\sim 1$ and significant   velocity of the inflaton at the bifurcation point), where the dynamical scales are of the same order. The resulting gravity wave spectra were located at frequencies too high to be observable, as in chaotic inflation models. Refs. \\cite{juan1, juan2} also displayed a conjecture for a low frequency spectrum, but the dependence on the parameters was not studied, so it remains  unclear which model, if any, could lead to an observable signal. In \\cite{frag}, it was conjectured that the peak frequency and amplitude of the gravity wave spectrum depends in a relatively simple way on the typical scale amplified during preheating, which is usually a known function of the parameters. In \\cite{DBFKU} we verified this conjecture  numerically for a model of chaotic inflation, and used it to estimate analytically the peak of the gravity wave spectrum produced in two different models of preheating after hybrid inflation. In particular, we noticed that models with low  velocity of the inflaton at the critical point are observationally more promising, but it is also more challeging to model this case with  numerical simulations, as noted above.  In the present paper, we investigate in detail gravity waves produced from preheating after hybrid inflation, focusing in particular on their dependence on various model parameters such as $g^2/\\lambda$ and the field velocities around the bifurcation point.  We study GW production in several qualitatively different settings of tachyonic preheating, some of which, including the most interesting cases, were completely missed in the literature. We use a combination of analytical estimates and improved numerical simulations to compute the resulting gravity wave spectra and to identify the regions of the parameter space which may be relevant for upcoming gravity wave experiments.   The rest of the paper is organized as follows. In Section \\ref{model}, we briefly review the hybrid inflation models that we will consider. In Section \\ref{analytics}, we identify three qualitatively different dynamical regimes of preheating in these models and we make analytical estimates of the resulting gravity wave spectra. In Section \\ref{gwcalculations} we describe  and advance further our numerical method for calculating gravity wave production from preheating. We apply this method to preheating after hybrid inflation in Section \\ref{numerics} and we study the dependence of the gravity wave amplitudes and frequencies on the parameters of the model. We then discuss which regions of the parameter space may lead to a detectable signal. Finally, we discuss the implications of our results and directions for future work. Some details of our numerical calculations are given in the appendices. ",
        "conclusions": "\\label{Discussion}  In this paper, we studied the stochastic background of gravitational waves produced from tachonic preheating after hybrid inflation in the very early universe. The present-day frequencies and amplitudes of these gravity waves may cover a wide range of values, depending on the main phenomenological parameters introduced in Section II: the VeV $v$ and self-coupling $\\lambda$ of the symmetry breaking fields, their coupling to the inflaton $g^2$ and the (unitless) initial velocity of the inflaton at the critical point $V_c$. We developed analytical and numerical tools to calculate the resulting GW spectra and to study in detail how they depend on these parameters. We identified three dynamical regimes leading to qualitatively different results for the GW spectra produced from tachyonic preheating: (i) the case with $g^2 \\gtrsim \\lambda$ and the onset of preheating driven by the classical rolling of the inflaton, (ii) the case with $g^2 \\gtrsim \\lambda$ and quantum diffusion at the onset of preheating (corresponding to a negligible initial velocity of the inflaton at the bifurcation point) and (iii) the case with $g^2 << \\lambda$ where GW are produced in successive bursts. Only the first case was considered in previous works about GW from tachyonic preheating \\cite{juan1, juan2}, where the resulting GW spectrum was computed in the specific range $g^2 \\sim \\lambda \\sim V_c \\sim \\mathcal{O}(1)$.  Based on our results, we can determine the range of parameters of hybrid inflation models for which preheating may lead to a GW signal that is potentially observable. Let us first discuss the frequencies and amplitudes of stochastic GW backgrounds that are relevant for GW astronomy. \\begin{figure}[htb] \\centering \\leavevmode \\epsfxsize=18cm \\epsfbox{sensit2.eps} \\caption{Expected sensitivities of interferometric experiments (adapted from \\cite{buonanno}) compared to gravity wave spectra from tachyonic preheating in the cases of significant intial velocity with $g^2 \\gtrsim \\lambda$ (green), negligible initial velocity with $g^2 \\gtrsim \\lambda$ (red), and $g^2 << \\lambda$ (blue). Also shown are the BBN and ms pulsar bounds, the stochastic background from extragalactic White Dwarfs binaries (taken from \\cite{phinney}), the inflationary background for two values of the tensor to scalar ratio, and the spectrum from preheating in a $\\lambda \\, \\phi^4$ model of chaotic inflation (double spectra in grey at the right, calculated with our previous code \\cite{DBFKU}). The GW spectra from tachyonic preheating, for the peak from left to right, correspond to the following values of the parameters: (1) $\\lambda = 2 g^2 = 10^{-14}$, $v = 3\\,10^{-7}\\,\\Mp$ and negligible initial velocity ; (2) $\\lambda = 0.1$, $g^2 = 10^{-16}$ and $v = 3\\,10^{-10}\\,\\Mp$ (independent of $V_c$) ; (3) $\\lambda = 2 g^2 = 10^{-14}$, $v = 10^{-3}\\,\\Mp$ and $V_c = 10^{-3}$ ; (4) $\\lambda = 0.1$, $g^2 = 10^{-4}$ and $v = 6\\,10^{-4}\\,\\Mp$ (independent of $V_c$) ; (5) $\\lambda = 2 g^2 = 10^{-5}$, $v = 10^{-3}\\,\\Mp$ and negligible initial velocity.} \\label{sensi} \\end{figure} We consider the present-day frequencies $f$ and the spectrum of energy density per logarithmic frequency interval $h^2\\,\\Omega_{\\mathrm{gw}}(f)$. Fig.~\\ref{sensi} shows the sketch of expected sensitivities of planned and future interferometric experiments: Advanced LIGO, LISA, the Einstein Telescope, Big Bang Observer and DECIGO (see \\cite{buonanno} for the relevant references). Note also the recent atomic interferometric sensor proposal of \\cite{agis}, which may be sensitive to GW backgrounds with frequencies in the range $1-10$ Hz and in the LISA range. We also show the BBN and ms pulsar bounds, and predictions for the stochastic GW background generated from inflation, for different values of the parameter $r$ (the ratio of the tensor and scalar amplitudes of the inflationary cosmological perturbations). Cosmological GW signals will be obscured by several expected astrophysical foregrounds, notably from White Dwarf binaries \\cite{phinney}, see Fig.~\\ref{sensi}.  There are also several bar and spherical resonant detectors (see e.g. \\cite{minigrail}) operating in the kHz range. Other experiments have been proposed at higher frequencies, up to $100$ MHz \\cite{Nishizawa:2007tn}. However, the sensitivity to $h^2\\,\\Omega_{\\mathrm{gw}}$ \\footnote{The concept of $h^2\\,\\Omega_{\\mathrm{gw}}$ for stochastic backgrounds is not applicable for GW signals peaked within a very narrow frequency band.} drops dramatically when the frequency increases. Indeed, we also have to take into account the current ``Standard Quantum Limit'' \\begin{equation}\\label{SQL} h^2\\,\\Omega_{\\mathrm{gw}} \\sim \\frac{10^{-17}}{x^2}  \\, \\left(\\frac{f}{Hz}\\right)^3 \\ , \\end{equation} shown in Fig.~\\ref{sensi}. The region on the right of this line is not observable with interferometric experiments due to the shot noise fluctuations of photons. Here $x$ is the squeezing parameter, which experimentalists are trying to push from one to a few.  Exemples of GW spectra from tachonic preheating calculated in this paper are shown in Fig.~\\ref{sensi}. We also show in this figure a GW spectrum from preheating after chaotic inflation calculated in \\cite{DBFKU} (the double-spectrum in grey), which is located in the high-frequency side and is not observable. GW from tachyonic preheating may cover a much wider range of frequencies, spreading from the highest frequencies shown in the figure to the region in between the SQL limit and the WD binaries, which is in principle observable \\footnote{For the  sake of curiosity, one can notice the analogy with the $2.73$K CMB radiation which is observable in the strip of frequencies between the borders of the foreground dust and synchrotron radiations, as was outlined in early 1960.}. In each case shown in the figure, a lower Higgs VeV would lead to lower amplitudes and smaller coupling constants would lead to smaller frequencies.  Our analytical and numerical results for the peak frequency $f_*$ and the peak amplitude $h^2\\,\\Omega_{\\mathrm{gw}}^*$ of the GW spectra are well described by Eqs.~(\\ref{fstarVc})-(\\ref{OmegastarVc}) in the case with $g^2 \\gtrsim \\lambda$ and the onset of preheating driven by the classical rolling of the inflaton, Eqs.~(\\ref{fstarV0})-(\\ref{OmegastarV0}) in the case with $g^2 \\gtrsim \\lambda$ and quantum diffusion at the onset of preheating (negligible initial velocity of the inflaton at the critical point) and Eqs.~(\\ref{fstargnum})-(\\ref{Omegastargnum}) in the case with $g^2 << \\lambda$. The peak frequency depends essentially on the coupling constants and is independent of the Higgs VeV $v$. As a result, $\\lambda \\sim g^2 \\sim \\mathcal{O}(1)$ leads to GW at very high frequencies, independent of the energy scale during inflation, and the only way to lower the frequency is to lower the coupling constants. On the other hand, $h^2\\,\\Omega_{\\mathrm{gw}}^* \\propto v^2$, so that at any frequency the amplitude can be relatively high for sufficiently high $v$, roughly up to the maximal bound~\\footnote{This bound follows from the approximate equation (\\ref{omegaR}) with the constraint $k_* > H$ (meaning that the scalar fields and the gravity waves are produced inside the Hubble radius). The same constraint, $k_* > H$, is also roughly the condition for the expansion of the universe to be negligible and for the waterfall condition to be satisfied. It also implies that $v << \\Mp$ and (for the small coupling constant we consider) that the inflationary model is not ruled out by the non-observation of inflationary gravity waves.} $h^2 \\Omega_{\\mathrm{gw}}^* < 10^{-6}\\,$. Such a high energy density in GW implies in particular that these GW may already be observable by Advanced LIGO, but this generally requires very small coupling constant(s), see Fig.~\\ref{sensi}. For illustration, we show in Fig.~\\ref{constraints}, for each of the three cases discussed above, the range of parameters of the hybrid inflation models such that the peak frequency of the GW produced from preheating satisfy $f_* < 10^3$ Hz. This is basically the condition for these gravity waves to be observable if $v$ is set to the maximum possible value consistent with the waterfall condition. In the first case of Fig.~\\ref{constraints}, corresponding to $g^2 \\gtrsim \\lambda$ and significant initial velocity, the coupling of interest $\\lambda$ typically has to be extremely small, down to $\\lambda < 10^{-30}$ for $V_c \\sim 1$. In the second case, with negligible initial velocity, we may have $g^2 \\sim \\lambda \\sim 10^{-11}$. In the third case ($g^2 << \\lambda$), the most interesting regime corresponds to $\\lambda \\sim 1$, but this still requires $g^2 < 10^{-14}$.  \\begin{figure}[hbt] \\centering \\leavevmode \\epsfxsize=18cm \\epsfbox{constraints.eps} \\caption{ The regions of parameter space for which the peak of the gravity wave signal satisfies $f_* < 10^3$ Hz. The panels show: Left - significant initial velocity and $g^2 \\gtrsim \\lambda$. Middle - negligible initial velocity and $g^2 \\gtrsim \\lambda$. Right - $g^2 \\ll \\lambda$.} \\label{constraints} \\end{figure}  Above we focused on the peak frequency $f_*$ of the GW spectra. At lower frequencies, the spectra have a power-law tail, $\\Omega_{\\mathrm{gw}} \\propto f^{\\alpha}$. For tachyonic preheating, we found in this paper an intermediate regime with a relatively steep slope, $\\alpha = 2 - 2.5$, before the IR tail with $\\alpha = 3$ corresponding to the modes produced outside the Hubble radius. The infra-red part of the spectrum does not significantly improve the constraints on the parameters shown in Fig. \\ref{constraints}. As outlined above, the range of parameters leading to the maximum amplitude in GW correspond to $k_*$ slightly below $H$. In this case, the spectra decrease as $\\Omega_{\\mathrm{gw}} \\propto f^3$ for frequencies just below the peak. In particular, since this is the same dependence as the SQL limit (\\ref{SQL}), if the peak of the GW spectrum is located at the right of this line, the same is true for the IR tail.  In sum, our sober conclusion is that GW from tachyonic preheating after hybrid inflation can fall in the observable range of the $h^2 \\Omega_{\\mathrm{gw}} -- f$ plane that is currently conceivable only for uncomfortably small values of the coupling constants of the model.  We should note, however, that particle physics models of hybrid inflation with such small coupling constants have been considered. First, if the field $\\sigma$ in the model (\\ref{hybrid}) corresponds to a flat direction, we may expect its negative mass squared at the origin to be of the order $\\lambda v^2 \\sim \\mathrm{TeV}^2$, and the true minimum of the potential to be located at a VeV at a much higher scale, as high as $v \\sim M_\\mathrm{Pl}$. This gives a coupling constant as low as $\\lambda \\sim 10^{-30}$. Hybrid inflation models along this line have been proposed in \\cite{randall, stewart}. Another model which, for the same reason, naturally involves very small coupling constants is thermal inflation \\cite{thermal}. Preheating mechanisms in this model have not been studied in detail yet, but tachyonic amplifications may lead to an interesting GW signal, see also \\cite{wanil}. Another model of hybrid inflation in supergravity is P-term inflation \\cite{pterm}, which includes F-term and D-term models as special cases and may also be realized in D3/D7 models in string theory. Satisfying the observational bounds on cosmic strings in P-term inflation may require small couplings ($\\lambda \\sim g^2 < 10^{-10}$), see \\cite{pterm}. Finally, we note that the case $g^2 << \\lambda$ may have connections with brane-anti-brane inflation in a warped throat \\cite{braninf}, which is another prototype of hybrid inflation in string theory. According to \\cite{neil}, the end of inflation in this scenario should exhibit similarities with the model (\\ref{hybrid}) with $g^2 << \\lambda$ (see in particular Eqs.~(30)-(32) of the second paper in \\cite{neil}). In this case, a very small $g^2$ would emerge naturally because of the exponential warping of the throat.  More generally, the hybrid inflation model (\\ref{hybrid}) that we considered in this paper is only one of the much broader class of inflationary models where preheating occurs through tachyonic effects. GW production from other models involving tachyonic preheating is a subject for further investigations.    \\vskip 0.5in \\vbox{ \\noindent{ {\\bf Acknowledgments} } \\\\ \\noindent We thank Juan Garcia-Bellido, Daniel Figueroa and Alessandra Buonanno for useful discussions.  We also acknowledge the use of the CITA and IFT computation clusters, and thank CITA and KITP for hosting G.F. during part of this work. The work of J.F.D. was supported by the Spanish MEC, via FPA2006-05807 and FPA2006-05423. G.F. was supported by PHY-0456631. L.K. was supported by NSERC and CIFAR.} \\vskip 0.5in"
    },
    "0812/0812.1313_arXiv.txt": {
        "abstract": "Radial-velocity measurements and sine-curve fits to the orbital radial velocity variations are presented for the last eight close binary systems analyzed the same way as in the previous papers of this series: QX~And, DY~Cet, MR~Del, HI~Dra, DD~Mon, V868~Mon, ER~Ori, and Y Sex. For another seven systems (TT~Cet, AA~Cet, CW~Lyn, V563~Lyr, CW~Sge, LV~Vir and MW~Vir) phase coverage is insufficient to provide reliable orbits but radial velocities of individual components were measured. Observations of a few complicated systems observed throughout the DDO close-binary program are also presented; among them an especially interesting is the multiple system V857~Her which -- in addition to the contact binary -- very probably contains one or more sub-dwarf components of much earlier spectral type. All suspected binaries which were found to be most probably pulsating stars are briefly discussed in terms of mean radial velocities and projected rotation velocities ($v \\sin i$) as well as spectral type estimates. In two of them, CU~CVn and V752~Mon, the broadening functions show a clear presence of non-radial pulsations. The previously missing spectral types for the DDO~I paper are given here in addition to such estimates for most of the program stars of this paper. ",
        "introduction": "\\label{sec1}  This is the last in a series of papers presenting results of spectroscopic observations taken within the program of radial velocity (hereafter RV) orbits of close binary stars at the David Dunlap Observatory (DDO); it contains a discussion of the all remaining targets of this program. For full references to the previous papers, see the last paper by \\citet[Paper XIV]{ddo14}; for technical details and conventions, for preliminary estimates of uncertainties, and for a description of the broadening functions (BFs) technique, see the interim summary paper \\citet[Paper VII]{ddo7}.  Most of data used in the present paper were obtained -- as in most of the previous fourteen papers -- using the broadening function approach applied to the spectral region of the Mg~I triplet (always centered at 5184~\\AA). Up to August 2005 we used a diffraction grating with 1800 lines/mm, after which a new one with 2160 lines/mm was used. A small number of observations were taken in the red region centered at 6290\\AA\\ which includes a telluric band to resolve our concerns on the stability of the RV system (see \\citet[Paper XI]{ddo11}). The RV observations reported in this paper have been collected between October 1996 and July 2, 2008, the day when the David Dunlap Observatory ceased to operate. The ranges of dates for individual systems can be found in Table~\\ref{tab1}. We note that this program utilized the efficient code of \\citet{pych2004} for removal of cosmic rays from 2-D images.  Throughout our program, selection of the targets was quasi-random: At a given time, we observed a few dozen close binary systems with periods usually shorter than one day, brighter than 10 -- 11 magnitude and with declinations $>-20^\\circ$; we published the results in groups of ten systems as soon as reasonable orbital elements were obtained from measurements evenly distributed in orbital phases. This paper is an exception as we have not been able to collect ten orbits before the observatory closure and we were left with material for a few stars already started. For that reason, in this last paper of the series, we present all remaining spectroscopic observations. Thus, they include: (i)~radial velocity orbits for eight close binaries observed and fully analyzed in the same way as in previous publications of this series (QX~And, DY~Cet, MR~Del, HI~Dra, DD~Mon, V868~Mon, ER~Ori, and Y~Sex); (ii)~close binaries with just one quadrature covered or with a few spectra available but not sufficient in number to define a reasonable orbit (TT~Cet, AA~Cet, CW~Lyn, V563~Lyr, CW~Sge, LV~Vir, and MW~Vir); (iii)~complicated or faint binaries/triples not providing reliable RVs (GO~Cyg, V857~Her, V752~Mon, V353~Peg, and MS~Vir), (iv)~additional data and improved orbits for a few multiple systems already presented within this series (ET~Boo, XY~Leo, and TV~UMi); (v)~radial velocities and projected rotational velocities, $v \\sin i$, for several stars found to be pulsating rather than being close binaries, and (vi)~the previously missing spectral-type estimates for close binaries in \\citet[Paper I]{ddo1}. All the RVs are given in Table~\\ref{tab1}. Among the present targets, only QX~And, ER~Ori, and Y~Sex, have had spectroscopic orbits previously published while preliminary orbits of AA~Cet and TT~Cet, based on limited data, were given by \\citet{south2}.  The RVs for the eight short period binaries reported in this paper were determined by fitting the double rotational profiles to extracted BFs, as explained in \\citet{ddo11}. Similarly as in our previous papers dealing with multiple systems (here the cases of MR~Del, V563~Lyr, DD~Mon, ER~Ori, Y~Sex, and LV~Vir), the RVs for the eclipsing pair were obtained after removal of the slowly rotating component, as was described most recently in \\citet{ddo14}.  As in other papers of this series, whenever possible, we estimated spectral types of the program stars using new classification spectra centered at 4200 \\AA\\ or 4400 \\AA. Additional classification spectra were obtained for part of the systems published in \\citet[Paper I]{ddo1} where no classifications were given. The estimated spectral types were compared with the mean $(B-V)$ color indices usually taken from the Tycho-2 catalog \\citep{Tycho2} and the photometric estimates of the spectral types using the relations of \\citet{Bessell1979}. In this paper we also made use of infrared colors determined from the 2$\\mu$ All Sky Survey (2MASS, \\citet{2MASS}). Especially useful is the $J-K$ color index, which is monotonically rising from the early spectral types to about M0V \\citep{cox2000}; this index is relatively less affected by the interstellar absorption than $B-V$. Parallaxes cited throughout the paper were adopted from the new reduction of the Hipparcos raw data \\citep{newhipp} which supersedes the original reductions \\citep{hip}.  The RV data for the binaries are given in Table~\\ref{tab1}. The preliminary sine-curve solutions for the eight binaries are in Table~\\ref{tab2} while the phase diagrams are shown in Figures~\\ref{fig1} and \\ref{fig2}. Section~\\ref{sec2} of the paper contains summaries of previous studies for individual systems with reliable orbits and comments on the new data. Systems with the insufficient orbit coverage are discussed in Section~\\ref{sec3}, while problematic systems for which RVs could not be determined are discussed in Section~\\ref{sec4}. Targets which were found to be pulsating are presented in Section~\\ref{sec5}. Examples of BFs of individual  systems are shown in Figures~\\ref{fig3} -- \\ref{fig5}.  The data for the eight close binaries in Table~\\ref{tab2} are organized in the same manner as in the previous papers of this series. In addition to the parameters of spectroscopic orbits, the table provides information about the relation between the spectroscopically observed upper conjunction of the more massive component, $T_0$ (not necessarily the primary eclipse) and the recent photometric determinations of the primary minimum in the form of the $O-C$ deviations for the number of elapsed periods $E$. The reference ephemeris of HI~Dra was taken from \\citet{gome1996} and for V868~Mon from \\citet{oter2004}. For the rest of systems the ephemerides were adopted from the on-line version of ``An Atlas of O-C diagrams of eclipsing binary stars''\\footnote{http://www.as.wsp.krakow.pl/ephem/} \\citep{Kreiner2004}. Because the on-line ephemerides are frequently updated, we give those used for the computation of the $O-C$ residuals below Table~\\ref{tab2} (the status as of October 2008). The deeper eclipse in W-type contact binary systems corresponds to the lower conjunction of the more massive component; in such cases the epoch in Table~\\ref{tab2} is a half-integer number. ",
        "conclusions": "\\label{summary}  With RV orbits for the last eight short-period binaries, this last paper of the DDO series brings the number of the systems studied at the David Dunlap Observatory to 141. There have been 138 orbits contained in 14 of 15 papers of this series (the seventh paper \\citep{ddo7} did not contain any new data) plus the separate investigations of W~Crv \\citep{RL2000}, AW~UMa \\citep{prib2008} and GSC~1387~475 \\citep{rucpri2008}.  This paper -- in addition to the eight systems with the reliable orbits -- summarizes all remaining, unpublished spectroscopic material obtained within the close-binary project. The included material consists of systems with insufficient phase coverage, problematic systems (too faint, too southern for DDO or too complex multiples), and systems found or suspected to be pulsating variables. For most of the binaries of this paper, we are presenting the first spectroscopic observations. The highlights of the current study are: (1)~the discovery of four triple systems GO~Cyg, V563~Lyr, DD~Mon, V353~Peg; (2)~the discovery of an under-luminous, early-type component in V857~Her; (3)~detection of prograde-rotating ripples in BFs of CU~CVn and V752~Mon indicating non-radial pulsations.  Numerous discoveries and reliable solutions of triple and quadruple systems show that the BF de-convolution approach utilizing the SVD method is a powerful and reliable technique which can be applied not only to close binaries but also to other types of spectral line broadening due to geometrical effects. A good example of the power of this technique is our detection of non-radial oscillations in two pulsating variables, CU~CVn and V752~Mon. As far as we know, the non-radial pulsations have been detected only in the variability of line profiles of selected spectral lines \\citep{zima2006} or in cross-correlation (CCF) studies; the latter obviously with a loss of information because the CCF technique reduces the effective spectral resolution. The only time the BF technique has been applied to a $\\delta$~Sct star (EE~Cam: \\citet{breger2007}) indicated a potential for discovery of fine structure in line profiles, probably akin to the ripples seen in CU~CVn and V752~Mon.  With the closure of the DDO observatory, this series has come to an end; only a dozen or so known W~UMa-type eclipsing binaries brighter than about $V = 10$ and potentially accessible from DDO remain unobserved. Our analysis for several stars presented here may feel incomplete and unsatisfactory. We still hope that the material presented will still lead to or will supplement future follow-up observations of several binary stars described here."
    },
    "0812/0812.1255_arXiv.txt": {
        "abstract": "Array-based, direct-sampling radio telescopes have computational and communication requirements unsuited to conventional computer and cluster architectures. Synchronization must be strictly maintained across a large number of parallel data streams, from A/D conversion, through operations such as beamforming, to dataset recording. FPGAs supporting multi-gigabit serial I/O are ideally suited to this application. We describe a recently-constructed radio telescope called ETA having all-sky observing capability for detecting low frequency pulses from transient events such as gamma ray bursts and primordial black hole explosions. Signals from 24 dipole antennas are processed by a tiered arrangement of 28 commercial FPGA boards  and 4 PCs with FPGA-based data acquisition cards, connected with custom I/O adapter boards supporting InfiniBand and LVDS physical links. ETA is designed for unattended operation, allowing configuration and recording to be controlled remotely. ",
        "introduction": "\\label{sec:intro}  The Eight-meter-wavelength Transient Array (ETA) is a new radio telescope designed to observe a variety of postulated but as-yet undetected astrophysical phenomena which are suspected to produce single pulses detectable at relatively long wavelengths. Potential sources for these pulses include the explosion of primordial black holes (PBHs) and prompt emission associated with gamma ray bursts (GRBs). PBHs are postulated to arise from density fluctuations in the early universe rather than the gravitational collapse of a star. If black holes evaporate as suggested by specific combinations of general relativity and quantum mechanics, those with masses below about $10^{14}$~g should be approaching a state of runaway evaporation, terminating in a massive burst of radiation \\cite{hawking-nature}. Also, prior to a possible terminal explosion, a primordial black hole may undergo a topological phase transition producing an observable outburst of energy if there exists a compact, extra spatial dimension in addition to the three known spatial dimensions [\\citeNP{Kavic1a}b].  A number of models have been proposed to explain short-duration GRBs, including the merger of closely-separated compact objects such as neutron stars and/or black holes. The formation of a single black hole would release an immense amount of energy over a few seconds \\cite{grb}. Prompt radio emission from GRBs would help to increase our understanding of these events. Recently, a dispersed pulse of duration $<$ 5 ms was detected during the analysis of archival pulsar survey data \\cite{Lorimer}. The brightness (30~Jy, where 1~Jy = $10^{-26}$~W~m$^{-2}$~Hz$^{-1}$) and singularity of the burst suggest the source was not a rotating neutron star. Although pulses can be quite strong by astronomical standards, they are difficult to detect due to their transient and unpredictable nature, and the narrow field of view provided by existing telescopes.  ETA, in contrast, is designed to provide roughly uniform (albeit very modest) sensitivity over most of the visible sky, all the time. The complete array consists of 12 dual-polarized dipole-like elements (i.e., 24 radio frequency inputs) at the Pisgah Astronomical Research Institute (PARI), located in a rural mountainous region of Western North Carolina ($35^{\\circ}$ $11.98' $ N, $82^{\\circ}$ $52.28'$ W). Each dipole is individually instrumented and digitized. The digital signals are combined to form fixed ``patrol beams'' which cover the sky, and the output of each beam is searched for the unique time-frequency signature expected from short pulses which have been dispersed by the ionized interstellar medium. ETA has the computational resources and communication bandwidth to implement up to eight beamforming datapaths across all antennas, permitting it to operate as eight independent telescopes simultaneously recording data from different regions of the sky. A patrol beam can be quickly reoriented by updating its beamforming coefficients. Additional information about ETA and its science objectives are available at \\cite{eta-website}.  This paper examines the science, engineering and computing aspects of the ETA project, and is organized as follows: Section~\\ref{sec:science} discusses some astrophysical objects of interest. The ETA architecture, with particular emphasis on the digital back-end, is described in Section~\\ref{sec:design}. Section~\\ref{sec:validation} discusses system validation issues such as determining bit error rates and ensuring channel synchronization. The offline steps performed to mitigate radio frequency interference (RFI) and dedisperse pulses are given in Section~\\ref{sec:analysis}. Current status and conclusions are summarized in Sections~\\ref{sec:status} and \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  FPGAs with integrated multi-gigabit serial transceivers are ideal platforms for the signal processing requirements of radio telescope arrays. Synchronous streaming computations may be implemented over many boards without requiring system-wide clock distribution. As in ETA, the physical communication topology can be tailored to the dataflow requirements, simplifying the interfaces and development effort. These characteristics also make FPGAs containing transceivers an essential technology in particle accelerators such as the LHC \\cite{LHC-ATLAS}. In contrast, software processors and their communication protocols are difficult to use in the lower tiers of telescope and accelerator instruments due to variable latencies.  Our goal to obtain results within the first year necessitated the use of commercial-off-the-shelf (COTS) FPGA boards. This approach distinguishes ETA from many other radio telescope arrays using custom PCBs. The design of complex, high speed boards reduces FPGA development time and cost advantages. Although a single board integrating the receiver and beamforming functions is desirable, the effort required to design a custom PCB would be excessive even if no respins were required. The low-cost evaluation boards meet the ETA system specifications remarkably well, and all interfacing is accomplished with standard cables and simple adapter boards containing only connectors and traces.  An additional benefit of COTS is straightforward upgrades to new FPGA families. ETA2 can use a newer generation of FPGA evaluation boards: the ML410 \\cite{ml410} and ML510 \\cite{ml510} contain XC4VFX60 and XC5VFX130T FPGAs, respectively, while retaining the PC ATX form factor and personality module interfaces, and the Stratix-II DSP Development Kit \\cite{dsp-2s60} contains a 2S60 device and preserves the MICTOR interface. ML410 development boards are used in the LHC ATLAS beam conditions monitor \\cite{LHC-ATLAS}. Evaluation boards further extend the low cost and rapid development virtues of FPGAs, with the new boards costing roughly the same as their predecessors despite having greater logic capacity, higher performance, and expanded interfaces such as PCI Express and Serial ATA. Common platforms such as BEE2 \\cite{bee2-rcs} facilitate integration and IP reuse on large scale instruments developed by geographically separate teams. Even if custom PCBs are justified, evaluation boards may still serve useful prototyping and validation roles.  We are developing an ML510 adapter board with eight Serial ATA connectors in order to have the FPGA write data directly to disks in a raw and interleaved manner. This should improve sustained data recording rates by removing the variable latencies arising from a PC's operating system, file system, interrupts, multiplexed processors, caches and buses. System cost, complexity and power are also reduced by having banks of inexpensive SATA disks replace server-class PCs containing data acquisition cards, SCSI disks and controllers. After an acquisition, data can be retrieved through the ML510's Gigabit Ethernet interface. Similar to the ML310-based cluster described in \\cite{RCC-project}, inter-node communication can use the SATA ports rather than ETA's more expensive InfiniBand cables and connectors. For streaming architectures, FPGAs with high-speed serial transceivers are often both necessary and sufficient for all computation, communication and data collection functions.  \\begin{acks} This work was supported by the National Science Foundation under  Grant No.\\ AST-0504677, a SCHEV ETF grant, and by the Pisgah Astronomical Research Institute. PARI personnel constructed the antenna masts and underground coaxial cable system. Graduate student Vivek Venugopal performed data streaming tests on the Aurora links and data acquisition PCs. We are grateful for early access to the ML310 board and the assistance of Vince Eck, Rick LeBlanc, Punit Kalra and Saeid Mousavi in Xilinx's Systems Engineering Group. \\end{acks}"
    },
    "0812/0812.1580_arXiv.txt": {
        "abstract": "We present the largest publicly available catalogue of compact groups of galaxies identified using the original selection criteria of Hickson, selected from the Sixth Data Release (DR6) of the Sloan Digital Sky Survey (SDSS). We identify 2297 compact groups down to a limiting magnitude of $r = 18$ ($\\sim0.24$\\,groups\\,degree$^{-2}$), and 74791 compact groups down to a limiting magnitude of $r = 21$ ($\\sim6.7$\\,groups\\,degree$^{-2}$). This represents $0.9\\%$ of all galaxies in the SDSS~DR6 at these magnitude levels. Contamination due to gross photometric errors has been removed from the bright sample of groups, and we estimate it is present in the large sample at the $14\\%$ level. Spectroscopic information is available for 4131 galaxies in the bright catalogue (43\\,\\% completeness), and we find that the median redshift of these groups is ${z_{med}} = 0.09$. The median line-of-sight velocity dispersion within the compact groups from the bright catalogue is $\\sigma_{LOS} \\simeq 230$\\,\\kms and their typical inter-galactic separations are of order $50 - 100$\\,kpc. We show that the fraction of groups with interloping galaxies identified as members is in good agreement with the predictions from our previous study of a mock galaxy catalogue, and we demonstrate how to select compact groups such that the interloper fraction is well defined and minimized. This observational dataset is ideal for large statistical studies of compact groups, the role of environment on galaxy evolution, and the effect of galaxy interactions in determining galaxy morphology. ",
        "introduction": "Galaxies cluster on all scales, and a basic tenet of the $\\Lambda-$CDM cosmological paradigm is that structure formation is hierarchical. Consequently, both weak and strong galaxy interactions are expected to be part of the typical life-cycle of a galaxy, and are believed to be among the main drivers in determining a galaxy's star formation history, chemical evolution and morphology.  The majority of $L_\\star$ galaxies at $z = 0$ are not found in rich environments and, while many are members of loose groups with a low velocity dispersion, few have nearby large companions with which they are interacting.  In contrast, galaxy clusters have a high galaxy number density. However, the considerable mass of these structures means that their velocity dispersion is of order $1000$\\kms, and thus individual galaxy-galaxy mergers are rare. On the other hand, compact groups (CGs) of galaxies generally have the low velocity dispersions typical of galaxy groups, but with galaxy-galaxy separations small enough for interactions to be significant (e.g., \\citealt{hickson1997} and references therein).  Compact groups of galaxies were first quantitatively defined by \\cite{hickson1982}, and there now exist several catalogues of compact groups identified using a variety of criteria (e.g., the Digitized POSS CG catalog, \\citealt{iovino2003}; the Southern CG catalogue, \\citealt{prandoni1994,iovino2002}; CGs in the UZC galaxy catalogue, \\citealt{focardi2002}; the CfA2 redshift survey, \\citealt{barton1996}; the Las Campanas redshift survey, \\citealt{allam2000}; the SDSS Commissioning Data, \\citealt{lee2004}). The identification of genuine groups with only a few members is a difficult problem, and considerable effort has been spent on determining the three dimensional properties of the identified systems (particularly the Hickson Compact Groups) and the effect of interlopers (e.g., \\citealt{mamon1986,hickson1988b,walke1989,hernquist1995,ponman1996}).  In this series of papers, we are conducting a homogeneous and systematic study of compact groups using a combination of theoretical and observational galaxy catalogues to provide a statistically-powerful platform from which to probe the effect of galaxy interactions on all aspects of galaxy evolution. In the first paper of this series (\\citealt{mcconnachie2008a}, hereafter Paper~I), we identified compact groups in a mock catalogue constructed from the Millennium Simulation (\\citealt{springel2005,delucia2007}) and developed a robust understanding of their spatial and dynamical properties. This allowed us to refine the compact group selection criteria to more efficiently identify these systems and understand the sources of contamination in any observed sample. In the second paper of this series (\\citealt{brasseur2008}, hereafter Paper~II), we compared the physical properties of compact group galaxies in our mock catalogue to previous observational studies. We demonstrated consistency with these earlier studies, and we concluded that interloping galaxies misidentified as compact group members have potentially affected our understanding of the influence of interactions on galaxy evolution and morphology.  In this paper, the third in the series, we create the largest observational catalogue of compact group galaxies to date by identifying all such groups in the Sloan Digital Sky Survey Data Release 6, consisting of more than 29\\,million galaxies to a limiting magnitude of $r = 21$. We make the resulting catalogue of compact groups publicly available through online tables associated with this paper. This catalogue is ideal for large statistical studies of compact groups, the role of environment on galaxy evolution, and the effect of galaxy interactions in determining galaxy morphology. Section 2 details our procedure for the identification of compact groups, Section 3 presents some of the basic observable properties of these systems, and Section 4 summarises our results. ",
        "conclusions": "In this paper we have presented a publically available catalogue of compact groups of galaxies identified in the SDSS~DR6 (\\citealt{adelmanmccarthy2008}) using the original Hickson criteria (\\citealt{hickson1982}).  We identify 2297 (74791) compact groups down to a limiting magnitude of $r = 18$ ($21$). 0.9\\% of all galaxies at both magnitude limits are identified as members of compact groups, although once interlopers are accounted for this fraction will decrease.  Figures~\\ref{top12} and \\ref{interest} show a selection of compact groups from our catalogue and illustrates the diversity of the groups in this sample. Many galaxies in these groups appear to be early-type, in agreement with previous studies (e.g. \\citealt{hickson1988a,palumbo1995}, Paper~II), although there are many examples of groups with late-type members (many of them visually spectacular). A later paper in this series will investigate the detailed morphology of galaxies in these compact groups.  Spectroscopic information is available for 43\\,\\% (5\\,\\%) of compact group galaxies to a limiting magnitude of $r = 18$ (21); we find that the median redshift of the compact groups identified in Catalogue A is ${z_{med}} = 0.09$. The median line-of-sight velocity dispersions within groups from Catalogue A are $\\sigma_{LOS} \\simeq 220 - 250$\\,\\kms. and the typical inter-galactic separations are of order $50 - 100$\\,kpc. The fraction of groups with interloping galaxies as members is significant, and is shown to be in good agreement with the predictions from the mock galaxy catalogue from Paper~I, despite the latter over-predicting the number density of compact groups by $\\sim 50\\%$. We empirically show that the selection of groups by group surface brightness can reduce the interloper fraction significantly; this will be a powerful tool for future studies of these catalogues.  The catalogue of compact groups derived in this paper is publically available: Tables~1 -- 4 contain basic information for all the groups and member galaxies we identify, including the identifiers necessary to extract full information on each galaxy from the main SDSS~DR6 database. Full versions of these tables are available for download from the MNRAS website."
    },
    "0812/0812.0968_arXiv.txt": {
        "abstract": "{The r-process constitutes one of the major challenges in nuclear astrophysics. Its astrophysical site has not yet been identified but there is observational evidence suggesting that at least two possible sites should contribute to the solar system abundance of r-process elements and that the r-process responsible for the production of elements heavier than $Z=56$ operates quite robustly producing always the same relative abundances. From the nuclear-physics point of view the r-process requires the knowledge of a large number of reaction rates involving exotic nuclei. These include neutron capture rates, beta-decays and fission rates, the latter for the heavier nuclei produced in the r-process. We have developed for the first time a complete database of reaction rates that in addition to neutron-capture rates and beta-decay half-lives includes all possible reactions that can induce fission (neutron-capture, beta-decay and spontaneous fission) and the corresponding fission yields.  In addition, we have implemented these reaction rates in a fully implicit reaction network.  We have performed r-process calculations for the neutrino-driven wind scenario to explore whether or not fission can contribute to provide a robust r-process pattern.}  \\FullConference{10th Symposium on Nuclei in the Cosmos\\\\ July 27 - August 1 2008\\\\ Mackinac Island, Michigan, USA}  \\begin{document} ",
        "introduction": "The r-process is a series of rapid neutron-capture reactions and beta-decays in explosive scenarios with high neutron densities. It is responsible for the synthesis of at least half of the elements heavier than Fe.  Its astrophysical site has not yet been identified, but there is observational evidence suggesting the contribution of at least two sites, namely core collapse supernovae from stars of different masses and neutron-star mergers.  Observed abundances of r-process elements with $Z>56$ in metal-poor stars follow a very robust pattern, that is consistent with scaled solar abundances\\cite{metalpoor}.  Our aim is to explore whether or not fission cycling can provide an explanation for the robustness. ",
        "conclusions": "Using a full set of reaction rates that include all possible fission channels, supplemented with the corresponding fission yield distribution, we have explored the influence of fission in the r-process. Our calculations show that neutron induced fission is the dominating channel. Similar results were obtained in ref. \\cite{panov} in studies of r-process in neutron star mergers. The importance of fission has been quantified computing the number of fission cycles for each r-process calculation. This shows that already at relatively low initial neutron-to-seed ratios a noticeable part of the final abundance is due to fission events. Moreover, calculations with different neutron-to-seed ratios give very similar final abundance distributions for the mass model used in this work."
    },
    "0812/0812.0473_arXiv.txt": {
        "abstract": "{ We consider the fuelling of the central massive black hole in Active Galactic Nuclei, through an inhomogeneous accretion flow. Performing simple analytical treatments, we show that shocks between elements (clumps) forming the accretion flow may account for the UV and X-ray emission in AGNs. In this picture, a cascade of shocks is expected, where optically thick shocks give rise to optical/UV emission, while optically thin shocks give rise to X-ray emission. The resulting blue bump temperature is found to be quite similar in different AGNs. We obtain that the ratio of X-ray luminosity to UV luminosity is smaller than unity, and that this ratio is smaller in massive objects compared to less massive sources. This is in agreement with the observed $L_{X}/L_{UV}$ ratio and suggests a possible interpretation of the $\\alpha_{OX}-l_{UV}$ anticorrelation. } ",
        "introduction": "According to the standard paradigm of Active Galactic Nuclei (AGN) physics, black hole fuelling occurs steadily via a geometrically thin, optically thick, accretion disc. The outward transport of angular momentum is set by viscosity, generally attributed to magnetic fields and turbulent motions, characterized by the $\\alpha$ parameter of the accretion disc (\\citet{S_S_1973}). The thermal emission from the disc, peaking in the optical/UV domain, has been associated with the blue bump component (\\citet{S_1978}, \\citet{M_1983}). \\\\  Much less is known on the kinematics of the accreting gas at large distances, where most (up to $99.9\\%$) of the angular momentum must be lost for matter to reach the $r \\lesssim 100 \\, R_{S}$ region where the bulk of the gravitational energy is released as radiation. In addition, standard accretion disc models are known to face several problems when confronted with detailed observations of AGNs (\\citet{C_2001}, \\citet{S_et_2008}). One of the major problems is given by the quasi-simultaneity of the UV and optical continuum variations that is not compatible with viscous time scales within discs characterized by temperature gradients (\\citet{C_C_1991}, \\citet{C_1991}). Another difficulty is the observed similarity in UV spectral properties over several orders of magnitude in luminosity, and in particular the stability of the blue bump temperature (\\citet{W_et_1994}). Moreover, a `bare' thin disc cannot account for the X-ray component of AGN emission, and therefore additional elements such as optically thin coronae, surrounding the disc and producing X-ray photons through inverse Compton processes, are required.  \\\\   One useful tool to study the accretion flow and radiative processes is provided by the relationship between the UV and X-ray emission components. In the majority of objects, observations indicate a small ratio of X-ray luminosity to blue bump luminosity. In particular, the relative importance of the hard X-ray component is found to be weaker in high luminosity AGNs (QSOs) compared to less luminous Seyfert galaxies (\\citet{K_B_1999}). This relative strength is generally quantified by the optical/UV to X-ray index $\\alpha_{OX}$. It is now observationally confirmed that the slope of the $L_{X}(\\mathrm{2 keV}) - L_{UV}(\\mathrm{2500 \\AA})$ relation is smaller than unity, indicating that the X-ray luminosity is smaller than the UV luminosity, and leading to the well-known $\\alpha_{OX}-l_{UV}$ anticorrelation. (\\citet{S_et_2005}, \\citet{S_et_2006}). This anticorrelation implies that luminous AGNs emit less energy in the X-rays relative to the optical/UV compared to less luminous objects (\\citet{J_et_2007}, \\citet{K_et_2008}). It seems that there is currently no satisfying theoretical study able to account for the small ratio of X-ray to UV luminosity, and to predict the observed $\\alpha_{OX}-l_{UV}$ anticorrelation.  \\\\  In contrast to binary systems where the angular mometum of the accreting material is constrained by the binary geometry, accretion in AGNs proceeds in a less regular and more chaotic manner, with matter arriving from a wide range of directions. This leads to a distribution of angular momentum in the accretion flow that leads to a wide distribution of angular velocities and complex phenomenology that in turn induces numerous shocks within the accretion flow. In view of the interest of widening the study of accretion flows and of difficulties met by disc models, we study the expected emission from the shocked material and show that several fundamental observations of AGN phenomenology can be well described by generic properties of the shocked material. In a previous paper, we considered an inhomogeneous flow in the form of interacting clumps of matter (\\citet{C_T_2005}, hereafter C\\&T 2005). We showed that shocks between clumps and the subsequent gas expansion are at the origin of the radiation and, at the same time, provide a physical mechanism of angular momentum transport. The survival of clumps in the deep gravitational well of the central black hole was also briefly discussed. \\\\  The present paper is organized as follows. The relevant features of the cascades of shocks model are briefly recalled in Sect. 2. We then further develop the model and estimate the UV luminosity arising from the optically thick shock, which is used to determine the blue bump temperature in terms of collision parameters (Sect. 3). We then calculate the X-ray luminosities resulting from the optically thin shocks, taking into account the filling factor of the expanding clouds and considering characteristic time scales, which define different classes of objects (Sect. 4). We compute the luminosity ratio $L_{X}/L_{UV}$ in Sect. 5, and try to identify the different cases with different AGN classes (Sect. 6). We compare the model $L_{X}/L_{UV}$ ratio, as given by the relative strength of optically thin shocks compared to optically thick shocks, with observational measurements of the $\\alpha_{OX}$ index (Sect. 7) and discuss resulting implications in Sect. 8. ",
        "conclusions": "A thorough study of the relationship between different emission components is a required step towards a theoretical understanding of energy generation mechanisms in AGNs. It is hard to explain all the observed AGN properties within standard accretion disc models, the main difficulties including: the similarity in spectral features in the UV and X-ray domains observed in sources with huge differences in central luminosity, the origin of the X-ray emission and its relative importance compared to UV emission. \\\\  In the optical/UV domain, the spectral shape of the blue bump component is very similar in objects varying by 6 orders of magnitude in luminosity (\\citet{W_F_1993}). In particular, the cut-off temperature of the bump is stable and does not vary by more than a factor of two in sources changing by a factor of $10^{4}$ in luminosity (\\citet{W_et_1994}). However, standard accretion disc models do not predict such similarity in spectral features. The relatively universal value of the blue bump temperature on source luminosity is difficult to explain in the framework of standard discs in which the optical/UV emission explicitly depends on the black hole mass and accretion rate.  On the contrary, in the picture presented here, the optical/UV emission arises from optically thick shocks and the resulting blue bump temperature is independent of the black hole mass and only weakly dependent on collision parameters. Optically thick shocks would then naturally account for the observed similarity of the blue bump temperature in objects of very different luminosities, hence different central masses. Our simple model may therefore explain the relatively universal value of the blue bump temperature in objects varying by several orders of magnitude in luminosity.   \\\\  In the X-ray domain, standard accretion discs cannot correctly account for the AGN emission. More complex models with additional components are therefore required, such as irradiated discs and disc-corona models. In the irradiated disc model, the disc emits as a result of both internal viscous heating and external radiative heating due to an X-ray point-like source located above it. But the origin and the location of the X-ray source irradiating the disc is a priori arbitrary and not physically justified. Moreover, in order to reproduce variations of similar amplitude in the UV and X-rays, the X-ray luminosity $L_{X}$ should be of the same order of the UV luminosity $L_{UV}$, which is contrary to  observations.  A more physically plausible picture is given by the disc-corona model (\\citet{H_M_1993}) in which the accretion disc is surrounded by a hot corona: the corona emits X-rays by Compton upscattering of the soft UV photons from the disc, while the disc reprocesses the X-ray photons from the corona into UV photons. But this model gives a larger $L_{X}/L_{UV}$ ratio than the observed value, as the corona and the disc luminosities are assumed to be of the same order. This led \\citet{H_M_G_1994} to propose a variant consisting in a `patchy' corona, which leads to a decrease in the X/UV ratio.  An alternative model is given by the cloud model (\\citet{C_et_1996}, \\citet{C_D_1998}, \\citet{A_C_2000}) in which a central X-ray source is surrounded by a number of Compton thick clouds in quasi-spherical geometry with a large coverage factor; this medium emits the blue bump and reprocesses the X-rays. Contrary to disc models, this latter model predicts a $L_{X}/L_{UV}$ ratio smaller than unity without any ad-hoc hypothesis due to the large coverage factor of the clouds. \\\\  In our picture, optically thick shocks give rise to optical/UV radiation, while the optically thin shocks are at the origin of the X-ray emission. The production of X-rays does not require any additional component, since X-rays are emitted as a consequence of the Compton cooling process of electrons heated in the optically thin shocks. The volume filling factor of the post-shock configuration plays an important role in determining the location of the optically thin shocks, and in defining two classes of objects, distinguished by the central mass and hence central luminosity, that we have identified with quasars and Seyfert galaxies. The competition between cooling and dynamical time scales suggests that there are two additional sub-classes for a given central mass, divided into high accretion rate and low accretion rate objects.  Computing the ratio of X-ray luminosity to UV luminosity, $L_{X}/L_{UV}$, we obtain that this ratio is always smaller than unity, in agreement with the small X-ray to UV ratio observed in the majority of objects. There are only few objects with $L_{X}/L_{UV}$ ratios exceeding unity in the total sample, and the majority of sources have luminosity ratios falling within the predicted range ($0.01 \\lesssim L_{X}/L_{UV} \\lesssim 0.8$). Our model is thus able to predict the observed range of $L_{X}/L_{UV}$ ratios, or equivalently the range of $\\alpha_{OX}$ indices. The observed $\\alpha_{OX}-M_{BH}$ correlation in the sample of \\citet{K_et_2008} implies that high mass objects have smaller $L_{X}/L_{UV}$ ratios, while low mass objects have larger $L_{X}/L_{UV}$ values. We also obtain here smaller $L_{X}/L_{UV}$ ratios in Class Q objects and larger $L_{X}/L_{UV}$ ratios in Class S objects, the predicted trend is thus consistent with the observational relation. Our model may therefore suggest a possible explanation for the observed $\\alpha_{OX}-l_{UV}$ anticorrelation. \\\\   Peculiar values of the $L_{X}/L_{UV}$ ratio, may be attributed to different factors such as additional X-ray emission from a jet component in radio-loud AGNs or absorption in BAL QSOs (considering residual sources not correctly removed in the sample selection). We should note that the observed $L_{X}/L_{UV}$ ratio in 3C 273 is one order of magnitude higher than the predicted value for massive objects. This discrepancy could be explained by an additional X-ray emission, probably associated with the jet component.\\\\   The existence of bulk relativisitc outflows in a preferred direction (collimated jets), cannot be accounted for within the proposed framework, as we have no privileged direction. However, the fraction of matter ejected outward following the expansion of the optically thick shock, may be associated with matter outflows observed in a number of AGNs. In order to obtain an outflow, the wind velocity should reach the local escape velocity and thus the launch radius should lie close to the escape radius. In our picture, the gas expansion velocity is given by the initial free-fall velocity (on the order of $\\sim$ 0.1c at $100 \\, \\mathrm{R_{S}}$), which value coincides with the local escape velocity. In addition, the mass outflow rate is comparable to the mass accretion rate, as indicated by observations. In the case of the QSO PG 1211+143, the outflow velocity is in the range 0.13-0.15c at an escape radius of $130 \\, \\mathrm{R_{S}}$ and with a mass outflow rate of $\\dot{M}_{out} \\sim 3.5 M_{\\odot}/\\mathrm{yr}$ (\\citet{P_P_2006}). The location of the phenomenon and the outflow speed are compatible with shocks occuring at $\\sim 100 \\, \\mathrm{R_{S}}$ and outflowing at the expansion velocity. The expanding matter will be later slowed down in the outer regions, and may give rise to the line emitting clouds of the Broad Line Region."
    },
    "0812/0812.2640_arXiv.txt": {
        "abstract": "\\baselineskip 11pt Holography is 3D imaging which can record intensity and phase at the same time. The importance of construct hologram is holographic recording and wavefront reconstruction. It is surprised that holography be discovered in study interstellar scintillation for pulsar provide a coherent light source recently. I think that is speckle hologram and speckle interference(i.e. intensity interference), and use modern technique which include phased array,CCD, digital signal processing and supercomputer can achieve that digital and computer holography from radio to X-ray astronomy. This means we can use it to image the universe and beyond the limited of telescope for cosmos provide much coherent light from pulsar,maser, black hole to 21cm recombination line. It gives a probe to the medium of near the black hole et al. From those coherent light sources in the sky, we can uncover one different universe that through astronomical quantum observation which use intensity interference. ",
        "introduction": "\\label{sec:intro}  After the invention of the telescope in 1608 in the Netherlands, by Hans Lippershey, Zacharias Janssen and Jacob Metius, it was the refinement of the telescope the following year by Galileo that triggered one of the greatest scientific revolutions of all time. Now astronomical observation has extend to all waveband of electromagnetic wave from radio to $\\gamma$-ray, cosmic rays, neutrino and gravitational wave. In usually,astronomer just measures celestial signal in the spatial or temporal that involves imaging, spectroscopy, polarimetry and variation in time.  The hologram is invented by Gabor,D in 1947, which base Bragg's wavefronts reconstruction in microscopy and Zernike's phase contrast. It originates in classic wave optic and fourier optic but thrives for maser and laser development that provide a excellent coherent referenced light. Now holography has been widely applied in many areas which include 3D displays, optical memories, diffuser display screen, nodestructive measurement and testing et al.\\cite{lgr02}.  It is surprised that holography be discovered pulsar astronomy recently\\cite{wkss08}. I think it is speckle hologram for pulsar provide a coherent light source. Speckle is a random intensity pattern produced by the mutual interference of a set of wavefronts. This phenomenon has been investigated by scientists since the time of Newton to explain that twinkle stars, but speckles have come into prominence since the invention of the laser and have now found a variety of applications which include optic coherence tomography,optical projection display, projection microlithography, laser radar(LADAR), Imaging Coherent Optical Radar (ICHOR) and Infrared Coherent infrared imaging camera (CIRIC) et al\\cite{gjw06}. The speckle effect as a result of the interference of many waves, having different phases, which add together to give a resultant wave whose amplitude and intensity varies randomly. It was considered to be a severe drawback in using lasers to illuminate objects, particularly in holographic imaging because of the grainy image produced. It was later realized that speckle patterns could carry information about the object's surface deformations, and this effect is exploited in holographic interferometry that a technique enables static and dynamic displacements of objects with optically rough surfaces to be measured to optical interferometric precision. These measurements can be applied to stress, strain and vibration analysis, as well as to non-destructive testing. It can also be used to detect optical path length variations in transparent media, which enables, for example, fluid flow to be visualised and analysised. And in astronomy, it widely is studied in stellar speckle of optical\\cite{gjw06,lgr02}. ",
        "conclusions": ""
    },
    "0812/0812.1828_arXiv.txt": {
        "abstract": "The hard X-ray emission of active galactic nuclei (AGN) is believed to originate from the hot coronae above the cold accretion discs. The hard X-ray spectral index is found to be correlated with the Eddington ratio $L_{\\rm bol}/L_{\\rm Edd}$, and the hard X-ray bolometric correction factor $L_{\\rm bol}/L_{\\rm X,2-10keV}$ increases with the Eddington ratio. {The Compton reflection is also found to be correlated with the hard X-ray spectral index for Seyfert galaxies and X-ray binaries.} These observational features provide very useful constraints on the accretion disc-corona model for AGN. We construct an accretion disc-corona model and calculate the spectra with different magnetic stress tensors in the cold discs, in which the corona is assumed to be heated by the reconnection of the magnetic fields generated by buoyancy instability in the cold accretion disc. Our calculations show that the magnetic stress tensor $\\tau_{r\\varphi}=\\alpha p_{\\rm gas}$ fails to explain all these observational features, while the disc-corona model with $\\tau_{r\\varphi}=\\alpha p_{\\rm tot}$ always leads to constant $L_{\\rm bol}/L_{\\rm X,2-10keV}$ independent of the Eddington ratio. The resulted spectra of the disc-corona systems with $\\tau_{r\\varphi}=\\alpha \\sqrt{p_{\\rm gas}p_{\\rm tot}}$ show that both the hard X-ray spectral index and the hard X-ray bolometric correction factor $L_{\\rm bol}/L_{\\rm X,2-10keV}$ increase with the Eddington ratio, which are qualitatively consistent with the observations. We find that the disc-corona model is unable to reproduce the observed very hard X-ray continuum emission from the sources accreting at low rates (e.g., $\\Gamma\\sim 1$ for $L_{\\rm bol}/L_{\\rm Edd}\\sim 0.01$), which may imply the different accretion mode in these low luminosity sources. We suggest that the disc-corona system transits to an advection-dominated accretion flow$+$disc corona system at low accretion rates, which may be able to explain all the above-mentioned correlations. ",
        "introduction": "It is believed that active galactic nuclei (AGN) are powered by accretion of matter on to massive black holes, and the observed UV/optical emission of AGN is thought to be a thermal emission from the standard geometrically thin, optically thick accretion discs \\citep*[e.g.,][]{s78,ms82,sm89}. On the other hand, the power-law hard X-ray spectra of AGN are most likely due to the inverse-Compton scattering of soft photons on a population of hot electrons \\citep{grv79,hm91,hm93}. In the accretion disc-corona model, such soft photons are from the cold disc, a fraction of which are Compton scattered by the hot electrons in the corona above the cold disc to the hard X-ray energy band. The disc-corona model was extensively explored in many previous works \\citep*[e.g.,][]{sz94,hm91,hm93,ksm01,lms02,lmo03}. In this disc-corona scenario, most gravitational energy is generated in the cold disc, probably through turbulence produced by the magneto-rotational instability \\citep{bh91}. The magnetic fields generated in the cold disc are strongly buoyant, and a substantial fraction of magnetic energy is transported vertically to heat the corona above the disc with the reconnection of the fields \\citep*[e.g.,][]{d98,dcf99,mf01,mf02}. It was found that the temperature of the hot electrons in the corona is roughly around $10^9$~K, which can successfully reproduce a power-law hard X-ray spectrum as observed \\citep*[e.g.,][]{lmo03}.    The physical processes of turbulence triggered by the amplified magnetic fields in the disc are very complicated, and quite unclear, though they are revealed to some extent with numerical magneto-hydrodynamical (MHD) simulations \\citep{bh91,bh98}. The so-called ``$\\alpha$-prescription\" is widely adopted in most of the works on accretion discs \\citep{ss73}, in which the magnetic stress tensor is assumed to be proportional to the total pressure ($p_{\\rm tot}=p_{\\rm gas}+p_{\\rm rad}$), gas pressure $p_{\\rm gas}$, or $\\sqrt{p_{\\rm gas}p_{\\rm tot}}$ \\citep{sc81,sr84,tl84}. The magnetic stress tensor $\\tau_{r\\varphi}=\\alpha p_{\\rm tot}$ was suggested in \\citet{ss73}, however, the disc is thermal unstable if the radiation pressure dominates over the gas pressure \\citep{ss76}\\citep*[but also see][]{hkb08}. It was argued that the magnetic field amplification is likely to be limited to that the magnetic pressure is less than the gas pressure even in the inner radiation-pressure-dominated regions, i.e., $\\tau_{r\\varphi}=\\alpha{p}_{\\rm gas}$ \\citep{sc81}. The disc with this stress tensor is thermal stable \\citep{sc81}. The stress tensor $\\tau_{r\\varphi}=\\alpha \\sqrt{p_{\\rm gas}p_{\\rm tot}}$ was initially suggested by \\citet{tl84} based on the idea that the viscosity is proportional to the gas pressure, but the size of turbulence should be limited by the disc scaleheight that is given in terms of the total pressure. This is also supported by the analysis on the local dynamical instabilities in magnetized, radiation-pressure-supported accretion discs \\citep{bs01}, which also leads to $p_{\\rm m}=B^2/8\\pi\\simeq \\beta_0 \\sqrt{p_{\\rm gas}p_{\\rm tot}}$, where $p_{\\rm m}$ is the magnetic pressure, and $\\beta_0$ is a constant of the order of unity \\citep*[see][for a detailed discussion]{mf02}. \\citet{nrm00} argued that the accretion disc model with $\\tau_{r\\varphi}=\\alpha p_{\\rm gas}$ is too stable to explain the unstable behaviour observed in GRS~1915+105, and they further suggested that it is possible to couple the radiation to the particles through collisions and thereby allow the radiation pressure to contribute to the stress tensor to some extent. As the complexity of the physics in the radiation-pressure-dominated fluids, any one of these stress tensors should only be regarded as a possible option in accretion discs.  The hard X-ray observations on AGN may provide useful clues on the accretion disc-corona models. \\citet{wwm04} compiled a sample of radio-quite AGN, and found a strong correlation between $L_{\\rm 2-10keV}/L_{\\rm bol}$ and $L_{\\rm bol}/L_{\\rm Edd}$, which was conformed by \\citet{vf07} with a different AGN sample. This correlation was then used to constrain the disc-corona models with different magnetic stress tensors, and they found that the model with magnetic stress tenor $\\tau_{r \\varphi}=\\alpha p_{\\rm gas}$ is favored by the correlation of $L_{\\rm 2-10keV}/L_{\\rm bol}-L_{\\rm bol}/L_{\\rm Edd}$. It was also found that the spectral index of hard X-ray continuum emission is correlated with the Eddington ratio \\citep{ly99,wwm04,s06,s08}. {\\citet{zls99} found a strong correlation between the Compton reflection and the hard X-ray spectral index for Seyfert galaxies and X-ray binaries.} In the accretion disc-corona scenario, the hard X-ray emission originates from the Compton scattering of the soft photons by the hot electrons in the corona. Therefore, the correlations of the hard X-ray spectral index with the Eddington ratio/Compton reflection, together with the correlation of $L_{\\rm 2-10keV}/L_{\\rm bol}-L_{\\rm bol}/L_{\\rm Edd}$, provide important constraints on the accretion disc-corona model.  {The parallel plane homogeneous corona is unable to produce an X-ray spectrum with $\\Gamma<2$ in 2--10~keV. The X-ray photons radiated from the corona are reprocessed in the cold disc, and the reprocessed photons irradiate the corona, which cool the corona and lead to rather soft X-ray spectra with $\\Gamma\\ga 2$. In the patchy corona model proposed by \\citet{hmg94}, the corona appears as individual blobs above the cold disc. Most of the reprocessed photons do not enter the blob, so the cooling is significantly reduced and the blobs are hotter than the parallel plane homogeneous corona, which leads to harder X-ray spectra. This model can explain the observed hard X-ray spectra with $\\Gamma<2$ in some AGN \\citep*[e.g.,][]{z96}, however, it is still unable to explain the correlation between the Compton reflection and the hard X-ray spectral index for Seyfert galaxies and X-ray binaries \\citep{zls99}, which is also the case for the parallel plane homogeneous corona model. \\citet{zls99} suggested that a central hot plasma surrounded by a cold disc may explain the correlation of the Compton reflection with the hard X-ray spectral index. In this scenario, the cold disk is truncated at a certain radius $d$, within which the hot plasma may possibly correspond to an advection dominated accretion flow (ADAF) \\citep*[e.g.,][]{ny95}. The radiation of the hot plasma is dominated by the inverse Compton scattering of the soft photons from the outer cold disc, while the cold disc is irradiated by the inner hot plasma. Thus, the Compton reflection component decreases with increasing the inner radius of the cold disc, which leads to less soft seed photons from the cold disc entering the hot plasma and then the harder X-ray spectrum from the hot plasma. The correlation between the Compton reflection and the hard X-ray spectral index can be reproduced by this model \\citep*[see][for the details]{zls99}. An alternative model was proposed by \\citet{b99} to explain the correlation between the Compton reflection and the X-ray spectral index. In this model, the hot plasma above the cold disc is assumed to move away from the cold disc at a mild relativistic velocity. Such an outflow reduces the downward flux, and then reduces both the reflection and reprocessing in the cold disc. The reduction of the reprocessing leads to less incident soft seed photons and cooling in the hot plasma above the cold disc, and in turn leads to harder X-ray spectrum \\citep{b99,mbp01}.  }  In this work, we take these different magnetic stress tensors as candidates in our disc-corona model calculations, which could be tested with the X-ray observations of AGN. We summarize the disc-corona model in Sect. 2, and the numerical results of the model calculations are given in Sect. 3. In Sect. 4, we discuss the physical implications of the results. ",
        "conclusions": "It is believed that the power generated in the disc is transported vertically with the buoyancy of the magnetic fields in the disc, and the fraction of the power dissipated in the corona to the total is mainly regulated by the magnetic fields \\citep*[e.g.,][]{d98,mf02,wwm04}. \\citet{vf07} found that the fraction $L_{\\rm cor}/L_{\\rm bol}$ is about 0.5 for the sources with low Eddington ratio $L_{\\rm bol}/L_{\\rm Edd}\\sim 0.01$, while it decreases to 0.1 for $L_{\\rm bol}/L_{\\rm Edd}\\sim 1$, for a sample of AGN. From Fig. \\ref{fig1}, we find that the model with $\\tau_{r\\varphi}=\\alpha p_{\\rm tot}$ always predict a constant $L_{\\rm cor}/L_{\\rm bol}$, which seems to be inconsistent with the observation. The models with other two magnetic stress tensors can roughly reproduce the trend of $L_{\\rm cor}/L_{\\rm Edd}$ decreasing with accretion rate $\\dot{m}$, though the model with $\\tau_{r\\varphi}=\\alpha p_{\\rm gas}$ always underpredicts the $L_{\\rm cor}/L_{\\rm bol}$ at high Eddington ratio end, even if $\\alpha=1$ is adopted. It seems that the model with $\\tau_{r\\varphi}=\\alpha \\sqrt{p_{\\rm gas}p_{\\rm tot}}$ is better than the other two models.   The ratio $L_{\\rm cor}/L_{\\rm bol}$ decreases with accretion rate $\\dot{m}$, which means the fraction of the power radiated from the cold disc increases with $\\dot{m}$. The cooling mechanism of the hot corona is dominated by the inverse Compton scattering of the soft photons from the cold disc by the hot electrons in the corona. There are more soft photons supplied by the cold disc for high-$\\dot{m}$ cases, and the fraction of the power radiated from the corona decreases with $\\dot{m}$. Thus, it can be easily understood that the temperature of the hot electrons in the corona should decrease with accretion rate $\\dot{m}$ (see Fig. \\ref{fig2}). It is found that the temperature of the electrons in the corona decreases to $\\sim 3-4\\times 10^8$~K for the models with $\\tau_{r\\varphi}=\\alpha{p}_{\\rm gas}$ or $\\tau_{r\\varphi}=\\alpha \\sqrt{p_{\\rm gas}p_{\\rm tot}}$ when $\\dot{m}=0.5$, which implies that the hot corona is nearly to be suppressed when $\\dot{m}$ is as high as $\\ga 0.5$.  We calculate the spectra of the accretion disc-corona systems, and find that their hard X-ray continuum emission indeed exhibits a power-law spectrum (see Fig. \\ref{fig3}). In Fig. 4, we plot the photon spectral indices $\\Gamma$ in 2--10~keV and the X-ray correction factors $L_{\\rm bol}/L_{\\rm X,2-10keV}$ as functions of accretion rate $\\dot{m}$ with different magnetic stress tensors. We find that the photon spectral index $\\Gamma$ predicted by the model with $\\tau_{r\\varphi}=\\alpha{p}_{\\rm tot}$ is within $\\sim 2-2.2$ for different values of $\\dot{m}$, and it increases slightly with $\\dot{m}$. The X-ray correction factor $L_{\\rm bol}/L_{\\rm X,2-10keV}$ remains constant with accretion rate $\\dot{m}$ for this model. The photon spectral index $\\Gamma$ increases with $\\dot{m}$ while $\\dot{m}\\la 0.02$, and it then decreases with $\\dot{m}$, for the model with $\\tau_{r\\varphi}=\\alpha{p}_{\\rm gas}$. We find that the bremsstrahlung emission {contributes some to} the hard X-ray energy band when $\\dot{m}$ is high, and the photon spectral index $\\Gamma$ is {slightly affected by the bremsstrahlung emission} (see the middle panel of Fig. 3). We find that the X-ray correction factor $L_{\\rm bol}/L_{\\rm X,2-10keV}$ for this model becomes extremely high for high-$\\dot{m}$ cases, which seems to be inconsistent with that derived from the observations \\citep*[see Fig. 12 in][]{vf07}. The X-ray spectra of the disc-corona systems with $\\tau_{r\\varphi}=\\alpha \\sqrt{p_{\\rm gas}p_{\\rm tot}}$ show that both the photon spectral index $\\Gamma$ and the X-ray correction factor $L_{\\rm bol}/L_{\\rm X,2-10keV}$ increase with accretion rate $\\dot{m}$, which are qualitatively consistent with the observations, except for the low $\\dot{m}$ end. The temperature of the electrons in the corona decreases with $\\dot{m}$ (see Fig. \\ref{fig2}), and the inverse Compton scattered X-ray spectrum therefore becomes softer for a higher $\\dot{m}$. {Similarly, anticorrelations between the X-ray spectral indices and X-ray fluxes were found in some X-ray binaries \\citep*[e.g.,][]{y03,z04,d08}, which seem to be roughly consistent with our present disc corona model. }  The photon spectral index $\\Gamma$ is found to increase with the Eddington ratio, and it could be as low as $\\sim 1$ when $\\dot{m}\\sim 0.01$ \\citep*[see][]{s06,s08}. The hard X-ray photon spectral index $\\Gamma\\sim 2$ derived from our model calculations with any magnetic stress tensor when $\\dot{m}\\sim 0.01$, which are significantly higher than the observed $\\Gamma\\sim 1$ \\citep*[e.g.][]{s06,s08}. {Although the patchy corona model proposed by \\citet{hmg94} can produce the observed hard X-ray spectra with $\\Gamma\\sim 1$, this patchy corona model is unable to explain the correlation between the Compton reflection and the hard X-ray spectral index. In an alternative model, the hot plasma above the cold disc is assumed to move away from the cold disc at a mild relativistic velocity \\citep{b99}, which reduces both the reflection and reprocessing in the cold disc. This naturally leads to harder X-ray spectrum  (see the discussion in Sect. 1). However, the detailed physics for producing such mildly relativistic outflows is still unclear.}   It was argued that the central engines in these low-luminosity sources with very hard X-ray emission may be different from their high-luminosity counterparts, i.e., the advection dominated accretion flows (ADAFs) may be present in these low-luminosity sources \\citep{ly99}. The soft incident photons for Comptonization are mainly due to the synchrotron$+$bremsstrahlung emission in the ADAFs. The energy density of the soft photons in the ADAF is much lower than that in the disc-corona model, which leads to inefficient cooling and relatively higher electron temperature in the ADAFs \\citep*[e.g.,][]{ny95}. Thus, the X-ray spectra of the ADAFs can be much harder than those of the disc-corona systems. The very hard X-ray spectra observed in these low luminosity sources can be well modelled with the ADAFs accreting at low rates \\citep*[see e.g.,][for the spectral modeling for the hard X-ray emission from the low luminosity AGN with ADAFs]{q99,xc08}. Our present model calculations are limited to the disc-corona model, in which the cold discs extend to the marginal stable orbits. The model of a disc-corona connecting with inner ADAF may resolve the photon index discrepancy at the low $\\dot{m}$ end. {The geometry of this ADAF$+$disc/corona scenario is quite similar to the hot plasma$+$cold disc model proposed by \\citet{zls99}. When the accretion rate decreases to a critical value $\\dot{m}_{\\rm crit}$, the inner cold disc may be truncated at radius $R_{\\rm tr}$ and it transits to an ADAF within this radius. The transition radius $R_{\\rm tr}$ may be close to the marginal stable orbits soon after the accretion mode transition \\citep*[e.g.,][]{q99,yn04,xc08}, and then it may increase with decreasing accretion rate $\\dot{m}$ \\citep*[e.g.,][]{l99,rc00,sd02,yn04}. The X-ray spectrum of such an ADAF$+$disc/corona system consists of emission from the inner ADAF and outer corona. The ratio of the X-ray emission from the ADAF to that from the corona increases with decreasing $\\dot{m}$, and therefore the photon spectral index $\\Gamma$ may decrease smoothly with decreasing  $\\dot{m}$ provided the initial truncated radius of the cold disc is not very large compared with  the marginal stable orbits. Similar to \\citet{zls99}'s model, the correlation between the Compton reflection and the hard X-ray spectral index can also be naturally explained by this ADAF$+$disc/corona model, if the truncated radius $R_{\\rm tr}$ increases with decreasing accretion rate $\\dot{m}$. This model can also successfully explain the spectral behaviours of X-ray binaries \\citep*[see][for a review and references therein]{dgk07}. } The detailed calculations on such ADAF$+$corona systems will be reported in our future work.                   In this work, we simply assume $T_{\\rm i}=0.9T_{\\rm vir}$, motivated by the previous work on the disc-corona model calculations \\citep{lms02,lmo03}. {We also check how $T_{\\rm i}$ may affect our results by tuning the value of $T_{\\rm i}$, and find that the X-ray spectra of the disc/corona systems change very little if all other disc parameters are fixed (see Fig. \\ref{fig5}). The cooling of the corona is dominated by the inverse Compton radiation, which is roughly proportional to Compton $y$-parameter. Thus, it is not surprising that our calculations show the Compton $y$-parameter varying little with $T_{\\rm i}$ if all other parameters are fixed (see equation \\ref{cor_energy} and Fig. \\ref{fig6}). The thickness of the corona is mainly regulated by the ion temperature $T_{\\rm i}$, and therefore the electron density decreases with increasing $T_{\\rm i}$, which leads to significant difference in synchrotron/bremsstrahlung spectra for different $T_{\\rm i}$ (see Fig. \\ref{fig5}). } We find that the resulted $\\dot{m}-\\Gamma$ relations are almost not changed if a different value of $T_{\\rm i}$ is adopted. Unlike the ADAFs, almost all the power dissipated in the hot corona with magnetic reconnection is radiated away locally. This means the radiated power in the corona is independent of the value $\\delta$, and the temperature and density of the electrons in the corona are almost insensitive with the value of $\\delta$. We also perform the same model calculations for different black hole masses (e.g., $M_{\\rm bh}=10^9{\\rm M}_\\odot$). It is found that our results are almost independent of $M_{\\rm bh}$ for massive black holes.  In all our calculations, we assume the magnetic fields to be equipartitioned with the gas pressure in the corona. As the cooling of the corona is dominated by the inverse Comptonization of the soft photons from the cold disc, the structure of the disc-corona and its spectrum (except in the radio wavebands) is almost independent of the magnetic field strength in the corona. Recently, \\citet{lb08} found there is a strong correlation between the radio luminosity ($L_{\\rm R}$) and X-ray luminosity ($L_{\\rm X}$) with $L_{\\rm R}\\sim 10^{-5}L_{\\rm X}$, for the radio quiet Palomar-Green (PG) quasar sample. The spectra of our disc-corona model show that the radio emission is correlated with the X-ray emission, which is roughly consistent with the correlation between $L_{\\rm R}$ and $L_{\\rm X}$ discovered by \\citet{lb08}."
    },
    "0812/0812.4585_arXiv.txt": {
        "abstract": "The primordial curvature fluctuation spectrum is reconstructed by the maximum likelihood reconstruction method using the five-year Wilkinson Microwave Anisotropy Probe data of the cosmic microwave background temperature anisotropy. We apply the covariance matrix analysis and decompose the reconstructed spectrum into statistically independent band powers. The prominent peak off a simple power-law spectrum found in our previous analysis turn out to be a $3.3\\sigma$ deviation. From the statistics of primordial spectra reconstructed from mock observations, the probability that a primordial spectrum including such excess is realized in a power-law model is estimated to be about $2$\\%. ",
        "introduction": "The cosmic evolution during the inflationary stage of the early Universe \\cite{inf1,inf2,inf2a,inf3} is recorded in the primordial fluctuation spectrum \\cite{yuragi,yuragi2,yuragi3}, which can be revealed by high-precision observations of the cosmic microwave background (CMB) and large-scale structures. If we found some nontrivial features exceeding a level expected by the cosmic variance after proper analysis of observational errors, we would have to seriously consider its generation mechanism, namely, an inflation model that generates such a nontrivial fluctuation spectrum. This would not only be a challenge to particle-physics model building but also a boon that provides an important clue to high-energy physics. The CMB anisotropy data provided by the Wilkinson Microwave Anisotropy Probe (WMAP) mission \\cite{WMAPBASIC,WMAPTEMP,WMAPCOSMO,WMAPINF,WMAPTTTE,loglike,WMAP3TEMP,WMAP5,WMAP5ONLY,WMAP5COSMO} is already on the precision level sufficient for such purpose.  There have been persistent controversy on the existence of nontrivial features in the CMB temperature anisotropy spectrum. The angular spectrum of the first-year WMAP data \\cite{WMAPBASIC,WMAPTEMP,WMAPCOSMO,WMAPINF,WMAPTTTE} exhibited a sign of a running spectral index, oscillatory behaviors on intermediate scales, and lack of power on large scales, which cannot be explained by a power-law primordial spectrum that is a generic prediction of simplest inflation models. These features may provide clues to unnoticed physics of inflation. Some of these anomalous structures disappeared on the three-year spectrum, but several anomalies are still remaining \\cite{WMAP3TEMP}. To explain these features, a number of inflation models have been proposed \\cite{CPKL03,CCL03,FZ03,KT03,HM04,LMMR04,KYY03,YY03,KS03,YY04,MR04A,MR04C,KTT04,HS04,HS07,CHMSS06}.  To reconstruct the primordial spectrum using CMB anisotropy data without specific prior assumptions about the shape of the primordial spectrum, various observational methods have been proposed \\cite{WSS99,BLWE03,MW03A,MW03B,MW03C,MW03D,SH01,SH04,BLH07,SVJ05,WMAP3COSMO,LSV08,VP08}. Among such attempts, nonparametric methods \\cite{MSY02,MSY03,KMSY04,KSY04,KSY05,NY08,SS03,TDS04,THS05,SS06,SS07}, which can reconstruct the primordial spectrum as a continuous function without any {\\it ad hoc} filtering scale, have a merit for investigating a detailed feature localized in a narrow wave number range. Recently, we reconstructed the primordial spectrum by the cosmic inversion method \\cite{MSY02,MSY03,KMSY04,KSY04,KSY05,NY08} from the temperature anisotropy spectrum of the five-year WMAP data. The reconstructed spectrum accurately preserved the information of detailed features observed on the angular spectrum. To distinguish possible true signal from the cosmic variance appropriately, we applied the covariance matrix analysis and decomposed the reconstructed spectrum into statistically independent band powers. Although the reconstructed spectrum was basically consistent with a simple power-law spectrum, it exhibited anomalous deviation localized on $\\approx 700$Mpc scale, which was unfortunately near the border of the wave number domain where accurate reconstruction was possible by the cosmic inversion method. To investigate the possible fine structure, we need to adopt another reconstruction method which can probe a detailed feature of the primordial spectrum in the corresponding wave number range.  The purpose of this work is to revisit the possibility of fine structure in the primordial spectrum. We apply the nonparametric reconstruction method, which is originally developed by Tocchini-Valentini et al.\\cite{TDS04,THS05}, to the five-year WMAP temperature anisotropy spectrum \\cite{WMAP5} and perform the band-power analysis introduced in our previous work. Because of the arbitrariness of the primordial spectrum, we inevitably incorporate infinite degree of freedom to our analysis, which results in degeneracy among spectral shape and cosmological parameters \\cite{KNN01,SBKET,SS07}. However, the fine structure on which we focus in this paper is much narrower than the structure dependent on the conventional cosmological parameters whose characteristic width is principally the inverse of the horizon length at the recombination epoch. As shown in \\cite{KMSY04}, different choices of cosmological parameters affect only the overall shape and the fine structures of the reconstructed spectrum remain intact. In this paper, we consider the concordance adiabatic $\\Lambda$CDM model, where the cosmological parameters (except for the ones characterizing the primordial spectrum) are those of the WMAP team's best-fit power-law model \\cite{WMAP5ONLY,WMAP5COSMO}, and instead focus on the detailed shape of the primordial spectrum.  This paper is organized as follows: In Sec. II, the overview of our analysis is described. In Sec. III, we show the reconstructed primordial power spectrum from the five-year WMAP data and discuss its implication. Finally, Sec. IV is devoted to the conclusion. ",
        "conclusions": "We have reconstructed the primordial power spectrum of curvature fluctuation by the maximum likelihood reconstruction method, assuming the absence of isocurvature modes and the best-fit values of cosmological parameters for the power-law $\\Lambda$CDM model. In a wide range of scales, where the cosmic variance currently dominates over measurement errors ($30 \\lesssim kd \\lesssim 390$), we can reproduce the fine structures of modulations off a simple power-law with which we can recover the highly oscillatory features observed in $C_\\ell$.  The statistical significance of the oscillatory structure in the reconstructed power spectrum is difficult to quantify due to the strong correlation among the neighboring wave number modes. We have therefore performed the covariance matrix analysis to calculate the window matrix, which diagonalizes the covariance matrix into statistically independent modes.  In regard to the cumulative deviation ($\\chi^2$), the reconstructed spectrum is a common realization of a power-law spectrum. In particular, on the scales of $kd \\gtrsim 130$, the best-fit power-law spectrum for the five-year WMAP data fits the reconstructed spectrum slightly excessively well. We have focused on the structure on the scale of $kd \\simeq 120$, where previously we found the prominent deviation from the power-law prediction. Although the corresponding band exhibits some deviation from a power-law spectrum, we have found that the feature is marginally consistent with a power-law model.  Finally, we briefly comment on reconstruction from the polarization angular spectrum. The WMAP temperature-polarization cross correlation spectrum exhibits remarkable deviation from the power-law prediction around $\\ell \\simeq 120$ \\cite{LAMBDA}. The enhancement of (negative) correlation observed there seems to agree with the anomalous peak found in the reconstructed primordial spectrum in this work, and besides they have apparently similar coherent width. Although the current polarization data is still suffering large measurement errors, the reconstruction from the precise polarization data, which will be available in near future \\cite{PLANCK}, promisingly brings it to a conclusion."
    },
    "0812/0812.0791_arXiv.txt": {
        "abstract": "SDSS J080434.20+510349.2 is the WZ type binary that displayed rare outburst in 2006 (Pavlenko et al., 2007). During the long-lasting tail of the late stage of the outburst binary shown the two-humped or four-humped profile of the orbital light modulation. The amplitude of orbital light curve decreased while the mean brightness decreased, more over that occurred $\\sim$ 10 times faster during the fast outburst decline in respect to the late quiet state of slow outburst fading.  There were no white dwarf pulsations detected neither 1 - 1.5 months prior to the outburst nor in 1.5 - 2 months after the 2006 outburst in this system. However the strong non-radial pulsations with period 12.6 minutes and mean amplitude of $0.05^{m}$ were first detected in V band with 2.6-m Shajn mirror telescope of the Crimean astrophysical observatory in $\\sim$ 8 months after the outburst. The evolution of pulsations over two years in 2006 - 2008 is considered. It is supposed that pulsations first appeared when the cooling white dwarf (after the outburst) entered the instability strip although the possibility of temporary lack of pulsations at some occasions also could not be excluded. ",
        "introduction": "12 cataclysmic variables (CVs) with orbital period close to the evolutional orbital period minimum, containing the pulsating white dwarf (WD), were known up to now (Mukadam et al., this volume). SDSS J080434.20+510349.2 (hereafter SDSS J0804) is the 13th  system showing such pulsations (Pavlenko and Malanushenko, 2009) which are generally believed are due to non radial g-mode pulsations of the white dwarf (Warner and Robinson, 1972).  P. Szkody et al. reported that there are 4 dwarf novae: PQ And, GW Lib, V455 And, REJ 1255+26 containing the accreting WD pulsator that undergone the outburst (is available from www.narit.or.th/conference/prcsa2008/).  SDSS J0804 (Pavlenko et al., 2007) is the 5th such binary.  SDSS J0804 is the WZ Sge type star displayed the outburst in 2006 with 11 rebrightenings. Superhumps that have been observed both during the main outburst and rebrightenings  later on were replaced by the orbital two-humped light variations (Zharikov et al., 2008; Pavlenko, 2007) typical to the WZ-Sge stars in quiescence. The phenomenon of a second series of rebrightenings or \"`mini-outburst\"' with longer recurrence time and less amplitude was found by Zharikov et al (2008). In quiescence before outburst optical spectrum of SDSS J0804 displayed a blue continuum with broad absorption lines from a white dwarf surrounding the double-peaked Balmer emission lines formed in accretion disk (Szkody et al., 2006). Authors also noted the particular behavior of the binary during one night in 2005 quit state - sudden increase of brightness ($0.5^{m}$ in a few minutes) accompanied by increase of amplitude from $0.05^{m}$ to $0.2^{m}$ that never was detected after the outburst (Zharikov et al., 2008).  Despite J0804 was promissing CV for search for the WD non-radial g-mode pulsations, it did not show them before the 2006 outburst  (Szkody et al., 2006). We had pointed attention to the short-term light variations with most significant  12.6 min. oscillation  in 2008, January 1, when its amplitude at some pulses was commensurable with orbital light modulation $\\sim 0.05^{m}$ (Pavlenko and Malanushenko, 2009).  Here the result of searching for the possible pulsations on the available base of our observations both before and after the 2006 outburst  in order to detect their first appearance is presented. ",
        "conclusions": "The orbital light curves of SDSS J0804 having typically two-humped profile, often displayed the splitting of one or both humps. The color-index V-R continued to decrease with time during the long brightness decay indicating to the decrease of the accretion disk contribution to the total light and, so, increase of the white dwarf contributing. The J0804 is localized in Galaxy at the low interstellar extinction, its E(V-R) could be calculated following Schlegel et al. (1998) and roughly is E(V-R) = $\\sim 0.037^{m}$, so corrected value $(V-R)_{0}$ = $0.073^{m}$. That corresponds to the black-body temperature $\\sim$ $15000^{\\circ}$ K with accuracy of a few of thousands K. Taking into account that in $\\sim$ 2 years after the outburst disk also could still contribute to the total light, one could expect that the temperature of the white dwarf  was hotter than $15000^{\\circ}$ K.  The dependence of the orbital amplitude on the mean brightness could be caused by  the decreasing brightness of the bright detail on the accretion disk due to the decrease of the mass transfer rate over the source of outburst that occured $\\sim$ 10 times faster during the fast outburst decline in respect to the late quiet state of slow outburst fading. Such dependence obviously is caused by decreasing initially more enhanced mass transfer rate, where decreasing itself fas $\\sim$ 10 times faster  during the  fast outburst decline (after the end of rebrightenings) than during the late stage of outburst.  The variable and sometimes for-humped orbital modulation points to the complex and variable morfology of accretion disk.  We did not find any significant pulsations in 1 - 1.5 months before the outburst. The first detection of the most stable 12.6-min. pulsations was in $\\sim$ 8 months after the expected start of the outburst. It is not clear whether the lack of pulsations at JD 2454063 was their temporary disappearance or caused by the insufficient data statistics. It is already known that the pulsators in dwarf novae could stop their pulsations by some reason  contrary to the ZZ Ceti stars (Southworth et al., 2008). So the lack of pulsations in SDSS J0804 in two occasions before outburst at JD 2453384 (Szkody et al., 2006) and  at JD 2453856 could not be the argument that before outburst the white dwarf never pulsated. However it is possible suggest that the compressional heating during the outburst and further fast cooling entered the white dwarf in the SDSS J0804 into the instability strip.  It is impossible also to consider confidently the 16.7 min. periodicity at JD 2453856 as the white dwarf pulsations because of the  periodogram is different from these performed in Fig. 7. \\subsection"
    },
    "0812/0812.2088_arXiv.txt": {
        "abstract": "We study the dependence of galaxy properties on the clustercentric radius and the environment attributed to the nearest neighbor galaxy using the Sloan Digital Sky Survey (SDSS) galaxies associated with the Abell galaxy clusters. We find that there exists a characteristic scale where the properties of galaxies suddenly start to depend on the clustercentric radius at fixed neighbor environment. The characteristic scale is $1\\sim 3$ times the cluster virial radius depending on galaxy luminosity. Existence of the characteristic scale means that the local galaxy number density is not directly responsible for the morphology-density relation in clusters because the local density varies smoothly with the clustercentric radius and has no discontinuity in general. What is really working in clusters is the morphology-clustercentric radius-neighbor environment relation, where the neighbor environment means both neighbor morphology and the local mass density attributed to the neighbor. The morphology-density relation appears working only because of the statistical correlation between the nearest neighbor distance and the local galaxy number density. We find strong evidence that the hydrodynamic interactions with nearby early-type galaxies is the main drive to quenching star formation activity of late-type galaxies in clusters. The hot cluster gas seems to play at most a minor role down to one tenth of the cluster virial radius. We also find that the viable mechanisms which can account for the clustercentric radius dependence of the structural and internal kinematics parameters are harassment and interaction of galaxies with the cluster potential. The morphology transformation of the late-type galaxies in clusters seems to have taken place through both galaxy-galaxy hydrodynamic interactions and galaxy-cluster/galaxy-galaxy gravitational interactions. ",
        "introduction": "In galaxy clusters the average morphology of galaxies changes with clustercentric radius or local density. The central region of clusters at the present epoch is dominated by early-type galaxies. This is known as the morphology-radius or morphology-density relation (hereafter MRR and MDR, respectively). Since the local density is on average a monotonically decreasing function of clustercentric radius, they appear to convey us the same information \\citep{hh31,dre80,dre97,treu03,smi05,post05,wei06}. However, there is discussion about which one plays a critical role in determining the galaxy morphology between the clustercentric radius and the local density \\citep{whi93,dml01,goto03,tk06}. There is also a remaining issue what physical parameters of galaxies (morphology, color, star formation rate, or stellar mass) correlate fundamentally with the environment (e.g., \\citealt{cz05,qui06,pog08,ski08}).  A number of physical mechanisms have been proposed to explain this morphology-environment relation. It is suggested that tidal interactions between individual galaxies can change mass profile and transform disk galaxies into spheroidals \\citep{bv90}. Frequency of galaxy-galaxy interaction depends strongly on the clustercentric radius, and thus galaxy interactions can result in the MRR or MDR. However, the duration of tidal interactions and thus the total tidal energy deposit are expected to be too small to change galaxy properties significantly for cluster galaxies that are moving at high speeds. Due to the fast orbital motions the frequency of galaxy mergers is also small in clusters. A series of frequent high-speed tidal interactions among cluster galaxies, called harassment, can produce impulsive heating and cause morphology transformation \\citep{moo96}. Tidal interaction with the whole cluster potential well can also change the structure and activity of galaxies significantly \\citep{mw00,gne03}.  Existence of the hot X-ray emitted by intracluster gas makes the mechanisms relying on hydrodynamic processes appear attractive. Stripping of cold gas from infalling late-type galaxies by a ram pressure of the hot intracluster gas, was proposed to explain the increase of the early-type fraction toward the cluster center \\citep{gg72}. Hot gas removal from infalling galaxies through hydrodynamic interactions with the intracluster medium can shut off gas supply and star formation activity (SFA) after the already existing cold gas is consumed. The interstellar medium of a galaxy traveling in the hot intracluster medium can be stripped from the disk due to a viscosity momentum transfer \\citep{nul82} or evaporate as the temperature of the interstellar medium rises \\citep{cs77}. These hydrodynamic processes have a common drawback that they can turn spirals only to S0's and do not produce ellipticals. The major effect of the hydrodynamic processes is quenching the SFA by removing or ionizing the cold gas in the disk of late-type galaxies. The bulge-to-disk ratio is not expected to increase by these processes. Therefore, none of the mechanisms proposed so far is able to fully account for the MRR or MDR observed in clusters.  Recently, \\citet{park08} and \\citet{pc09} have found that galaxy properties like morphology and luminosity depend strongly on the distance and morphology of the nearest neighbor galaxy. This dependence was found even when the large-scale background density is fixed, and thus is completely different from the commonly known morphology-local density relation. Most importantly, the effects of the nearest neighbor change at the characteristic scale given by the virial radius of the neighbor galaxy. When a galaxy is located within the virial radius of its nearest neighbor, its morphology tends to be the same as that of the neighbor. But such tendency disappears when the separation is larger than the virial radius. This fact strongly suggests that galaxies interact hydrodynamically when the separation to their nearest neighbor is smaller than the virial radius of the neighbor, and that the effects of such interaction are significant enough to change their morphology and SFA.  Outside the virial radius galaxy morphology still depends on the distance to the nearest neighbor, but is suddenly independent of neighbor's morphology. The probability for a galaxy to be an early type monotonically decreases as the separation increases. This strongly supports the idea that the conformity in morphology of galaxy pairs is not primordial but an acquired one through interactions since there should be no reason for the break in morphology conformity to be at the current virial radius of the neighbor if it is initially given. \\citet{park08} proposed that the mechanism responsible for this be tidal interactions causing both gravitational and hydrodynamic effects. Outside the virial radius the interactions are purely gravitational between the galaxy plus dark halo systems, but within the virial radius hydrodynamic (and radiative) effects must be involved as well.  A series of tidal interactions and mergers will keep the total cold gas contents in bright galaxies decreasing, producing more early-type galaxies as time passes \\citep{park08,pc09,hp08}. The speed of this process is an increasing function of the large-scale density. Namely, even though the direct physical processes affecting galaxy morphology are the gravitational and hydrodynamic interactions between neighboring galaxies, the large-scale background density appears to control galaxy morphology through its statistical correlation with the frequency and strength of galaxy-galaxy interactions. Given the knowledge that the MDR in most region of the universe (note that Park et al. did not resolve the cluster regions) is the result of galaxy-galaxy interactions, one can naturally suspect the MRR and MDR in clusters are also due to the interactions between individual galaxies. It is the purpose of this paper to explore this possibility. ",
        "conclusions": "We have studied the dependence of various galaxy properties on the clustercentric radius and the environment attributed to the nearest neighbor galaxy using the SDSS galaxies associated with the Abell galaxy clusters. Our major findings are as follows.  1. There exists a characteristic scale where the galaxy properties such as morphology, color gradient, and structural parameters suddenly start to depend on the clustercentric radius at fixed nearest neighbor environment.  2. The characteristic scale depends on galaxy luminosity; the faint galaxies with $-17.0\\ge M_r > -19.0$ has the scale at $\\sim 3 r_{\\rm 200,cl}$ while the scale is $\\sim r_{\\rm 200,cl}$ for the brighter galaxies with $-20.5\\ge M_r > -22.5$.  3. The hydrodynamic interactions with nearby early-type galaxies seem to be the main drive to quenching star formation activity of late-type galaxies in clusters. We do not find evidence that the hot cluster gas is the main drive down to the clustercentric radius of $\\sim 0.1 r_{\\rm 200,cl}$.  4. The interaction with the cluster potential and harassment are the viable mechanisms that can account for the clustercentric radius dependence of the structural and internal kinematics parameters.  Existence of the characteristic scale means that the local galaxy number density is not responsible for the MDR in cluster because the local density is a smooth function of the clustercentric radius and has no discontinuity in general. The MDR appears working in clusters and also in the field because of the statistical correlation between the local density and the nearest neighbor distance. What is really working in clusters is the morphology-clustercentric radius-neighbor environment relation, where the neighbor environment means both neighbor morphology and the mass density attributed to the neighbor.  According to the third and fourth conclusions, the morphology transformation of the late-type galaxies in clusters is not due to a single mechanism and has been taken place through both hydrodynamic and gravitational processes."
    },
    "0812/0812.0102_arXiv.txt": {
        "abstract": "We calculate neutrino processes involving two nucleons at subnuclear densities using chiral effective field theory. Shorter-range noncentral forces reduce the neutrino rates significantly compared with the one-pion exchange approximation currently used in supernova simulations. For densities $\\rho < 10^{14} \\gcmiq$, we find that neutrino rates are well constrained by nuclear interactions and nucleon-nucleon scattering data. As an application, we calculate the mean-square energy transfer in scattering of a neutrino from nucleons and find that collision processes and spin-dependent mean-field effects dominate over the energy transfer due to nucleon recoil. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.4517_arXiv.txt": {
        "abstract": "The young stars near the supermassive black hole at the galactic center follow orbits that are nearly random in orientation and that have an approximately ``thermal'' distribution of eccentricities, $N(<e)\\sim e^2$. We show that both of these properties are a natural consequence of a few million years' interaction with an intermediate-mass black hole (IBH), if the latter's orbit is mildly eccentric and if its mass exceeds approximately 1500 solar masses. Producing the most tightly-bound S-stars requires an IBH orbit with periastron distance less than about $10$ mpc. Our results provide support for a model in which the young stars are carried to the galactic center while bound to an IBH, and are consistent with the hypothesis that an IBH may still be orbiting within the nuclear star cluster. ",
        "introduction": "Observations of the inner parsec of the Milky Way reveal two groups of massive young stars near the supermassive black hole (SBH). A group of roughly 40 stars move in approximately circular orbits that extend inward to a tenth of a parsec from the SBH \\citep{PaumardApJ2006,LuApJ2008}. Another group of roughly 20 stars, the S-stars, follow eccentric orbits well inside the disk stars \\citep{EckartMN1997}. The orbital periods of the S-stars are as short as 15 years and the orbits of these stars provide strong constraints on the mass and size of the central dark object \\cite{GillessenApJ2008,GhezApJ2008}.  The presence of young stars so near the galactic center is a puzzle, since giant molecular clouds, which are the sites of star formation elsewhere in the galaxy, would be unable to collapse and fragment in the tidal field of the SBH \\citep{MorrisApJ1993}. An orbiting gas cloud would instead be sheared into a set of rings; shocks between the rings would dissipate energy and the gas would settle into a thin disk.  If such a disk reaches a critical density, its self-gravity can overcome the tidal forces from the SBH allowing stars to form \\citep{LevinApJ2003,NayakshinAA2005}. This model has been shown to successfully reproduce the observed properties of the young disk stars, including their top-heavy mass function and their mildly noncircular orbits \\citep{BonnellScience2008}. But it fails to account for the S-stars, which move on eccentric, randomly-oriented orbits much closer to the SBH.  A number of models have been proposed to explain the S-stars but none is completely satisfactory \\citep{AlexanderPhysRept2005,Paumard2008}. The S-stars could be old stars that migrated inward and were ``rejuvenated\" by tidal heating or collisions with other stars. However their relatively normal spectra argue against such an exotic history \\citep{Figer2008}. Capture of young stars onto tightly-bound orbits via three-body exchange interactions involving compact remnants (e.g. stellar-mass black holes) may explain some of the S-stars \\citep{AlexanderApJ2004}. A similar model assumes that the S-stars were originally in binary systems; the more massive component of the binary was ejected during a close passage to the SBH, leaving the lower-mass star behind \\citep{GouldApJ2003}. Both of these models require an ad hoc reservoir of new stars at large radii, as well as some mechanism for placing these stars onto plunging orbits soon after their birth, so that they can pass near the SBH where the chance of a three-body exchange is appreciable. The models also have difficulty reproducing the observed distribution of S-star orbital eccentricities.  An alternative scenario posutlates that the young stars formed in a giant molecular cloud, far enough from the SBH that tidal forces did not preclude collapse and fragmentation \\citep{GerhardApJ2001}. The newly-formed star cluster then migrated inward via dynamical friction before tidal forces from the SBH dispersed it. The presence of an intermediate-mass black hole (IBH) at the center of the cluster would assist in the transport \\cite{HansenApJ2003}; in the absence of the IBH, the cluster would be completely disrupted by tidal stresses at a distance of ~one parsec from the SBH \\citep{ZwartApJ2003}. This model requires a relatively high mass ($\\gap 10^4 M_\\odot$) and density for the cluster, although not much greater than what is observed in existing galactic center star clusters like the Arches and the Quintuplet \\citep{Figer2008}. While the timescale for formation of the IBH is uncertain, simulations suggest that collisions between stars could form a massive remnant in a time shorter than the time for cluster inspiral and dissolution \\citep{ZwartApJ2002,FreitagMN2006}. This model is also appealing since IBHs provide potential solutions to a number of other outstanding problems, including the origin of the hyper-velocity stars \\citep{LevinApJ2006}, the structure of the stellar disks \\citep{LevinApJ2005,BerukoffApJ2006}, and the growth of SBHs \\cite{ZwartApJ2006}.  Here we show that the infalling star cluster model can also reproduce the peculiar orbits of the S-stars. ",
        "conclusions": "We have found that the presence of an IBH orbiting within the nuclear star cluster at the center of the Milky Way can efficiently randomize the orbits of stars near the SBH, converting an initially thin, co-rotating disk into a nearly isotropic distribution of stars moving on eccentric orbits around the SBH. Randomization of the orbital planes requires $\\sim 1$ Myr if the IBH mass exceeds $\\sim 1500\\msun$ and if the orbital eccentricity $\\sim 0.5$ or greater. Evolution to a ``thermal'' eccentricity distribution occurs on an even shorter time scale. The final distribution of stellar semi-major axes depends on the assumed size of the IBH orbit, but stars with apastron distances as small as the {\\it peri}astron distance of the IBH are naturally produced.  Our simulations contribute only a small piece to the bigger unsolved puzzle of the origin of the young stars at the galactic center. If models for the genesis of the S-stars that invoke an inspiralling IBH are deemed otherwise viable, our results show how the same models can also naturally reproduce the random and eccentric character of the stellar orbits, and all in a time that is less than stellar evolutionary time scales -- thus providing an essentially complete explanation for the ``paradox of youth\" of the S-stars.  Although we chose parameters for the IBH such that gravitational radiation would not alter its orbit appreciably in 5 Myr, the rapidity with which the IBH modifies the stellar orbits in our simulations suggests that even an IBH on a decaying orbit might be able to randomize and ``thermalize\" the S-star distribution before coalescing with the SBH. Thus, our model does not necessarily imply that an IBH is present, at the current time, within the S-star cluster. But if the IBH is there now, its presence might be detected in a number of ways:  1. The IBH will induce a motion of the SBH \\citep{HansenApJ2003}. Upper limits on the astrometric wobble of the radio source Sgr A$^*$ are so far consistent with the presence of an IBH with mass and semi-major axis in the range considered here \\citep{GillessenApJ2008}.  2. Stars can remain bound to the IBH if its Hill sphere is larger than its tidal disruption sphere; this condition is satisfied for SBH-IBH separations greater than $\\sim 0.05$ mpc. The motion of a star bound to the IBH would be the superposition of a Keplerian ellipse around the SBH and an additional periodic component due to its motion around the IBH; the latter would have a velocity amplitude $\\sim 0.1-10$ times the IBH orbital velocity and an orbital frequency from several hours to a few years, potentially accessible to astrometric monitoring.  3. In favorable circumstances, a near encounter of the IBH with a star unbound to it could produce observable changes in the star's orbit over month- or year-long time scales.  4. In our simulations with $q = 0.001$ and $e = (0.5,0.7)$, a fraction $13\\%$ of the stars were ejected from the SBH-IBH system in 5 Myr. A star ejected at $\\sim 10^3$ km s$^{-1}$ requires $\\sim 10^2$ yr to move beyond 0.1 pc implying a probability $\\sim 0.2 (N/10^4)$ of observing an escaping star at any given time in the Galactic center region, where $N$ is the number of stars subject to ejection. One S-star in fact appears to be on an escaping trajectory \\citep{GillessenApJ2008}.  Finally we note that some models for IBH formation predict that a large number of IBHs might co-exist in the galactic center region \\citep{ZwartApJ2006}. If so, the rate of orbital evolution of the young stars might be even higher than what we observe in our simulations with a single IBH."
    },
    "0812/0812.1618_arXiv.txt": {
        "abstract": "We study internal extinction of late-type galaxies in the Sloan Digital Sky Survey. We find that the degree of internal extinction depends on both the concentration index $c$ and $K_s$-band absolute magnitude $M_K$. We give simple fitting functions for internal extinction. In particular, we present analytic formulae giving the extinction-corrected magnitudes from the observed optical parameters. For example, the extinction-corrected $r$-band absolute magnitude can be obtained by $M_{r,0}=-20.77 +(-1+\\sqrt{1+4\\Delta (M_{r,obs}+20.77+4.93\\Delta)})/2\\Delta$, where $\\Delta = 0.236\\{1.35(c-2.48)^2-1.14\\}\\log(a/b)$, $c=R_{90}/R_{50}$ is the the concentration index, and $a/b$ is the isophotal axis ratio of the 25 mag/arcsec$^2$ isophote in the $i$-band. The $1\\sigma$ error in $M_{r,0}$ is $0.21{\\rm log}(a/b)$. The late-type galaxies with very different inclinations are found to trace almost the same sequence in the $(u-r)$-$M_r$ diagram when our prescriptions for extinction correction are applied. We also find that $(u-r)$ color can be a third independent parameter that determines the degree of internal extinction. ",
        "introduction": "Dust in spiral galaxies causes internal extinction. For example, due to dust, edge-on spiral galaxies in general look fainter than face-on spirals with the same intrinsic luminosity in short wavelength optical bands, in particular (see Fig. 12 of  Choi, Park, \\& Vogeley 2007). The study of internal extinction is important for determination of distances and absolute magnitude of galaxies. Giovanelli et al.~(1995) demonstrated that inadequate treatment of internal extinction has strong effects on the distances estimated using the Tully-Fisher relation (Tully \\& Fisher 1977). Internal extinction also causes biased measurements of galaxy star formation rates (see, for example, Bell \\& Kennicutt 2001; Sullivan et al. 2000) and affects determination of galaxy luminosity function (Shao et al.~2007). Earlier studies of internal extinction include Giovanelli et al.~(1994, 1995), Tully et al. (1998), and Masters, Giovanelli, \\& Haynes (2003). Recent works that contain relevant discussions on the topic are Rocha et al.~(2008), Unterborn \\& Ryden (2008), and Maller et al.~(2008).  The amount of obscuration by dust is larger when the observing wavelength is shorter. Therefore, the effect of inclination is more pronounced in short-wavelength bands, such as $u$ or $r$ band, and it may be much smaller in $K_s$ (2.17$\\mu$m) band. As a consequence, when the viewing angle changes for a given galaxy, $u$- or $r$-band magnitude changes while $K_s$-band magnitude does not change much. Therefore, $u-K_s$ or $r-K_s$ color tends to be larger when the galaxy is viewed more edge-on.  Internal extinction depends on many factors. Perhaps, the most important factor is the amount of dust. How dust is distributed can also be an important factor. The study of extinction versus inclination will ultimately reveal how dust is distributed and how much dust is contained in spiral galaxies. Then, what determines the amount and distribution of dust in spiral galaxies?  Earlier works have discussed dependence of internal extinction on luminosity. Giovanelli et al.~(1995) found that the amount of internal extinction in $I$ band depends on the galaxy luminosity. Tully et al.~(1998) also found a strong luminosity dependence using a magnitude-limited samples drawn {}from the Ursa Major and Pisces Clusters. Masters, Giovanelli, \\& Haynes (2003) studied internal extinction in spiral galaxies in the near infrared and also found a luminosity dependence. On the other hand, there are suggestions that internal extinction depends on galaxy type (de Vaucouleurs et al. 1991; Han 1992).  In this paper, we study the internal extinction in SDSS late-type galaxies. We investigate the dependence of the inclination effects on luminosity, concentration index $c$, and $u-r$ color. In \\S2, we describe the data set used in this paper. In \\S3, we study dependence of internal extinction on the concentration index $c$ and $K_s$-band luminosity separately. In \\S4, we measure dependence of internal extinction on the concentration index $c$ and $K_s$-band luminosity simultaneously. In \\S5, we derive dependence of inclination effects on $r$-band luminosity. The result in \\S5 is useful when $K$-band magnitude is not available. In \\S6, we present dependence on $u-r$ color. We give discussion in \\S7 and conclusion in \\S8. ",
        "conclusions": "Our major conclusions are as follows. \\begin{enumerate} \\item{} We have shown that the slope $\\gamma_{\\lambda}$ depends on both $K_s$-band absolute magnitude $M_K$ and the concentration index $c$. The fitting functions for the relations are given in Table 1 and 2. \\item{} We have also shown that the slope $\\gamma_{\\lambda}$ depends on both the concentration index $c$ and $r$-band absolute magnitude $M_{r,0}$, where the subscript `$0$' denotes the value after inclination correction. The relations are also given in Table 1 and 2. \\item{} We have derived analytic formulae giving the extinction-corrected magnitudes from the observed optical parameters (see, for example, Equation (\\ref{eq:RfromR})). \\item{} We have shown that $(u-r)_0$ can be a third parameter that determines the slope $\\gamma_{\\lambda}$. \\end{enumerate}"
    },
    "0812/0812.4721_arXiv.txt": {
        "abstract": "There is no direct evidence for radiation domination prior to big-bang nucleosynthesis, and so it is useful to consider how constraints to thermally-produced axions change in non-standard thermal histories. In the low-temperature-reheating scenario, radiation domination begins as late as $\\sim 1$ MeV, and is preceded by significant entropy generation. Axion abundances are then suppressed, and cosmological limits to axions are significantly loosened. In a kination scenario, a more modest change to axion constraints occurs. Future possible constraints to axions and low-temperature reheating are discussed. ",
        "introduction": "If the axion has mass $m_{\\rm a}\\gsim 10^{-2}~{\\rm eV}$, it will be produced thermally, with cosmological abundance $\\Omega_{{\\rm a}}h^{2}=\\left(m_{{\\rm a}}/130~{\\rm eV}\\right)\\left(10/g_{*\\rm F}\\right)\\label{thermalfreeze},$ where $g_{*{\\rm F}}$ is the effective number of relativistic degrees of freedom when axions freeze out \\cite{crewther, baluni,pq,changchoi,notinvisible,kt,murayama}. If $m_{\\rm a}\\lsim 1~{\\rm eV}$, axions free-stream to erase density perturbations while they are relativistic, and thus suppress the matter power spectrum on small scales, much like neutrinos \\cite{hu1,hu2,kt,haneutrino5,changchoi,crotty,tegmarksdss,barger,seljakneutrino,fukugitawmap3,spergelwmap3,pierpaoli}. Data from large-scale structure (LSS) surveys and cosmic microwave-background (CMB) observations impose the constraint $m_{\\rm a}\\lsim1~{\\rm eV}$ to light hadronic axions \\cite{hanraffelt,raffeltwong,slosar}. We restrict our attention to hadronic axions in this work.  These constraints were determined in the standard radiation-dominated scenario. The transition to radiation domination after inflation might be gradual \\cite{kamion}. In a modified thermal history, relic abundances may change, due to modified freeze-out temperatures and suppression from entropy generation.  The universe could have reheated to a temperature as low as $1~{\\rm MeV}$ \\cite{kawasaki2,ichikawa3,giudice2,hannestadreheat,kolbnotario}. This low-temperature reheating (LTR) scenario may be modeled simply through the decay of a massive particle $\\phi$ into radiation, with fixed rate $\\Gamma_{\\phi}$. This decay softens the scaling of temperature $T$ with cosmological scale factor $a$, increasing the Hubble parameter $H(T)$ and leading to earlier freeze-out for certain relics. Entropy generation then highly suppresses these relic abundances. Kination models offer another alternative to the standard thermal history, without entropy production, and cause more modest changes in abundances \\cite{kination}.  Past work has determined the relaxation in constraints to neutrinos, weakly interacting massive particles, and non-thermally produced axions are relaxed in LTR \\cite{giudice1,giudice2,yaguna}. Here, we present new constraints to thermally-produced axions in the LTR scenario. We point the reader to Ref. \\cite{reh} for a discussion of the more modest changes to axion constraints in the kination scenario, and for additional details relevant to the following discussion. We conclude by discussing the impact of future LSS surveys and CMB measurements of the primordial helium abundance on the allowed parameter space for axions. ",
        "conclusions": "LTR suppresses the abundance of thermally-produced axions, once  $T_{\\rm rh}\\sim 50~{\\rm MeV}$, as a result of dramatic entropy production. The cosmologically allowed window for $m_{\\rm a}$ is extended as a result. Future probes of the matter power spectrum or the primordial helium abundance may definitively explore some of this parameter space.   \\begin{theacknowledgments} D.G. was supported by the Gordon and Betty Moore Foundation and thanks the organizers of DM08. T.L.S. and M.K. were supported by DoE DE-FG03-92-ER40701 and the Gordon and Betty Moore Foundation. \\end{theacknowledgments}"
    },
    "0812/0812.3271_arXiv.txt": {
        "abstract": "We report on our time-resolved CCD photometry during the 2005 June superoutburst of a WZ Sge-type dwarf nova candidate, ASAS 160048-4846.2. The ordinary superhumps underwent a complex evolution during the superoutburst. The superhump amplitude experienced a regrowth, and had two peaks. The superhump period decreased when the superhump amplitude reached to the first maximum,  successively gradually increased until the second maximum of the amplitude, and finally decreased again. Investigating other SU UMa-type dwarf novae which show an increase of the superhump period, we found the same trend of the superhump evolution in superoutbursts of them. We speculate that the superhump regrowth in the amplitude has a close relation to the increase of the superhump period, and all of SU UMa-type dwarf novae with a superhump regrowth follow the same evolution of the ordinary  superhumps as that of ASAS 160048-4846.2. ",
        "introduction": "\\label{intro}  Cataclysmic variables (CVs) are close binaries containing a white dwarf (primary) and a late-type star (secondary). The secondary fills its Roche-lobe and transfers gas to the primary, so that an accretion disk is formed around it (for a review, see e.g. \\cite{war95book}; \\cite{hel01book}; \\cite{smi07review}). Dwarf novae are a subclass of CVs.  SU UMa-type stars are a subgroup of dwarf novae which show two types of outbursts: normal outbursts, and superoutbursts. During the superoutbursts, repetitive modulations with an amplitude of 0.1-0.3 mag, called superhumps, are shown. The period of the superhumps is a few percent longer than the orbital period of the system. The thermal-tidal instability model is the most acceptable one for explaining the general behavior of SU UMa-type dwarf novae \\citep{osa89suuma}. According to the tidal-instability theory, an accretion disk becomes unstable due to the gravitational interaction with the secondary star when it reaches a critical radius of the 3:1 resonance \\citep{whi88tidal}. Superhumps can be explained by a beat phenomenon of the precession of a tidally-distorted disk and the orbital motion.  WZ Sge-type dwarf novae are one of the subtypes of SU UMa-type dwarf novae with the shortest orbital periods among SU UMa stars (see e.g. \\cite{kat01hvvir}). Their observational properties are (1) the extremely long supercycle (over 5 years), which is a period between two successive superoutbursts, (2) the large amplitude of superoutbursts over 6 mag (4-5 mag in many SU UMa-type stars), (3) the absence of normal outbursts, and (4) the presence of double-peaked humps, called early superhumps, before emergence of the ordinary superhumps. Although the clear definition of the WZ Sge type is still controversial, its representative members where at least two outbursts have been observed so far and the supercycle was measured are AL Com (\\cite{kat96alcom}, \\cite{pat96alcom}), EG Cnc (\\cite{kat04egcnc}), HV Vir (\\cite{ish03hvvir}), WZ Sge itself (\\cite{kat04vsnet}), and GW Lib (Imada et al. in preparation).   On 2005 June 9, ASAS 160048-4846.2 (hereafter ASAS 1600) was initially discovered by the All Sky Automated Survey (\\cite{poj02asas3}) as an eruptive object. This is the only one recorded superoutburst. In the light curves of this superoutburst, \\citet{ima06asas1600letter}, hereafter called Paper I, found double-peaked humps with a period of 0.063381(41) days, before ordinary superhumps emerged with a period of 0.064927(3) days. Based on evidence for early superhumps, they identified ASAS 1600 as a promising candidate for WZ Sge-type dwarf novae.  In this paper, we reanalysed the data reported in Paper I, and report on the evolution of the ordinary superhumps. The details of the observations are described in the next section. The results of the observations and the analyses are summarized in the section 3. We will discuss the superhump evolution and the WZ Sge nature of ASAS 1600 in the section 4, and put conclusions in the last section 5. ",
        "conclusions": "We photometrically studied the 2005 superoutburst of a dwarf nova ASAS 1600. Our conclusion is summarized below:  \\begin{enumerate}  \\item  ASAS 1600 showed a regrowth of the ordinary superhumps in the amplitude during the 2005 superoutburst, which is frequently found in SU UMa-type dwarf novae with an increase of $P_{SH}$.  \\item  During the 2005 superoutburst, the superhump period decreased when the superhump amplitude reached to the first maximum, successively gradually increased until the second maximum of the amplitude, and finally decreased again. This course of the superhump evolution seems to be common in SU UMa-type dwarf novae with a superhump regrowth.  \\item  ASAS 1600 does not seem to be a typical WZ Sge-type dwarf nova, but, at least, a very close to typical WZ Sge-type dwarf nova.  \\end{enumerate}  \\bigskip  We acknowledge with thanks the variable star observations from the AAVSO and VSNET International Database contributed by observers worldwide and used in this research. Thanks are also due to the anonymous referee for useful comments. This work was supported by the Grant-in-Aid for the Global COE Program \"The Next Generation of Physics, Spun from Universality and Emergence\" from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan."
    },
    "0812/0812.1797_arXiv.txt": {
        "abstract": "\\noindent We present projections for reconstruction of the inflationary potential expected from ESA's upcoming Planck Surveyor CMB mission.  We focus on the effects that tensor perturbations and the presence of non-Gaussianities have on reconstruction efforts in the context of non-canonical inflation models. We consider potential constraints for different combinations of detection/null-detection of tensors and non-Gaussianities.  We perform Markov Chain Monte Carlo and flow analyses on a simulated Planck-precision data set to obtain constraints.  We find that a failure to detect non-Gaussianities  precludes a successful inversion of the primordial power spectrum, greatly affecting uncertainties, even in the presence of a tensor detection.  In the absence of a tensor detection, while unable to determine the energy scale of inflation, an observable level of non-Gaussianities provides correlations between the errors of the potential parameters, suggesting that constraints might be improved for suitable combinations of parameters.  Constraints are optimized for a positive detection of both tensors and non-Gaussianities. ",
        "introduction": "A faithful reconstruction of the inflaton potential is paramount to understanding the epoch of inflation.  Recent cosmological observations \\cite{Hinshaw:2008kr,Gold:2008kp,Nolta:2008ih,Dunkley:2008ie,Komatsu:2008hk} have revealed much about the possible forms of the potential \\cite{Kinney:2006qm,Lesgourgues:2007gp,Peiris:2008be,Lorenz:2008je}.  Future CMB missions, such as ESA's Planck Surveyor \\cite{planck}, and ground-based polarization experiments such as BICEP \\cite{Yoon:2006jc}, promise to further hone our understanding of the early universe.  The Planck Surveyor will provide improved measurements of the temperature and E-mode polarization anisotropies of the CMB out to $\\ell = 3000$, and will detect the B-mode polarization characteristic of gravity waves if the tensor/scalar ratio, $r$, is greater than around  0.05.  Planck will also provide the first high-quality measurement of any departure from Gaussianity exhibited by the CMB temperature fluctuations, with a projected sensitivity $|f_{NL}| \\gtrsim 5$ \\cite{Komatsu:2002db,Creminelli:2003iq}.  Recent times have also seen exciting progress in inflationary model building, and much is now understood about how to implement inflation within string theory (see \\cite{Cline:2006hu,Kallosh:2007ig,Burgess:2007pz,McAllister:2007bg} for reviews).  One particularly exciting realization is provided by the dynamics of D-branes evolving in a higher-dimensional, warped spacetime.  A novel aspect of these models is the existence of a speed limit on the field space, resulting from causality restrictions on the motion of the branes in the bulk spacetime.   This allows for slow roll to be achieved even in the presence of a steep potential.  In terms of effective field theory, the speed limit is enforced by non-canonical kinetic terms in the DBI action describing the motion of the D-branes.  These models, termed DBI inflation \\cite{Silverstein:2003hf}, are further distinguished by the fact that scalar perturbations propagate at a sound speed $c_s \\leq 1$.  In the limit $c_s \\ll 1$, these fluctuations exhibit strong non-Gaussianities \\cite{Alishahiha:2004eh,Chen:2006nt,Spalinski:2007qy,Bean:2007eh,LoVerde:2007ri}.  Non-Gaussianities of this type were first popularized in the context of {\\it k-inflation} \\cite{ArmendarizPicon:1999rj}, a model-independent construction in which higher-derivative kinetic terms drive inflation in the absence of any potential energy.   This phenomenon is a general characteristic of models with non-canonical kinetic terms in the Lagrangian.  With the advent of these more sophisticated theoretical constructions, and the increasingly accurate cosmological probes planned to study them, more generalized methods of inflationary reconstruction have been developed.  Bean {\\it et al.} \\cite{Bean:2008ga} present a formalism for reconstructing a general action -- both potential and non-canonical kinetic energy functions can be linked to observations.  The new functional degree of freedom presented by the non-canonical form of the kinetic energy typically gives rise to additional observational signatures.  For example, the speed of sound, $c_s$, becomes an additional dynamical degree of freedom in k-inflation and DBI, giving rise to non-Gaussian density fluctuations. However, $c_s$ also affects the form of the primordial power spectra.  A successful reconstruction of the inflaton potential, $V(\\phi)$, thus requires a measurement of $c_s$, since otherwise the expressions giving the spectral parameters cannot be uniquely inverted to obtain $V(\\phi)$.  An accurate reconstruction must also include a detection of the tensor power spectrum.  Currently, neither of these quantities have been measured: present bounds on any non-Gaussianities lie within $-9 < f_{NL}^{\\rm local} < 111$ and $-151 < f_{NL}^{\\rm equil} < 253$ (95\\% CL) \\cite{Komatsu:2008hk}, and $r < 0.2$ (95\\% CL) with WMAP+BAO+SN \\cite{Komatsu:2008hk} and $r<0.35 $ (95\\% CL) for WMAP+ACBAR \\cite{Kinney:2008wy} for power-law spectra, but might be larger if more exotic spectra are considered \\cite{Powell:2007gu}.  In this paper, we make reconstruction projections for the Planck satellite within the context of this larger inflationary parameter space, namely, one that contains $c_s$ as an additional degree of freedom.  We seek to determine what combination of observations will be required to most successfully reconstruct the potential.  In the best of all worlds, future observations will detect both tensors and non-Gaussianities, telling us not only that exotic physics is responsible for inflation, but also allowing for a successful reconstruction program.  In this paper, we focus particularly on cases in which only one of these observables is positively detected.  In Section 2, we review potential reconstruction for the case of non-canonical inflation. In Section 3, we perform a Bayesian analysis on a simulated Planck-precision data-set for which there is a positive tensor detection, $r = 0.1$, but no detection of non-Gaussianities, $|f_{NL}| < 5$.  This corresponds to a universe in which canonical inflation took place (with $c_s = 1$), or, for example, DBI inflation with $c_s \\in [0.25,1]$.  We obtain constraints on $V(\\phi_0)$, $V'(\\phi_0)$, and $V''(\\phi_0)$ for both possibilities and find that the error bars are significantly larger if $c_s$ is taken as a free parameter within the above range. Our inability to resolve $c_s$ prohibits us from using the shape of the power spectra alone to reconstruct $V(\\phi)$. In Section 4, we investigate this problem using inflationary flow methods.  The flow method is a natural extension of the MCMC analysis, allowing higher dimensional parameter spaces to be explored with relative ease. We reproduce the results obtained in Section 3, and then investigate the effect of allowing the speed of sound to vary over the course of inflation.  We find that constraints on $V_0''$ loosen considerably as this new degree of freedom is introduced.  Next, we use flow methods to investigate the observational possibility of a positive detection of $f_{NL}$ and a null detection of $r$. We find that constraints on the individual potential parameters are worse in this case than in the case where tensors are detected, but that the errors are highly correlated, and tight bounds can be placed on combinations of these parameters. Finally, we consider the case for which {\\it both} $r$ and $f_{NL}$ are measured. This is clearly the best case scenario for near future experiments, since it will not only place bounds on the potential parameters from above and below, but it will also be a smoking gun for exotic physics. In Section 5 we present conclusions. ",
        "conclusions": "We have considered potential reconstruction forecasts for the upcoming Planck Surveyor CMB mission in the context of non-canonical inflation.  We have focused on the following observational outcomes: a detection of a tensor signal and a null detection of non-Gaussianities, a null detection of tensors and a positive detection of non-Gaussianities, and a positive detection of both tensors and non-Gaussianities.  To explore the first case, we perform a Markov Chain Monte Carlo analysis on a simulated Planck-precision data-set.  We obtain constraints on the potential parameters $V_0$, $V_0'$, and $V_0''$, and the sound speed $c_s$.  A null detection of non-Gaussianities at Planck resolution corresponds to $|f_{NL}| \\lesssim 5$, with the result that $c_s$ is only constrained to lie within the range $[0.25,1]$ for the case of DBI inflation.  Since $c_s$ also affects the form of the power spectrum, a failure to constrain it precludes us from inverting our spectrum observation to obtain the inflaton potential.  The errors on $V_0'$ and $V_0''$ increase by an order of magnitude relative to constraints obtained on canonical inflation.  We next used the flow formalism to investigate the effect of allowing the speed of sound to vary over the course of inflation.  This requires the introduction of terms governing its evolution.  Unlike MCMC methods, the flow formalism is well suited to studying reconstructions to high-order.  The strongest effect of a variable speed of sound is a much increased error on the second derivative of $V(\\phi)$.  A variable speed is generic to models such as DBI inflation.  The worsening of the constraint on $V_0''$ results from the fact that there is additional freedom in tuning the geometry in order to get the same observables. (We note that this has been a very conservative analysis as far as our constraints on $f_{NL}$ are concerned.  If one considers $f_{NL}^{\\rm equil}$, the non-Gaussianities associated with equilateral triangles in $k$-space, the Planck detection threshold could be as large as $f_{NL}^{\\rm equil} \\lesssim 60$ \\cite{Liguori:2008vf}.)  We also used the flow approach to investigate the observational outcome of a detection of different levels of non-Gaussianity. We first consider the case of a null measurement of $r$. Because neither $V_0$ or $V_0'$ are bounded from below, and since a larger level of non-Gaussianity is generated by larger values of $V_0'$, we find that the errors on the individual parameters $V_0$, $V_0'$, and $V_0''$ grow relative to the case where neither are detected. However, these errors are strongly correlated, and it is nonetheless possible to tightly constrain combinations of these parameters. The parameters $r$ and $f_{NL}$ therefore play complementary roles in potential reconstruction. Finally, we consider the case where {\\it both} $r$ and $f_{NL}$ are detected. The complementary role that $r$ and $f_{NL}$ play in potential reconstruction is again clear, where gravity waves place bounds on $V_{0}$ and $V'_{0}$ from above and below, and non-Gaussianities determine the width of the bands in the $V$-$V'$ and $V$-$V''$ planes.  The approach taken in this paper has been phenomenological.  We have treated the potential and $c_s$ as freely tunable parameters. We have purposefully kept the study general to encompass the larger class of $k$-inflation and other non-canonical models based on effective field theory. However, it is very likely that many of the potentials and sound speeds sampled in this analysis do not represent realistic compactification scenarios in string theory, and might not correspond to realizations of DBI inflation. What we have shown in this study is that, if the observational outcomes considered here are reflected in the actual data, then our ability to reconstruct the potential is significantly weakened. While we ultimately hope that the best-case scenario will be realized with a detection of both tensors and non-Gaussianities, in the meantime we can only wait with guarded anticipation as to what Planck will reveal about the universe."
    },
    "0812/0812.0178_arXiv.txt": {
        "abstract": "We present observations of the CO(1-0) emission in the central 750~pc (10\\arcsec) of the counter-rotating disc galaxy NGC~4550, obtained at the Institut de Radioastronomie Millim\\'etrique (IRAM) Plateau de Bure Interferometer. Very little molecular gas is detected, only $1\\times 10^{7}$ $\\rm M_{\\sun}$, and its distribution is lopsided, with twice as much molecular gas observed at positive relative velocities than at negative relative velocities. The velocity gradient in the CO(1-0) emission shows that the molecular gas rotates like the thicker of the two stellar discs, which is an unexpected alignment of rotations if the thinner disc was formed by a major gas accretion event. However, a simulation shows that the gas rotating like the thicker disc naturally results from the coplanar merger of two counter-rotating disc galaxies, demonstrating the feasibility of this scenario for the formation of NGC~4550. We investigate various star formation tracers to determine whether the molecular gas in NGC~4550 is currently forming stars. UV imaging data and optical absorption linestrengths both suggest a recent star formation episode; the best-fitting two population model to the UV-optical colours yields a mass of young stars of $5.9 \\times 10^{7}$ M$_{\\sun}$ with an age of 280 Myr. The best information on the current star formation rate is a far infrared-based upper limit of only 0.02 M$_{\\sun}$ yr$^{-1}$. We are thus witnessing NGC~4550 either in a dip within a bursty star formation period or during a more continuous low-level star formation episode. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.0980_arXiv.txt": {
        "abstract": "We begin an exploration of the physics associated with the general CP-conserving MSSM with Minimal Flavor Violation, the pMSSM. The 19 soft SUSY breaking parameters in this scenario are chosen so as to satisfy all existing experimental and theoretical constraints assuming that the WIMP is a conventional thermal relic, \\ie, the lightest neutralino. We scan this parameter space twice using both flat and log priors for the soft SUSY breaking mass parameters and compare the results which yield similar conclusions. Detailed constraints from both LEP and the Tevatron searches play a particularly important role in obtaining our final model samples. We find that the pMSSM leads to a much broader set of predictions for the properties of the SUSY partners as well as for a number of experimental observables than those found in any of the conventional SUSY breaking scenarios such as mSUGRA. This set of models can easily lead to atypical expectations for SUSY signals at the LHC. ",
        "introduction": "The LHC will soon begin operations in earnest and will thereafter begin a detailed study of the Terascale. Amongst the many prospective new physics signals that may be realized at the LHC, the possible observation of an excess of events with missing energy is awaited with great anticipation. Given the strong collection of evidence for the existence of dark matter \\cite{Bertone:2004pz} and the naturalness of the WIMP paradigm \\cite{Primack:1988zm}, it is no surprise that many beyond the Standard Model (SM) scenarios predict the appearance of missing energy signatures at the LHC \\cite{missinglhc}. Perhaps the most well explored of these scenarios is R-Parity conserving, N=1 Supersymmetry (SUSY) where the WIMP is identified as the lightest supersymmetric particle (LSP), which most commonly corresponds to the lightest neutralino or  gravitino. The simplest manifestation of this picture is the Minimal Supersymmetric Standard Model (MSSM), where the gauge group of the SM is maintained and the matter content of the SM is augmented only by an additional Higgs doublet.  While the MSSM has many nice features, perhaps its least understood aspect is the nature of the mechanism which breaks the Supersymmetry. There is an ever growing list of possible candidate scenarios, including mSUGRA \\cite{msugrab}, GMSB \\cite{gmsbb}, AMSB \\cite{amsbb}, and gaugino mediated supersymmetry breaking \\cite{ggmsbb}.  Each of these mechanisms strongly influences both the mass spectra and decay patterns of the various sparticles and hence it is highly non-trivial to make model-independent statements about the properties of the LSP or about SUSY signatures at colliders such as the LHC.  This puts forth the question of whether there is some way to study the MSSM in a broad fashion while simultaneously adopting as few simplifying assumptions as possible.  The purpose of this paper is to address this issue. In trying to answer this question, one is immediately faced with the daunting task of dealing with the over 120 parameters, arising mainly in the soft-breaking sector, which describe SUSY breaking in the MSSM \\cite{mssmrev}.  Clearly the phenomenological study of any model with such a large number of parameters is totally hopeless. Obviously {\\it some} assumptions need to be made in order to reduce this large number of independent parameters to a viable subset and there are many possible avenues one could follow. Most studies assume a specific SUSY breaking mechanism, employing theoretical assumptions at the GUT scale, and thus reduce the number of independent model parameters to 4 or 5. In this investigation, we will not make any assumptions at the high scale. We will, however, restrict ourselves to the CP-conserving MSSM (\\ie, no new phases) with minimal flavor violation (MFV)\\cite{mfv}. Moreover, to simplify matters further and help soften the impact of the experimental constraints arising from the flavor sector, we will require that the first two generations of sfermions be degenerate. This leaves us with 19 independent, real, weak-scale, SUSY Lagrangian parameters to consider. These include the gaugino masses $M_{1,2,3}$, the Higgsino mixing parameter $\\mu$, the ratio of the Higgs vevs $\\tan \\beta$, the mass of the pseudoscalar Higgs boson $m_A$, and the 10 squared masses of the sfermions (5 for the assumed degenerate first two generations and a separate 5 for the third generation). Finally, due to the small Yukawa couplings for the first two generations, independent $A$-terms are only phenomenologically relevant for the third generation and will be included in our study for the $b,t$ and $\\tau$ sectors. This set of 19 parameters leaves us with what has been called the phenomenological MSSM (pMSSM) in the literature{\\cite{Djouadi:2002ze}} and is the scenario that we will study in some detail below. We note that for this study we will make the further assumption that the LSP is the neutralino and is a conventional thermal relic within the standard cosmology with a radiation dominated era.  The primary goal of the first phase of this study is to obtain a large set of MSSM parameter points that can be used for future studies of SUSY signatures, \\eg, signals at the LHC or in cosmic rays.  To accomplish this, we will perform a scan over this 19-dimensional parameter space and subject the resulting specific sets of parameter choices (which we denote as models) to a large and well-known (and some perhaps, not so well-known) set of experimental and theoretical constraints{\\cite {Freitas:2008uk}}. Furthermore, we will perform two independent scans which will employ different priors for choosing the above mass parameters; this is important since some bias may be introduced into the analysis by how one choose parameter space points. By choosing more than one prior and comparing results this sensitivity can be estimated. We will show that while both scans yield qualitatively similar results, quantitative differences will be observed in some cases.  We will perform our comparison to the experimental data with some care to identify sets of parameter space points which are at present phenomenologically viable. As we will see below, applying these constraints can be complicated and is not always as straightforward as employed in some other analyses.  This is particularly true in the case where experimental bounds are predicated on specific SUSY breaking mechanisms. As we are not a priori interested in any particular region of parameter space (\\eg, we are enforcing the WMAP dark matter density{\\cite {Komatsu:2008hk}} observations only as an upper bound) or in obtaining a set of points that provide a `best fit' to the data, we will not need to make use of Markov-chain Monte Carlo techniques{\\cite {MCMC,orig}}.{\\footnote {The MCMC approach asks a different question than we do here, namely `What regions of the parameter space provide the best fits to the data?', by using, for example, a chi-squared analysis. This requires one to assume precise knowledge of the probability distributions for both experimental results as well as theoretical predictions. For example, given the uncertainties associated with the g-2 of the muon, is it justified to treat the error associated with the apparent difference with the SM prediction as Gaussian?  These types of questions introduce prejudice into the analysis, which we try to avoid, and so we have instead adopted the scanning approach.}} We will then perform a detailed study of the properties of these surviving models and, in particular, highlight those aspects which are not generally encountered when only a given set of SUSY breaking scenarios is considered. This allows us to get a perspective on  the wide range of predictions for various observables that can arise in the more general pMSSM framework in comparison to those found in specific SUSY breaking scenarios.  In the next Section, we will discuss our scanning techniques, the ranges of the soft breaking parameters that we employ, as well as the values of the SM input parameters that we use. In Section 3, we will detail the various theoretical and experimental constraints at some length and will explain exactly how they are employed in our analysis. Section 4 presents our analysis and a survey of the results, while Section 5 contains an overview discussion of the lessons learned from this analysis. Our conclusions can be found in Section 6. ",
        "conclusions": "In this paper we begin a detailed study of the 19 parameter, CP-conserving pMSSM under the assumptions of ($i$) minimal flavor violation, ($ii$) the LSP being identified with the lightest neutralino, ($iii$) only the Yukawa couplings of the third generation are important and ($iv$) the WMAP measurement of the dark matter relic density is taken only as an upper limit to that arising from the LSP. Carefully subjecting numerous combinations of parameter space points, with mass terms randomly generated assuming either flat or log priors, to a large number of theoretical and experimental constraints we arrive at two sets of models that survive all of the restrictions.  With these two sets of models, we first examined the resulting distributions of expected Higgs and sparticle masses; qualitatively and even semi-quantitatively either choice of prior was found to return quite similar results except for trivially obvious differences due to how the parameters were generated. In both cases, we found some instances where surprisingly light sparticles are still allowed by the data. Next, we studied the identity of the nLSP and its mass splitting with the LSP. We found that for either choice of priors, over a dozen sparticles can play the role of the nLSP with more or less the same probabilities in the two model samples. The most common nLSP identities were found to be $\\tilde \\chi_1^\\pm$ and $\\tilde \\chi_2^0$ with other possibilities being $\\tilde e_R, \\tilde \\tau_1$ and $\\tilde t_1$. Less familiar nLSP identities were also found to occur such as the gluino or the first/second generation squarks. We also showed that the mass splitting between these two states, $\\Delta m\\equiv m_{nLSP}-m_{LSP}$, which is a critical parameter for collider signals and studies, was found to vary over a very large range of six orders of magnitude with values in excess of $\\sim 100$ GeV and as small as $\\sim 100$ keV or less.  The greatest range was found in the case where the lightest chargino was the nLSP. Small values of $\\Delta m$ were found to be quite common particularly in the case of the lightest chargino or the second lightest neutralino being the nLSP. This would imply that searches for long-lived and (essentially) stable particles at the LHC will be of particular importance. In some cases, squarks and gluinos were found to be relatively light and not widely separated in mass from the corresponding LSP. The decays of such states would likely lead to soft jets in the final state which may be difficult to observe above SM backgrounds at the LHC. The mass spectra of the models we obtained were such as to substantially diminish the needed amount of fine tuning in comparison to, \\eg, mSUGRA-based scenarios.  Subsequently, we examined the postdictions arising from both model sets for a number of observables such as the branching fraction for $b\\to s\\gamma$ and the shift in the value of the $g-2$ of the muon; these were found to be reasonably similar for both cases of priors. Since the predicted dark matter density was allowed to float to any value at or below that obtained by WMAP, we found that the average contribution of the LSP to dark matter peaked at $\\sim 4\\%$, with long tails, implying that most of the time some other source of dark matter must typically dominate. However, we do see that there is a reasonable fraction of models where the LSP contribution to the dark matter density is, in fact, quite appreciable in comparison to the WMAP measurement. These results were found to hold for either choice of priors. We then went further and examined the implications of these models for the direct detection of dark matter; we found that the range of possible cross sections is far larger than that obtained in any of the well-known SUSY breaking frameworks, such as the cMSSM or mSUGRA. Furthermore these ranges extended into the regions one would have expected to only arise from either the Universal Extra Dimensions or Little Higgs models. This implies that these observables alone cannot be used to discriminate between different Terascale models.  It is important to note that while the two sets of priors have led to qualitatively very similar pictures, some {\\it quantitative} differences do appear as can be seen from the figures above. For example, the mass spectra of the SUSY particles in both cases, while qualitatively similar, are different in detail since the log prior case generally leads to smaller masses.  Clearly the detailed study of the pMSSM models is quite exciting as well as complex; we are just at the initial stages of this analysis and there are many things remaining to be done which we will address in subsequent publications.  In addition, we plan to make our model sets public in the near future, in order to encourage a wider set of possible analyses. Hopefully we will not have to wait too long until we actually have real data from the LHC to examine."
    },
    "0812/0812.1809_arXiv.txt": {
        "abstract": "We report on the design and estimated performance of a balloon-borne hard X-ray polarimeter called HX-POL. The experiment uses a combination of Si and Cadmium Zinc Telluride detectors to measure the polarization of 50 keV-400 keV X-rays from cosmic sources through the dependence of the angular distribution of Compton scattered photons on the polarization direction. On a one-day balloon flight, HX-POL would allow us to measure the polarization of bright Crab-like sources for polarization degrees well below 10\\%. On a longer (15-30 day) flight from Australia or Antarctica, HX-POL would be be able to measure the polarization of bright galactic X-ray sources down to polarization degrees of a few percent. Hard X-ray polarization measurements provide unique venues for the study of particle acceleration processes by compact objects and relativistic outflows. In this paper, we discuss the overall instrument design and performance. Furthermore, we present results from laboratory tests of the Si and CZT detectors. ",
        "introduction": "\\IEEEPARstart{X}{-ray} astronomy has made major contributions to modern astronomy and cosomology for the last three to four decades. To name a few examples, X-rays allowed us to study the hot gas in which the galaxies of galaxy clusters are embedded, to study gas briefly before it plunges into stellar mass and supermassive black holes, and to study mass accreting neutron stars. Presently, the National Aeronautics and Space Administration (NASA) and the European Space Agency (ESA) fly the Chandra and XMM-Newton soft X-ray telescopes, which set world records in terms of angular resolution (Chandra: 0.5'') and collection area (XMM-Newton: 4,300 cm$^2$ at 1.5 keV). Even though X-ray astronomy has been extremely successful over several decades, there are still largely unexplored areas which future space-borne X-ray telescopes can pioneer. The NuSTAR hard X-ray telescope\\footnote{http:$//$www.nustar.caltech.edu$/$} is scheduled for launch in 2011 and will image 6-80 keV X-rays with an angular resolution of 40''. The International X-ray Observatory (IXO)\\footnote{http:$//$ixo.gsfc.nasa.gov$/$} would enable high-throughput excellent energy resolution (E/$\\Delta$E$>$2400 at 6 keV) X-ray spectroscopy in the 0.3-10 keV energy band. The EXIST black hole finder probe\\footnote{http:$//$exist.gsfc.nasa.gov$/$} would use coded mask imaging to scrutinize the entire sky every 170 min in the 5 keV to 600 keV energy band.  This paper focuses on X-ray polarimetry, another very promising and largely unexplored area. X-ray polarimetry would increase the parameter space for the study of compact objects like black holes and neutron stars from two dimensions (time variability and energy spectra) to four dimensions, by adding two qualitatively new parameters: polarization degree and polarization direction \\cite{Rees:75,Shap:75,Shap:76,Mesz:88}. The polarization measurements can be used to identify the processes responsible for the observed X-ray emission, to constrain viewing perspectives and opacities in accretion disk systems, and to gain access to the properties of the magnetic field (order and orientation) in the emission regions. Recently, concepts have been developed to measure the polarization of soft X-rays by tracking the photo-effect electrons that tend to be ejected parallel to the electric field vector of the X-rays with gas pixel detectors \\cite{2008arXiv0810.2700C} and with time projection chambers \\cite{2007SPIE.6686E..29H}. At $>$20 keV energies, the Compton effect can be used to measure the polarization, as photons are preferentially scattered in the direction perpendicular to the electric field vector. Early soft X-ray polarization measurements were done with the OSO~8 experiment which used mosaic crystals of graphite to yield polarization-sensitive Bragg reflection of 2.6 keV and 5.2 keV X-rays \\cite{Weis:78}. The Ge detectors on board the INTEGRAL $\\gamma$-ray observatory were used to measure the polarization of gamma-rays based on Compton effect polarimetry \\cite{2008Sci...321.1183D}. Recently, NASA approved the Gravity and Extreme Magnetism SMEX (GEMS) mission as a ``small explorer mission''. GEMS will combine grazing incident X-ray mirrors with time projection chamber detectors to measure the polarization of 2-10 keV X-rays, and will be launched in the next decade \\cite{GEMS}. The japanese ASTRO-H mission (formerly called ``NeXT'') will have a ``soft $\\gamma$-ray telescope'' which uses Si and CdTe detector stacks and will have polarization sensitivity in the $\\gamma$-ray energy regime \\cite{next}.  \\begin{figure}[t] \\begin{center} \\includegraphics[angle=-90,width=3.5in]{Fig1.eps} \\vspace*{0.3cm} \\caption{\\label{gondola} Conceptual design of the {HX-POL} balloon payload. The polarization is measured with a Si-CZT detector configuration inside a 2~m high collimator-shield assembly. The collimator is pointed towards a source by rotating the gondola and by changing the elevation.} \\end{center} \\end{figure}  The OSO~8 observations revealed a polarization degree of $\\sim$20\\%  of the 2.6~keV and 5.2~keV X-ray emission from the Crab Nebula, and a polarization angle aligned around 30 degrees oblique to the X-ray jet \\cite{Weis:78}. In the 0.1-1 MeV energy range, the INTEGRAL observations showed a high polarization degree of 46\\%$\\pm$10\\%, and a polarization direction aligned with the orientation of the X-ray jet \\cite{2008Sci...321.1183D}. The result indicates that the X-ray emission comes from the inner jet. The fact that the polarization angle differs from the classic X-ray band to the soft gamma-ray band provides a significant diagnostic on the jet-nebula environment. It may be a general trend that polarization fractions increase with the energy of the X-rays, as the electrons emitting higher-energy X-rays lose their energy faster than the electrons responsible for the lower-energy X-rays; as a consequence, the regions responsible for harder X-rays should be more compact and thus more uniform than the regions responsible for soft X-rays. A more uniform magnetic field would lead to a higher degree of polarization. However, such trends would critically depend on the relative order of the magnetic field on different spatial scales, a salient property for MHD and plasma models. Even at balloon altitudes of $\\sim$40 km, the atmosphere is only transparent for $>$20 keV X-rays. Soft X-ray polarimeters have thus to be space borne. In the hard X-ray energy band, balloon-borne experiments with $\\sim$500 cm$^2$ detection areas can measure polarization fractions of between 10\\% and a few percent for bright, mostly galactic sources with fluxes between 10\\% and 100\\% of the flux from the Crab Nebula. The measurement of the hard X-ray polarization of very short transient events (Gamma Ray Bursts, GRBs) and most extragalactic sources requires larger, most likely, space-borne experiments. \\begin{figure}[t] \\begin{center} \\vspace*{2cm} \\includegraphics[angle=+90,width=3.5in]{Fig2.eps} \\vspace*{2.26cm} \\caption{\\label{SiCZT} ``Exploded view'' of the {HX-POL} detector assembly. A total of 32 Si detectors (each 0.2~cm$\\times$10~cm$\\times$10~cm, active detector area: 81 cm$^2$) are tiled and stacked to make a 1.6 cm thick detector module with an active detection area of 324 cm$^2$. The Si detectors are positioned ``above'' one hundred CZT detectors (each  0.5~cm$\\times$2~cm$\\times$2~cm) with an active detection area of 400 cm$^2$. The ASIC readout of the Si detectors is located at the perimeter of the Si detectors and the readout of the CZT detectors is located below the CZT detectors. The Si-CZT detectors are positioned above a 2~cm$\\times$25~cm$\\times$25~cm CsI active rear shield. } \\end{center} \\end{figure}   In Sect.\\ \\ref{HX-POL}, the concept of the hard X-ray polarimeter HX-POL is presented, a balloon-borne experiment that uses Si and Cadmium Zinc Telluride (CZT) detectors. In Sect.\\ \\ref{science}, the science potential is discussed, and in Sect.\\ \\ref{discussion} the results are summarized. ",
        "conclusions": "\\label{discussion} This paper describes the concept of a balloon-borne hard X-ray polarimeter. The detector module combines a low-Z Compton scatterer (Si) with a high-Z photoeffect absorber (CZT). The detector combination achieves excellent polarization sensitivity in the 50 keV to 400 keV energy band. Based on 3+3 hrs of ON-OFF observations during a one-day balloon flight, the polarization of strong Crab-like sources can be measured for polarization degrees well below 10\\%. On a 15-30 day balloon flight, polarization degrees of a few percent can be measured. The 50-400 keV energy band is well suited for a balloon-borne experiment: atmospheric absorption hinders observations below 20 keV and observations above 400 keV have to deal with very low photon fluxes and require relatively large experiments and/or long integration times, both better suited for space borne experiments. One of the strengths of the Si/CZT detector combination is the excellent energy resolution that makes it possible to measure the energy dependence of the polarization properties. Compared to experiments that use Si detectors only, the Si/CZT combination avoids the neutron backgrounds that plague  experiments exclusively made of low-Z detectors. Compared to experiments that use CZT only, the Si/CZT combination achieves a lower energy threshold. In CZT detectors, photoeffect interactions dominate over Compton interactions up to energies of $\\sim$250 keV; furthermore, $<$100 keV photons are absorbed quickly in CZT, making it necessary to use CZT detectors with very small pixel pitches."
    },
    "0812/0812.1038_arXiv.txt": {
        "abstract": "\\\\ I discuss some theoretical expectations for the synchrotron emission from a relativistic blast-wave interacting with the ambient medium, as a model for GRB afterglows, and compare them with observations. An afterglow flux evolving as a power-law in time, a bright optical flash during and after the burst, and light-curve breaks owing to a tight ejecta collimation are the major predictions that were confirmed observationally, but it should be recognized that light-curve decay indices are not correlated with the spectral slopes (as would be expected), optical flashes are quite rare, and jet-breaks harder to find in Swift X-ray afterglows.  The slowing of the early optical flux decay rate is accompanied by a spectral evolution, indicating that the emission from ejecta (energized by the reverse shock) is dominant in the optical over that from the forward shock (which energizes the ambient medium) only up to 1 ks. However, a long-lived reverse shock is required to account for the slow radio flux decays observed in many afterglows after $\\sim 10$ day.  X-ray light-curve plateaus could be due to variations in the average energy-per-solid-angle of the blast-wave, confirming to two other anticipated features of GRB outflows: energy injection and angular structure. The latter is also the more likely origin of the fast-rises seen in some optical light-curves. To account for the existence of both chromatic and achromatic afterglow light-curve breaks, the overall picture must be even more complex and include a new mechanism that dominates occasionally the emission from the blast-wave: either late internal shocks or scattering (bulk and/or inverse-Compton) of the blast-wave emission by an outflow interior to it. ",
        "introduction": "A relativistic motion of GRB sources was advocated by \\cite{paczynski86,goodman86} from that the energies released exceed by many orders of magnitude the Eddington luminosity for a stellar-mass object, especially if GRBs are at cosmological distances (see also \\cite{shemi90,meszaros92}).  The detection by CGRO/EGRET of photons with energy above 1 MeV during the prompt burst emission (e.g. \\cite{hurley94}) shows that GRB sources are optically thin to such photons. Together with the sub-MeV burst isotropic-equivalent output of $10^{52}-10^{54}$ ergs (e.g. \\cite{bloom01}) and the millisecond burst variability timescale, the condition for optical thickness to high energy photons gives another reason why GRBs must arise from ultra-relativistic sources, moving at Lorentz factor $\\Gamma \\simg 100$ (e.g. \\cite{fenimore93,lithwick01}).  The same conclusion is enforced by the measurement of a relativistic expansion of the radio afterglow source. That expansion was either measured directly, as for GRB 030329 ($z=0.17$), whose size increased at an apparent speed of 5c, indicating a source expanding at $\\Gamma \\siml 6$ at 1--2 months \\cite{taylor04}, or was inferred from the rate at which interstellar scintillation \\cite{goodman97} quenches owing to the increasing source size, as for GRB 970508, whose expansion speed is inferred to be close to $c$ at 1 month \\cite{waxman98}. The adiabatic dynamical evolution of a blast-wave, $\\Gamma^2 M = const$, where $M$ is the mass of the ambient medium, leads to $\\Gamma \\propto t^{-3/8}$ for a homogeneous medium ($t$ being the observer-frame photon arrival time). Then, $\\Gamma (30\\,{\\rm d}) = 2$ extrapolated to the burst time implies $\\Gamma(100\\,{\\rm s}) \\sim 100$. Extrapolating to such early times is justified by that most optical afterglow light-curves display a power-law decay starting after the burst, which sets a lower limit on the source Lorentz factor at that time.  Whether the GRB ejecta are a cold baryonic outflow accelerated by the adiabatic losses of fireball's initial thermal/radiation energy (e.g. \\cite{meszaros93,piran93}), or relativistic pairs formed through magnetic dissipation in a Poynting outflow, as in the electromagnetic model of \\cite{lyutikov06}, their interaction with ambient medium will drive two shocks: a reverse shock crossing the ejecta and a forward-shock sweeping the circumburst medium, as illustrated in Figure \\ref{bw}. Both shocks energize their respective media, accelerate relativistic particles and generate magnetic fields through some plasma-instability related process, such as the two-stream Weibel instability driven by an anisotropic particle distribution function \\cite{medvedev99}. The original magnetic field of the fireball at $10^7$ cm becomes too weak by the time the fireball reaches the $10^{15}-10^{17}$ cm radius (where the burst and afterglow emissions are produced) for the synchrotron emission to account for the sub-MeV burst emission and for the longer-wavelength, ensuing afterglow emission, even if the fireball was initially magnetically dominated \\cite{meszaros93,medvedev99}.  \\begin{figure} \\parbox[h]{7cm}{ \\includegraphics[height=6.4cm,width=7cm]{blast.eps} } \\hspace*{2mm} \\parbox[h]{8cm}{ \\includegraphics[height=6cm,width=8cm]{syic.eps} } \\caption{ \\footnotesize {\\sl Left panel}: particle density distribution in a relativistic blast-wave (moving from left to right) sweeping the ambient medium (of particle density $1\\,{\\rm cm^{-3}}$) located ahead of the forward shock (FS). The ejecta are located behind the contact discontinuity (CD) and are energized by the reverse shock (RS). Solid lines are for lab-frame density, dotted lines show comoving density. Locations are measured in deceleration radii, defined by the first reverse shock crossing the ejecta shell (other reverse shocks can develop later). {\\sl Right panel}: as the blast wave decelerates, the synchrotron and inverse-Compton emissions from the reverse and forward shocks become weaker and softer (assuming constant electron and magnetic fields parameters). For a brief release of the ejecta, the reverse shock is only mildly relativistic, radiates at lower frequencies, and accelerates electrons for a shorter duration (while it crosses the ejecta shell), after which electrons cool adiabatically and their characteristic synchrotron frequency may fall below that of observations. Results shown here are based on 1-dimensional hydrodynamical simulations of the ejecta--medium interaction \\cite{wen97}. } \\label{bw} \\end{figure}  The evolution of the synchrotron and inverse-Compton fluxes produced by the blast-wave at a fixed frequency (i.e. the light-curve) is determined by how the characteristics of the spectrum (break frequencies and peak flux) change with time. Figure \\ref{spek} shows the expected afterglow synchrotron spectrum, whose characteristics depend on the blast-wave radius, number of radiating electrons, their distribution with energy, magnetic field strength, and Lorentz factor.  \\begin{figure} \\parbox[h]{8cm}{ \\includegraphics[height=6cm,width=8cm]{spek.eps} } \\hspace*{2mm} \\parbox[h]{8cm}{ \\vspace*{-5mm} \\footnotesize Broadband spectrum of the synchrotron emission showing its three break frequencies: self-absorption (in radio), injection ($\\nu_i$, corresponding to the energy of the electrons currently accelerated by the shock), and cooling ($\\nu_c$, for electrons whose radiative cooling timescale is equal to the dynamical timescale). The energy distribution of the shock-accelerated electrons is assumed to be a power-law above the injection break -- $dN_e/d\\epsilon \\propto \\epsilon^{-p}$ -- which leads to a power-law afterglow spectrum ($F_\\nu \\propto \\nu^{-\\beta}$) of slope $\\beta = (p-1)/2$ between $\\nu_i$ and $\\nu_c$ and $\\beta = p/2$ above $\\nu_c$. For a forward-shock afterglow origin, the three break frequencies and the peak flux determined from a \"snap-shot\" afterglow spectrum provide four constraints for four unknown parameters: ejecta kinetic energy (per solid angle), ambient density, and fractional energies of electrons and magnetic field. Thus, afterglow spectra extending down to radio frequencies can be used to determine afterglow fundamental physical parameters (e.g. \\cite{wijers99,granot99}). } \\caption{} \\label{spek} \\end{figure}  If the typical electron energy and magnetic field energy correspond to some fixed fraction of the post-shock energy, or if they start from such a fixed fraction and then evolve adiabatically (as for adiabatically colling ejecta), then the afterglow light-curve depends on (1) the evolution of the blast-wave Lorentz factor, the blast-wave radius being $R \\simeq \\Gamma^2 ct$ (with $t$ the photon-arrival time measured since burst trigger) (2) the spectrum of the blast-wave emission (i.e. the distribution of electrons with energy), and, in the case of the reverse-shock, (3) the evolution of the incoming mass.  The power-law deceleration of the blast-wave ($\\Gamma \\propto r^{-(3-s)/2} \\propto t^{-(3-s)/(8-4s)}$ for $s<3$, where $n \\propto r^{-s}$ is the radial stratification of the ambient medium density) and the power-law afterglow spectrum ($F_\\nu \\propto \\nu^{-\\beta}$) are two factors which lead to a power-law afterglow light-curve ($F_\\nu \\propto t^{-\\alpha}$), with the decay index $\\alpha$ being a linear function of the spectral slope $\\beta$. These are the only two factors at work for the forward-shock emission and the ejecta emission during the adiabatic cooling phase (which starts when the reverse shock has crossed the ejecta shell), the two models that yield power-law afterglow light-curves in the most simple and natural way.  In contrast, for the reverse-shock emission (i.e. the ejecta emission while the shock exists), the light-curve depends also on the radial distribution of ejecta mass and of their Lorentz factor, thus the observed power-law light-curves require additional properties to be satisfied by the relativistic ejecta. Such properties seem {\\sl ad-hoc} when it comes to explaining single power-law afterglows whose flux displays an unchanged decay over 2--4 decades in time (such as the X-ray afterglows of GRBs 050801, 050820A, 06011B, 060210, 060418, 061007), but they also provide the flexibility required to account for the prevalent X-ray afterglow light-curves that exhibit one or more breaks. ",
        "conclusions": "The temporal and spectral properties of GRB afterglows are, in general, consistent with those predicted for the synchrotron emission from the blast-wave produced when highly relativistic ejecta (initial Lorentz factor above 100) interact with the ambient medium.  As expected from shocks accelerating particles with a power-law distribution with energy, power-laws are observed in the optical and X-ray afterglow continua. A spectral softening (i.e. decrease of peak frequency) is expected owing to the blast-wave deceleration and is observed in radio afterglows spectra and in the behaviour of afterglow radio light-curves, which rise slowly until the synchrotron spectrum peak reaches the radio domain and fall-off afterward.  Rising afterglow light-curves are seen at early times in the optical, and are consistent with the pre-deceleration emission from the forward shock, although a structured outflow with a hot-spot that gradually becomes visible to the observer is also possible, in which case the rising afterglow emission could also be explained with the reverse shock.  Much more often, afterglow light-curves display power-law decays of indices that are not correlated with the spectral slopes, as would be expected for the external-shock model. That inconsistency could be due to a substantial intrinsic scatter in decay indices and spectral slopes, owing to more than one variant of the blast-wave model occurring in GRB afterglows, combined with a small range of those indices being realized.  Bright optical flashes accompanying the burst emission were predicted to arise from the reverse shock, owing to the larger number of ejecta electrons than in the forward shock. Fast-falling optical light-curves have been observed at 100--1000s in two afterglows (991023 and 021211), followed by a slower of the decay, which was taken as evidence for the reverse-shock emission dominating the early afterglow, although spectral information was not available to test that hypothesis. More recently, the early optical spectral slopes were measured for two afterglows (061126 and 080319B), the decay of the former being consistent with that from adiabatically-cooling ejecta, while the later is faster than expected and consistent with it being the large-angle emission released at an earlier time.  For the above two optical afterglows with spectral information at early times, the slowing of the optical flux decay at 1 ks is accompanied by a spectral evolution, which indicates the transition from one mechanism to another. Most naturally, that is the transition from ejecta emission to forward-shock emission, with the reverse-shock emission being relevant for the optical afterglow only during its early phase.  Evidence for a reverse-shock emission is also provided by the slow radio flux decays observed after 10 day in several afterglows. Adiabatically cooling ejecta would yield a decay faster than observed, particularly if the synchrotron cooling frequency were to fall below the radio, thus a reverse-shock accelerating ejecta electrons is required to operate for days and produce a radio emission decaying much slower than the optical at the same time, the latter being attributed to the forward shock. Together with the above conclusion regarding the contribution of the reverse shock to the early optical afterglow, this suggests that the reverse shock is the main afterglow source for a duration that decreases with observing frequency, perhaps never being dominant in the X-rays and having no connection with the chromatic X-ray light-curve breaks seen at $\\sim 10$ ks in most afterglows.  That the X-ray-to-optical flux-ratio is larger (by a factor 5) for afterglows with chromatic X-ray light-curve breaks than for those with coupled optical and X-ray light-curves (i.e. with achromatic breaks or similar power-law decays), indicates the existence of a different mechanism producing the X-ray emission of afterglows with chromatic X-ray breaks, coupled light-curves resulting when the emission from that novel mechanism is negligible. Long-lived internal shocks or scattering of the blast-wave emission by an outflow located behind it are two possibilities that could explain chromatic X-ray breaks, as well as the flares seen in many X-ray light-curves. The latter mechanism is also related to energy injection in the blast-wave, whose cessation accounts naturally for achromatic light-curve breaks. Thus, in the scattering model, the diversity of optical and X-ray light-curve relative behaviours is attributed to the interplay between the scattered and direct blast-wave emissions, combined with the changing dynamics of the blast-wave produced when the scattering outflow brings into the shock more energy than already existing.  On energetic grounds, GRB outflows should be collimated into jets narrower than 10 degrees. The steepening of the afterglow flux decay when the jet boundary becomes visible to the observer was another major prediction confirmed by observations. Just as for the pre jet-break phase, more than one jet model is required to account for the measured decay indices (given the observed spectral slopes). Swift X-ray afterglows display light-curve jet-breaks less often than pre-Swift optical afterglows, which could be due to that the former afterglows (being dimmer) arise from wider jets whose jet-breaks occur later and could, thus, be missed more often."
    },
    "0812/0812.4337_arXiv.txt": {
        "abstract": "We know that the galactic magnetic field possesses a random component in addition to the mean uniform component, with comparable strength of the two components. This random component is considered to play important roles in the evolution of the interstellar medium (ISM). In this work we present numerical simulations associated with the interaction of the supersonic flows located at high latitude in our Galaxy (High Velocity Clouds, HVC) with the magnetized galactic ISM in order to study the effect that produces a random magnetic field in the evolution of this objects. ",
        "introduction": "Numerical simulations of the evolution of HVC collisions with the Milky Way have been performed for more than two decades by different authors. The details of resulting supersonic flows depend of on the model assumptions, and the intensity and initial configuration of the magnetic field. Santill\\'an et al. (1999) made models that illustrate the effects of magnetic pressure, and differentiate them from those due to magnetic tension. The evolution of the interaction of a HVC with a magnetized interstellar medium, is studied by setting a random magnetic field that satisfies the divergence--free constraint ($\\nabla~\\bullet$ $\\bf B$ =0) at all times. To mimic the average galactic magnetic field, that is oriented parallel to the disk, the horizontal component dominates over the vertical component. ",
        "conclusions": ""
    },
    "0812/0812.3451_arXiv.txt": {
        "abstract": "We investigate shattering and coagulation of dust grains in turbulent interstellar medium (ISM). The typical velocity of dust grain as a function of  grain size has been calculated for various ISM phases based on a theory of grain dynamics in compressible magnetohydrodynamic turbulence. In this paper, we develop a scheme of grain shattering and coagulation and apply it to turbulent ISM by using the grain velocities predicted by the above turbulence theory. Since large grains tend to acquire large velocity dispersions as shown by earlier studies, large grains tend to be shattered. Large shattering effects are indeed seen in warm ionized medium (WIM) within a few Myr for grains with radius $a\\ga 10^{-6}$ cm. We also show that shattering in warm neutral medium (WNM) can limit the largest grain size in ISM ($a\\sim 2\\times 10^{-5}~\\mathrm{cm}$). On the other hand, coagulation tends to modify small grains since it only occurs when the grain velocity is small enough. Coagulation significantly modifies the grain size distribution in dense clouds (DC), where a large fraction of the grains with $a<10^{-6}$ cm coagulate in 10 Myr. In fact, the correlation among $R_V$, the carbon bump strength, and the ultraviolet slope in the observed Milky Way extinction curves can be explained by the coagulation in DC. It is possible that the grain size distribution in the Milky Way is determined by a combination of all the above effects of shattering and coagulation. Considering that shattering and coagulation in turbulence are effective if dust-to-gas ratio is typically more than $\\sim 1/10$ of the Galactic value, the regulation mechanism of grain size distribution should be different between metal-poor and metal-rich environments. ",
        "introduction": "Dust grains absorb stellar ultraviolet (UV)--optical light and reprocess it into far-infrared (FIR), thereby affecting the energetics of interstellar medium (ISM) \\citep[e.g.][]{hirashita02}. For this process concerning the interaction between grains and radiation, the optical properties of grains are important. The grain optical properties are determined not only by dust species but also by the grain size \\citep[e.g.][]{draine84}. In fact, the grain species and size distribution are derived from the observed interstellar extinction curve \\citep[][hereafter MRN]{mathis77}. These grain properties also affect the FIR spectrum of dust emission \\citep[e.g.][]{takeuchi05}.  The grain size distribution is known to be affected by various processes in the interstellar space. Grains are supplied from stars at their death \\citep{gehrz89} with a certain grain size distribution \\citep{dominik89,todini01,nozawa03,bianchi07,nozawa07}. These grains dispersed from stars are processed in the interstellar space. They grow in dense environments such as molecular clouds by grain-grain coagulation {\\citep*{chokshi93}} and accretion of heavy elements \\citep{spitzer78}. They are also destroyed by supernova shocks by gas-grain sputtering and by grain-grain collision (shattering) \\citep{dwek80,tielens94,borkowski95}. In particular,  {\\citet[hereafter JTH96]{jones96}} show that the grain size distribution can be significantly modified by shattering. Such a change of grain size distribution would significantly modify the extinction curve and the infrared spectral energy distribution of dust emission.  Potential importance of shattering and coagulation in ISM has often been pointed out, since relative velocity between grains is naturally expected if ISM is turbulent \\citep{kusaka70,volk80,draine85,ossenkopf93,lazarian02}. Because turbulence is ubiquitous in ISM {\\citep[e.g.][]{mckee07}}, the relative grain motion induced by turbulence is of general importance in the grain evolution in ISM. Moreover, the ISM is known to be magnetized \\citep{arons75}. Thus, dust motion in magnetohydrodynamic (MHD) turbulence should be considered. \\citet[hereafter YLD04]{yan04} calculate the relative grain velocity in compressible MHD turbulence, taking into account gas drag (hydrodrag) and gyroresonance. The basis of their theory can be seen in \\citet{lazarian02} and \\citet{yan03}. According to their results, grains can be accelerated to a velocity larger than a few km s$^{-1}$ in diffuse medium by gyroresonance. At such a high velocity, grains can be shattered (JTH96). On the other hand, small grains can obtain velocities small enough for coagulation to occur, especially in dense medium (YLD04).  The size distribution of grains processed in various ISM phases was investigated by \\citet{odonnell97}. Their models incorporate coagulation by turbulent motion in clouds, accretion of gas-phase metals onto grains, and shattering and sputtering in interstellar shocks. They also take into account the phase exchange of the ISM. They show that both the observed extinction curve and the observed depletion of refractory elements are reproduced by considering phase exchange among diffuse clouds, warm neutral intercloud gas, and molecular clouds. In their models, shattering mainly occurs in ISM shocks. However, \\citet{yan03} shows that grains can be accelerated to velocities large enough for shattering in MHD turbulence. Thus, it is worth focusing on the effect of turbulence on the grain size distribution. In addition, the treatment of shattering can be revised by including the framework of JTH96 to take into account the velocity dependence of fragment production. It is also useful to compare the difference between the size distribution modified by supernova shocks as treated by JTH96 and that by interstellar turbulence as examined in this paper.  The aim of this paper is to examine quantitatively whether or not shattering and coagulation in turbulent ISM really modify the grain size distribution. We focus on the effects of MHD turbulence as treated in \\citet{yan03}. The other types of dust processing such as shattering and sputtering in supernova shocks and dust condensation in stellar mass loss are not treated in this paper, in order to make our discussions concentrated and clear. For an observational comparison, we adopt the Milky Way extinction curve following \\citet{odonnell97}.  This paper is organized as follows. First, in Section~\\ref{sec:model}, we describe the model adopted to treat shattering and coagulation in turbulent ISM. Then, in Section~\\ref{sec:result}, we overview the results. In Section~\\ref{sec:discussion}, we discuss our results, focusing on the regulation mechanism of the grain size distribution in turbulent ISM. Section~\\ref{sec:sum} is devoted to the summary. ",
        "conclusions": "\\label{sec:discussion}  The important features found in the previous section can be summarized as follows.  (i) The largest effect of shattering is seen in WIM, where grains with $a\\ga 10^{-6}$ cm are efficiently shattered.  (ii) The largest effect of coagulation is observed in DC around $a\\la\\mbox{a few}\\times 10^{-6}$ cm.  (iii) Grains with $a\\ga\\mbox{a few}\\times 10^{-5}$ cm can be shattered in WNM and graphite grains with $a\\ga 10^{-5}$ cm may be quite efficiently destroyed in CNM. These destructions could affect the upper limit of grain size in ISM.  (iv) On the other hand, the lower limit of grain size may be determined by coagulation in DC and MC.  The features (i) and (ii) indicate that once grains are included in WIM or DC, the grain size distribution is significantly modified. It is interesting to note that the shattered grains in (i) and the coagulated grains in (ii) have a similar size. Thus, it is worth investigating if the MRN size distribution can be realized as a balance between (i) and (ii). This point is investigated in Section \\ref{subsec:equilibrium}.  Regarding the feature (iii), as mentioned in Section \\ref{subsubsec:cnm}, the results in CNM are sensitive to the grain velocities and the shattering thresholds. We leave more careful treatment of shattering in CNM for future work. Shattering in WNM is interesting to investigate, since turbulence in WNM accelerates grains with $a\\ga \\mbox{a few}\\times 10^{-5}$ cm much above the threshold velocity for shattering. This size really matches the upper limit of the grain size distribution (MRN). This point is investigated in Section \\ref{subsec:limit}.  The issue (iv) has already been investigated and discussed in Section \\ref{subsubsec:dc}.  \\subsection{Grain size distributions in diffuse-dense phase exchange}\\label{subsec:equilibrium}  In ISM, mass is exchanged between various phases \\citep{mckee77,ikeuchi88}. Thus, it is important to investigate the effects of multi-phase ISM on the evolution of grain size distribution, although the main aim of this paper is to examine the dust processing in individual phases. The largest shattering and coagulation effects are seen in WIM and DC, respectively, and we here examine the dust processing in both WIM and DC to address a possible importance of multi-phase ISM in determining the grain size distribution. For DC, we adopt DC2 because of the success in explaining the trend of $R_V$ in terms of the UV slope and the 2175~\\AA\\ bump (Section \\ref{subsubsec:dc}).  We start from the size distribution of grains processed in DC2 for 10 Myr. Then, we apply the condition of WIM. In Fig.\\ \\ref{fig:dc_wim}, we show the results at $t=3~\\mathrm{Myr}$ and 5 Myr in WIM. Around 5 Myr, the number of small grains is recovered to the level of the MRN distribution. In other words, if grains pass their lifetimes in WIM more than in DC, the grains are shattered too much to be consistent with the MRN distribution. This implies a short lifetime of WIM. Combining this short lifetimes of WIM with a theoretically implied timescale for the phase exchange ($\\mbox{a few}\\times 10^7$--$10^8$ yr; Section \\ref{subsec:timescale}), we obtain a picture that a large fraction of warm medium is in a neutral form and a certain small fraction is ionized. It is interesting to point out that such a short timescale is consistent with the recombination timescale as mentioned in Section \\ref{subsubsec:wim}.  The corresponding extinction curves are shown in Fig.\\ \\ref{fig:extinc_exchange}. The Milky Way extinction curve is indeed recovered by the phase exchange. This demonstrates that it is really possible to reproduce the Milky Way extinction curve by considering dust grains processed in multiphase medium.  The above phase exchange model is too simple, and the realistic ISM has more continuous density distribution and more complicated structure of turbulence \\citep{wada01}. Such complexity should tend to eliminate the specific features such as accumulation of grains around $a\\sim \\mbox{a few}\\times 10^{-6}$ cm in DC and selective grain destruction at $a\\sim 10^{-6}$ cm in WIM. Thus, we expect that the grain size distribution becomes smoother in realistic ISM than we calculate in this paper.  \\begin{figure} \\begin{center} \\includegraphics[width=7.5cm]{ext_dc_wim.eps} \\end{center} \\caption{The extinction curves calculated for the size distributions in Fig.\\ \\ref{fig:dc_wim}. The solid line represents the initial size distribution (10 Myr in DC2), and the dashed and dot-dashed lines show the extinction curves at $t=3$ Myr and 5 Myr in WIM, respectively. \\label{fig:extinc_exchange}} \\end{figure}  \\begin{figure*} \\begin{center} \\includegraphics[width=6.5cm]{sil_wnm_large.eps} \\includegraphics[width=6.5cm]{gra_wnm_large.eps} \\end{center} \\caption{The grain size distributions of (a) silicate and (b) graphite after $t=50$ Myr (solid line) and $t=100$ Myr (dashed line) in WIM for the MRN size distribution extending up to $a=10^{-4}$ cm as the initial condition (dotted line). \\label{fig:wnm_large}} \\end{figure*}  \\subsection{Upper and lower limits of grain size} \\label{subsec:limit}  According to MRN, the upper grain radius is $\\sim 0.25~\\mu$m (see also \\citealt{kim94}). Coagulation has negligible influence on grains larger than $a\\sim 0.2~\\mu$m both for silicate and for graphite because they generally obtain larger velocity than the coagulation thresholds. Thus, if there is no grain with $a\\ga 0.2~\\mu$m initially, it is not possible to make such large grains by coagulation in ISM.  Even if grains larger than $a\\sim 0.2~\\mu$m form by condensation in stellar ejecta, shattering could destroy such large grains. \\citet{nozawa03} show that silicon grains with $a>0.2~\\mu$m form in Type II supernovae. In the outflows from evolved late-type stars, the grain radius is expected to become of order $\\sim 0.1~\\mu$m \\citep{gail99}, and grains with $a\\ga 0.2~\\mu$m may have a chance to form. It is interesting to note that grains with $a\\ga 0.2$--0.3 $\\mu$m are accelerated above the shattering threshold in WNM. Thus, shattering in WNM may play a central role in determining the upper limit of the grain size in ISM.  In order to examine whether or not shattering in WNM really plays a role in determining the upper limit of the grain size, we perform a test by adopting an initial grain size distribution extending up to $a=1~\\mu$m with the total mass of grains conserved. Then the evolution of the grain size distribution is calculated by applying the conditions in WNM. Fig.\\ \\ref{fig:wnm_large} shows the results. We observe that the grains with $a\\ga 0.2$--0.3 $\\mu$m are significantly shattered in 50 Myr. Thus, shattering in WNM is a strong candidate for the determining mechanism of the upper limit of grain size.  In Fig.\\ \\ref{fig:extinc_large}, we show the corresponding extinction curves. The initial extinction curve is significantly lower than the observed one because large grains tend to have low mass absorption coefficients. However, after 50 Myr, the level of the extinction is already consistent with the Milky Way curve. This means that shattering of large grains in WNM is efficient enough to reproduce the upper grain size consistent with the observed Milky Way extinction curve.  \\begin{figure} \\begin{center} \\includegraphics[width=7.5cm]{ext_wnm_large.eps} \\end{center} \\caption{The extinction curves calculated for the size distributions in Fig.\\ \\ref{fig:wnm_large}. The dotted line represents the initial extinction curve (MRN size distribution extending up to $a=10^{-4}$ cm). The solid and dashed lines show the extinction curves of grains at $t=50$ Myr and $t=100$ Myr in WNM, respectively. \\label{fig:extinc_large}} \\end{figure}  \\begin{figure*} \\begin{center} \\includegraphics[width=6.5cm]{sil_wim_metallicity.eps} \\includegraphics[width=6.5cm]{gra_wim_metallicity.eps} \\end{center} \\caption{The grain size distributions of (a) silicate and (b) graphite at $t=10$ Myr in WIM. The solid, dashed, dot-dashed lines show $Z=1$, 1/3 and 1/10~$\\mathrm{Z}_\\odot$, respectively. The dust abundance (vertical axis) of the dashed and dot-dashed lines are multiplied by 3 and 10, respectively, to offset the low dust abundances. \\label{fig:wim_metallicity}} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[width=6.5cm]{sil_dc_metallicity.eps} \\includegraphics[width=6.5cm]{gra_dc_metallicity.eps} \\end{center} \\caption{Same as Fig.\\ \\ref{fig:wim_metallicity} but in DC2. \\label{fig:dc_metallicity}} \\end{figure*}  \\subsection{In the context of galaxy evolution}  The efficiencies of shattering and coagulation are affected by the grain abundance. This indicates that metal-poor galaxies, which are generally poor in dust content \\citep{issa90}, have different grain size distributions. Here we examine the metallicity dependence of shattering and coagulation. We assume that the dust-to-gas ratio is proportional to the metallicity $Z$; that is, we adopt ${\\cal R}=4.0\\times 10^{-3}Z/\\mathrm{Z}_\\odot$ and $3.4\\times 10^{-3}Z/\\mathrm{Z}_\\odot$ for silicate and graphite, respectively, in equation (\\ref{eq:norm_grain}). In other words, the dust density in the ISM is proportional to the metallicity, and we expect that the effects of shattering and coagulation become weak as the metallicity decreases. {The turbulence model and the grain velocities are not changed, which means that we implicitly assume that the parameters listed in Table \\ref{tab:yan} are fixed.}  We test WIM and DC, where shattering and coagulation, respectively, are the most efficient among the various phases. In Fig.~\\ref{fig:wim_metallicity}, we show the grain size distributions in WIM at $t=10$ Myr. We apply a longer timescale than adopted in the other part of this paper to enhance the effect of shattering. We observe that the shattering effect is significantly reduced at $1/10~\\mathrm{Z}_\\odot$. The same is true for coagulation in DC as shown in Fig.\\ \\ref{fig:dc_metallicity}, where we adopt DC2 because of the success in reproducing the trend of $R_V$ in terms of the UV slope and the 2175 \\AA\\ bump (Section \\ref{subsubsec:dc}). Thus, as the metallicity decreases, the relative importance of processing by interstellar turbulence becomes minor in determining the grain size distribution. This indicates that the initial grain size distribution at the grain formation in stellar ejecta is relatively preserved in metal-poor galaxies (typically $Z<1/10~\\mathrm{Z}_\\odot$), although we should keep in mind that there are other processes, such as interstellar shocks by supernovae, which could modify the grain size distribution in any metallicity.  \\subsection{Toward the grain evolution in protoplanetary discs}  {The condition of turbulence in the circumstellar discs is still unclear. Let us consider protoplanetary discs. According to \\citet{nomura06}, turbulence is very weak ($\\delta V\\sim 0.01\\mbox{--}0.1c_\\mathrm{s}$, where $c_\\mathrm{s}$ is the sound speed). The acceleration by turbulence will be marginal in this case, and grain motions are more likely to be Brownian. As a result coagulation is at least as efficient as in DC.  As shown in Figs.\\ \\ref{fig:sil} and \\ref{fig:gra}, small grains with $a\\la 10^{-6}$ cm are strongly depleted in DC because of coagulation. Thus, we can justify that the grain size distribution in protoplanetary discs is biased to radii $\\ga 10^{-6}$ cm. Moreover, because grain velocities are expected to be lower than the coagulation threshold even at $a>10^{-6}$ cm, grains grow further.}  {As shown by \\citet{sano00}, the grain size in protoplanetary discs is important in determining the unstable regions for magnetorotational instability, which induces MHD turbulence \\citep{balbus98}. Consequently the grain size distribution is further affected by the presence/absence of the turbulent motion determined by the instability/stability condition. The coupling between turbulence and grain size is interesting to investigate as a future work.}"
    },
    "0812/0812.3673_arXiv.txt": {
        "abstract": "We present radio continuum polarimetry observations of the nearby edge-on galaxy NGC\\,253 which possesses a very bright radio halo. Using the vertical synchrotron emission profiles and the lifetimes of cosmic-ray electrons, we determined the cosmic-ray bulk speed as $300\\pm30\\,\\rm km\\,s^{-1}$, indicating the presence of a galactic wind in this galaxy. The large-scale magnetic field was decomposed into a toroidal axisymmetric component in the disk and a poloidal component in the halo. The poloidal component shows a prominent X-shaped magnetic field structure centered on the nucleus, similar to the magnetic field observed in other edge-on galaxies. Faraday rotation measures indicate that the poloidal field has an odd parity (antisymmetric). NGC\\,253 offers the possibility to compare the magnetic field structure with models of galactic dynamos and/or galactic wind flows. ",
        "introduction": "Gaseous halos are a common property of star forming galaxies. The abundance of hot coronal gas in the halo requires heating from the star forming disk and a transport of energy from the disk into the halo. This energy transport has been discussed as galactic fountains or chimneys \\cite[(Field et al.\\ 1969, Norman \\& Ikeuchi 1989)]{field_69a, norman_89a}. The relativistic cosmic-ray gas has an adiabatic index of 4/3 compared to 5/3 of the (single atomic) hot gas; this results in a larger pressure scaleheight of the cosmic-ray gas. The magnetic field has an even larger pressure scaleheight \\cite[(Beck 2009)]{beck_09}. The pressure contributions of the cosmic rays and the magnetic field are thus dominating in the halo and can accelerate the outflow. This led to the model of a cosmic-ray driven wind that causes acceleration of the flow accelerates further away from the disk \\cite[(Breitschwerdt et al.\\ 1991)]{breitschwerdt_91a}.  The existence of radio halos around many nearby galaxies shows the presence of cosmic rays and magnetic fields. Cosmic rays spiral around magnetic field lines and follow them from the disk into the halo. Hence, a vertical component of the large-scale magnetic field can enhance the vertical cosmic-ray transport. How to generate vertical fields is still an open question. \\cite[Parker (1992)]{parker_92} suggested that field lines start to overlap and reconnect to form ``open'' field lines that connect the disk with the halo. Another possibility is the mean field $\\alpha\\Omega$-dynamo, which creates dipolar or quadrupolar field configurations.  Furthermore, the galactic wind may shape the magnetic field and align with it. This may particularly hold for \\emph{superwinds}, which originate in a region of very active star formation (starburst).  As shown by X-ray observations, NGC\\,253 has a superwind and is an ideal object to study the interaction between the superwind and the magnetic field. There is a nuclear plume of H$\\alpha$ and X-ray emitting gas \\cite[(Strickland et al.\\ 2000, Bauer et al.\\ 2007)]{strickland_00a,bauer_07a} where the outflow velocity was directly measured by spectroscopy as $390\\,\\rm km\\,s^{-1}$ \\cite[(Schulz \\& Wegner 1992)]{bauer_07a,schulz_92a}. The connection between the heated gas in the halo, which extends far from the disk, and the the superwind is likely as suggested by numerical simulations of a centrally driven wind \\cite[(Suchkov et al.\\ 1994)]{suchkov_94a}. A detailed study of this galaxy is possible because it is nearby ($D=3.94\\,\\rm Mpc$, \\cite[Karachentsev et al.\\ 2003]{karachentsev_03a}, $30''=600\\,\\rm pc$). With an inclination angle of $78^\\circ$ the galaxy's disk is only mildly edge-on. But it contains a very bright nucleus which complicates the data reduction of radio continuum observations. ",
        "conclusions": "NGC\\,253 possesses a bright radio halo. The scaleheight of the synchrotron emission depends on the lifetime of the cosmic-ray electrons. This requires a vertical cosmic-ray transport from the disk into the halo. The disk wind has an average velocity of $300\\pm30\\,\\rm km\\,s^{-1}$ and is surprisingly constant over the extent of the disk. The disk wind transports material from the disk into the halo. If this material is the origin of the luminous, heated gas in the halo, it can explain the asymmetry between the northeastern and southwestern halo, as the cosmic-ray transport is much more efficient in the convective northeastern halo than in the diffusive southwestern halo.  In the disk the magnetic field is parallel to the disk. A model of an axisymmetric spiral disk field can explain the observed asymmetries in polarization by a geometrical effect. After subtracting the disk field model from the observations we obtain the poloidal magnetic field in the halo. This field has a prominent X-shaped structure that is also seen in several other edge-on galaxies. For the first time the parity of the large-scale field in any galaxy could be determined, which is odd (antisymmetric) in NGC\\,253. The X-shape may be explained by interaction of the halo gas transported by the disk wind with the superwind from the starburst center. However, several galaxies without a superwind also show X-shaped halo fields \\cite[(Krause 2008)]{krause_08a}, which needs to be further investigated. \\begin{figure}[tb] \\begin{minipage}[b]{0.66\\textwidth} \\resizebox{\\hsize}{!}{\\includegraphics[bb=5 60 535 562,clip=]{heesenfig5.pdf}} \\end{minipage} \\hfill \\begin{minipage}[b]{0.33\\textwidth} \\caption{Polarized emission and magnetic field orientation of the poloidal magnetic field overlaid onto diffuse X-ray emission. Contours and vectors as in Fig.\\,\\ref{fig:n253cm6ve_piaf_b30}. The X-ray map is from XMM observations in the energy band $0.5-1.0\\,\\rm keV$ (\\cite[Bauer et al.\\ 2008]{bauer_08a}).} \\label{n253_pol_XMM2_col} \\end{minipage} \\end{figure}"
    },
    "0812/0812.4579_arXiv.txt": {
        "abstract": "% Astrophysical fluids are generically turbulent, which means that frozen-in magnetic fields are, at least, weakly stochastic. Therefore realistic studies of astrophysical magnetic reconnection should include the effects of stochastic magnetic field. In the paper we discuss and test numerically the Lazarian \\& Vishniac (1999) model of magnetic field reconnection of weakly stochastic fields. The turbulence in the model is assumed to be subAlfvenic, with the magnetic field only slightly perturbed. The model predicts that the degree of magnetic field stochasticity controls the reconnection rate and that the reconnection can be fast independently on the presence or absence of anomalous plasma effects. For testing of the model we use 3D MHD simulations. To measure the reconnection rate we employ both the inflow of magnetic flux and a more sophisticated measure that we introduce in the paper. Both measures of reconnection provide consistent results. Our testing successfully reproduces the dependences predicted by the model, including the variations of the reconnection speed with the variations of  the injection scale of turbulence driving as well as the intensity of driving. We conclude that, while anomalous and Hall-MHD effects in particular circumstances may be important for the initiation of reconnection, the generic astrophysical reconnection is fast due to turbulence, irrespectively of the microphysical plasma effects involved. This conclusion justifies  numerical modeling of many astrophysical environments, e.g. interstellar medium, for which plasma-effect-based collisionless reconnection is not applicable. ",
        "introduction": "%   Magnetic reconnection is a fundamental problem of astrophysical MHD (see e.g., Priest \\& Forbes 2000 and references therein). Indeed, while the condition of magnetic field being frozen is well satisfied in the bulk of astrophysical plasma, one should wonder what happens when magnetic flux tubes try to push through each other. This is a very important question. For instance, the commonly employed {\\it dynamo} theory allows a large-scale magnetic field to grow exponentially at the expense of small-scale turbulent energy. However, it has been known for some time that the back reaction of the field can suppress magnetic diffusion and defeat the dynamo. Indeed, if a parcel of fluid is threaded by a field loop in equipartition with the turbulence, the loop acts like a rubber band returning the parcel to its origin. As a result, small-scale magnetic fields accumulate most of the energy, while the mean magnetic field saturates far below equipartition with the surrounding fluid. To enable free motions of the fluid parcels, as required by the mean field dynamo, there should be a mechanism for cutting the ``rubber bands''. This mechanism is magnetic reconnection. Unfortunately, if reconnection happens at the rate allowed by generally accepted Sweet-Parker model (Parker 1957, Sweet 1958), it is far too slow. Turbulence would cause many magnetic reversals per parsec within the interstellar medium. Observations, on the contrary, show that magnetic field is coherent over the scales of hundreds of parsecs. This fact, as well as direct observations of Solar flares (which are generally accepted to be fed by reconnection), suggest that the rate of reconnection is many orders of magnitude more rapid than the Sweet-Parker theory suggests.  In general, if the reconnection were slow, the change of magnetic topology did not happen and the interacting flux tubes would create knots, redistributing magnetic energy to small scales. This is the expected outcome for the Sweet-Parker reconnection, which is extremely slow and inadequate for nearly all astrophysical situations.  Petscheck (1964) reconnection is an attempt to remedy the problem assuming that magnetic flux tubes get into contact over areas determined by microphysics and the reconnection regions open up. So far, the feasibility of Petscheck reconnection has been demonstrated in very restricted circumstances (see Shay \\& Drake 1998), which, for instance, exclude interstellar medium and other important astrophysical environments like stars and many types of accretion disks. One of the sources of these restrictions arises from the requirement that the medium should be collisionless in a very special sense, namely, that the extension of the current sheet should not exceed several dozens of electron mean free paths (see Yamada et al. 2006).  Lazarian \\& Vishniac (1999, henceforth LV99) proposed a model that naturally generalizes Sweet-Parker reconnection scheme for weakly turbulent magnetic fields. LV99 consider the case in which there exists a large scale, well-ordered magnetic field, of the kind that is normally used as a starting point for discussions of reconnection\\footnote{The LV99 model has nothing to do with the concept of turbulent magnetic diffusivity, which by some authors to justify the mean field dynamo. The latter concept requires bending of magnetic fields on very small scales, which is prohibited on energetic grounds for any dynamically important field.} (see Fig.~\\ref{recon}). The difference with the Sweet-Parker reconnection arises both  from the fact that in 3D generic configuration (note that 3D is essential and the effect of fast reconnection is absent in 2D) many independent patches of magnetic field get into contact and undergo reconnection. In addition, the outflow of plasma and shared magnetic flux happens not over a microscopically narrow region determined by Ohmic diffusion, but through a substantially wider region determined by field wandering. LV99 showed that magnetic reconnection gets fast when these two effects are taken into account. In fact, the LV99 reconnection rate is \\begin{equation} V_{rec} = V_A \\min \\left[ \\left( \\frac{L}{l} \\right)^{1/2}, \\left( \\frac{l}{L} \\right)^{1/2} \\right] \\left( \\frac{V_l}{V_A} \\right)^{1/2}, \\label{eq:constraint} \\end{equation} where $V_A$ is the Alfv\\'en speed, $L$ is the size of the system, $l$ is the injection scale, and $V_l$ is the velocity amplitude at the injection scale. In this relation, the reconnection speed is determined by the characteristics of turbulence, namely, its strength and injection scale. Most importantly, this rate is determined by magnetic field wandering and contains no explicit dependence on the Ohmic or anomalous resistivity. For isotropic injection of energy when the injection velocity is less than $V_A$ the injection power $P_{inj}$ is proportional to $V_l^4$ (see LV99), which means that Eq.~(ref{eq:constraint}) predicts $V_{rec}\\sim P_{inj}^{1/2} l^{1/2}$, assuming that  $l<L$. This is the dependence that we are testing in the paper.  \\begin{figure}[!t] \\includegraphics[width=0.7\\columnwidth]{ethan.eps} \\caption{{\\it Upper plot}: Sweet-Parker model of reconnection. The outflow is limited by a thin slot $\\Delta$, which is determined by Ohmic diffusivity. The other scale is an astrophysical scale $L\\gg \\Delta$. {\\it Middle plot}: Reconnection of weakly stochastic magnetic field according to LV99. The model that accounts for the stochasticity of magnetic field lines. The outflow is limited by the diffusion of magnetic field lines, which depends on field line stochasticity. {\\it Low plot}: An individual small scale reconnection region. The reconnection over small patches of magnetic field determines the local reconnection rate. The global reconnection rate is substantially larger as many independent patches come together.} \\label{recon} \\end{figure}   Within this short paper we do not have space to discuss in detail neither LV99 model nor other approaches to reconnection attempted through years by various reseachers. More on these issues can be found in the paper by \\citet{lazarian08}, where also the implications of the LV99 model of reconnection are also discussed. We instead concentrate here on numerical testing of the model. The initial numerical testing of LV99 model is presented in \\cite{kowal08}, where the reconnection rate was measured as the inflow rate of unreconnected flux. Here we present a more sophisticated and more theory-justified approach to measuring of the reconnection rate and compare the results of the two approaches. ",
        "conclusions": "\\label{sec:summary}  In this article we introduced a new way of measuring of reconnection rate and we tested the dependence of the LV99 reconnection rate on the injection power and the injection scale of MHD turbulence.  We found that the old and new reconnection measures produce consistent results and both confirm the predictions of the LV99 reconnection model.  We note that the numerical testing of the LV99 reconnection model is far from trivial. The model is intrinsically three dimensional. In order to develop a turbulent cascade and minimize the role of numerical diffusion we have to use high resolution simulations. Another problem is the choice of proper boundary conditions. Our model requires open boundaries in order to allow the ejection of the reconnected flux.  This property is crucial for the global reconnection constraint, since the reconnection stops when the outflow of the reconnected flux is blocked.  If the turbulence level increases, LV99 model predicts that this accelerates magnetic reconnection. At very low levels of turbulence, the width over which magnetic fields wander gets smaller than the thickness of the Sweet-Parker current sheet and the Sweet-Parker reconnection takes over.  The successful testing of LV99 model is good news for many areas of astrophysical numerical modeling. Unlike the models that rely on particular plasma properties to induce fast reconnection, the LV99 model is robust  and fast in any type of fluid, provided that the fluid is, at least, weakly turbulent. If the fluid is laminar initially, the LV99 model predicts that such magnetic configurations should be prone to bursts of reconnection, when the outflow increases the level of turbulence in the system. The latter may provide an appealing explanation of Solar Flares (see more in Lazarian \\& Vishniac 2008).  The LV99 model was used as a starting point for development of the model of reconnection of magnetic field in a partially ionized gas in Lazarian, Vishniac \\& Cho (2004). Although the latter model has to be tested, our successes in testing of LV99 model is an encouraging sign in terms of understanding of reconnection in partially ionized interstellar gases. The latter process can serve provide removal of magnetic field during star formation, for instance (cf. Shu et al. 2007).  Other important implications of the fast reconnection of the weakly stochastic field include the acceleration of cosmic rays as magnetic fields lines shrink as a result of reconnection (see  de Gouveia dal Pino, E. \\& Lazarian, A. 2003, 2005). This process is similar to acceleration of energetic particles considered in a more restrictive case of collisionless Hall-MHD reconnection by Drake et al. (2006)."
    },
    "0812/0812.1789.txt": {
        "abstract": "We report Mopra (ATNF), Anglo-Australian Telescope, and Atacama Submillimeter Telescope Experiment observations of a molecular clump in Carina, BYF73 = G286.21+0.17, which give evidence of large-scale gravitational infall in the dense gas.  From the millimetre and far-infrared data, the clump has mass $\\sim$ 2$\\times$10$^4$\\,M\\solar, luminosity $\\sim$ 2--3$\\times$10$^4$\\,L\\solar, and diameter $\\sim$\\,0.9 pc.  From radiative transfer modelling, we derive a mass infall rate $\\sim$ 3.4$\\times$10$^{-2}$ M\\solar yr$^{-1}$.  If confirmed, this rate for gravitational infall in a molecular core or clump may be the highest yet seen.  The near-infrared $K$-band imaging shows an adjacent compact HII region and IR cluster surrounded by a shell-like photodissociation region showing H$_2$ emission.  At the molecular infall peak, the $K$ imaging also reveals a deeply embedded group of stars with associated H$_2$ emission.  The combination of these features is very unusual and we suggest they indicate the ongoing formation of a massive star cluster.  We discuss the implications of these data for competing theories of massive star formation. ",
        "introduction": "Many details of massive star formation in dense molecular clouds are still unclear \\citep{Chu02}, despite much recent progress \\cite[e.g.,][]{SBS02,FWS05,LBB07}.  For example, it is still debated whether massive stars can form by a scaled-up version of the accretion thought to occur with low-mass protostars \\cite[e.g.][]{MT03}, or rather form by collective processes in a clustered environment \\cite[e.g.][]{BBV03}.  Consequently, examples of massive star formation showing evidence of either behaviour can be informative to this debate, especially since there are still relatively few examples known of true massive protostars.  As part of the {\\it Census of High- and Medium-mass Protostars} (CHaMP, \\S\\ref{strategy}), we identified the massive dense clump\\footnote{Here we use the \\citet{WBM00} terms for ``core''  (that part of a molecular cloud which will collapse to form an individual star or binary) and ``clump'' (which will collapse and fragment, via many cores, to form a star cluster).} G286.21+0.17 as showing striking evidence of large-scale gravitational infall, which we report here.  This source \\cite[hereafter referred to as BYF73, from the master CHaMP source list;][]{BYF09} has been included in some previous surveys \\citep{BNM96,DBS03,FBG04,YAK05} and is an Infrared Astronomical Satellite (IRAS) point source, but has not previously been shown to be remarkable.  The precise location is ($l$,$b$) = (286\\degree.208, +0\\degree.169) or ($\\alpha$,$\\delta$)$_{J2000}$ = ($10^{h}38^{m}32^{s}\\hspace{-1mm}.2, -58^{\\circ}19'12''$), about 1\\degree.5 northwest of $\\eta$ Carinae and 12$'$ north of the rim of the 15$'$-diameter HII region/bubble NGC\\,3324/IC\\,2599, at an assumed distance of 2.5 kpc. ",
        "conclusions": "From Mopra and ASTE HCO$^+$ observations, the Galactic source G286.21+0.17 (which we also call BYF73 from the CHaMP survey master list) has been found to be a massive dense molecular clump exhibiting clear signs of gravitational infall.  The size and scale of this infall, $dM/dt \\sim 3.4\\times10^{-2}$\\,M\\solar\\,yr$^{-1}$ over $\\sim$1\\,pc, is either a record or close to it, and may indicate the global formation of a massive protocluster.  AAT near-IR imaging confirms the existence of unusual spectral signatures and a deeply embedded cluster of stars in the infall zone, as well as an adjacent compact HII region and young star cluster.  Higher-resolution mm-wave and FIR/MIR observations of this source are encouraged, since it appears to be an exemplary test case for confronting competing theories of massive star formation.  %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "0812/0812.2219_arXiv.txt": {
        "abstract": "The amount of detected baryons in the local Universe is at least a factor of two smaller than measured at high redshift. It is believed that a significant fraction of the baryons in the current Universe is ``hiding'' in a hot filamentary structure filling the intergalactic space, the Warm-Hot Intergalactic Medium ($WHIM$). We found evidence of the missing baryons in the $WHIM$ by detecting their signature on the angular correlation of diffuse X-ray emission with the XMM-Newton satellite. Our result indicates that $(12\\pm 5)$\\% of the total diffuse X-ray emission in the energy range 0.4-0.6~keV is due to intergalactic filaments. The statistical significance of our detection is several sigmas ($\\chi ^2>136$ $N=19$). The error bar in the X-ray flux is dominated, instead, by cosmic variation and model uncertainties. ",
        "introduction": "Several methods have been used to measure the quantity of baryons in the Universe. Measurements at high redshift all point to about 4\\% of the matter-energy density of the Universe to be in the form of baryons, while the rest consists of dark matter and dark energy \\cite{rauch, weinberg, burles,kirkman,bennet}. In contrast, the amount of baryons measured in the local Universe in the form of stars, galaxies, groups, and clusters is less than 2\\% \\cite{fukugita}.  Hydrodynamic simulations tracking the evolution of the Universe suggest that much of the ``missing'' material in the local Universe lies in a hot filamentary gas filling the intergalactic space, the Warm-Hot Intergalactic Medium ($WHIM$). Theoretical estimates indicate that the $WHIM$ filaments, have typical density 20 to 1000 times the average baryon density of the universe and temperature greater than $5\\times 10^5$~K \\cite{cen99,borgani,cen06}. At these temperatures and densities the baryons are in the form of highly ionized plasma, making them invisible to all but low energy X-ray and UV observations. Recent UV observations \\cite{danforth} confirm the presence of ``warm'' gas in the intergalactic medium that should account for half of the missing baryons. However, the ``hot'' core of the filaments can only be detected in the soft X-ray band, mostly through excitation lines of highly ionized heavy elements, such as O{\\tiny VII}, O{\\tiny VIII}, Fe{\\tiny XVII}, C{\\tiny VI}, and N{\\tiny VII} with energy between 0.2 and 0.7~keV.  Different approaches have been proposed to detect and study the missing baryons in X-rays. Absorption features can be spotted along the line of sight of distant bright sources. Evidence of the filament existence comes from an absorption measurement made in the direction of the blazar Markarian 421 in a period of maximum brightness \\cite{nicastro05,nicastro05b}, however such data are still considered controversial \\cite{kaastra,williams,rasmussen}. Unfortunately, with the capability of current X-ray satellites, Markarian 421 is practically the only source sufficiently bright to be used as beacon for absorption measurements.  Another option is to identify red-shifted X-ray emission lines emitted by the filaments \\cite{ursino}, but the emission is very weak and must be detected on top of a larger background from galactic emission and unresolved point sources. Dedicated missions are on the drawing board for the study of the WHIM  (e.g., den Herder et al. - 2007), but, if approved, they are still several years away. With the capability of current missions, searching for individual filaments in emission is impractical, except in a few specific directions where filaments are expected. Clusters of galaxies form in the higher density regions (or nodes) of the filamentary structure and are connected by lower density filaments identified with the $WHIM$. Following the cluster network could therefore point to the location of the filaments. The complexity of the network and the distance between clusters makes this approach very difficult, except in few cases where clusters are sufficiently close together. Recently, an observation between the the clusters Abell 222 and Abell 223 has found evidence of excess emission attributed to a filament connecting them \\cite{werner}.  A different way to detect the $WHIM$ in emission is to use a statistical approach. The plasma filaments are expected to have a characteristic angular structure that can be identified and studied using the angular autocorrelation function \\cite{ursino,croft}.  Hydrodynamic simulations have been used to predict the angular distribution of soft X-rays emitted by the filaments. The results show that the filaments should have typical angular scales of a few arcminutes or smaller, leaving a clear signature in the angular distribution of their X-ray emission \\cite{ursino}. The difficulty in detecting such signature lies in the fact that the filament emission only represents a fraction of the total diffuse X-ray emission, which is dominated by unresolved point sources and diffuse galactic emission, such as Local Bubble and Galactic Halo \\cite{galeazzi,mushotzky}. Data from the ROSAT All-Sky Survey ($RASS$) and ROSAT pointed observations have been used to calculate the angular autocorrelation function on scales greater than 2 arcmin, but the results were dominated by emission from unresolved point sources \\cite{soltan,kuntz}. The angular autocorrelation function has also been used to quantify the possible detection of a filament using the ROSAT $PSPC$ detector \\cite{scharf}.  The XMM-Newton satellite, launched in 1999, is currently the best platform to investigate the angular distribution of X-rays emitted by the missing baryons. The CCD detectors aboard XMM-Newton ($PN$ and $MOS$) have a field of view of about 30' and an angular resolution of about 14~arcsec, covering the range of angular scales of interest. Moreover, the angular resolution of the detectors allows good source identification, strongly reducing the effect of unresolved point sources. The detectors also maintain a reasonable response down to about 0.2~keV, covering the energy range of interest. ",
        "conclusions": "We studied the angular distribution of the diffuse X-ray emission from 6 targets available in the XMM-Newton public archive. After removing all knows sources of contamination we found a clear signal that we attribute to emission from the Warm-Hot Intergalactic Medium. A comparison with simulations indicate that, in the energy range 0.4-0.6~keV, the WHIM emission correspond to $(12\\pm 5)$\\% of the total diffuse X-ray emission. The statistical significance of our detection is several sigmas ($\\chi ^2>136$ $N=19$). The error bar in the X-ray flux is dominated, instead, by cosmic variation and model uncertainties."
    },
    "0812/0812.2396_arXiv.txt": {
        "abstract": "{Photoelectric heating is a dominant heating mechanism for many phases of the interstellar medium. We study this mechanism throughout the Large Magellanic Cloud.} {We aim to quantify the importance of the [\\cii] cooling line and the photoelectric heating process of various environments in the LMC and to investigate which parameters control the extent of photoelectric heating.} {We use the BICE [\\cii] map and the Spitzer/SAGE infrared maps. We examine the spatial variations in the efficiency of photoelectric heating: photoelectric heating rate over power absorbed by grains, i.e. the observed [\\cii] line strength over the integrated infrared emission. We correlate the photoelectric heating efficiency and the emission from various dust constituents and study the variations as a function of H$\\alpha$ emission, dust temperatures, and the total infrared luminosity. The observed variations are interpreted in a theoretical framework. From this we estimate radiation field, gas temperature, and electron density.} {We find systematic variations in photoelectric efficiency. The highest efficiencies are found in the diffuse medium, while the lowest coincide with bright star-forming regions ($\\sim$1.4 times lower). The [\\cii] line emission constitutes 1.32\\% of the far infrared luminosity across the whole of the LMC. We find correlations between the [\\cii] emission and ratios of the mid infrared and far infrared bands, which comprise various dust constituents. The correlations are interpreted in light of the spatial variations of the dust abundance and by the local environmental conditions that affect the dust emission properties. As a function of the total infrared surface brightness, $S_\\mathrm{TIR}$, the [\\cii] surface brightness can be described as: $\\mathrm{S_\\mathrm{[\\cii]}=1.25~S_\\mathrm{TIR}^{0.69}~[10^{-3}~erg~s^{-1}cm^{-2}sr^{-1}]}$, for $\\mathrm{S_\\mathrm{TIR}~\\gtrsim~3.2\\cdot~10^{-4}erg~s^{-1}cm^{-2}sr^{-1}}$. We provide a simple model of the photoelectric efficiency as a function of the total infrared. We find a power-law relation between radiation field and electron density, consistent with other studies. The [\\cii] emission is well-correlation with the 8~$\\mu$m emission, suggesting that the polycyclic aromatic hydrocarbons play a dominant role in the photoelectric heating process.} {} ",
        "introduction": "\\label{sec:intro} The structure and evolution of the interstellar medium (ISM) is largely dependent upon the thermal processes taking place \\citep{goldsmith:1969, dejong:1977, mckee:1977, draine:1978, ferriere:1988}. This, in turn, shapes the evolution of galaxies as a whole, as the constituents of the ISM are responsible for the characteristics of incipient stellar generations. Therefore, an understanding of the agents which dominate the heating and cooling of interstellar gas is of fundamental importance.  A dominant heating source of the ISM of galaxies is the photoelectric (PE) emission of interstellar dust grains. Absorption of a far-ultraviolet (FUV) photon by a dust grain may result in the ejection of an energetic electron which heats interstellar gas via collisions. Photoelectric emission as a heating mechanism of the ISM was first proposed by \\citet{spitzer:1948} and later revisited by \\citet{watson:1972} and \\citet{dejong:1977}. Since then, it has been found that the process dominates the heating of a range of interstellar media: neutral atomic gas clouds, the photo-dissociation regions (PDRs), and the warm inter-cloud medium \\citep[e.g.][]{1982A&A...106....1M,weingartner:2001}  The PE heating process has received much theoretical attention \\citep[e.g.][]{watson:1972, dejong:1977, draine:1978, tielens:1985a, bakes:1994, dwek:1996, weingartner:2001}. Due to grain charging, PE heating efficiency is highly dependent on the physical conditions which determine the ionisation and recombination rates. Specifically, it depends on the FUV radiation field, gas temperature and electron density. In turn, the extent of grain charging and therefore the efficiency of PE heating is also highly dependent on the grain species involved.  The $\\mathrm{C^{+}}$ fine structure transition ($\\mathrm{^{2}P_{3/2}-^{2}P_{1/2}}$) is the dominant coolant of the diffuse ionised and diffuse atomic gas as well as in PDRs \\citep{dalgarno:1972, tielens:1985a, tielens:1985b, stacey:1991, madden:1993, petuchowski:1993, heiles:1994}. Several reasons account for its dominance: carbon is the fourth most abundant element and it has an ionisation potential of 11.3 eV, less than that of H. This $\\mathrm{C^{+}}$ transition is also easy to excite as it has a relatively low excitation temperature ($\\sim$92K). Thus, it is able to cool warm neutral gas ($\\mathrm{T\\cong30-10^{4}K}$) whereas other species can not \\citep{tielens:1985a, tielens:1985b, wolfire:1990}. The efficiency of the $\\mathrm{C^{+}}$ as a coolant is dependent upon environment. When temperatures or densities are high, other lines, primarily [\\oi]~63$\\mu$m, participate in the cooling process \\citep{tielens:1985a, tielens:1985b, hollenbach:1991}. The critical density of the $\\mathrm{C^{+}}$ transition is relatively low ($\\mathrm{\\sim 3 \\cdot 10^{3}~cm^{-3}}$). At densities above the critical density or temperature above 92\\,K the cooling by the [\\cii] line saturates and its importance as coolant diminishes.  An observational study of PE heating and gas cooling requires [\\cii]\\footnote{Throughout this paper, we utilise the following notation: we refer to the fine structure transition with $\\mathrm{C^{+}}$ and to the emission line that it produces with [\\cii].} and infrared (IR) observations, covering wavelengths tracing the variety of dust components in the ISM, and ideally of sufficient spatial resolution to separate environments, because of the dependence of the process on environment and composition. Because of its proximity, the Large Magellanic Cloud (\\object{LMC}) is an obvious candidate for a study of these processes. We undertake such a study using the Spitzer legacy program: Surveying the Agents of a Galaxy's Evolution \\citep[SAGE][]{meixner:2006, bernard:2008}, and the Balloon-borne Infrared Carbon Explorer mission \\citep[BICE][]{mochizuki:1994}. SAGE fully mapped the \\object{LMC} at high spatial resolution from 3 to 160$\\mu$m while the latter offers a [\\cii] map of the entire \\object{LMC}. These datasets offer advantages over previous work because of their enhanced spatial resolution, wavelength coverage, and sensitivity. Prior studies using NASA's Kuiper Airborne Observatory (KAO) and the Long Wavelength Spectrometre (LWS) aboard the Infrared Space Observatory (ISO) examined [\\cii] emission integrated across whole galaxies \\citep[e.g.][]{stacey:1991, madden:1993, malhotra:2001}, or in specific regions within galaxies \\citep[e.g.][]{ stutzki:1988, meixner:1992, poglitsch:1995, israel:1996, madden:1997}. There have not been many [\\cii] studies which probe the range of different phases of the ISM. Moreover, even though PE efficiency is believed to be a strong function of grain size, there have not been many prior observational studies of the effect of PE heating due to distinct grain populations. Most studies have considered the total PE heating across all grain populations because they did not have access to bands which trace the emission from distinct grain populations.  The composition of the ISM of the \\object{LMC} makes it an interesting laboratory because it has a low metallicity; $\\mathrm{Z \\cong0.3-0.5Z_{\\sun}}$ \\citet{westerlund:1997} and $\\mathrm{Z \\cong0.25Z_{\\sun}}$ \\citet{dufour:1984}. The presence of metals, in the form of dust, is integral to the PE effect. How does the low metallicity lower dust abundance affect PE heating?  Moreover, it is known from observational studies that the \\object{LMC} and other low metallicity galaxies have a dearth of polycyclic aromatic hydrocarbons (PAHs) compared to Galactic values \\citep{sakon:2006, vermeij:2002, madden:2006, wu:2006, galliano:2007, engelbracht:2008}. This condition raises another question: since it is found theoretically that PAHs should play the greatest role in the PE heating process, how is the PE heating affected in a galaxy with a prominent dearth of PAHs?  In contrast to the metallicity argument given above, it is found from observations that the $\\mathrm{C^{+}}$ coolant plays an even more important role at low metallicity than it does at high metallicity , contributing up to $\\sim$10 times more to the far-infrared (FIR) emission \\citep{ poglitsch:1995, israel:1996, madden:1997}. The larger relative strength has been explained by the clumpy nature of the ISM in low metallicity galaxies. Observations of [\\cii]/CO show that for low metallicity galaxies, this ratio can be 10-30 times higher than for normal galaxies or active galaxies \\citep[e.g.][]{madden:2000}. The far UV radiation penetrates deeper into the molecular cloud at low Z, for the same $A_{V}$, leaving a smaller CO core and larger $\\mathrm{C^{+}}$ emitting envelopes. The result is a preponderance of CO cores clumps with larger exposed surface area. As the $\\mathrm{C^{+}}$ resides primarily in the envelopes of the clouds the increase in surface area results in a higher ratio of $\\mathrm{[\\cii]/CO}$ emission (e.g. \\citet{roellig:2007}).  The aim of this paper is to explore the qualitative behaviour and observationally quantify the extent of PE heating and [\\cii] cooling in relation to environment. The format of this paper is as follows: Sec.~\\ref{sec:observations} reviews the observational details of the data used in this study. Section~\\ref{sec:convolved_data} presents data treatment process and the final images.  Section~\\ref{sec:Cii_across_LMC} quantifies the importance of the [\\cii] coolant globally across the \\object{LMC}. In Sec.~\\ref{sec:diffuse_and_non_diffuse} we examine PE heating as a function of environment based on the H$\\alpha$ surface brightness and Sec.~\\ref{sec:Cii_within_LMC} explores the variations in [\\cii] emission within the \\object{LMC} using the H$\\alpha$ criterion and dust temperature. To explore the correlation between [\\cii] emission and the emission from the various grain components in the \\object{LMC}, Sec.~\\ref{sec:dist_of_grain_emission} compares the spatial distribution of the [\\cii] emission and the emission at various Spitzer bands. The key-concepts of PE heating are reviewed in Sec.~\\ref{sec:PE_heating_and_eff}, which are applied in Sec.~\\ref{sec:pe_eff_and_phys_cond} to describe the PE heating as a function of radiation field within the \\object{LMC}. Using the observed relation we calculate electron densities, and find a correlation between radiation field and electron density (Sec.~\\ref{sec:model_for_eff}). The paper concludes with Sec.~\\ref{sec:dust_and_PE_heating}, an analysis of the dependence of PE heating on grain population; we study the qualitative behaviour of the extent of PE heating on grain population, and we also quantify the contributions to PE heating from the various grain populations. ",
        "conclusions": "Using the MIR to FIR SAGE maps, and the BICE [\\cii] map of the \\object{LMC}, we have, for the first time, conducted an observational study of PE heating and [\\cii] cooling in relation to spatially resolved grain emission throughout the \\object{LMC}.  Integrated throughout the entire \\object{LMC}, the [\\cii] line accounts for $0.64 \\pm 0.01\\%$ of the total infrared ($\\sim 1.2\\%$ of the FIR). Applying a correction for the pixels below $2\\sigma_\\mathrm{[\\cii]}$, we find that the [\\cii] line accounts for $1.32 \\pm 0.01\\%$ of the FIR. This value is greater than that of normal and gas rich galaxies (with values typically from 0.1-1\\%), as found in other low metallicity galaxies.  Distinguishing environments by H$\\alpha$ surface brightnesses and by location, we find that the [\\cii] line contributes significantly less to the TIR emission in SF regions versus diffuse ISM regions: \\begin{itemize} \\item $\\mathrm{L_\\mathrm{[\\cii]}/L_\\mathrm{TIR}=0.57 \\pm 0.01 \\%}$ for the SF regions \\item $\\mathrm{L_\\mathrm{[\\cii]}/L_\\mathrm{TIR} = 0.42 \\pm 0.01\\%}$ for \\object{30 Doradus} \\item $\\mathrm{L_\\mathrm{[\\cii]}/L_\\mathrm{TIR} =0.56 \\pm 0.03\\%}$ for \\object{N11} \\item $\\mathrm{L_\\mathrm{[\\cii]}/L_\\mathrm{TIR}= 0.79 \\pm 0.01 \\%}$ for the diffuse regions. \\end{itemize} We also calculate the contribution of the total output of [\\cii] emission from the \\object{LMC} from the same regions. We find that, although the SF regions have the highest surface brightness values, most of the \\object{LMC}'s [\\cii] emission originates from the diffuse medium: \\begin{itemize} \\item $\\mathrm{L_\\mathrm{[\\cii]}/L_{[\\cii], LMC}= 31.5 \\pm 0.1 \\%}$ for the SF regions \\item $\\mathrm{L_\\mathrm{[\\cii]}/L_{[\\cii], LMC} = 5.58 \\pm 0.1 \\%}$ for \\object{30 Doradus} \\item $\\mathrm{L_\\mathrm{[\\cii]}/L_{[\\cii], LMC} =1.53 \\pm 0.1 \\%}$for \\object{N11} \\item $\\mathrm{L_\\mathrm{[\\cii]}/L_{[\\cii], LMC} = 68.5 \\pm 0.1 \\%}$ for the diffuse regions. \\end{itemize}  To examine variations in PE efficiency within the \\object{LMC}, we use $\\mathrm{S_\\mathrm{[\\cii]}}$ as a proxy for the total PE heating rate. We estimate that this assumption is valid for our 15$^{\\prime}$ beam for all but the most active SF regions which might underestimate the total heating rate by at most 20\\%. We study how PE efficiency varies with environment using the observed values of $\\mathrm{S_\\mathrm{TIR}}$ as an indicator of the local radiation field. As a function of $\\mathrm{S_\\mathrm{TIR}}$, $\\mathrm{S_\\mathrm{[\\cii]}}$ flattens at the highest value of $\\mathrm{S_\\mathrm{TIR}}$. The flattening trend is interpreted as a decrease in the PE efficiency for the most illuminated regions in the \\object{LMC}. We provide a prescription for $\\mathrm{S_\\mathrm{[\\cii]}}$ as a function of $\\mathrm{S_\\mathrm{TIR}}$ (Eq.~\\ref{eqn:CII_as_function_of_TIR}) by fitting a power law to the data. Such a decrease in efficiency is theoretically expected due to grain charging effects.  Previous studies have found a correlation between these two parameters, with $\\mathrm{G_{o} \\propto n_{e}^{4/3}}$. This relation has been explained by invoking a simple model assuming pressure-balance between \\hii~regions and the adjacent PDRs. Theoretically, the PE efficiency depends strongly on the recombination-rate, and thus on the ratio of $\\mathrm{G_o/n_e}$. We thus calculate values for $\\mathrm{G_o}$ and $\\mathrm{n_e}$ using the observed efficiencies. We convert the observed $\\mathrm{S_\\mathrm{TIR}}$ to $\\mathrm{G_o}$, assuming illumination by a central source in each 45$\\times$45~pc pixel. We find that a similar scaling-relation between $\\mathrm{G_o}$ and $\\mathrm{n_{e}}$ holds for the \\object{LMC}.It is unclear why the Str\\\"{o}mgren sphere argument should hold on the large $\\sim$45~pc scale that we probe.  We analyse observed PE efficiencies in relation to the grain component emission from each Spitzer band. We note that this is the first such analysis utilising spatially resolved grain emission components throughout an entire galaxy. From the correlations between observed efficiency and the grain component emission, and from a calculation of the PE efficiencies for each population, we show that the PAH emission is the most spatially correlated with the PE heating rate. We therefore conclude that it is the PAH population that dominates the PE heating process.  The efficiency of the PE heating process is dependent on both environmental conditions, such as radiation field, and grain abundances. Our study has examined PE efficiency within the \\object{LMC} without fully disentangling the extent of PE heating due to existence of grain populations favourable to PE heating and the extent due to environmental conditions favourable to PE heating. Disentangling the effects of both on the observed PE efficiency, however, is difficult as the regions where PAH populations are destroyed naturally have intense radiation fields which also suppress the extent of PE heating. To break this degeneracy we will undertake detailed SED modelling of the different regions in the LMC in an upcoming paper. Using the SED models we can independently solve for radiation field values and PAH abundances."
    },
    "0812/0812.0670_arXiv.txt": {
        "abstract": "{ The disks of spiral galaxies commonly show a lopsided mass distribution, with a typical fractional amplitude of 10\\% for the Fourier component m=1. This is seen in both stars and gas, and the amplitude is higher by a factor of two for galaxies in a group. The study of lopsidedness is a new topic, in contrast to the extensively studied bars and two-armed spirals (m=2). Here, first a brief overview of the observations of disk lopsidedness is given, followed by a summary of the various mechanisms that have been proposed to explain its physical origin. These include tidal interactions, gas accretion, and a global instability. The pattern speed of lopsidedness in a real galaxy has not been measured so far, the various issues involved will be discussed. Theoretical studies have shown that the m=1 slow modes are long-lived, while the modes with a moderate pattern speed as triggered in interactions, last for only about a Gyr. Thus a measurement of the pattern speed of lopsided distribution will help identify the mechanism for its origin. ",
        "introduction": "It is known that the light and hence the mass distribution in disks of spiral galaxies is not strictly axisymmetric, as for example in M101 or in NGC 1637, where the isophotes are elongated in one half of the galaxy. This phenomenon was first highlighted in the paper by Baldwin, Lynden-Bell, \\& Sancisi (1980) for the atomic hydrogen gas in the outer regions in the two  halves of some galaxies, and they called these  `lopsided' galaxies. A galaxy is said to be lopsided if it displays a non-axisymmetric mass distribution of type $m=1$ where $m$ is the azimuthal wavenumber, or a cos$\\phi$ distribution where $\\phi$ is the azimuthal angle in the plane of the disk. From the HI global velocity profiles compiled for a large sample of galaxies, Richter \\& Sancisi (1994) concluded that nearly half the galaxies show lopsidedness.  \\begin{figure}[h] \\centering \\includegraphics[height=1.7in,width=2.0in]{fig1.eps} \\caption{\\footnotesize The histogram showing the number of galaxies  vs. the lopsided amplitude, measured in 149 galaxies at the inclination of $< 70^0$ from the OSU sample (Bournaud et al. 2005). The typical normalized lopsided amplitude $<A_1>$ measured over the radial range between 1.5-2.5 disk scalelengths is = 0.1. Thus most spiral galaxies show significant lopsidedness.} \\label{fig1} \\end{figure}  The near-IR observations show that lopsidedness is common in stars as well. The stellar disks in 30 \\% of galaxies studied are significantly lopsided, with an amplitude A$_1$ $> 0.1 $. This is measured as the Fourier amplitude of the m=1 component normalized to the average value (Rix \\& Zaritsky 1995). The lopsidedness is stronger at larger radii. This was confirmed for a larger sample of 149 galaxies from the Ohio State University  (OSU) database by Bournaud et al. (2005), see Fig. 1 here. Thus, lopsidedness is a typical feature, hence it is important to study its dynamics.   \\medskip  In a first such study, a similar Fourier analysis for the interstellar atomic hydrogen gas (HI) has been done by analyzing the two-dimensional surface density plots for 18 galaxies in the Eridanus group (Angiras et al. 2006). Using HI as the tracer allows the asymmetry to be measured upto several times the disk scalelengths, see Fig. 2. This is more than twice the radial distance covered for the stars, since the sky background in the near-IR limits the Fourier analysis to $\\sim 2.5$ disk scalelengths. Note that, during the Fourier decomposition, the same centre has to be used for all the annular radial bins (Rix \\& Zaritsky 1995, Jog \\& Combes 2008).  \\begin{figure}[h] \\centering \\includegraphics[height=1.25in,width=2.5in]{fig2a.eps} \\medskip \\includegraphics[height=1.25in,width=2.5in]{fig2b.eps} \\caption{\\footnotesize The lopsided amplitude and phase of the HI surface density distribution vs. radius (in units of the disk scalelength) for the galaxy UGC 068 in the Eridanus group, taken from Angiras et al. (2006). The amplitude increases with radius (Fig. 2a), and the phase is nearly constant indicating that m=1 is a global mode (Fig. 2b).} \\end{figure}  \\bigskip  \\begin{figure}[h] \\centering \\includegraphics[height=1.8in,width=2.0in]{fig3.eps} \\caption{\\footnotesize The histogram showing the number of galaxies vs. the lopsided amplitude, A$_1$, measured in the 1.5-2.5 disk scalelength range for the Eridanus group galaxies, from Angiras et al. (2006). The group galaxies are more lopsided,  compare with Fig. 1.} \\end{figure}    Since these galaxies belong to a group, this serendipitously allows a study of the effect of group environment on lopsidedness. The average lopsided amplitude A$_1$ for the group galaxies is higher $\\sim 0.2$, see the histogram in Fig. 3. The fraction of galaxies showing lopsidedness is higher in groups. For example, all the group galaxies have A$_1$ $> 0.1$, which is the mean value for the field case. The frequent interactions in galaxies in groups could explain the higher fraction of galaxies showing lopsidedness.   \\section {Origin of disk lopsidedness}   A non-axisymmetric feature will tend to get wound up in a few dynamical timescales or  in a few $\\times$ 10$^8$ yr, due to the differential rotation in the galactic disk. This is the well-known winding dilemma. However, a high fraction of galaxies is observed to be lopsided. To get around the limited lifetime implied by the winding problem, a number of models have been proposed.   \\subsection {Kinematical model} Starting with aligned orbits, Baldwin et al. (1980) calculated  the winding up time while taking account of the epicyclic motion of the particles. This decreases the effective radial range over which the differential rotation is effective, hence it increases the lifetime to be about 5 times longer than for the usual material arms, or $\\sim$ a Gyr. This is still much smaller than the Hubble time so they concluded that the lopsidedness seen could not be of primordial origin.   \\subsection {Dynamical models}  The various mechanisms that have been proposed to explain the physical origin of the disk lopsidedness are discussed next.  \\subsubsection {Disk response to lopsided potential}  The basic features of the orbits, isophotes and kinematics in a lopsided galaxy can be understood  by treating the disk response to an imposed lopsided perturbation potential (Jog 1997, Schoenmakers et al. 1997). The origin of lopsided halo potential was assumed to be of tidal origin. The results from Jog (1997, 2000) are summarized below:  The unperturbed axisymmetric potential, $\\psi_0$, and the first-order perturbation potential, $\\psi_{lop}$, are taken to be respectively:  $$ \\psi_0 = V_c^2 \\: ln R   \\eqno(1) $$  $$ \\psi_{lop} = V_c^2 \\: {\\epsilon_{lop}}  cos {\\phi}  \\eqno(2)$$  \\noindent where V$_c$ is the rotation velocity in the region of flat rotation, and $\\epsilon_{lop}$ denotes the small perturbation parameter in the potential which is taken to be constant with radius. The perturbation potential is taken to be non-rotating for simplicity.   The equations of motion are solved using the first order epicyclic theory. The perturbed closed loop orbits are given by:  $$ R = R_0 ( 1 - 2 {\\epsilon_{lop}} cos {\\phi} ) \\eqno (3) $$ $$ V_R = 2 V_c  \\: {\\epsilon_{lop}} sin {\\phi}  \\eqno (4) $$ $$ V_{\\phi} = V_c ( 1 + {\\epsilon_{lop}} cos {\\phi} ) \\eqno (5) $$   Thus an orbit is elongated along $\\phi$ = 180$^0$ or the minimum of the perturbation potential, and it is shortened along the opposite direction (along $\\phi$ = 0$^0$).  The observations measure isophotes rather than orbits, hence the isophotal shapes are obtained for an exponential galactic disk in a lopsided potential. {\\it The resulting isophotes have an egg-shaped oval appearance}. This agrees with the observed isophotal shapes in M101. On solving the equations of perturbed motion (eqs.[3]-[5]), the equation of continuity, and the effective surface density together, one gets:  $$ {\\epsilon_{lop}} = {\\frac {A_1}{(\\frac{2 R}{R_{exp}}) - 1}}   \\eqno (6)$$  \\noindent From this, the typical value of the perturbation potential parameter $\\epsilon_{lop}$ is obtained to be $\\sim 0.05$ for the average observed lopsided amplitude. This denotes the global distortion of the halo potential. Thus, the disk lopsidedness can be used as a diagnostic to study the lopsidedness of the halo potential.  The disk response to a distorted halo that gives spatial lopsidedness also results in kinematical lopsidedness in the disk. The resulting rotation curve is asymmetric (eq.[5]). The rotational velocity is maximum along $\\phi = 0^0$ and minimum along the opposite direction. The typical difference in the two halves of a galaxy is $ 2 \\epsilon_{lop} V_c$ or $\\sim$ 10 \\% of the rotation velocity, $\\sim$ 25 km s$^{-1}$ (Jog 2002).  Thus the observers should give the full two-dimensional data for the azimuthal velocity, rather than an azimuthally averaged value. The observed asymmetry in the rotation curve then can be used as a tracer to determine the magnitude of the lopsided potential.  \\subsubsection {Tidal interactions}  Tidal interactions between galaxies have long been suggested as a possible mechanism for the origin of disk lopsidedness. The magnitude and age of lopsidedness thus generated have been studied by N-body simulations by Bournaud et al. (2005) for various mass ratios of the galaxies, orbits, and inclinations. The simulations include stars, gas and the dark matter halo, and all three are taken to be live.  The lopsidedness  generated in a prograde encounter between two spiral galaxies with a mass ratio 2:1 is shown in Fig. 4. A strong lopsidedness is triggered during the interaction, but it lasts for only $\\sim 2$ Gyr. Thus this mechanism cannot explain the high A$_1$ seen in many isolated galaxies. The latter may be explained due to an interaction with or accretion of a satellite with a much smaller mass, but this can also cause a vertical puffing up of the disk. This idea has to be explored systematically. In a recent paper, Mapelli, Moore \\& Bland-Hawthorn (2008) show that a fly-by encounter with a neighbouring galaxy UGC 1807 is the preferred mechanism for the origin of lopsidedness seen in NGC 891.  \\begin{figure}[h] \\centering \\includegraphics[height=3.0in,width=2.2in]{fig4.eps} \\caption{\\footnotesize Plot of A$_1$ vs. time in the middle radial range of 1.5-2.5 disk scalelengths, generated in a distant interaction between galaxies (Bournaud et al. 2005). The peak value is $\\sim 0.2$, higher than the average value observed for the field case, but it drops rapidly to 0.05 in a few Gyr. The lower panel shows the same in terms of Q$_1$, the cumulative potential from the disk. } \\end{figure}   Tidal interactions appear to be the dominant mechanism for the generation of lopsidedness in the group galaxies. First, tidal interactions are frequent in a group because of the high number density of galaxies. Second, the early type galaxies show a higher lopsidedness in group galaxies, as would be expected if lopsidedness arises due to tidal encounters (Angiras et al. 2006).   \\subsubsection {Gas accretion}  \\begin{figure}[h] \\centering \\includegraphics[height=1.4in,width=2.65in]{fig5.eps} \\caption{\\footnotesize NGC 1637 - the lopsidedness due to gas accretion (l.h.s.) matches the near-IR data (r.h.s.), from Bournaud et al. (2005).} \\end{figure}   External gas accretion onto a galaxy from a cosmological filament can give rise to disk lopsidedness, as shown by Bournaud et al. (2005). This can explain the lopsidedness seen in many isolated galaxies. They used reasonable values of accretion rates of 4 and 2 M$_{\\odot}$ yr$^{-1}$ on two sides of a disk, which would double the galaxy mass in a Hubble time. The accretion was shown to explain the main features including the lopsidedness observed in the near-IR map of NGC 1637, see Fig. 5. Here once the accretion stops, the lopsidedness disappears within $\\sim 2$ Gyr.    \\subsubsection {Global m=1 instability}  An internal mechanism involving a global mode instability as the generation mechanism was explored by Saha, Combes \\& Jog (2007). They treated a slowly rotating, global m=1 mode in a purely exponential galactic disk, and showed that the inclusion of self-gravity makes it long-lasting. This model was motivated by the nearly constant phase of lopsidedness that is observed (see e.g., Fig. 2), which indicates a global mode. A unique feature of the m=1 mode is that it shifts the centre of mass of the disturbed galaxy away from the original centre of mass. This further acts as an indirect force on the original centre of mass, which results in long-lived m=1 modes.  Using the linearized fluid equations and the softened self-gravity of the perturbation, a self-consistent quadratic eigenvalue equation was derived and solved. The resulting mode gives the basic observed feature, namely the fractional Fourier amplitude A$_1$ increases with radius. Fig. 6 shows the resulting isodensity contours, clearly the centres of isocontours are progressively more disturbed in the outer parts.  \\begin{figure}[h] \\centering {\\rotatebox{270}{\\includegraphics[height=2.0in,width=2.0in]{fig6.eps}}} \\caption {Contours of constant surface density for a global m=1 mode, from Saha et al. (2007). Here the x and y axes are given in units of the disk scalelength. The maximum surface density occurs at (0,0). The outer contours show a progressive deviation from the undisturbed circular distribution, indicating a more lopsided distribution in the outer parts- as observed. } \\end{figure}    \\begin{figure}[h] \\centering {\\rotatebox{270}{\\includegraphics[height=2.0in,width=2.0in]{fig7.eps}}} \\caption{The precession rate for the global lopsided mode vs. the size of the disk in units of the disk scalelength in a galactic disk (shown as the line with circles), from Saha et al. (2007). The precession rate is very low, thus the mode is long-lived. The dashed line denotes the free precession rate $(\\kappa - \\Omega)$. } \\end{figure}    The self-gravity of the mode results in a significant reduction in the differential precession, by a factor of $\\sim 10$ compared to the free precession. This leads to persistent m=1 modes, as shown in Fig. 7. The lifetime or the winding up time is large $\\sim 12$ Gyr. These results are confirmed by N-body simulations. Future work on this topic should include a live halo to get a more realistic picture. ",
        "conclusions": "Lopsidedness is a common feature of spiral galaxies, seen in both stars and gas, and in the field as well as the group galaxies. The lopsidedness has an amplitude of $\\sim 10$ \\% in the radial range of 1.5-2.5 disk scalelengths, and is higher at larger radii.  Various physical mechanisms have been proposed to explain the origin of lopsidedness. These  include external ones such as tidal interactions and gas accretion, or an internal one such as a global m=1 instability. These could all play a role to a varying degree, and which particular mechanism dominates depends on the specifics of the case. For example, in group galaxies, the tidal interactions seem to play a dominant role in generating lopsidedness.  The pattern speed of lopsidedness is an important parameter, but it has not yet been measured in any galaxy.  For a recent review of lopsided spiral galaxies, see Jog \\& Combes (2008).  \\bigskip"
    },
    "0812/0812.3265_arXiv.txt": {
        "abstract": "In 1986 Alex Dalgarno published a paper entitled {\\it Is Interstellar Chemistry Useful?}\\cite{Da86} By the middle 1970s, and perhaps even earlier, Alex had hoped that astronomical molecules would prove to: possess significant diagnostic utility; control many of the environments in which they exist; stimulate a wide variety of physicists and chemists who are at least as fascinated by the mechanisms forming and removing the molecules as by astronomy. His own research efforts have contributed greatly to the realization of that hope. This paper contains a few examples of: how molecules are used to diagnose large-scale dynamics in astronomical sources including star forming regions and supernovae; the ways in which molecular processes control the evolution of astronomical objects such as dense cores destined to become stars and very evolved giant stars; theoretical and laboratory investigations that elucidate the processes producing and removing astronomical molecules and allow their detection. ",
        "introduction": "Jane Fox's eloquent comments about Alex Dalgarno's qualities as a friend, made at the September 2008 symposium honoring him, provided a fitting tribute to Alex's generosity, kindness and thoughtfulness. I am privileged to have Alex as a friend and also to have benefitted from his professional and intellectual support.  Many of Alex's former students and postdocs, as well as other colleagues, have stories about him similar to those that I will relate. We should all wonder how many letters of recommendation Alex has written. That activity must have consumed a tremendous amount of time. However, Alex's professional support of others has often gone far beyond letter writing. He has visited many of us shortly after our moves to new positions and has helped us make good impressions on our new colleagues. In Leeds, as elsewhere, he served on a visiting panel providing advice on how the local physics and astronomy research effort might be developed. The other two panelists visited one day each. In contrast, over a three day period, he spoke one-on-one with every available permanent member of our Physics and Astronomy academic staff to gain a complete and thorough overview of our activity. He produced an insightful report, which due to its concise, reasoned and incisive nature, carried considerable impact. It contributed very positively to our subsequent, successful efforts to establish a new group, with four permanent academic posts, conducting theoretical research on fundamental quantum processes. Alex continued to help Leeds after his 2003 stay. In 2007 he served as an external member of the committee that appointed Paola Caselli as our Professor of Astronomy.  In 1975 when I, as a postgraduate student, first worked with Alex, I soon developed interests in hydromagnetics and plasma kinetics. Rather than encourage me to focus only on molecular processes, Alex supported my other interests. One problem that I wished to pursue concerns the screening of molecular clouds from ionizing cosmic rays by scattering on Alfv\\'en waves. Alex kindly arranged for me to spend substantial fractions of the summers of 1976 and 1977 in Cambridge, England. While there in 1976, I received a letter from him drawing my attention to a recently published paper by Skilling and Strong\\cite{Sk76} on that topic. Alex has often astounded others with his encyclopedic knowledge of a tremendous range of literature. His awareness of so much has often been of great help to others. Alex, Holly Doyle and I made use of results in Ref.~\\refcite{Sk76} in a paper on ionization rates inferred for diffuse molecular clouds.\\cite{Ha78b}  The title of the present paper echoes that of a paper by Alex\\cite{Da86} to which John Black and Ewine van Dishoeck also referred during their talks at the September 2008 symposium. The remainder of this paper is divided into six sections, half of which mention selected contributions of Alex primarily in molecular astrophysics. My coauthors and I have divided his contributions into those concerning molecular diagnosis, those showing how molecular processes control astrophysical environments and those in the investigation of quantum processes relevant to astrophysics. Sections addressing selected recent works on these themes by other researchers interleave with those summarizing Alex's efforts. I apologize to Alex and to others for the fact that the selections cannot be comprehensive. We have had to neglect a great deal of excellent work done by Alex and by others. Ewine gave a talk at the symposium in which she summarized some of Alex's work in astrochemistry and some recent work of others. The overlaps and the differences between what she said and what I said demonstrate the strength and breadth of Alex's work and of the current state of molecular astrophysics. ",
        "conclusions": ""
    },
    "0812/0812.3924_arXiv.txt": {
        "abstract": "I review five of Bohdan Paczy{\\'n}ski's ideas on black hole accretion disk theory. They formed my understanding of the subject and often guided intuition in my research. They are fundamentally profound, rich in physical consequences, mathematically elegant and clever, and in addition are useful in several technically difficult practical applications. ",
        "introduction": "\\label{section-introduction}   \\subsection{Motivation: the five ``easy'' pieces}  This review article does not give a full (or even coherent) description of black hole accretion disk theory. Instead, it only reflects briefly on a few of Bohdan Paczy{\\'n}ski's particularly important contributions to the subject. Bohdan had several brilliant ideas about black hole accretion.  Today, most of them are mainstream and firmly accepted. One is still considered controversial.  {\\it 1. The Potential:~} For free particles, both Newton's and Einstein's orbital dynamics are described by the same principle, and indeed the same equation: the orbital frequency follows from the first derivative, and the epicyclic frequencies follow from the second derivatives of the ``effective potential''. Therefore, as Bohdan pointed out, by a proper definition of an artificial Newtonian potential, one should be able to accurately describe (formally in Newton's theory!) the fully general relativistic orbital motion. He then guessed the form of the ``Paczy{\\'n}ski-Wiita'' potential which immediately became a great success. The potential is simple and very practical in numerous applications.  {\\it 2. The Doughnut:~} One is often interested in phenomena that occur on a ``dynamical'' timescale $t_{*}$ much shorter than the ``viscous'' time $t_{\\cal L}$ needed for angular momentum redistribution, and the ``thermal'' time $t_{\\cal S}$ needed for entropy redistribution. In modeling such cases, Bohdan noticed, the angular momentum and entropy distributions may be considered {\\it free functions}. Therefore, when studying processes with $t_* \\ll t_0 = {\\rm Min}(t_{\\cal L}, t_{\\cal S})$, one may ad hoc {\\it assume} angular momentum and entropy distributions. A constant angular momentum is the simplest possible assumption. It was used by Bohdan and his Warsaw team to construct the ``Polish Doughnuts'', i.e. super-Eddington thick accretion disks.  {\\it 3. The Funnel:~} In Warsaw we found that a long, narrow empty funnel forms along the axis of rotation of a Polish doughnut. It collimates radiation to hyper-Eddington fluxes.  {\\it 4. The Roche Lobe:~} In equilibrium, the isobaric surfaces of the accretion disk coincide with the surfaces of constant effective potential, called the ``equipotential'' surfaces. Near a black hole one of the equipotential surfaces, called the ``Roche lobe'', self-crosses along the ``cusp'', i.e. a circle $r=r_{cusp}$ located between the innermost stable circular orbit at $r = r_{ISCO}$, and the marginally bound orbit at $r=r_{mb}$. Bohdan's knowledge of close binaries helped him to realize that the black hole Roche lobe overflow must induce dynamical mass loss from the disk. Thus, for $r < r_{cusp}$ accretion is not caused by ``stresses'' that remove angular momentum, but by the strong gravity. The ``inner edge'' of an accretion disk is at the ISCO only for accretion flows with very small ${\\dot M}/{\\dot M}_{Edd}$ that are radiatively efficient (Shakura-Sunyaev). He understood that the efficiency of accretion $\\eta$ (defined by $L = \\eta {\\dot M}c^2$) equals the Keplerian binding energy at the location of the cusp: when $r_{cusp} = r_{ISCO}$, efficiency takes its maximal possible value, but it goes to zero when $r_{cusp}$ approaches $r_{mb}$. All radiatively inefficient accretion flows (RIAFs) like advection-dominated accretion flows (adafs) and slim disks have their inner edges close to $r_{mb}$.  {\\it 5. The Inner Torque:~} Against the view of an influential part of the black hole community, Bohdan argued that when the accretion disk is vertically very thin, for very small accretion rates the viscous or magnetohydrodynamic (MHD) torque at the inner edge should be vanishingly small.  \\begin{figure} \\includegraphics[width=.95\\textwidth] {analytic-numerical.eps} \\caption{Isobaric surfaces in a black hole accretion flow. Left: according to a very simple, analytic Polish doughnut model. Right: according to a state-of-art, full 3D MHD numerical simulation. Figure taken from the {\\it Living Review on black hole accretion disks} by Abramowicz \\& Fragile, in preparation.} \\label{fragile} \\end{figure}  \\subsection{Analytic models and computer simulations}  Bohdan was a master of massive supercomputing methods in astrophysics. Some of the computer codes that he developed, e.g. for stellar structure and evolution, are the standard tools in the field. However, although supercomputer MHD simulations are a large part of black hole accretion research today, Bohdan was not himself involved in this activity. Indeed, all Bohdan's ideas that I discuss here are based on simple analytic models. One may ask --- why? This I do not know, but I can think of a possible reason. The present-day MHD simulations do not address the fundamental problems that the theory of black hole accretion faces. Our understanding of black hole accretion rests on analytic models. Indeed, all the fundamental features of black hole accretion flows that have been calculated in MHD simulations, were {\\it previously well understood} in terms of analytic models. It seems that today, and in the foreseeable future, the super-computer simulations will keep confirming rather than solving or discovering. Today, the most important message of the supercomputer simulations seems to be that the approximations and simplifications adopted in the analytic models have been rather wisely chosen (see Figure \\ref{fragile} to illustrate this point).  \\subsection{Notation and units}  In this review, I use the Kerr metric in the standard spherical Boyer-Lindquist coordinates $t, \\phi, r, \\theta$ with the $+---$ signature. The Kerr metric is stationary and axially symmetric, \\begin{equation} \\label{the-metric} ds^2 = g_{tt}\\,dt^2 + 2g_{t\\phi}\\,dt\\,d{\\phi} + g_{\\phi\\phi}\\,d{\\phi}^2 + g_{rr}\\,dr^2 + g_{\\theta\\theta}\\,d{\\theta}^2, \\end{equation} i.e. its metric tensor components are (known) functions of radial and polar coordinates  $r$, $\\theta$, the mass $M$ and the dimensionless spin parameter $\\vert a \\vert < 1$, \\begin{equation} \\label{the-kerr-metric} g_{ik} = g_{ik}(r, \\theta; M, a). \\end{equation} However, in all specific examples, I will use a much simpler Schwarzschild metric that describes spacetime geometry of a non-rotating ($a=0$) black hole, \\begin{equation} \\label{the-schwarzschild-metric} ds^2 = \\left(1 - \\frac{r_G}{r}\\right)dt^2 - \\left(1 - \\frac{r_G}{r}\\right)^{-1}dr^2 - r^2\\left[ d\\theta^2 + \\sin^2\\theta\\,d{\\phi}^2\\right], \\end{equation} with the gravitational radius $r_G$ is defined by, \\begin{equation} \\label{gravitational-radius} r_G \\equiv \\frac{2GM}{c^2} \\end{equation} My apologies to Bohdan; here I often but not always use the $c = 1 = G$ units he greatly disliked. ",
        "conclusions": ""
    },
    "0812/0812.2922_arXiv.txt": {
        "abstract": "We present a study on the effects of the intracluster medium (ICM) on the interstellar medium (ISM) of 10 Virgo cluster galaxies using {\\it Spitzer} far-infrared (FIR) and VLA radio continuum imaging. Relying on the FIR-radio correlation within normal galaxies, we use our infrared data to create model radio maps which we compare to the observed radio images. For 6 of our sample galaxies we find regions along their outer edges that are highly deficient in the radio compared with our models. We also detect FIR emission slightly beyond the observed radio disk along these outer edges. We believe these observations are the signatures of ICM ram pressure. For NGC~4522 we find the radio deficit region to lie just exterior to a region of high radio polarization and flat radio spectral index, although the total 20~cm radio continuum in this region does not appear strongly enhanced. These characteristics seem consistent for other galaxies with radio polarization data in the literature. The strength of the radio deficit is inversely correlated with the time since peak pressure as inferred from stellar population studies and gas stripping simulations, suggesting the strength of the radio deficit is good indicator of the strength of the current ram pressure. We also find that galaxies having {\\it local} radio {\\it deficits} appear to have {\\it enhanced global} radio fluxes. Our preferred physical picture is that the observed radio deficit regions arise from the ICM wind sweeping away cosmic-ray (CR) electrons and the associated magnetic field, thereby creating synchrotron tails as observed for some of our galaxies. We propose that CR particles are also re-accelerated by ICM-driven shocklets behind the observed radio deficit regions which in turn enhances the remaining radio disk brightness. The high radio polarization and lack of precisely coincident enhancement in the total synchrotron power for these regions suggests shearing, and possibly mild compression of the magnetic field, as the ICM wind drags and stretches the leading edge of the ISM. ",
        "introduction": "The physical processes associated with interactions between the intracluster medium (ICM) and the interstellar medium (ISM) play critical roles driving the evolution of spiral galaxies in clusters \\citep[e.g.][]{gg72,ltc80,amb99,ss01,bv01}. Many basic effects related to ICM-ISM interactions (e.g., ram pressure stripping) are still not well understood. These effects include the fate of star-forming molecular clouds, the rates of triggered star formation  versus gas removal, and the possible reconfiguration of a galaxy's large-scale magnetic field and/or cosmic-ray (CR) particles. It is also not known how far from cluster centers the stripping occurs; for example, gas-deficient galaxies are found in cluster outskirts \\citep{sol01} yet it remains unclear if they are stripped in cluster cores or locally in the outskirts as seems to be the case in NGC~4522 \\citep{kvv04}.  If we can find clear tracers of ongoing ram pressure and identify galaxies in different stages of the stripping process, then we can further understand the consequences of stripping. Since ram pressure will more easily affect lower density constituents of the ISM, it seems that the diffuse radio continuum halos of galaxies may be sensitive tracers of active ICM pressure as indicated by observations for a number of Virgo cluster spirals including NGC~4522 \\citep{bv04}, NGC~4402 \\citep{hc05}, NGC~4254 \\citep{kc07}, and a few others \\citep{bv07}.  % More specifically, large regions of enhanced polarized radio continuum emission have been found within a number of cluster spirals \\citep{bv07, kc07}; the maxima of the polarized radio continuum within these galaxies are located along outer edges and thought to arise from external influences of the cluster environment. However, without a good idea of the galaxy's unperturbed appearance in the radio, these observations alone make it difficult to quantify the ram pressure effects. We now address this by predicting the radio distribution expected in the absence of ram pressure by using the nearly universal correlation between the far-infrared (FIR) and non-thermal radio continuum emission of normal galaxies \\citep[e.g.][]{de85,gxh85}.  Studies of the FIR-radio correlation within field galaxies \\citep[e.g.][]{hoer98,hip03,ejm06a} have found that the dispersion in the FIR/radio ratios across a galaxy disk is as small as the dispersion of the global correlation (i.e. $\\la$0.3~dex) on scales  down to $\\sim$100~pc. While a number of other authors have studied how global FIR/radio ratios behave among cluster galaxies \\citep[e.g.][]{mo01,ry04}, finding them to be systematically lower than those for field galaxies, it is only now by using {\\it Spitzer} that we can study the influence of the cluster environment on the FIR/radio behavior {\\it within} galaxies allowing us to better explain the global FIR/radio observations.  The FIR-radio correlation is driven by the process of massive star formation. Massive stars both heat dust, yielding FIR emission, and end their lives as supernovae (SNe) whose remnants (SNRs), via shocks, accelerate CR electrons in a galaxy's magnetic field bringing about synchrotron radiation. Using a phenomenological smearing model, in which a galaxy's FIR map is smoothed by a parameterized kernel to compensate for the fact that the mean free path of dust heating photons is much shorter than the diffusion length of CR electrons, \\citet{ejm08} has shown for a sample of 15 non-Virgo spiral galaxies that the dispersion in the FIR/radio ratios on sub-kiloparsec scales {\\it within} galaxies can be reduced by a factor of $\\sim$2, on average. Accordingly, it is possible to derive a good first order approximation of an undisturbed galaxy's non-thermal radio continuum morphology from its FIR image.  Assuming that dust and gas are well mixed within galaxies, the FIR emission distribution provides a reasonable picture for the molecular gas distribution as modulated by heating by starlight. The non-thermal radio continuum emission, on the other hand, is a function of a galaxy's magnetic field and CR electron distributions (i.e. the relativistic phase of the ISM). Systematic departures from the nominal FIR/radio ratios across galaxy disks undergoing ICM-ISM interactions can provide insight into differential influence of these interactions on the gaseous and relativistic phases of the ISM.  % These differences may help to quantify the ICM pressure acting on the ISM of such galaxies.  With the capabilities of current observatories the Virgo cluster remains an ideal laboratory allowing for many resolution elements across a large number of galaxy disks, especially at FIR wavelengths. Coupling FIR observations taken by the {\\it Spitzer} Space Telescope with VLA radio continuum imaging, obtained as part of the VLA Imaging of Virgo in Atomic gas (VIVA; A. Chung et al. 2008, in preparation) survey, we study how the relativistic and gaseous phases of the ISM are affected by ICM-ISM interactions for a sample of Virgo cluster galaxies. This is done using FIR {\\it Spitzer} maps to predict how the radio morphology should appear if the galaxy were not experiencing an interaction; in this paper we make the case that significant deviations from such an appearance are directly related to ICM-ISM interactions.  The paper is organized as follows: In $\\S2$ we briefly describe the galaxy sample and observations. We then present our analysis techniques for comparing the FIR and radio continuum imaging in $\\S$3. Results from this comparison are presented in $\\S$4. In $\\S$5 we compare our results in detail with other observational evidence of ICM effects for NGC~4522, a well studied galaxy experiencing ram-pressure stripping, and put together a physical scenario for our observations. Then, in $\\S$6, we compare and discuss this physical picture for the rest of the sample. Finally, we summarize our conclusions in $\\S$7.   \\begin{deluxetable*}{ccccccccccccc} \\tablecaption{Basic Galaxy Data\\label{tbl-galdat}} \\tablehead{ \\colhead{} & \\colhead{R.A.} & \\colhead{Decl.} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{$M_{B}$} & \\colhead{$V_{r}$} & \\colhead{$D_{25}$} & \\colhead{} & \\colhead{PA} & \\colhead{$d_{\\rm M~87}$} & \\colhead{H~{\\sc i}~Def.}\\\\ \\colhead{Galaxy} & \\colhead{(J2000)} & \\colhead{(J2000)} & \\colhead{Type} & \\colhead{H$\\alpha$~Type} & \\colhead{Nuc.} & \\colhead{mag} & \\colhead{(km s$^{-1}$)} & \\colhead{(arcmin)} & \\colhead{$b/a$} & \\colhead{($\\degr$)} & \\colhead{($\\degr$)} & \\colhead{(dex)} \\\\ \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} & \\colhead{(7)} & \\colhead{(8)} & \\colhead{(9)} & \\colhead{(10)} & \\colhead{(11)} & \\colhead{(12)} & \\colhead{(13)}} \\startdata NGC~4254&12 18 49.4 &+14 25 07 & SAc\t  &N\t   & \\nodata&-20.66 & 2407& 5.4& 0.87&   0   & 3.6&-0.10\\\\%$\\pm$0.02\\\\ NGC~4321&12 22 55.2 &+15 49 23 & SABbc    &T/N     & L/H    &-21.05 & 1571& 7.4& 0.85&  30   & 4.0& 0.35\\\\%$\\pm$0.12\\\\ NGC~4330&12 23 16.5 &+11 22 06 & Scd\t&T/N $^{a}$& \\nodata&-18.01 & 1565& 4.5& 0.20&  59   & 2.1& 0.80\\\\%$\\pm$0.04\\\\ NGC~4388&12 25 47.0 &+12 39 42 & SAb &T/N~[s]$^{a}$& Sy2    &-19.34 & 2524& 5.6& 0.23&  92   & 1.3& 1.16\\\\%$\\pm$0.12\\\\ NGC~4396&12 25 59.3 &+15 40 19 & SAd\t  &N$^{a}$& \\nodata&-18.04 & -128& 3.3& 0.29& 125   & 3.5& 0.30\\\\%$\\pm$0.04\\\\ NGC~4402&12 26 07.9 &+13 06 46 & Sb\t &T/N$^{a}$& \\nodata&-18.55 &  232& 3.9& 0.28&  90   & 1.4& 0.74\\\\%$\\pm$0.12\\\\ NGC~4522&12 33 40.0 &+09 10 30 & SBcd\t  &T/N~[s] & \\nodata&-18.11 & 2307& 3.7& 0.27&  33   & 3.3& 0.86\\\\%$\\pm$0.02\\\\ NGC~4569&12 36 50.1 &+13 09 48 & SABab    &T/N~[s] & L/Sy   &-20.84 & -235& 9.6& 0.46&  23   & 1.7& 1.47\\\\%$\\pm$0.20\\\\ NGC~4579&12 37 44.2 &+11 49 11 & SABb\t  &T/N     & L/Sy1.9&-20.62 & 1519& 5.9& 0.79&  95   & 1.8& 0.95\\\\%$\\pm$0.20\\\\ NGC~4580&12 37 48.4 &+05 22 09 & SABa pec &T/N~[s] & \\nodata&-19.27 & 1034& 2.1& 0.78& 165   & 7.2& 1.53% \\enddata \\tablecomments{Col. (1): ID. Col. (2): The right ascension in the J2000.0 epoch. Col. (3): The declination in the J2000.0 epoch. Col. (4): RC3 type. Col. (5): Star formation class of \\citet{kk04}: N: Normal; A: Anemic; E: Enhanced; T/A: Truncated/Anemic;  T/C: Truncated/Compact; T/E: Truncated/Enhanced; T/N: Truncated/Normal; T/N~[s] Truncated/Normal (Severe); $^{a}$: classified for this paper. Col. (6): Nuclear type: H: H~{\\sc ii} region; L: LINER; Sy: Seyfert (1, 1.9, 2); SB: Starburst. Col. (7): Absolute $B$-band magnitudes. Col. (8): Heliocentric velocity. Col. (9): Major axis diameters. Col. (10): Semi-minor/semi-major axis ratio. Col. (11): Position Angle in degrees. Col. (12): Distance from Virgo cluster center \\citep[$12^{\\rm h}30^{\\rm m}47^{\\rm s}$, $+12\\degr 20\\arcmin 13\\arcsec$;][]{he98} in degrees. Col. (13): Type-independent H~{\\sc i} deficiencies, having a typical uncertainty of 0.2~dex arising from the uncertainty in morphological classifications, taken from A. Chung et al. (2008, in preparation) (deficiencies logarithmically increase with number). } \\end{deluxetable*} ",
        "conclusions": "We have studied the interstellar medium (ISM) of 10 Virgo galaxies included in the VLA Imaging of Virgo in Atomic Gas (VIVA) survey using {\\it Spitzer} MIPS and VLA 20~cm imagery. By comparing the observed radio continuum images with modeled distributions (i.e. appropriately smoothed FIR images), created using a phenomenological image-smearing model described by \\citet{ejm06a, ejm08}, we find that the edges of many cluster galaxy disks are significantly radio deficient. These radio deficit regions are consistent with being leading-edge regions affected by intracluster medium (ICM) induced ram pressure as suggested by the location of H~{\\sc i} and radio continuum tails. From our results we are able to conclude the following: \\begin{enumerate} \\item The distributions of radio/FIR ratios within cluster galaxies thought to be experiencing ICM-ISM effects are systematically different from the distributions in field galaxies. Radio/FIR ratios are found to be systematically low, due to a deficit of radio emission, along the side of the galaxy experiencing the ICM wind.  \\item In the case of NGC~4522, a clear example of ongoing strong stripping, we find that the radio deficit region lies exterior to a region of high radio polarization and a flat radio spectral index. We interpret this to suggest that CR electrons in the halos of galaxies are being swept up by the ICM wind. The ICM wind drives shocklets into the ISM of the galaxy which re-accelerate CR particles interior to the working surface at the ICM-ISM interface. Some CR electrons may also be redistributed downstream creating synchrotron tails as observed for NGC~4522. The high radio polarization is probably mostly the result of shear as the ICM wind stretches the magnetic field; compression may also play a role, though a modest one, since the total radio continuum in these regions does not appear significantly enhanced.   \\item The severity of the current ICM-ISM interaction, as measured by the normalized difference between the modeled and observed radio flux densities in the radio deficit regions, appears to be strongly correlated with the time since peak pressure as measured by stellar population studies and gas stripping simulations. Thus the radio deficit appears to be a good tracer for the strength of current ram pressure.   \\item Using the identified radio deficit regions we are able to get a quantitative estimate of the minimum strength of the ICM pressure required to affect a galaxy disk; we find values in the range of $\\sim$$2-4\\times 10^{-12}$~dyn~cm$^{-2}$. These pressures are generally smaller than, but similar to, those estimated for typical values of the ICM gas density and galaxy velocities, as well as the range of ram pressures calculated by 3D hydrodynamical simulations. Therefore, our estimates are consistent with the scenario of ram pressure creating the observed deficit regions.   \\item The internal relativistic plasma pressures range between $\\sim$2 and $16\\times 10^{-12}$~dyn~cm$^{-2}$ and appear to be lowest in some of the galaxies with the strongest radio deficits. Thus the strength of the radio deficit may also be affected by the internal relativistic plasma pressures of the galaxies as well as external ram pressure. Further work will be required to learn how much the radio deficits depend on these 2 variables.  \\item The global radio/FIR ratios of these cluster galaxies are systematically higher than the average value found for field galaxies and appear to increase with increasing severity of the local radio deficit, which may trace current ram pressure.  % We attribute this to an increase in the CR energy density which is greater when the interaction between the ICM-ISM is more energetic. We also show that the amount of mechanical luminosity available from a galaxy's motion through the ICM is significantly larger than what is required to heat the CR plasma.   \\end{enumerate}   This paper describes a preliminary study comparing the FIR and radio emission distributions of Virgo cluster galaxies using the first 10 out of $\\sim$40 galaxies for which {\\it Spitzer} data was available. We have clearly shown that by comparing the FIR and radio continuum distributions of these cluster galaxies we can address a number of physically important characteristics affecting their evolution including the strength and directions of the ICM wind. The inclusion of the remaining $\\sim$30 sample galaxies along with other ancillary data sets will help to better distinguish among scenarios presented here and better illuminate our understanding of how the ICM affects the relativistic phase of the ISM."
    },
    "0812/0812.2325_arXiv.txt": {
        "abstract": "Over the last few years enormous progress has been made in the numerical description of the inspiral and merger of binary black holes. A particular effort has gone into the modelling of the physical properties of the final black hole, namely its spin and recoil velocity, as these quantities have direct impact in astrophysics, cosmology and, of course, general relativity. As numerical-relativity calculations still remain computationally very expensive and cannot be used to investigate the complete space of possible parameters, semi-analytic approaches have been developed and shown to reproduce with very high precision the numerical results. I here collect and review these efforts, pointing out the relative strengths and weaknesses, and discuss which directions are more promising to further improve them. ",
        "introduction": "\\label{sect0}  Despite the almost unnatural simplicity with which the problem can be formulated (black holes are after all the simplest macroscopical objects we know), the final evolution of a binary system of black holes is an impressively complex problem to solve. At the same time, this very simple process plays a fundamental role in astrophysics, in cosmology, in gravitational-wave astronomy, and of course in general relativity. Recent progress in numerical relativity initiated by the works in refs.~\\cite{Pretorius:2005gq, Campanelli:2005dd, Baker:2006yw}, have made it now possible to compute the different stages of the evolution, starting from the inspiral at large separations, for which post-Newtonian (PN) calculations provide an accurate description, through the highly relativistic merger, and finally to the ringdown.  As long as the two black holes are not extremal and have masses which are not too different from each other (see however~\\cite{Gonzalez:2008} for simulations with a mass-ratio of $1\\!\\!:\\!\\!10$), no major technical obstacle now prevents the solution of this problem in full generality and with an overall error which can be brought down to less than $1\\%$~ (see ref.~\\cite{Pollney:2007ss, Baker:2008, Campanelli:2008, Hinder:2008kv, Gonzalez:2008} for some recent examples). Yet, obtaining such a solution still requires a formidable computational power sustained over several days. Even for the simplest set of initial data, namely those considering two black holes in quasi-circular orbits, the space of parameters is too vast to be explored entirely through numerical-relativity calculations. Furthermore, many studies of astrophysical interest, such as many-body simulations of galaxy mergers, or hierarchical models of black-hole formation, span a statistically large space of parameters and are only remotely interested in the evolution of the system during the last few tens of orbits and much interested in determining the properties of the final black hole when the system is still widely separated.  In order to accommodate these two distinct and contrasting needs, namely that of sampling the largest possible space of parameters and that of reducing the computational costs, a number of analytical or semi-analytic approaches have been developed over the last couple of years.  In most of these approaches the inspiral and merger is considered as a process that takes, as input, two black holes of initial masses $M_{1}$, $M_{2}$ and spin vectors {\\boldmath $S$}$_{1}$, {\\boldmath $S$}$_{2}$ and produces, as output, a third black hole of mass $M_{\\rm fin}$, spin {\\boldmath $S$}$_{\\rm fin}$ and recoil velocity {\\boldmath $v$}$_{\\rm kick}$. Mathematically, therefore, one is searching for a mapping between the initial 7-dimensional space of parameters (\\ie~the one containing the six spin components $S^j_{1,2}$ and the mass ratio $q \\equiv M_2/M_1$) to two distinct 3-dimensional spaces (\\ie~the ones containing the three components of the final spin and of the recoil velocity)\\footnote{The space of final parameters could in principle be larger if the final mass of the black hole were to be modelled. Two attempts to do so have been recently proposed~\\cite{Kesden:2008,Tichy:2008} (see also the discussion in Sec.~\\ref{sect1}), but this quantity is not usually modelled.}. Clearly this is a degenerate mapping (two different initial configurations can lead to the same final one) and it would seem a formidable task to accomplish given the highly nonlinear features of the few last orbits. Yet surprisingly, or perhaps not so surprisingly given the underlying simplicity of the problem, all of these studies have shown that the final spin vector and the final recoil velocity vector, can be estimated to remarkably good accuracy if the initial parameters are known~\\cite{Rezzolla-etal-2007, Buonanno:07b, BoyleKesdenNissanke:07, BoyleKesden:07, Marronetti:2007wz, Rezzolla-etal-2007b, Rezzolla-etal-2007c, Kesden:2008, Tichy:2008}.  This paper is dedicated to review the different approaches developed so far to tackle this problem, pointing out the relative strengths and weaknesses, and finally comparing them against the numerical data. The manuscript is organized as follows: in Sect.~\\ref{sect1} I review the different methods that have been suggested to model the final spin, distinguishing those that are purely analytic from those that employ, with different amounts, the results of numerical simulations. Sect.~\\ref{sect2} is instead dedicated to the specific approach developed at the AEI to obtain a simple and robust algebraic expression for the final spin vector. In Sect.~\\ref{sect2} I also exploit the advantages of such analytic formula to explore those configurations that may be astrophysically more interesting and then validate the accuracy of the formula against all of the available data and across alternative approaches. Finally, Sect.~\\ref{sect3} is dedicated to the modelling of the recoil velocity vector and a summary is presented of the different contributions and of the debate still present on one them. Section~\\ref{sect4} will present the conclusions and highlight the prospects of future work in the modelling of the final state.  As a final but important remark I note that all of the considerations made here apply to binary systems that inspiral from very large separations and hence through quasi-circular orbits. Such configurations are the ones more likely to occur astrophysically since any residual eccentricity is lost quickly by the gravitational-radiation reaction~\\cite{Peters:1964}. Furthermore, at least for nonspinning equal-mass black holes, recent work has shown that the final spin does not depend on the value of the eccentricity as long as the latter is not too large~\\cite{Hinder:2007qu}. ",
        "conclusions": "\\label{sect4}  The determination of the final spin and recoil velocity from the from the knowledge of the initial properties of the initial black holes is of great importance in several fields. In astrophysics, it provides information on the properties of isolated stellar-mass black holes produced at the end of the evolution of a binary system of massive stars. In cosmology, it can be used to model the distribution of masses and spins of the supermassive black holes produced through the merger of galaxies~\\cite{Berti2008,Arun:2008zn} and whether massive galaxies should be expected to host a massive black hole at their centre (but see also refs.~\\cite{McNamara:2008dy,Sikora:2008fb}). In addition, in gravitational-wave astronomy, the a-priori knowledge of the final spin can help the detection of the ringdown~\\cite{Ajith:2007kx}.  I have reviewed those semi-analytic approaches to model the final state from binary black-hole coalescences and that have been derived to bypass the still very expensive numerical-relativity calculations and obtain a complete and (reasonably) accurate description of the full space of parameters. Although following different routes and approximations, the several approaches proposed so far have provided a rather accurate description of the final state from binary coalescences. For the first time after decades, these approaches and the numerical-relativity calculations behind them, have placed us in the position of \\textit{knowing} the properties of the final black hole. However, our \\textit{understanding} of these properties has not progressed equally rapidly. It is now possible to compute the second and third-order dependences on the initial spins and mass ratios but many of these dependences are still unclear. Because knowledge without understanding is often sterile, the works reported here should serve as stimulus to other studies where the numerical results represent the ``background'' to be used by perturbative or PN studies to go from knowledge to understanding. The recent works in refs.~\\cite{DamourNagar:07a, Berti:2007snb, Berti:2007nw, Baker:2008, Mino:2008, Favata:2008yd} are encouraging and useful steps in this direction.    \\ack It is a pleasure to thank Peter Diener, Nils Dorband, Sascha Husa, Denis Pollney, Christian Reisswig, Erik Schnetter and Jennifer Seiler who have been my collaborators in this work. I am particularly indebted to Enrico Barausse, who has been the source of many discussions and ideas. This work was supported in part by the DFG grant SFB/Transregio~7."
    },
    "0812/0812.0785_arXiv.txt": {
        "abstract": "{} { To improve the parameters of the HD~17156 system (peculiar due to the eccentric and long orbital period of its transiting planet) and constrain the presence of stellar companions. } { Photometric data were acquired for 4 transits, and high precision radial velocity measurements were simultaneously acquired with SARG@TNG for one transit. The template spectra of HD~17156 was used to derive effective temperature, gravity, and metallicity. A fit of the photometric and spectroscopic data was performed to measure the stellar and planetary radii, and the spin-orbit alignment. Planet orbital elements and ephemeris were derived from the fit. Near infrared adaptive optic images was acquired with ADOPT@TNG. } { We have found that the star has a radius of R$_S = 1.44\\pm 0.03$R$_\\odot$ and the planet R$_P = 1.02\\pm 0.08$R$_J$. The transit ephemeris is T$_c = 2\\,454\\,756.73134 \\pm 0.00020 + N\\cdot 21.21663 \\pm 0.00045$ BJD. The analysis of the Rossiter-Mclaughlin effect shows that the system is spin orbit aligned with an angle $\\lambda = 4.8^\\circ\\pm 5.3^\\circ$. The analysis of high resolution images has not revealed any stellar companion with projected separation between 150 and 1\\,000 AU from HD~17156. } {} ",
        "introduction": "The discovery of transiting extrasolar planets (TESP) is of special relevance for the study and characterization of planetary systems. The combination of photometric and radial velocity measurements allows to measure directly the mass and radius of an exoplanet, and hence its density, which is the primary constraint on a planet's bulk composition. Dedicated follow-up observations of TESP during primary transit and secondary eclipse at visible as well as infrared wavelengths allow direct measurements of planetary emission and absorption (e.g., \\citet{2007prpl.conf..701C}, and references therein). Transmission spectroscopy during primary eclipse in recent years has been successful in characterizing the atmospheric chemistry of several Hot Jupiters (e.g., \\citet{2002ApJ...568..377C} and \\citet{2007Natur.448..169T}). Infrared measurements gathered at a variety of orbital phases, including secondary eclipse, have permitted the characterization of the longitudinal temperature profiles of nearby TESP. The quickly increasing amount of high-quality data obtained for TESP has provided the first crucial constraints on theoretical models describing the physical structure and the atmospheres of gas and ice exoplanets. The detailed characterization of TESP ultimately is of special relevance to test several proposed formation and orbital evolution mechanisms of close-in planets.  The planet \\object{HD~17156b}, detected by \\cite{2007ApJ...669.1336F}(hereinafter F07) using the radial velocity method, was  shown to transit in front of its parent star by \\cite{2007A&A...476L..13B}. Additional photometric measurements were presented in follow-up papers by \\cite{2008A&A...485..871G}, \\cite{2008PASJ...60L...1N}, \\cite{2008ApJ...681..636I}, \\cite{2008arXiv0810.4725W}. This planet is unique among the known transiting systems in that its period (21.2 days) is more than 5 times longer than the average period for this sample, and it has the largest eccentricity ($e = 0.67$).  \\cite{1910PAllO...1..123S}, \\cite{1924ApJ....60...15R} and \\cite{1924ApJ....60...22M} showed that a transiting object, like a stellar companion, produces a distortion in the stellar line profiles due to the partial eclipse of the rotating stellar surface during the event, and thus an apparent anomaly in the measured radial velocity of the observed primary star. In the case of close-in giant planets transiting solar-type stars, the amplitude of the Rossiter-McLaughlin effect ranges typically between a few and $\\sim100$ m s$^{-1}$, depending on orbital period, stellar and planetary radii, and stellar rotational velocity. This is within the current instrument capabilities for the brightest transiting planets. Indeed, this effect has been previously observed on several TESP (e.g. \\citet{2008arXiv0807.4929W} and references therein). The observation of the Rossiter-McLaughlin effect allows to measure the relative inclination angle between the sky projections of the orbital plane and the stellar spin axis. Another important result derived from these observations is the possibility to determine whether short orbital period transiting planets move in the same direction of the stellar spin, indicating the existence of strong dynamical interactions between the parent star and its planet. Measurements of the Rossiter-McLaughlin effect thus provide relevant fossil evidence of planet formation and migration processes, as well as dynamical interactions with perturbing bodies \\citep{2002Icar..156..570M}   Up to now 8 of the 9 transiting planets for which the Rossiter-McLaughlin effect was measured are coplanar systems, as our Solar System (the orbital axes of all the planets and the Sun spin axis are aligned within a few degrees). This is compatible with a formation mechanism for close-in giant planets including migration via tidal interactions with the protoplanetary disk. Only the XO-3 planet seems to be a non aligned system \\citep{2008A&A...488..763H}.   The measurement of the Rossiter-McLaughlin effect is of special relevance for the \\object{HD~17156} system, given the very high eccentricity and orbital period, much longer than the other transiting planets. With a period of 21 days, \\object{HD~17156b} is well outside the peak in the period distribution of close-in planets at about 3 days, possibly implying a different migration history with respect to the other transiting planets. An eccentricity as high as that of \\object{HD~17156b} can be hardly explained by models of planet migration via tidal interactions with the protoplanetary disk. Alternatively, high-eccentricity planets might be the outcome of a variety of possible planet-planet dynamical interactions \\citep{2002Icar..156..570M}. In these scenarios, high relative inclinations between the stellar rotation axis and the planet orbit plane are possible. An additional way to get large relative inclinations and eccentricities is represented by Kozai resonances with close stellar companion. Interestingly, a rather large planetary mass coupled with a short orbital period might be an indication for binarity, as typically high-mass short-period planets occur in binary systems \\citep{2007A&A...462..345D}.  \\cite{2008PASJ...60L...1N} presented the first observations of the Rossiter-McLaughlin effect in \\object{HD~17156}. They found an angle between the sky projections of the orbital axis and the stellar rotation axis $\\lambda=62\\pm25^{\\circ}$. However, \\cite{2008ApJ...683L..59C} did not confirm this claim, suggesting instead well-aligned axes. Such a discrepancy calls for additional high-precision RV monitoring during planetary transit. These are presented in this work, along with additional photometric observations and a critical revision of stellar and planetary parameters. To further constrain the origin of the special properties of \\hdb, we also searched for possible wide stellar companions using adaptive optics observations.  The overall organization of this paper begins in Sec.~\\ref{s:hrs} with the description of high resolution spectroscopy data obtained with SARG@TNG. The following Sec.\\ref{s:star} covers the analysis of the stellar parameters. Sec.\\ref{s:photom} presents the photometric data collected during several planetary transits. Next Sec.\\ref{s:system} describes the analysis of the radial velocity and photometric data, and their results. In Sec.\\ref{s:binarity} we present the results for the search of additional stellar companions to \\hd. Finally in Sec.\\ref{s:concl} our conclusion are presented. ",
        "conclusions": "\\label{s:concl} In this paper, we studied the characteristics of HD 17156 and its transiting planets. Stellar parameters (mass, radius, metallicity) agree quite well with the previous study by F07.  Our measurement of the stellar radius of \\hd obtained through the analysis of the transit light-curve  is the same as the one obtained using  the Stefan-Boltzmann law or the  Kervella calibration (Tab.\\ref{t:pstar}). \\cite{2008A&A...485..871G} obtained a radius of 1.63 $\\pm 0.2$ R$_\\odot$, which is only marginally compatible  with our estimate 1.44 $\\pm 0.08$ R$_\\odot$. On the one hand, to explain such a large stellar radius and the observed visual magnitude, it would be necessary to add $\\sim0.3$ mag of interstellar absorption, that at the distance of  \\hd is not realistic  (suggesting a mean  extinction of 4 mag/kpc) because \\hd is located well inside of the Local Bubble, where no strong  absorption is present, and the  maximum expected absorption is few hundreds of mag. On the other hand, also the comparison of the \\cite{2008A&A...485..871G} radius estimate with stellar models does not appear satisfactory: it is not possible to find model  with a radius that agrees with the observed  temperature  and metallicity  of \\hd. We conclude that the determination of the stellar radius, and by inference planetary radius, \\cite{2008A&A...485..871G} is overestimated by 15\\%.   For  a  planet  of  3M$_{\\rm J}$ and an age of $\\sim2$~Gyr,  theoretical  models  of  planet evolution \\citep{2008A&A...482..315B} predict  a radius ranging between 0.9 and 1.1R$_{\\rm J}$  as a function of  chemical  composition  of  the planet.  Our determination of the radius of \\hdb is $R_p =1.02  \\pm 0.08$ R$_{\\rm J}$. This is in excellent agreement theoretical expectations. Thus, the strong tidal heating effects on the planet do not appear to contribute to significantly inflate its radius.  \\cite{2007ApJ...659.1661F} suggested that \\hdb,  due to its large orbital eccentricity, can change its spectral type  during the orbit from a warm pL type at apoastron to  a hotter pM type at  periastron. Their models suggest that for pM class planets the observed radius at {\\em R} band could be 5\\% larger  than the  radius measured  at {\\em B}  band (due  to  the increased opacity of TiO and VO in  the {\\em R} band). Comparing our radius measurements in the {\\em R}  band and  the measurements  obtained in the {\\em B}  band by \\cite{2008A&A...485..871G}, we find identical results. This result is however not significant because of the large error involved in radius measurement. In order to obtain a significant difference of the radius in {\\em B} and {\\em R}  band each measurement should be more accurate than 0.03~R$_{\\rm J}$.    Our  RV monitoring  of \\hd  during  the 2007 December 3  transit does  not confirm  the misalignment between  the stellar  spin and  planet orbit axes claimed  by \\cite{2008PASJ...60L...1N}, but it agrees instead with the opposite finding by \\cite{2008ApJ...683L..59C}. We think our results are more robust because the other most accurate dataset (HET in \\citealt{2008ApJ...683L..59C}) does not cover the full transit, leaving some uncertainties in their Rossiter modeling. We then conclude that the projection on the sky of the stellar spin and planet  orbit  axes are aligned to better than 10 deg.  Therefore  \\hd  joins  most of the other  exoplanet systems with available measurements of the Rossiter-McLaughlin effect in being compatible with coplanarity. The only possible exception is represented by the XO-3 system, for which \\cite{2008A&A...488..763H} found indications for a large departure from coplanarity ($\\sim 70^\\circ$). However, as acknowledged by the authors, this result should be taken as preliminary, because of the possibility of unrecognized systematic errors for observations taken at large airmass and with significant moonlight contamination.  Our result confirms that large deviations from coplanarity between stellar spin and planet orbit  axes are at most rather rare. Such a rarity had already been established at a high confidence level for the ``classical'' Hot Jupiters in  short-period circular orbits. For massive  eccentric  planets  the situation is less clear: \\object{HD  147506b} ($M=8.6$ M$_{\\rm J}$ , $P=5.6$ days, $e=0.52$)  and \\hd ($M=3.2$~M$_{\\rm J}$ , $P=21$ days, $e=0.67$) have projected inclinations below 10$^\\circ$ while the possible detection of spin-orbit misalignment in the XO-3 system ($M=12.5$~M$_{\\rm J}$, $P=3.2$  days, $e=0.29$) still awaits confirmation as discussed above.       The results of the spin-orbit alignment measurements for the \\hd system can be compared with the prediction of the planet scattering models. A range of alignments can be the outcome of planet-planet scattering \\citep{2002Icar..156..570M}. Therefore, our indication for coplanarity does not exclude planet-planet scattering in the HD~17156 system. A larger number of transiting planets with significant eccentricities have to be discovered and characterized to get more conclusive inferences.  We also searched for stellar companions using adaptive optics, to test the hypothesis of Kozai mechanism to explain the large eccentricity of HD~17156~b. We did not detect companions within 1\\,000 AU, and our detection limits allowed us to exclude main sequence companions with projected separations from about 150 to 1\\,000 AU. This result makes unlikely the occurrence of a companion inducing Kozai eccentricity oscillations on the planet, but this possibility can not be yet completely rule out (companions at small projected separation and faint white dwarfs and brown dwarfs companion still possible). Continuation of radial velocity and photometric monitoring will allow a more complete view on the existence of additional companions at small separations."
    },
    "0812/0812.0971_arXiv.txt": {
        "abstract": "{ Galactic nuclei are unique laboratories for the study of processes connected with the accretion of gas onto supermassive black holes. At the same time, they represent challenging environments from the point of view of stellar dynamics due to their extreme densities and masses involved. There is a growing evidence about the importance of the mutual interaction of stars with gas in galactic nuclei. Gas rich environment may lead to stellar formation which, on the other hand, may regulate accretion onto the central mass. Gas in the form of massive torus or accretion disc further influences stellar dynamics in the central parsec either via gravitational or hydrodynamical interaction. Eccentricity oscillations on one hand and energy dissipation on the other hand lead to increased rate of infall of stars into the supermassive black hole. Last, but not least, processes related to the stellar dynamics may be detectable with forthcoming gravitational waves detectors. ",
        "introduction": "Standard model of active galactic nuclei consists of a supermassive black hole (SMBH) of mass $\\mbh$ surrounded by an accretion disc which extends to $\\sim 10^4 - 10^5\\rg;\\;\\rg\\equiv G\\mbh/c^2$. Further away, at $r\\gtrsim10^6 \\rg$, there is assumed to be an obscuring massive molecular torus. Within the radius of $\\sim 10^6 - 10^7\\rg$, there is a dense stellar cusp, of the total mass comparable to $\\mbh$. In spite of its extreme density which may exceed $10^8$ stars per cubic parsec in the centre, the two-body relaxation time in this region is of the order of Gyrs. Hence, on the orbital (crossing) time scale, the stars follow nearly Keplerian orbits in the gravitational field of the SMBH. Interaction of stars with the axisymmetric gaseous structures leads to a secular orbit evolution. Its characteristic time may be well below the relaxation time and, therefore, it can lead to observable effects.  In this contribution we briefly introduce various modes of the mutual interaction of the stars and gas in galactic nuclei and discuss possible consequences. In Section~\\ref{sec:gc} we present our results within the context of the nucleus of the Milky Way. ",
        "conclusions": "Gas in the galactic nuclei can manifest itself not only by means of its own radiation. It forms complex environment together with the supermassive black hole and millions of stars. All these components mutually interact and influence themselves. In this contribution we have focused on the effects of the gravitational and hydrodynamical influence of the gas in the form of accretion disc or torus upon the stars. We have discussed the effect of Kozai oscillations and shown that they are likely to lead to enhanced rate of tidal stellar disruptions. Due to the relativistic pericentre advance, this process is strongly damped for $\\mbh\\gtrsim10^7\\msun$, while it is the most efficient for intermediate masses of the black holes, i.e. it can help them to grow from the IMBH to the SMBH stage.  Passages of the stars through the gas dissipate their kinetic energy. This leads to their continuous inflow towards the black hole. The mass flow in the form of stars has to be small, compared to the accretion rate of the gas, otherwise, the gaseous structure would be destroyed. The individual stars may, however, be detectable in the final phase of their inspiral due to the emission of gravitational waves.  Recent observations of the Galactic Centre indicate recent star formation at a distance of $10^5 - 10^6\\rg$ from the SMBH. This goes in line with previous theoretical considerations about fragmentation of the outer parts of the accretion discs due to self-gravity \\citep{collin99}. In addition, recent numerical simulations of the fragmenting accretion flow suggest that star formation may inhibit further inflow of the gas onto the central mass \\citep{nayakshin07}. Hence, the feedback from stars should be considered in realistic models of AGNs."
    },
    "0812/0812.0376_arXiv.txt": {
        "abstract": "The gravitational lensing distortion of distant sources by the large-scale distribution of matter in the Universe has been extensively studied. In contrast, very little is known about the effects due to the large-scale distribution of dark energy. We discuss the use of Type Ia supernovae as probes of the spatial inhomogeneity and anisotropy of dark energy. We show that a shallow, almost all-sky  survey can limit rms dark energy fluctuations at the horizon scale down to a fractional energy density of $\\sim 10^{-4}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.1000_arXiv.txt": {
        "abstract": "We present a determination of the mass of the supermassive black hole (BH) and the nuclear stellar orbital distribution of the elliptical galaxy Centaurus~A (NGC~5128) using high-resolution integral-field observations of the stellar kinematics. The observations were obtained with SINFONI at the ESO Very Large Telescope in the near-infrared ($K$-band), using adaptive optics to correct for the blurring effect of the earth atmosphere. The data have a spatial resolution of 0\\farcs17 FWHM and high $S/N\\ga80$ per spectral pixel so that the shape of the stellar line-of-sight velocity-distribution can be reliably extracted. We detect clear low-level stellar rotation, which is counter-rotating with respect to the gas. We fit axisymmetric three-integral dynamical models to the data to determine the best fitting values for the BH mass $\\bh=(5.5\\pm3.0)\\times10^7 \\msun$ ($3\\sigma$ errors) and $(\\ml)_K=(0.65\\pm0.15)$ in solar units. These values are in excellent agreement with previous determinations from the gas kinematics, and in particular with our own published values, extracted from the same data. This provides one of the cleanest gas versus stars comparisons of \\bh\\ determination, due to the use of integral-field data for both dynamical tracers and due to a very well resolved BH sphere of influence $R_{\\rm BH}\\approx0\\farcs70$. We derive an accurate profile of the orbital anisotropy and we carefully test its reliability using spherical Jeans models with radially varying anisotropy. We find an increase in the tangential anisotropy close to the BH, but the spatial extent of this effect seems restricted to the size of $R_{\\rm BH}$ instead of that $R_b\\approx3\\farcs9$ of the core in the surface brightness profile, contrary to detailed predictions of current simulations of the binary BH scouring mechanism. More realistic simulations would be required to draw conclusions from this observation. ",
        "introduction": "The existence of supermassive black holes (BHs) in normal galaxy nuclei was predicted forty years ago by \\citet{LyndenBell69}, but until fifteen years ago it was still considered an interesting possibility which had to be demonstrated. Nowadays BHs are regarded as a crucial ingredient for our understanding of how galaxies form. Key to this paradigm shift was the launch in 1990 of the Hubble Space Telescope (HST). It all started with the realisation that the mass of the BH is correlated to other global characteristics of the host galaxy as a whole. Initially a correlation  $\\bh-L$ was found between the mass of the BH and the luminosity of the host-galaxy stellar spheroid \\citep{kormendy95,magorrian98}. In 1997 the installation of the STIS long-slit spectrograph on HST allowed the spatially-resolved kinematical observations to probe inside the radius of the subarcsecond BH sphere of influence $R_{\\rm BH}\\equiv G \\bh/\\sigma^2$ in nearby galaxies ($\\sigma$ being the velocity dispersion of the stars in the galaxy). The increased accuracy in the \\bh\\ determinations contributed to the discovery of the tighter $\\bh-\\sigma$ correlation \\citep{ferrarese00,gebhardt00}.  As $\\sigma$ is a good predictor of galaxy properties, it is perhaps not surprising that similar correlations were found between \\bh\\ and respectively the galaxy concentration \\citep{graham01}, the dark-halo mass \\citep{ferrarese02,Baes03,pizzella05}, the bulge mass \\citep{McLure02,marconi03,haring04} and the stars' gravitational binding energy \\citep{aller07}. The existence of these correlations is broadly consistent with a scenario in which the BH regulates the galaxy formation, during the hierarchical galaxy merging, by shutting off the conversion of gas into stars via a feedback mechanism due to its powerful outflows \\citep{silk98,granato04,dimatteo05,Bower06}.  Even though the current scenario explains a number of observed facts, our interpretation of the role of BHs in galaxy formation is far from secure. One of the problems lies in the fact that the models have few observables to compare with, mainly the galaxy mass (or $\\sigma$) and \\bh. Moreover, even after a decade of HST spectroscopy and models, only about 40 secure BH determinations exist \\citep[compilations are given e.g. in][]{tremaine02,ferrarese05rev,Graham08}. With few exceptions the BH determinations have been performed with a single dynamical tracer (either gas or stars), so that no independent test of the two measurements methods could be made. Very little is known about the orbital distribution near the BHs, which is expected to contain key information on the BH accretion process \\citep{quinlan97,milosavljevic01}.  The advent of integral-field spectroscopy on all the 8--10-m class telescopes, combined with adaptive optics (AO) to reduce the blurring effect of the Earth's atmosphere is opening a new epoch for BH studies. The large mirrors allow for shorter exposure times and higher signal-to-noise ratios ($S/N$) of the observations, than what was possible with the 2.4-m mirror and the STIS long-slit spectrograph of HST. Observations at near-infrared wavelengths allow dust absorption effects to be virtually eliminated. The integral-field observations, due to the tight constraint on the orbital distribution, dramatically improve the accuracy of \\bh\\ determination, for a given spatial resolution and $S/N$ \\citep{verolme02}. Integral-field data are also needed for a unique recovery of the orbital distribution from the observations \\citep{cappellari05,krajnovic05,vandeVen08}. This can be understood from dimensional arguments, considering that most orbits in a stationary potential conserve three isolating integrals of motion. This three-dimensional orbital distribution cannot be recovered without the knowledge of at least another three-dimensional observed quantity. Motivated by these arguments, \\bh\\ determinations from AO-assisted integral-field observations in the near-infrared are starting to appear in the literature \\citep{Davies06,nowak07,Nowak08}.  The elliptical galaxy NGC~5128 (Centaurus~A) is a prime candidate for AO-assisted integral-field observations. At least nine independent distance $D$ determinations for Cen~A, based on different methods, are available in the literature\\footnote{See the NED-1D compilation by Barry F. Madore and Ian P. Steer at http://nedwww.ipac.caltech.edu/level5/NED1D/.} \\citep[e.g.][]{tonry01,rejkuba04,ferrarese07}. The median value is $D=3.8$ Mpc, with extreme ranges of 3.4 Mpc and 4.4 Mpc respectively. Here we adopt a value\\footnote{The choice of the distance $D$ does not affect our results but sets the scale of our models in physical units. Specifically, lengths and masses scale as $D$, while mass-to-light ratios scale as $D^{-1}$.} $D=3.5$ Mpc to be consistent with all the earlier papers on \\bh\\ determination on this galaxy \\citep{marconi01,marconi06,silge05,haring06,krajnovic07,Neumayer07}. At this close distance Cen~A is the nearest elliptical galaxy and one arcsec corresponds to 17 pc.  Cen~A is among the only $\\sim10$  galaxies on the whole sky with an observed $R_{\\rm BH}\\approx1\\arcsec$ \\citep[see][for a partial list]{kormendy04}. Moreover Cen~A possesses a nuclear gaseous disk in regular rotation from which a number of independent determinations of \\bh\\ have been performed, with ever increasing accuracy \\citep{marconi01,marconi06,haring06,krajnovic07,Neumayer07}. This allows for an accurate comparison between the BH mass derived with gaseous or stellar kinematics. The $K$-band central surface brightness of Cen~A is quite bright at $\\mu_K\\approx12.2$ mag \\citep{Jarrett03}, which allows a high $S/N$ in the stellar spectra to be achieved in reasonable exposure times. A bright star, sufficiently close to the nucleus, can be used as reference for the AO correction. All these facts make Cen~A a unique observational benchmark for \\bh\\ determinations in the near-infrared.  In Section~2 we describe our data and the extraction of the stellar kinematics. In Section~3 we present the stellar dynamical models. We discuss our results in Section~4 and we finally summarize them in Section~5. ",
        "conclusions": "Our determination, using the stellar kinematics, for the BH mass in the center of Cen~A $\\bh=(5.5\\pm3.0)\\times10^7 \\msun$ agrees very well with the determination $\\bh=(4.9\\pm1.4)\\times10^7 \\msun$ we performed using the gas kinematics in \\citet{Neumayer07} from the same SINFONI data. This gas-stars \\bh\\ comparison constitutes one of the most robust and accurate ones, due to a very well resolved BH sphere of influence ($R_{\\rm BH}\\approx0\\farcs70$ compared to a PSF FWHM$\\approx0\\farcs17$) and thanks to the use of high-resolution integral-field data for both kinematical tracers. Comparable and equally successful comparisons were done by \\citet{shapiro06} and \\citet{Siopis08}, while a less good agreement was found in \\citet{cappellari02}, likely due to the disturbed gas kinematics.  In \\citet{Neumayer07} we made a detailed comparison with all the numerous previous \\bh\\ determinations for Cen~A \\citep{marconi01,marconi06,silge05,haring06,krajnovic07}. In summary the only significant disagreement in \\bh\\ is with the previous stellar \\bh\\ determination by \\citet{silge05}. In \\refsec{sec:jeans} we showed that the disagreement is likely caused by a difference in the data quality and in the treatment of the contribution from the central non-thermal continuum in the kinematic extraction. It is not due to differences in the details of the modeling methods. As shown in \\citet{Neumayer07}, with our new gas and stars \\bh\\ determination Cen~A lies within the errors on the $\\bh-\\sigma$ relation as given by either \\citet{tremaine02} or \\citet{ferrarese05rev}.  The value $(\\ml)_K=0.65$ is well consistent with what one would expect from a $\\sim7$ Gyr luminosity-weighted age of the stellar population of Cen~A, almost independently of the assumed near-solar metallicity. This adopts the \\citet{kroupa01} {\\em normalization} of the initial mass function (IMF) and uses the population models of \\citet{maraston05}. These population models include a proper treatment of the TP-AGB stars which is essential for a reliable prediction of the near-infrared flux and $(\\ml)_K$. The $(\\ml)_K$ of the stellar population would decrease to 50\\% of the observed dynamical value if Cen~A had a mean age of $\\sim3$ Gyr. In that case dark-matter would be required to explain the observations. However the younger age is at the lower extreme of the estimated age of 3--8 Gyr, for the main body of Cen~A, from the analysis of its globular cluster system \\citep{Peng04gc}. Adopting the normalization of the \\citet{salpeter55} IMF, the mean age of Cen~A would have to be lower than $\\sim3$ Gyr for the population $(\\ml)_K$ not to over-predict the dynamical one. However this IMF normalization has been shown to be inconsistent with the observed dynamical \\ml\\ of early-type galaxies as a class (\\citealt{cappellari06}; see \\citealt{deJong07} for a review).  Cen~A possesses a clear core in the luminosity profile, with a break from a shallow central slope to a steeper outer one at a radius $R_b\\approx3\\farcs9$ \\citep{marconi00}, which is extremely well resolved by our kinematics. Every galaxy spheroid seems to host a supermassive BH \\citep{magorrian98,ferrarese00,gebhardt00}, and galaxies are expected to form by mergers. These cores have been interpreted as due to the scouring of the galaxy profile due to the ejection of stars in radial orbits passing close to the BH binary which forms shortly after mergers \\citep{faber97,quinlan97,milosavljevic01}. The models predicts that the anisotropy should show a bias towards tangential orbits inside $R_b$.  To test this prediction, in this paper we present the recovered stellar anisotropy profile $\\sigma_r/\\sigma_t$ for Cen~A (\\reffig{fig:anisotropy}). The profile is to first order flat around the value $\\sigma_r/\\sigma_t\\approx1$. The maximum deviations of the velocity ellipsoid from a sphere (isotropy) are on the order of 15\\%. The profile shows no sudden bias towards tangential orbits inside $R_b$, while a tangential bias appears only inside $R_{\\rm BH}$. This profile seems typical of the few nuclear anisotropy profiles that have been published so far using integral-field observations or spherical models (the only cases for which the orbital distribution can be robustly recovered): M32 \\citep{verolme02}, M87 \\citep{cappellari05}, NGC~3379 \\citep{shapiro06}, NGC~1399 \\citep{houghton06,Gebhardt07}. A similar anisotropy trend was observed with high-resolution HST long-slit data by \\citet{gebhardt03}, although in that case the profiles appear more noisy.  Although the predicted tangential bias is observed, its size does not seem related to that of the core. So these anisotropy determinations would not seem to support the scenario in which the cores have been produced by the core-scouring mechanism. However a big caveat is that the current simulations are not entirely realistic: they only consider binary mergers of cusp galaxies which are initially isotropic. Now that high-quality anisotropy measurements are starting to appear, it would be valuable if the numerical simulations could revisit their predictions in a more cosmological motivated setting, considering sequences of multiple mergers. For the moment we must refrain from making firm statements on what our anisotropy profile implies for BH formation."
    },
    "0812/0812.4413_arXiv.txt": {
        "abstract": "We compute the flux of linear momentum carried by gravitational waves emitted from spinning binary black holes at 2PN order for generic orbits. In particular we provide explicit expressions of three new types of terms, namely next-to-leading order spin-orbit terms at 1.5 PN order, spin-orbit tail terms at 2PN order, and spin-spin terms at 2PN order. Restricting ourselves to quasi-circular orbits, we integrate the linear-momentum flux over time to obtain the recoil velocity as function of orbital frequency. We find that in the so-called superkick configuration the higher-order spin corrections can increase the recoil velocity up to a factor $\\sim 3$ with respect to the leading-order PN prediction. Whereas the recoil velocity computed in PN theory within the adiabatic approximation can accurately describe the early inspiral phase, we find that its fast increase during the late inspiral and plunge, and the arbitrariness in determining until when it should be trusted, makes the PN predictions for the total recoil not very accurate and robust. Nevertheless, the linear-momentum flux at higher PN orders can be employed to build more reliable resummed expressions aimed at capturing the non-perturbative effects until merger. Furthermore, we provide expressions valid for generic orbits, and accurate at 2PN order, for the energy and angular momentum carried by gravitational waves emitted from spinning binary black holes. Specializing to quasi-circular orbits we compute the spin-spin terms at 2PN order in the expression for the evolution of the orbital frequency and found agreement with Mik\\'{o}czi, Vas\\'{u}th and Gergely. We also verified that in the limit of extreme mass ratio our expressions for the energy and angular momentum fluxes match the ones of Tagoshi, Shibata, Tanaka and Sasaki obtained in the context of black hole perturbation theory. ",
        "introduction": "\\subsection{Motivation and summary of results}  In the past few years the study of linear momentum carried by gravitational radiation and the subsequent recoil (or kick) velocity it imparts to a binary merger has received a lot of attention, as this recoil effect is astrophysically very relevant~\\cite{Merritt}. There has been a lot of effort devoted in quantifying the impact of gravitational recoil on stellar mass black hole population and supermassive black hole (SMBH) growth scenarios~\\cite{MQ,Haiman,TH,Volonteri,VP,LCFH,BL,MAS}, as well as on formation of galactic cores~\\cite{BKMQ,GM}. In addition the observation of a candidate black hole ejected from its host galaxy after a merger has recently been reported~\\cite{KZL}. The evidence for a recoiling black hole lies in the detection of broad blueshifted emission lines, presumably from gas carried by the recoiling hole, accompanied by a corresponding set of narrow emission lines from the gas left behind in the host object. The estimated recoil velocity from the line blueshift is $\\sim 2650\\, {\\rm km/s}$. However Refs.~\\cite{Bogdanovic,Dotti} have proposed alternative scenarios to explain the observations of Ref.~\\cite{KZL} based on massive black hole binary models.  Recent estimates from binary black hole merger simulations~\\cite{Baker2,Pollney,Koppitz,jena-no-spin,Herrmann,Gonzalez,Campanelli, recoilFAU,Bruegmann,Baker,Lousto1,Lousto2,Lousto3} indicate that for some special spin configurations recoil velocities of order $\\sim 4000 \\, {\\rm km/s}$ could occur in nature. Such high kicks are easily strong enough to eject the black hole remnant from its host galaxy. Therefore a precise understanding of the magnitude of kick velocities and their dependence on binary parameters is paramount for the development of accurate galactic population synthesis models and massive black hole formation scenarios. It is however currently impossible to simulate the number of mergers required to span the expected binary parameter space when spins are included, due to overwhelming computational cost. One must therefore develop analytical models for the recoil velocity as function of the masses and spins of the black holes.  The first computations of gravitational recoil in binary systems were performed by Fitchett~\\cite{Fitchett1} and later by Fitchett and Detweiler~\\cite{Fitchett2}. These papers relied upon earlier work by Peres~\\cite{Peres} and Bekenstein~\\cite{Bekenstein}, who independently computed leading-order expressions for linear momentum flux carried by gravitational waves in terms of interference between multipole moments of the radiation field. Fitchett's~\\cite{Fitchett1} result is limited to the regime where the binary's dynamics can be accurately described by Newtonian physics supplemented by dissipative terms due to emission of gravitational waves. It therefore becomes rather inaccurate when the binary is near merger, which is where most of the recoil is accumulated. Thus it is imperative to include post-Newtonian (PN) corrections within this particular framework. To obtain correct recoil velocities at higher PN order, one must use the equations of motion for the binary at the appropriate PN order, and also include additional couplings between multipole moments of the radiation field. This extension of the work of Fitchett at 1PN order for non-spinning binaries has been performed by Wiseman~\\cite{Wiseman}. More recently the computation of the recoil for non-spinning binaries at 2PN order has been reported by Blanchet, Qusailah and Will~\\cite{BQW} [henceforth BQW].  If the binary contains spinning black holes, then the spins of the holes contribute additional terms to the linear momentum flux. Kidder~\\cite{Kidder} has computed the leading-order (spin-orbit) contributions from the spins to the recoil. These contributions turn out to be 0.5PN order relative to Fitchett's calculation, i.e. to the leading Newtonian order, showing that spins, if large, play a crucial role in determining the recoil, especially since they introduce extra asymmetries in the binary. In this paper we improve the work of Kidder by computing the kick velocity including all spin effects up to 2PN order beyond Fitchett's leading-order computation. The new contributions we compute in this paper are the next-to-leading order spin-orbit terms at 1.5PN order, spin-orbit tail terms at 2PN order, and the leading spin-spin terms at 2PN order. These terms include in particular contributions from the quadrupole moment of each spinning black hole, which affect both the orbital equations of motion at 2PN order, and the time evolution of the spins themselves at 1.5PN, through precession induced by quadrupole-monopole coupling (see for example Refs.~\\cite{Racine,DamourSpinEOB}).  The PN computations just described can of course claim to provide a reliable estimate of the recoil velocity accumulated solely during the early inspiral phase preceding the plunge, merger and ringdown. However since a significant amount of the total recoil is generated during the plunge, merger and ringdown~\\cite{DG,anatomy}, resummation methods and the inclusion of quasi-normal modes must be invoked in order to provide a complete analytical model of the recoil. Preliminary attempts in this direction were pursued in Refs.~\\cite{DG,anatomy,SB} within the effective-one-body approach~\\cite{EOB1,EOB2,EOB3}. Our present work pushing the calculation of the PN-expanded linear-momentum flux at higher orders provides a foundation for constructing more accurate resummed versions of the linear-momentum flux.  Our paper is structured as follows. We first complete our introductory section with a summary of the notation and conventions employed throughout. Next in Sec.~\\ref{sec:spins} we provide a somewhat detailed overview of the treatment of spins in general relativity and in PN theory. We define carefully different spin variables appearing at various steps of our computations, e.g. spin variables of the PN source multipole moments and spin variables with constant magnitude. In Sec.~\\ref{sec:dPdt} we outline the main computation and give our main results for generic orbits. In Sec.~\\ref{sec:circ} we specialize our results to quasi-circular orbits and integrate the momentum flux to obtain the kick velocity. We provide numerical estimates of the kick velocity accumulated throughout the inspiral for specific configurations, namely equal mass binaries with spins equal in magnitude but opposite in direction. The spins are either collinear with the orbital angular momentum or lying in the orbital plane. In the collinear case we also provide an estimate of the kick for equal masses but unequal spins. Finally in Sec.~\\ref{sec:EJfluxes}, we provide the expressions for the fluxes of energy and angular momentum accurate at 2PN order for spinning binary black holes. We provide flux expressions for generic orbits and compare with existing results in the literature. More specifically we verify that in the extreme mass-ratio limit our fluxes match the formulas obtained by Tagoshi {\\it et al.}~\\cite{Tagoshi} in the framework of black hole perturbation theory. We also compute the spin-spin terms at 2PN order in the expression for the evolution of the orbital frequency derived from the usual balance argument, and verify that it matches the expression obtained by Mik\\'{o}czi, Vas\\'{u}th and Gergely~\\cite{MVG} when one substitutes the proper expression for the quadrupole moment of a Kerr black hole, and one neglects contributions from magnetic dipoles.  \\subsection{Conventions}  In this paper we consider black holes which can be nearly maximally spinning. To reflect this property, our PN counting for spin variables is defined as follows. We introduce spin variables for each body, say $\\Sp^A$, which are related to the true physical spins $\\Sp_{\\rm true}^A$ by \\be \\label{rescaledspins} \\Sp^A \\equiv c \\Sp_{\\rm true}^A. \\ee This rescaling stems from the fact that for maximally spinning compact objects, the physical spin scales as $\\Sp_{\\rm true} \\sim G M^2 / c$, where $M$ is the body's mass. The convention for PN order counting thus states that the physical spin is of order 0.5PN. On the other hand the rescaled spin variables $\\Sp^A$ of Eq.~\\eqref{rescaledspins} are of Newtonian order and do not contain any hidden power of the speed of light. Of course this scaling does not apply to slowly spinning objects, for which $\\Sp_{\\rm true} \\sim G M^2 v_{\\rm spin}/c^2$. For such bodies the rescaled spins~\\eqref{rescaledspins} are of 0.5PN order, but such slow spins are not targeted by this work. Throughout the body of the paper we use geometric units $G=c=1$. However we keep track of the PN order of a given term by assigning to it a multiplicative factor which is a power of $1/c$. This factor should not be thought of as carrying dimensions; it appears solely as part of our PN bookkeeping. We also denote a term of order $n$PN, i.e. scaling as $c^{-2n}$, as being ${\\cal O}(2n)$. The PN harmonic coordinates are denoted as $x^\\mu = (ct,x^i)$, latin indices being spatial. Note that the time coordinate $x^0$ carries a PN counting factor of $c$, to track the relative smallness of time variations compared to spatial variations in PN theory.  In addition we denote the antisymmetric permutation symbol by $\\{\\mu,\\nu,\\lambda,\\rho\\}$, with $\\{0,1,2,3\\} = +1$. The Levi-Civita tensor is then \\be \\varepsilon^{\\mu\\nu\\lambda\\rho} \\equiv \\frac{1}{\\sqrt{-g}}\\{\\mu,\\nu,\\lambda,\\rho\\}. \\ee The version with indices down is given by \\be \\varepsilon_{\\mu\\nu\\lambda\\rho} = -\\sqrt{-g}\\, \\{\\mu,\\nu,\\lambda,\\rho\\}, \\ee which can be derived by simply lowering the indices  with $g_{\\mu\\nu}$ on the upper-index version. ",
        "conclusions": "In this paper we computed the linear-momentum flux carried by gravitational waves emitted from spinning binary black holes at 2PN order for generic orbits, notably the next-to-leading order spin-orbit terms at 1.5 PN order, spin-orbit tail terms at 2PN order, and spin-spin terms at 2PN order. In addition, as far as we know, the 2PN non-spinning terms for generic orbits we provide do not seem appear in the literature. We also performed the reduction to quasi-circular orbits and integrated the simplified flux over time to obtain the kick velocity as function of orbital frequency. We specialized our formula for the kick velocity of equal mass binary configurations where the spins are equal in magnitude, opposite in direction and are either collinear with the orbital angular momentum or lying in the orbital plane. In particular, we found that in the so-called superkick configuration the higher-order spin corrections computed in this paper can increase the recoil velocity up to a factor $\\sim 3$ with respect to the leading-order PN prediction.  Comparisons between PN and numerical relativity results for the gravitational-wave energy flux have already shown that when the latter is computed for quasi-circular orbits in the adiabatic approximation, as done in this paper for the linear-momentum flux, it tends to overestimate the numerical energy flux during the last stages of inspiral and plunge~\\cite{Boyleetal08}. Similar conclusions can be drawn here for the linear-momentum flux.  The PN expressions for the linear-momentum flux can be used to grasp which asymmetry in the parameter space can produce the recoil velocity, or suggest phenomenological formulas describing numerical-relativity results~\\cite{Campanelli,Herrmann,Baker,Lousto3}. However, not surprisingly, the fast increase during the late inspiral and plunge, and the arbitrariness in determining until when those formula should be trusted, make the PN predictions not very accurate and robust for predicting the recoil velocity at the peak. By contrast, the computation of the linear-momentum flux at higher PN orders is crucial for building more reliable resummed expressions aimed at capturing the non-perturbative effects until merger~\\cite{EOB1,EOB2,DG,SB} and predict the total recoil velocity.  We also provided expressions valid for generic orbits, and accurate at 2PN order, for the energy and angular momentum carried by gravitational waves emitted from spinning binary black holes. Specializing to quasi-circular orbits we computed the derivative of the orbital frequency through 2PN order, and found agreement with results of Mik\\'{o}czi, Vas\\'{u}th and Gergely~\\cite{MVG}. We also verified that in the limit of extreme mass ratio our expressions for the energy and angular momentum fluxes match the results of Tagoshi {\\it et al.}~\\cite{Tagoshi} obtained in the context of black hole perturbation theory.  It would certainly be quite interesting to extend this computation to 3PN order to provide more refined estimates of the recoil velocity accumulated during the inspiral and build more accurate resummed expressions. This would require the computation of several new source multipole moments. In addition the 3PN acceleration and 3PN spin precession equations currently available in the literature~\\cite{PortoS1S2, PortoS1S1,SpinSchafer1,SpinSchafer2,SpinSchafer3} would need to be computed in harmonic gauge."
    },
    "0812/0812.3375_arXiv.txt": {
        "abstract": "We present high angular resolution observations of the HC$_{\\rm 3}$N J=5--4 line from the Egg nebula, which is the archetype of protoplanetary nebulae. We find that the HC$_{\\rm 3}$N emission in the approaching and receding portion of the envelope traces a clumpy hollow shell, similar to that seen in normal carbon rich envelopes. Near the systemic velocity, the hollow shell is fragmented into several large blobs or arcs with missing portions correspond spatially to locations of previously reported high--velocity outlows in the Egg nebula. This provides direct evidence for the disruption of the slowly--expanding envelope ejected during the AGB phase by the collimated fast outflows initiated during the transition to the protoplanetary nebula phase. We also find that the intersection of fast molecular outflows previously suggested as the location of the central post-AGB star is significantly offset from the center of the hollow shell. From modelling the HC$_3$N distribution we could reproduce qualitatively the spatial kinematics of the HC$_3$N J=5--4 emission using a HC$_3$N shell with two pairs of cavities cleared by the collimated high velocity outflows along the polar direction and in the equatorial plane. We infer a relatively high abundance of HC$_3$N/H$_2$ $\\sim$3x10$^{-6}$ for an estimated mass--loss rate of 3x10$^{-5}$ M$_\\odot$ yr$^{-1}$ in the HC$_3$N shell. The high abundance of HC$_3$N and the presence of some weaker J=5--4 emission in the vicinity of the central post-AGB star suggest an unusually efficient formation of this molecule in the Egg nebula. ",
        "introduction": "The rapid evolution of low and intermediate mass stars (1 to 10 M$_\\odot$) after the end of the Asymptotic Giant Branch (AGB) phase is accompagnied by a radical change in the morphology of their circumstellar envelopes around them. The circumstellar envelope created by the slow and dusty stellar wind during the AGB phase is known to be roughly spherically symmetric as the radiation pressure on dust grains is expected to be isotropic. On the other hand, a variety of morphologies ranging from spherical to multipolar shapes have been observed in post-AGB envelopes and planetary nebulae. It is also during this phase that collimated high velocity outflows are seen, such as in the Egg nebula (Cox et al. 2000) and in CRL 618 (S\\'{a}nchez-Contreras et al. 2004). The shaping of the complex envelope morphology and the mechanism that generates the high--velocity outflows of post-AGB stars are very poorly understood. Lee \\& Sahai (2003) suggest that envelopes with bipolar morphologies are shaped by the interactions with a fast collimated outflow or jet launched in the polar directions. Their hydrodynamic simulations are able to reproduce both the envelope morphology and other properites such intensity of emission lines excited in the shocked gas at the interface between the jet and the slow wind.  The Egg nebula (CRL 2688) is widely considered as the prototype of proto-planetary nebulae (PPN), a class of stars in the rapid transition phase between the AGB and planetary nebulae. During the PPN phase the central post-AGB star contracts and gradually becomes hotter, but is not yet hot enough to ionize its surrounding envelope. Using high resolution spectroscopic observations of the scattered stellar light Klokova et al. (1996) determine a spectral type F5Ia for the central star with an effective temperature of T$_{\\rm eff}$=6500K. The abundance analysis reveals the enrichment of C and N together with strong enhancement of slow neutron capture (s-process) elements, as expected for the carbon rich post-AGB star at the center of the Egg nebula.  Recently, using multi-epoch optical images from Hubble Space Telescope, Ueta et al. (2006) succesfully measured the spatial expansion of the nebula and thereby determined the distance to CRL 2688 of 420 pc. We shall henceforth use this distance for the Egg nebula  High spatial resolution optical images of CRL 2688 reveal a pair of bipolar lobes seen as search-light beams emanating from the central (obscured) star and oriented perpendicular to a dark lane (Sahai et al. 1998a). This conspicuous dark lane has been commonly interpreted as an equatorial disk of cold dust. High angular resolution images of hot molecular hydrogen gas emission at 2.2 $\\mu$m reveal shocked gas, tracing the strong interaction between the fast outflows and the slow wind (Sahai et al. 1998b). Surprisingly, hot molecular hydrogen emission is seen in both bipolar lobes as well as in the equatorial plane far beyond the dark lane, indicating the presence of multiple fast collimated outflows.  Indeed, high spatial resolution mapping of CO J=2--1 emission by Cox et al. (2000) reveals the presence of several pairs of collimated fast molecular outflows along both the bipolar axis and the equatorial plane. These molecular outflows can be traced back to a common origin, presumably the location of the central post-AGB star in the Egg nebula. Cox et al. (2000) suggest that the observed CO emission does not constitute the actual fast collimated outflows, but rather molecular gas swept up by even faster collimated outflows that comprise mainly atomic or ionized gas. By contrast to its appearance in the optical, no equatorial disk is the Egg nebula is seen in the CO J=2--1 by Cox et al. (2000). High spatial resolution observation of HCN J=1--0 by Bieging \\& Nguyen-Q-Rieu (1996) traces molecular gas at higher densities along both the bipolar lobes and the equatorial plane consistent with gas swept up gas by the fast outflows. The only observation that seems to directly trace the AGB wind is that by Yamamura et al. (1996) in $^{13}$CO J=1--0, which does not exhibit any high velocity wings. The limited spatial resolution (about 5\\arcsec) of the observation, however, does not provide any detailed information on the structure of the circumstellar envelope.  Emision lines from cyanopolyyne molecules such as HC$_3$N are very prominent in the centimeter and millimeter wavelengths region toward the Egg nebula. Cyanopolyyne molecules, as large as HC$_9$N, have been detected (Truong-Bach et al. 1993). The line profiles of cyanopolyyne molecules do not exhibit high velocity wings, suggesting that their emission originate from the slowly expanding shell ejected during the AGB phase. Current chemical models (Millar et al. 2000, Cherchneff et al. 1993) for carbon rich envelopes suggest that HC$_3$N molecules are formed by photochemistry in the outer part of the expanding envelope. Thus the spatial distribution of HC$_3$N is predicted to exhibit a hollow shell structure. In the case of the carbon rich envelope of IRC+10216, the spatial distribution of cyanopolyyne molecules has been imaged at high angular resolution (Bieging \\& Tafalla 1993, Lucas \\& Gu\\'{e}lin 1999, Dinh-V-Trung \\& Lim 2008) and shown to have the hollow shell structure as expected from chemical models. The relatively high anbundance and high electric diople moment of HC$_3$N make its rotational lines relatively strong in the circumstellar envelope. The HC$_3$N lines might be useful to probe not just the chemistry but also the structure of the envelope at high angular resolution. In the Egg nebula we also expect the molecule to form in the outer part of the envelope and its emission lines could be used to trace the spatial structure and kinematics of the outer envelope of the Egg nebula. Interferometric observation of HC$_3$N emission line in the 3 mm band by Nguyen-Q-Rieu \\& Bieging (1990) indicates a compact and centrally peaked emission, most likely due to the limited angular resolution of about 8\\arcsec.  In this paper we present high resolution observations of the HC$_3$N J=5--4 emission line, which traces the outer molecular envelope of the Egg nebula. The envelope is found to be disrupted by the passage of the fast collimated outflows along the polar direction and in the equatorial plane. We also present a simple model of the envelope to better understand its spatial kinematics. ",
        "conclusions": "We have imaged at high angular resolution the distribution of HC$_{\\rm 3}$N J=5--4 emission line in the Egg nebula  We find that in the approaching and receding portion of the envelope the HC$_{\\rm 3}$N emission traces a  clumpy hollow shell structure, similar to that seen in normal carbon rich envelopes. Near the systemic velocity, however, the hollow shell is fragmented into several large blobs or arcs. The missing portions of the hollow shell correspond spatially to locations of the high velocity outlows observed previously in the Egg nebula. We also find that the intersection of fast molecular outflows previously suggested as the location of the central post-AGB star is significantly offset from the center of the hollow shell. We interprete the observed spatial-kinematics of HC$_3$N J=5--4 emission as the direct evidence for the disruption of the slowly expanding envelope ejected during the AGB phase by the interaction with collimated high velocity outflows initiated during the transition. From modelling the HC$_3$N distribution we could reproduce qualitatively the spatial kinematics of the HC$_3$N J=5--4 emission using a HC$_3$N shell with two pairs of cavities with opening angle of $\\sim$40$^\\circ$ opened by the collimated high velocity outflows along the polar direction and in the equatorial plane. We infer a relatively high abundance of HC$_3$N/H$_2$ $\\sim$3x10$^{-6}$ for an estimated mass loss rate of 3 10$^{-5}$ M$_\\odot$ yr$^{-1}$. The high abundance of HC$_3$N and the presence of some weaker J=5--4 emission in the vicinity of the central post-AGB star suggest an unusually efficient formation of this molecule in the Egg nebula."
    },
    "0812/0812.4249_arXiv.txt": {
        "abstract": "The dynamics of infinite, asymptotically uniform, distributions of purely self-gravitating particles in one spatial dimension provides a simple and interesting toy model for the analogous three dimensional problem treated in cosmology. In this article we focus on a limitation of such models as they have been treated so far in the literature: the force, as it has been specified, is well defined in infinite point distributions only if there is a centre of symmetry (i.e. the definition requires explicitly the breaking of statistical translational invariance). The problem arises because naive background subtraction (due to expansion, or by ``Jeans swindle'' for the static case), applied as in three dimensions, leaves an unregulated contribution to the force due to surface mass fluctuations.  Following a discussion by Kiessling of the Jeans swindle in three dimensions, we show that the problem may be resolved by defining the force in infinite point distributions as the limit of an exponentially screened pair interaction. We show explicitly that this prescription gives a well defined (finite) force acting on particles in a class of perturbed infinite lattices, which are the point processes relevant to cosmological $N$-body simulations. For identical particles the dynamics of the simplest toy model (without expansion) is equivalent to that of an infinite set of points with inverted harmonic oscillator potentials which bounce elastically when they collide. We discuss and compare with previous results in the literature, and present new results for the specific case of this simplest (static) model starting from ``shuffled lattice'' initial conditions. These show qualitative properties of the evolution (notably its ``self-similarity'') like those in the analogous simulations in three dimensions, which in turn resemble those in the expanding universe. ",
        "introduction": "The development of clustering in initially quasi-uniform {\\it infinite} distributions of point particles evolving purely under their Newtonian self-gravity has been the subject of extensive numerical study in cosmology over the last several decades (see e.g. \\cite{bagla_review} for a review). This is the case because these ``N-body'' (particle) simulations of the Newtonian limit are believed to give a very good approximation to the formation of structure formation in current dark matter dominated models of the universe. The impressive growth in the size of these simulations has led essentially to phenomenological models of the associated dynamics. Analytical understanding, which would be very useful in trying to extend the numerical results and also control for their reliability, remains very limited. In attempts to progress in this direction it is natural to look to simplified toy models which may provide insight and qualitative understanding. Such models may also be interesting theoretically in a purely statistical mechanics setting, and specifically in the context of the investigation of out of equilibrium dynamics of systems with long-range interactions (see e.g. \\cite{Dauxoisetal, Assisi}).  An obvious toy model for this full $3-d$ problem is the analogous problem in $1-d$, i.e., the generalization to an infinite space (static or expanding) of the so-called ``sheet model'', which is formulated for finite mass distributions. In this latter model, which has been quite extensively investigated (see, e.g., \\cite{hohl+feix, reidl+miller, severne+luwel, rouet+feix, tsuchiya+gouda, miller_1dreview} and references therein), particles in $1-d$ experience pair forces independent of their separation, like those between parallel self-gravitating  sheets in $3-d$ of infinite extent. Several groups of authors \\cite{rouet_etal, yano+gouda, tatekawa+maeda, aurell_etal, miller+rouet_2002, aurell+fanelli_2002a, aurell+fanelli_2002b, miller+rouet_2006, valageasOSC_1, valageasOSC_2, miller_etal_2007} have then discussed different variants on this model to develop the analogy with the $3-d$ infinite space problem. Just as for the finite sheet model, these models have the particular interest of admitting exact solutions between sheet crossings, which means that they can be easily evolved numerically to machine precision, and at modest numerical cost for quite large numbers of particles.  In this article we revisit the basics of these toy models (in either static or expanding universes), addressing the problem of their general formulation for infinite distributions. Indeed, as we will discuss, previous discussions have required, in their implementation, the imposition of symmetry about a point, or finite extent of the considered density perturbations\\footnote{This is not true of the treatments in \\cite{yano+gouda, tatekawa+maeda}, which start directly from the fluid limit (rather than from a particle description). See further discussion below.}. Such a restriction on the class of point processes which can be considered, and notably the requirement that statistical translational invariance be broken, is not desirable. Indeed in the context of the cosmological problem, this latter property of the distributions usually considered as initial conditions for simulations is very important, because of the ``cosmological principle'' which supposes that there are no preferred centres (see e.g. \\cite{peebles, book}). Further the question of the extrapolation of the finite version of the model (which is what is simulated numerically) to the infinite system limit has, as we will discuss below, not been carefully examined. We will show that problems with the definition of the force (as used in these previous treatments) arise from a subtlety about how the so-called ``Jeans swindle'' is applied in one dimension. We draw here on the work of Kiessling in \\cite{kiessling}, where it has been shown that, in $3-d$, the usual formulation of the Jeans swindle --- subtraction of a compensating negative mass background in calculation of the potential --- may be more physically formulated as a prescription for the calculation of the {\\it force} in the infinite volume limit. It turns out, as we will see, that while in $3-d$ it is sufficient to prescribe that the force on a given particle is obtained by summing symmetrically about it (e.g. summing in spheres of radius $R$ with centre at the particle, and then sending $R$ to infinity), in $1-d$ this limiting procedure needs to be further specified. More specifically the force turns out to be defined in $1-d$ for a broader class of point distributions --- and notably for distributions without a centre --- when the summation is performed by taking the unscreened limit of the same sum for a screened version of the interaction, rather than as the limit of the sum truncated to a finite symmetric``top-hat'' interval.  In the next section we give a more detailed heuristic discussion of the problem of defining the force in the infinite volume limit, and then give the prescription we adopt. We then present in Sect.~\\ref{Forces in infinite perturbed lattices} a rigorous calculation showing that the force is indeed well defined for a certain class of infinite perturbed lattices, i.e., infinite configurations generated by perturbing particles off a perfect lattice. To do this we treat these infinite point distributions as stochastic point processes and study the probability distribution of the force on particles\\cite{chandra43, book}. For finite variance displacements off the lattice, either correlated or uncorrelated, which do not cause particles to cross one another, the result turns out to be extremely simple: the force on any particle is simply proportional to its displacement from its lattice site.  In the following section we turn to consider the definition of dynamical models using this force. While the most evident model is the simple one obtained by using the conservative Newtonian dynamics under the derived forces, there is also a simple variant with an additional damping term which is the natural toy model for the $3-d$ cosmological problem (with an expanding background).  Given that particle crossings are, up to interchanges of particle labels, equivalent, in $1-d$, to an elastic collision (with exchange of velocities) the evolution in these toy models, starting from infinite perturbed lattices in the class in which we have shown the force to be defined, is in fact equivalent {\\it at all times} simply to an infinite set of particles with inverted harmonic oscillator potentials centred at their original lattice sites, and which collide elastically when they meet. In the last subsection of this section,  we then give a detailed discussion of the relation of these two toy models to those which have been discussed previously in the literature, explaining that our formulation provides essentially both a simplification, and a generalization, of most of these previous treatments. In Sect.~\\ref{Case study} we present briefly results of numerical simulations for the simplest toy model (without expansion), and uncorrelated initial displacements to the lattice (the ``shuffled lattice''). We show that in this case the evolution of the clustering in time is qualitatively very similar to that which has been observed in the analogous $3-d$ system. Notably these static space simulations share the features of ``hierachichal'' structure formation and ``self-similarity'' which are well documented in full $3-d$ simulations. In the final section we summarize our findings and conclusions, and discuss some directions we envisage for further work. ",
        "conclusions": "We have revisited in this article a basic question concerning the definition of the gravitational force in $1-d$ infinite point distributions. Previous definitions of this quantity in the literature have required the assumption of the existence of a special point (centre) in the distribution, i.e., explicit global breaking of statistical translational invariance which is typically a feature of the infinite distributions one instead wishes to study. We have noted that the problem, associated with the non-converging surface fluctuations in such distributions, may be solved by employing a definition using a smooth screening which is sent to zero at the end of the calculation. We have then shown explicitly that this leads to a well defined force for a specific class of infinite perturbed lattices --- those subject to perturbations of finite variance which do not make particles cross. In this case, when the mean displacement of particles is also assumed to vanish, the force on each particle take a unique value which is simply proportional to its own displacement from its lattice site. We note that we have assumed also that variance of the displacement fields is finite, which restricts to initial density fluctuations which have a sufficiently rapidly decaying power spectrum at small wavenumber (specifically, such that $P(k \\rightarrow ) \\sim k^{n}$ where $n>1$, analogous to the same condition with $n > -1$  in $3-d$).  We have then discussed different dynamical toy models which incorporate this definition of the force --- the simple conservative Newtonian dynamics and then one which incorporates a damping term mimicking the effect of $3-d$ expansion. Since the crossing of particles is equivalent, up to labels, to elastic collisions with exchange of velocities, the configurations generated by such dynamics, at any finite time, are always in the class of infinite perturbed lattices for which the force is defined (provided such a configuration is the initial condition). This is the case because, at any finite time, collisions/crossings may only correlate particles up to a finite distance, and the correlation properties of displacements at asymptotically large separations therefore always obey the required conditions. The equations of motion are then simply those of an infinite set of inverted harmonic oscillators (with damping in the expanding case) with centres on the original lattice sites, and which bounce elastically when they collide.  In this context we have also discussed in detail the different formulations of these models in the previous literature. We have then presented some results for the development of clustering in the simplest model, for the simplest class of initial conditions. We have underlined the similarity of the results to those which have been obtained for the analogous system in $3-d$, which is itself a simplified model for the full cosmological model. The physical meaning of the prescription adopted for the force is also very clearly illustrated by these simulations: the dynamics we observe in this infinite system limit is simply that which would occur in a system with a screened gravitational force, in the regime in which the size of the structures is much less than the scale of the screening.  A few additional remarks on our prescription for the force, and the class of point distributions we have considered, are appropriate:  \\begin{itemize}  \\item While we have emphasized that our definition of the force does not require explicit breaking of translational invariance, we have in fact shown it to give a well defined answer only for a class of point processes which have discrete statistical translational invariance (by lattice vectors) rather than full statistical translational symmetry. We will consider in forthcoming work the more general question of the general conditions on point processes for the force to be defined. In this respect we note that it has, in fact, been shown in \\cite{aizenman_etal2001} that {\\it any} point process in $1-d$ with a variance of mass in an interval which is bounded  necessarily breaks continuous statistical translational invariance. While we have not shown here that the definedness of the $1-d$ force requires, in general, such boundedness of the variance, it is straightforward to show that this boundedness indeed characterizes the class of perturbed lattices for which we have found the force to be defined.  \\item We note that the results of \\cite{aizenman_etal2001} in fact generalize an earlier result (see \\cite{aizenman+martin_1980} and references therein), that the thermodynamic equilibria of the $1-d$ one component plasma (OCP, or``jellium'' model) break translational invariance. This model is in fact just the same system we are considering here up to the sign of the force, with the difference that the presence of a physical neutralizing negative background is specified. We therefore do not have, a priori, the freedom to use the regularization of the problem in the infinite system limit which we have exploited here. We note, however, that if we were to do so, i.e., define the model in the thermodynamic limit as the zero screening limit of a screened Coulomb interaction in $1-d$, we obtain, for the class of perturbed lattices we have determined, a quite trivial model in which all particles are simple harmonic oscillators about their lattice sites, which bounce elastically when they collide. It would be interesting to investigate further whether such a formulation of the OCP is in fact equivalent to the usual one.  \\item Even if the force itself is not defined, with our prescription, for a given point process in the infinite volume limit, it may still be possible to construct toy models of clustering in an infinite space, with only the weaker condition that {\\it differences} in forces at a finite separation be defined. This is related also to the restriction on the variance of the displacements we imposed, which we have noted corresponds to a restriction on the small wavenumber behaviour of the power spectra for which the force is defined. If one considers differences in the forces rather than forces, it is simple to show that any power spectrum of density fluctuations which is integrable at small wavenumber can be treated. We will discuss this issue, and the more general conditions for definition of the force, in forthcoming work.  \\item In cosmological simulations initial conditions are prepared by applying stochastic perturbations to a perfect lattice (see e.g. \\cite{couchman, discreteness1_mjbm}). These perturbations usually have a correlated Gaussian joint PDF, which means that in the $1-d$ analogy the no-crossing condition we have required cannot be satisfied. In practice, however, the initial amplitudes are always chosen sufficiently small so that no such crossing occurs in the finite simulation box, and indeed the approximation in which the algorithm for the initial conditions is valid corresponds to this limit. To establish rigorously that our results can be extended to incorporate this case would require a generalisation of the calculation we have given which incorporates the contribution to the force PDF from crossings.  This will lead to a non-zero variance in the PDF, but one would expect that its effect should indeed be negligible in the limit that the typical displacement (measured by the displacement variance) is small compared to the lattice spacing.  \\end{itemize}  In future work we aim to exploit further this model, and the expanding variant, as toy models for the cosmological problem. Despite the previous works in the literature which have explored various versions of these models in the regime in which the analogy to $3-d$ cosmological simulations is most direct, and found various simple behaviours in the clustering --- notably the studies of \\cite{rouet_etal, yano+gouda, tatekawa+maeda, miller+rouet_2002, miller+rouet_2006} --- they have led so far to little of the analytical insight one might hope to gain from a toy model. For example, although fractal behaviours have been documented in numerical simulations by \\cite{rouet_etal, miller+rouet_2002, miller+rouet_2006}, the relevant exponents remain unexplained. The very simple formulation of the models we have given may help in this respect.  These toy models may also be useful in understanding other aspects of numerical simulations in $3-d$, such as the question of the importance of discreteness effects: in cosmological simulations the aim is to reproduce as closely as possible through a particle simulation the Vlasov-Poisson limit, which corresponds to an appropriate infinite particle number limit. The corrections which arise due to finite particle number are poorly understood (see \\cite{discreteness3_mjbm} for a detailed discussion and references), leading to the absence of control on the associated errors in theoretical predictions. Although there are evidently important differences, the problem can be formulated and studied in the greatly simplified $1-d$ context. This requires firstly a clear formulation of the Vlasov-Poisson limit, which so far has been given rigorously only for finite systems \\cite{braun+hepp}. This then allows one to define an appropriate numerical extrapolation which should be used to study convergence. In $1-d$ there is the interesting possibility also of performing directly simulations of the Vlasov-Poisson system for comparison.  We thank D. Fanelli, M. Kiessling, B. Marcos, B. Miller and F. Sylos Labini for useful discussions and exchanges, and B. Jancovici for providing us with references \\cite{aizenman_etal2001, aizenman+martin_1980}."
    },
    "0812/0812.4555_arXiv.txt": {
        "abstract": "{ Recently the PAMELA satellite-based experiment reported an excess of galactic positrons that could be a signal of annihilating dark matter. The PAMELA data may admit an interpretation as a signal from a wino-like LSP of mass about 200 GeV, normalized to the local relic density, and annihilating mainly into W-bosons.  This possibility requires the current conventional estimate for the energy loss rate of positrons be too large by  roughly a factor of five.  Data from anti-protons and gamma rays also provide tension with this interpretation, but there are significant astrophysical uncertainties associated with their propagation. It is not unreasonable to  take this well-motivated candidate seriously, at present,  in part because it can be tested in several ways soon.  The forthcoming PAMELA data on higher energy positrons and the FGST (formerly GLAST) data, should provide important clues as to whether this scenario is correct.  If correct, the wino interpretation implies a cosmological history in which the dark matter does not originate in thermal equilibrium. } ",
        "introduction": "Recently, the PAMELA collaboration (a Payload for Anti-Matter Exploration and Light-nuclei Astrophysics) released preliminary results \\cite{Adriani:2008zr} indicating an excess of cosmic ray positrons above the $10$ GeV energy range. This confirms earlier results from HEAT \\cite{Barwick:1995gv,Barwick:1997ig} and AMS-01 \\cite{AMS01}, which had already received some initial interest from theorists, e.g.,  \\cite{Baltz:2002ua,Baltz:2001ir,Kane:2001fz}.  One possible explanation for the positron excess is the annihilation of weakly interacting massive particles (WIMPS) in the galaxy.  A spate of new particle physics models have been proposed, in part to fit the detailed features of the PAMELA data \\cite{NewModels,list2}.  However, it is worth exploring in detail whether a well-established candidate (such as the neutralino) could plausibly fit the data.  As discussed recently in \\cite{Grajek:2008jb},  for this explanation to be valid the neutralinos would need to have a larger cross section than dark matter of thermal origin.   In light of the preliminary findings of PAMELA, we revisit the non-thermal neutralino models considered in \\cite{Grajek:2008jb,Nagai:2008se,Moroi:1999zb,Kane:2001fz} to determine whether they could account for the excess\\footnote{See also a recent discussion in \\cite{Ishiwata:2008cv}.}.  We find a wino-like neutralino with mass roughly $200$ GeV comes close to accounting for the excess, but only if  unconventional assumptions about the underlying distribution of the dark matter or the propagation of its annihilation products are made.  Without such modifications, a light supersymmetric particle appears unable to account for the data. In part, the purpose of this paper is to point out the places where the sources of tension arise for this explanation, while simultaneously highlighting the types of astrophysical modifications that would need to be made to accomodate the data.  Such a candidate is well motivated theoretically.  For example, a wino LSP arises in theories where the anomaly--mediated\\cite{AMSB} contribution to the gaugino masses dominates, including simple realizations of split supersymmetry, and the string constructions where $M$-theory is compactified on a $G_{2}$ manifold\\cite{G2}.  A pure-wino 200 GeV neutralino annihilates dominantly to $W$-bosons, with a cross section $\\langle \\sigma v \\rangle = 2  \\times$ 10$^{-24}$ cm$^{3}$ s$^{-1}$.  It is remarkable that this cross section is approximately the correct one needed to explain the size of the signal in the data. Masses somewhat below 200 GeV could conceivably explain the spectrum from the positron data, but such candidates come into sharper conflict with the existing limits from anti-protons and gamma rays. Even at 200 GeV, a wino has tension with the existing data, a fact implicit in \\cite{Cirelli:2008pk,Donato:2008jk}. In fact, taking the data at face value, such a candidate is excluded.  In the following, we will show how close the 200 GeV neutralino comes to the current data, given the present understanding of the astrophysics.  Given the inherent astrophysical uncertainties, it is not unreasonable to think the $200$ GeV case might ultimately be consistent with existing positron, anti-proton and $\\gamma$-ray data.    Neutralino masses much larger than this give a bad fit to the PAMELA results unless very large astrophysical boost factors are employed \\cite{ZurekHooper}. ",
        "conclusions": "The PAMELA group \\cite{Adriani:2008zr} has reported a robust excess in galactic positrons with energy above about 10 GeV, compared to those expected from known astrophysical sources. This is consistent with earlier reported excesses from \\cite{Barwick:1995gv,Barwick:1997ig,AMS01}. The source that produces postirons will likely produce antiprotons and $\\gamma$-rays.  No unambiguous excess in these channels has been observed, which provides constraints on a dark matter interpretation.  Only if standard assumptions about the propagation of cosmic rays are relaxed can the excess in positrons arise from galactic annihilation of winos, the superpartners of W bosons. Although there are apparent conflicts with anti-protons and gammas, given the astrophysical uncertainties it is likely premature to assume the wino interpretation is completely excluded.  Further investigation of these uncertainties is a topic for future work.  Fortunately, data from FGST and further  data from PAMELA will help to clarify the picture soon.  Later data from AMS-02 will also help.  If the PAMELA excess is indeed due to a well-motivated wino LSP, the implications are remarkable.   We would learn what the dark matter of our universe is.  It would be the discovery of supersymmetry, telling us something about the resolution to the hierarchy problem.  It would imply a number of superpartners will likely be seen at LHC, confirming the result. And we would be learning that the universe had a non-thermal cosmological history  that we can probe.  \\vspace{-.2in}"
    },
    "0812/0812.3469_arXiv.txt": {
        "abstract": "A new method is proposed to compute the bulk viscosity in strange quark matter at high densities. Using the method it is straightforward to prove that the bulk viscosity is positive definite, which is not so easy to accomplish in other approaches especially for multi-component fluids like strange quark matter with light up and down quarks and massive strange quarks. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.0760_arXiv.txt": {
        "abstract": "{We present high-sensitivity, high-resolution images of the Ultraluminous Infrared Galaxies (ULIRG; L$_{\\mathrm{FIR}} > 10^{12}$ L$_\\odot$) IRAS 23365+3604 and IRAS 07251-0248, taken with the EVN at 6 and 18 cm. The images show a large number of compact components, whose luminosities are typical of Type IIL and Type IIn Radio Supernovae (RSNe). Further observations of these ULIRGs  will allow us to confirm, or to rule out, their nature. The present observations are part of a project that should result in a significant number of SN detections, providing a direct measurement of the Core Collapse Superova (CCSN) rate and allowing us to estimate the Star Formation Rate (SFR) in our sample of ULIRGs (see  \\cite{miguel}).}  \\FullConference{The 9th European VLBI Network Symposium on The role of VLBI in the Golden Age for Radio Astronomy and EVN Users Meeting\\\\ September 23-26, 2008\\\\ Bologna, Italy}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "We have obtained the deepest and highest resolution radio images ever of two of the most distant ULIRGs in the local Universe, using quasi-simultaneous observations with the EVN at 6 and 18 cm (see Figure \\ref{fig:maps}). We expect to get a better RMS and u-v coverage in our maps once we add the MERLIN observations.  We found a number of bright, compact components, some of which are suggestive of CCSNe exploding in the innermost regions of those ULIRGs. Indeed, the luminosities inferred for the compact components seen at 6 cm range from $\\sim 5 \\times 10^{27}$ in IRAS 23365+3604 up to $\\sim 16 \\times 10^{28}$ erg s$^{-1}$ Hz$^{-1}$ in IRAS 07251-0248 which correspond to luminosities typical of Type IIL and Type IIn SN (bright radio events). We also show overlays of the 6 and 18 cm EVN observations in Figures \\ref{fig:ir23cl} and \\ref{fig:ir07cl}. Not all of the 6 cm brightness peaks are coincident with those seen at 18 cm (see Table \\ref{ta:param}). This is consistent with an scenario where we observe very young SNe, so that their 6 cm emission would now be around their peak, while their 18 cm emission could still be rising. On the other hand, we have also found several 18 cm peaks without a clear counterpart at 6 cm. This can be explained if their emission arises from CCSNe that are already in their optically thin phase, as indicated by their two-point spectral indices.  Our results show that the EVN is a powerful tool to study the central regions of nearby ULIRGs, and can be efficiently used to determine the CCSN rate and SFR in those galaxies. A new observing epoch should allow us to disentangle the nature of the compact objects we have found.  \\begin{table} \\begin{center} {\\scriptsize \\begin{tabular}{|c|c|c|c|c|c|c|c|c|} \\hline Component & \\multicolumn{4}{|c|}{IRAS23365+3604}   & \\multicolumn{4}{c|}{IRAS07251-0248}                                                                                                           \\\\ \\cline{2-9} ~  & \\multicolumn{2}{c|}{Coordinates (J2000)} & $\\sigma_{\\mathrm{position}}$  & L$_{\\nu} \\times 10^{28}$ &  \\multicolumn{2}{c|}{Coordinates (J2000)} & $\\sigma_{\\mathrm{position}}$  &   L$_{\\nu} \\times 10^{28}$  \\\\ \\cline{2-3} \\cline{6-7} ~        &           RA     &      DEC       & (mas) &  (erg s$^{-1}$ Hz$^{-1}$)        &           RA     &      DEC       & (mas) &   (erg s$^{-1}$ Hz$^{-1}$)   \\\\ \\hline L1             & 23 39 01.2687    & 36 21 08.5230  & 1.3   &   1.1                      & 07 27 37.6154    & -02 54 54.2540 & 1.4   &     8.4                      \\\\ L2             & 23 39 01.2674    & 36 21 08.4840  & 1.4   &   1.1                      & 07 27 37.6129    & -02 54 54.2510 & 0.7   &    16.0                      \\\\ L3             & 23 39 01.2653    & 36 21 08.5260  & 1.2   &   1.2                      & 07 27 37.6113    & -02 54 54.2510 & 1.4   &     8.5                      \\\\ L4             & 23 39 01.2643    & 36 21 08.5690  & 1.0   &   1.4                      &  ~               & ~              &  ~    &     ~                        \\\\ L5             & 23 39 01.2624    & 36 21 08.5520  & 0.8   &   1.8                      &  ~               & ~              &  ~    &     ~                        \\\\ L6             & 23 39 01.2606    & 36 21 08.5720  & 0.7   &   1.9                      &  ~               & ~              &  ~    &     ~                        \\\\ L7             & 23 39 01.2582    & 36 21 08.5860  & 1.2   &   1.2                      &  ~               & ~              &  ~    &     ~                        \\\\ C1             & 23 39 01.2665    & 36 21 08.5830  & 0.8   &   0.7                     & 07 27 37.6144    & -02 54 54.2516 & 0.5   &    2.1                       \\\\ C2             & 23 39 01.2656    & 36 21 08.5895  & 1.1   &   0.5                     & 07 27 37.6129    & -02 54 54.2560 & 0.3   &    3.2                       \\\\ C3             & 23 39 01.2605    & 36 21 08.5895  & 1.1   &   0.5                     & 07 27 37.6125    & -02 54 54.2540 & 0.3   &    3.1                       \\\\ C4             & 23 39 01.2599    & 36 21 08.5855  & 1.1   &   0.5                     & 07 27 37.6120    & -02 54 54.2588 & 0.4   &    2.3                       \\\\ \\hline \\end{tabular} } \\end{center} \\caption{\\footnotesize{Parameters estimated from the EVN-observations  \\label{ta:param}}} \\end{table}  \\vspace{2.5cm}  \\begin{figure}[h!] \\begin{center} \\subfigure[]{ \\label{fig:ir23c} \\includegraphics[angle=0, width=4.5cm, height=3.7cm]{iras2336c-small.jpg}} \\subfigure[]{ \\label{fig:ir23l} \\includegraphics[angle=0, width=4.5cm, height=3.7cm]{iras2336-l.jpg}} \\subfigure[]{ \\label{fig:ir23cl} \\includegraphics[angle=0, width=4.5cm, height=3.7cm]{iras2336cl.jpg}} \\vspace{.1in} \\subfigure[]{ \\label{fig:ir07c} \\includegraphics[angle=0, width=4.5cm, height=3.7cm]{iras0725c-small.jpg}} \\subfigure[]{ \\label{fig:ir07l} \\includegraphics[angle=0, width=4.5cm, height=3.7cm]{iras0725-l.jpg}} \\subfigure[]{ \\label{fig:ir07cl} \\includegraphics[angle=0, width=4.5cm, height=3.7cm]{iras0725cl.jpg}} \\caption { \\footnotesize{IRAS 2336+3604 at 6 (a) and 18 cm (b), with estimated RMS of about 15 $\\mu$Jy in both maps. IRAS 07251-0248 at 6 (d) and 18 cm (e), with estimated RMS of 28 and 60 $\\mu$Jy respectively. Emission components are marked with C- and L- labels; note that the first component in the C-maps do not correspond with the first component in the L-maps and so on. IRAS 23365+3604 (c) and IRAS 07251-0241(f) contours from observations on 29 February 2008 at 18 cm overlaid on the 6 cm image from 11 March 2008 observations.} } \\label{fig:maps} \\end{center} \\end{figure}"
    },
    "0812/0812.3413_arXiv.txt": {
        "abstract": "\\noindent We present a set of formalisms for comparing, evolving and constraining primordial non-Gaussian models through the CMB bispectrum. We describe improved methods for efficient computation of the full CMB bispectrum  for any general (non-separable) primordial bispectrum, incorporating a flat sky approximation and a new cubic interpolation.   We review all the primordial non-Gaussian models in the present literature and calculate the CMB bispectrum up to $l <2000$ for each different model.   This allows us to determine the observational independence of these models by calculating the cross-correlation of their CMB bispectra.  We are able to identify several distinct classes of primordial shapes - including equilateral, local, warm, flat and feature (non-scale invariant) - which should be distinguishable given a significant detection of CMB non-Gaussianity. We demonstrate that a  simple shape correlator provides a fast and reliable method for determining whether or not CMB shapes are well correlated.  We use an eigenmode decomposition of the primordial shape to characterise and understand model independence.   Finally, we advocate a standardised normalisation method for $f_{NL}$ based on the shape autocorrelator, so that observational limits and errors $\\Delta f_{NL}$ can be consistently compared for different models. ",
        "introduction": "Constraints on non-Gaussianity arguably provide the most stringent observational tests of the simplest inflationary paradigm and, in the near future, these limits are set to tighten substantially.   Single-field slow-roll inflation predicts to high precision that the CMB will be a Gaussian random field and hence can be completely described by its angular power spectrum. However, if there was mechanism for generating some non-Gaussianity in the initial perturbations then its measurement would open up a wealth of extra information about the physics governing the early universe. Motivated by the first discussions of the local case which is dominated by squeezed states, non-Gaussianity is usually parametrised by $f_{NL}$, a quantity that can be extracted from CMB observations. The purpose of this paper is to apply and improve the methods detailed in ref.~\\cite{0612713} for calculating the CMB bispectrum for any general (non-separable) primordial non-Gaussian model and, on this basis, to determine the extent to which competing models are independent and can be constrained by present limits on $f_{NL}$.  At present, CMB observations directly constrain only three separable primordial non-Gaussian models because of the calculational difficulties of using bispectrum estimators in the general case.   These are the local model and the equilateral model   (a separable approximation to the DBI inflation) and warm inflation with the latest (model-dependent) CMB constraints on $f_{NL}$ becoming  \\cite{09012572,08030547,07071647} \\begin{align} \\label{eq:locallim} -4 &< f^{local}_{NL} < 80\\,, \\\\ \\label{eq:equilim} -151 &< f^{equi}_{NL} < 253\\,, \\\\ \\label{eq:warmlim} -375 &< f^{warm}_{NL} < 37\\, \\end{align} The above are 2$\\sigma$-limits on non-Gaussianity, though we note that there is a tentative claim \\cite{07121148} of a detection of $f^{local}_{NL}$ at almost $3\\s$, with limits $27 < f^{local}_{NL} < 147 \\;(2\\s)$.  Given the growing range of theoretical possibilities for other (non-separable) non-Gaussian bispectra, an important goal is the development of  methods which can be used to place direct constraints in the general case \\cite{inprep}.  However, our purpose here is to note the models which are tightly constrained by the present limits and, conversely, to identify those for which further investigation is warranted.  The bispectrum has been shown to be optimal for detecting primordial non-Gaussianity \\cite{0503375}, but alternative approaches seem to be able to produce comparable limits.  A wavelet analysis of the WMAP 5 year data yields   $-8 < f_{NL} < 111 \\;(2\\s)$ \\cite{08070231} and positional information from wavelets can be used to examine the likelihood of specific features.  Minkowski functionals provide a geometric characterisation of the temperature fluctuations in the CMB, yielding a slightly weaker constraint $-70 < f_{NL} < 91 \\;(2\\s)$ \\cite{08023677}. Of course, the effects of non-Gaussianity are not only felt in the CMB but could be detectable in a wide range of astrophysical measurements, such as cluster abundances and the large scale clustering of highly biased tracers. However, given the imminent launch of the Planck satellite with a projected constraint $|f_{NL}| \\le 5 \\;(2\\s)$  \\cite{0005036} (almost at the level of cosmic variance, $|f_{NL}| \\le 3 \\;(2\\s)$), the focus of this paper remains on the CMB bispectrum.  In section 2,  we review the relationship between the primordial bispectrum and its counterpart in the CMB, noting the importance of separability for present estimators.   We also present a new analytic solution for the CMB bispectrum from a \\it{constant} primordial model, against which it is useful to normalise other models. In section 3, we introduce the shape function relevant for nearly scale-invariant bispectra, and we review the current literature classifying current non-Gaussian models into centre-, corner- and edge-weighted categories, as well as those which are not scale-invariant.    We use a shape correlator to forecast the cross-correlation of the CMB bispectra for all these different models, identifying five distinct classes of shapes. This work takes forward earlier discussions of the non-Gaussian shape (see, for example, \\cite{0405356,0509029,0612571}) but here we are able to directly compare to the actual CMB bispectra for all the models.   Moreover, we propose a specific eigenfunction decomposition of the shape function which offers insight as to why particular shapes are related or otherwise.  In section 4, we describe important improvements to the numerical methods which we use to make accurate computations of the CMB bispectrum for all the models surveyed \\cite{0612713}.  The most important of these are the flat sky approximation and a cubic interpolation scheme for the tetrahedral domain of allowed multipoles.  What was previously regarded as an insurmountable computational problem has now become tractable, irrespective of separability. In section 5, we present the main results detailing the cross-correlations for all the different models, while confirming the different classes of independent shapes previously identified, some of which remain to be fully constrained (even by present data).   We compare the results of the shape and CMB bispectra correlators, noting the efficacy of the shape approach in identifying models for which quantitative CMB analysis is  required. Finally, in section 6, we propose an alternative normalisation procedure for $f_{NL}$ which brings the constraints into a more consistent pattern, allowing for a model-independent comparison of the true level of non-Gaussianity. ",
        "conclusions": "In this paper we have presented a comprehensive set of formalisms for comparing, evolving, and constraining primordial non-Gaussian models through the CMB bispectrum.  The primary goal was to directly calculate CMB bispectra from a general primordial model, enhancing methods previously outlined in ref.~\\cite{0612713}.   This was achieved by exploiting common features of primordial bispectra to reduce the dimensionality of the transfer functions required to evaluate the CMB bispectra.   The new innovations reported here include the use of the flat sky approximation when all $l_i\\ge 150$, greatly reducing computational times for most of the allowed region, and a cubic reparametrisation, significantly reducing the number of points required for accurate interpolation of the bispectrum. (We note that this CMB bispectrum code is being prepared for public distribution.) These methods make feasible the repetitive calculation of highly accurate CMB bispectra at Planck resolution without specific assumptions about the separability of the underlying primordial bispectrum.  Further, we have calculated the CMB bispectra for all the distinct primordial bispectrum shapes $S(k_1,k_2,k_3)$ currently presented in the literature, motivated by a range of inflationary and other cosmological scenarios.   We have presented these, plotted relative to the large-angle CMB bispectrum for the constant   primordial model ($S^ { const.}=1$) for which we presented an analytic solution (\\ref{eq:constbispectrum}).    The CMB bispectra from the different primordial models exhibit a close correspondence to their original shape modulated, of course, by the oscillatory transfer functions.  We were able to quantitatively determine the observational independence of the CMB bispectra by measuring their cross-correlations using the estimator (\\ref{eq:estimator}). These results revealed five independent classes of shapes which it should be possible to distinguish from one another with a significant detection of non-Gaussianity in future experiments such as Planck. These were the equilateral (\\ref{eq:equi}), local (\\ref{eq:local}), warm (\\ref{eq:warm}), and flat (\\ref{eq:flat}) shapes, all described in section~3, together with the feature model (\\ref{eq:feature}) which is basically the constant shape (\\ref{eq:constmodel}) plus broken scale-invariance.  Different models belonging within the same class will be difficult to distinguish, a fact best demonstrated by the 95\\% correlation of the equilateral shape with DBI (\\ref{eq:dbi}), ghost (\\ref{eq:ghost}) and single (\\ref{eq:single}) shapes.  We have also presented a shape correlator (\\ref{eq:shapecor}) which provides a fast and simple method for determining the independence of different shapes.   In particular, for highly correlated models, it yielded results accurately reflecting those of the full CMB correlator (\\ref{eq:cmbcor}), thus avoiding unnecessary calculation.   The shape correlator also reliably identified poorly correlated models, that is, new shapes for which a full CMB bispectrum analysis was warranted.   We also proposed a straightforward two-dimensional eigenmode analysis of shape functions, valid for nearly scale-invariant  models.   This allowed us to identify the shape correlations with products of eigenvalues of the non-orthogonal eigenmodes, immediately revealing, for example, why warm and local models are independent.  In principle, the analysis can guide the theoretical search for primordial models with distinctive non-Gaussian signatures.   It also revealed the constant eigenmode (or shape (\\ref{eq:constmodel})) as the primary cause of the cross-correlation between many models, suggesting strategies for distinguishing, for example, between local and equilateral models.  Finally, we proposed an alternative approach (\\ref{eq:normalise}) to the normalisation of the non-Gaussianity parameter $f_{NL}$ using the shape autocorrelator. This new normalisation allows us to employ $f_{NL}$ to systematically compare the true level of non-Gaussianity in different models. In contrast, with current methods, the constraints for competing models can vary by a factor of four or more, with bounds varying significantly for models even in the same class of highly correlated shapes.  A detection of non-Gaussianity would have profound consequences for our understanding of the early universe, uprooting the present simplest inflationary paradigm. The present work indicates that the next generation of CMB experiments (notably Planck) may be able to distinguish between different classes of shapes for primordial non-Gaussianity.   Delineating the bispectrum shape would provide important clues about viable alternative theories for the origin of large-scale structure in the universe."
    },
    "0812/0812.1416_arXiv.txt": {
        "abstract": "We present the result of near-infrared and optical observations of the BL Lac object S5~0716$+$714 carried out by the KANATA telescope. S5~0716$+$714 has both a long term high-amplitude variability and a short-term variability within a night. The shortest variability (microvariability) time-scale is important for understanding the geometry of jets and magnetic field, because it provides a possible minimum size of variation sources. Here, we report the detection of 15-min variability in S5~0716$+$714, which is one of the shortest time-scales in optical and near-infrared variations observed in blazars. The detected microvariation had an amplitude of $0.061{\\pm}0.005$~mag in $V$ band and a blue color of $\\Delta(V-J)=-0.025{\\pm}0.011$. Furthermore, we successfully detected an unprecedented, short time-scale polarimetric variation which correlated with the brightness change. We revealed that the microvariation had a specific polarization component. The polarization degree of the variation component was higher than that of the total flux. These results suggest that the microvariability originated from a small and local region where the magnetic field is aligned. ",
        "introduction": "Blazars are one of classes of active galactic nuclei (AGNs), whose relativistic jets are considered to be directed along the line of sight (Blandford \\& K\\\"onigl, 1979 , \\cite{Antonucci}). Blazars show a large and rapid variability, and their fluxes often vary by an order of magnitude in a time-scale of days. A short time-scale variability has also been reported, that is, so-called intra-day variability or microvariability, having a time-scale of $<1$~day (e.g. \\cite{Antonucci}).  The observation of the shortest time-scale variability gives us important information for understanding the variation mechanism, the emitting region and the jet geometry. An extremely large variability with a time-scale of tens of minutes has recently been reported in the optical region (e.g. \\cite{Xie} for eight blazars e.g. S5~0716$+$714). However, \\citet{Cellone07a} argued that they can be spurious results due to systematic errors arising from analysis. The shortest time-scale in blazar variability was, in general, believed to be several hours.  In gamma-ray region, however, a significant variability with a time-scale of minutes have recently been discovered by \\citet{Albert} in Mrk~501 and \\citet{Aharonian} in PKS2155$-$304. Also in the optical band, a 15-min microvariation has been reported by \\citet{Gupta} in Mrk~501. The amplitude of the microvariation was not so extremely large, but rather small, with 0.05~mag in $R$ band. Such a short time-scale indicates that the size of the emitting region of the variation is smaller than a Schwarzschild radius $R_{\\rm S}$ of the central black hole in AGN if we assume typical values for the black hole mass and the jet Doppler factor. It is, hence, possible that microvariations originate from a small confined region rather than a whole emitting area. On the other hand, observational features of those shortest time-scale variations have not been established due to the lack of detailed observations, for example, polarimetric and multi-wavelength observations.  The spectral energy distribution (SED) of blazars can be described by synchrotron radiation and inverse-Compton scattering. In the optical band, the synchrotron radiation is dominant. Variations of polarization can, hence, be observed in blazars (e.g. Angel \\& Stockman, 1980). The polarimetric observation provides an important clue for the magnetic field and its geometry. However, there is still few polarimetric observation of microvariability so far. \\citet{Cellone07b} observed hours-scale variations in the polarization of AO~0235$+$164. According to their observation, the temporal variation of polarization parameters showed no correlation with the variation of the total flux. For the shortest time-scale variations (${\\sim} 10$~min), there was no report studying possible correlation between the polarization and the total flux.  S5~0716$+$714 is a BL Lac object, and one of systems in which microvariability was observed. The variability has been observed on the time-scale of days and hours (e.g. \\cite{Wagner}, \\cite{Stalin06} and \\cite{Wu07}). \\citet{Nilsson08} estimated its redshift to be $z = 0.31 \\pm 0.08$.  In order to study the mechanism of the microvariability and to find the shortest time-scale variations, we performed photopolarimetric observations of S5~0716$+$714 in the optical and near-infrared (NIR) bands simultaneously. In this paper, we report our detection of 15-min variability with polarimetric variations. The 15-min variability has a small amplitude, but the light-curve is very well-sampled. The paper is arranged as follows: In section 2, we present the observation method and analysis. In section 3, we report the result of the photometry and polarimetry. We discuss the emitting region and the geometry of the magnetic field based on the observation in section 4. Finally, the conclusion is drawn in section 5. ",
        "conclusions": "We obtained multicolor photometric and polarimetric data of the shortest time-scale variation in S5~0716$+$714 for the first time. Our findings are summarized as below; first, the object became blue with ${\\Delta}(V-J)=-0.025$ during the bump. Second, the bump component has a specific polarization vector with the large polarization degree of $27\\%$, which is distinct from the overall brightening trend. It can be suggested that the emitting region of the microvariability in optical band is a small and local area compared with the whole optical emitting region in the jet. \\\\ \\\\ This work was partly supported by a Grant-in-Aid from the Ministry of Education, Culture, Sports, Science, and Technology of Japan (19740104)."
    },
    "0812/0812.1285_arXiv.txt": {
        "abstract": "Using Matrix Theory  as a concrete example of a fundamental holographic theory, we show that the emergent macroscopic spacetime displays a new macroscopic quantum structure, holographic geometry, and a new observable phenomenon, holographic noise, with phenomenology similar to that previously derived on the basis of a quasi-monochromatic wave theory. Traces of matrix operators on a light sheet with a compact dimension  of size $R$ are interpreted as transverse position operators  for macroscopic bodies. An effective quantum wave equation for spacetime is derived from the Matrix Hamiltonian.  Its solutions display eigenmodes that connect longitudinal separation and transverse position operators on macroscopic scales.   Measurements of transverse relative positions of macroscopically separated bodies, such as signals in Michelson interferometers, are shown to display holographic nonlocality, indeterminacy and noise, whose properties can be predicted with no parameters except $R$.  Similar results are derived using a detailed scattering calculation of the matrix wavefunction.  Current experimental technology will allow a definitive and precise test or validation of  this interpretation  of holographic fundamental theories.  In the latter case,  they will yield   a direct measurement of $R$ independent of the gravitational definition of the Planck length, and a direct measurement of the total number of degrees of freedom. ",
        "introduction": "If the world as a whole is a  quantum system, there should  be a quantum description of observables that correspond to properties of spacetime.  Arguments from gravitational thermodynamics \\cite{'tHooft:1993gx,Susskind:1994vu} and superstring theory \\cite{Maldacena:1997re,Witten:1995ex} suggest that the quantum description of spacetime may be holographic, that is, expressible as a local quantum theory with only two large spatial dimensions and a finite minimum length.  The large dimensions are interpreted as null surfaces or wavefronts in a dual, emergent three-dimensional world. Such theories are strongly motivated because by construction they respect holographic scaling of entropy and information \\cite{Bousso:1999xy, Bousso:2002ju}.  Holographic degrees of freedom statistically explain the laws of gravitational thermodynamics \\cite{Bekenstein:1972tm,Bardeen:gs,Bekenstein:1973ur, bek2, Hawking:1975sw,Jacobson:1995ab}: for example, the null surface represented by a black hole event horizon encodes the states of the hole, so that the total entropy is given by one quarter of the area in Planck units.   Recently   it  has been conjectured \\cite{Hogan:2007pk,Hogan:2008zw} that theories of this type generically create a macroscopic holographic quantum geometry with a new,  particular and precisely defined form of nonlocal  correlation:  an indeterminacy in transverse position that grows with scale $L$ like $\\sqrt{LR}$, where $R$ is some fundamental length scale.  Arguments based on wave mechanics, rather than a fundamental theory, were used to derive this generic phenomenology.  In this paper we show concretely how exactly such a macroscopic holographic quantum geometry arises in a specific candidate fundamental theory that includes quantum spacetime: Matrix Theory \\cite{Banks:1996vh}, a realization of M-theory, the strongly-coupled version of string theory \\cite{Witten:1995ex}.  We derive from Matrix Theory an effective wave theory, including an expression for  the fundamental   quantum eigenmodes of the spacetime wavefunction.   This macroscopic quantum structure is new; it shows explicitly how the quantum degrees of freedom of spacetime are holographically encoded, and how they manifest themselves in the three dimensional world.  They lead to new nonlocal correlations, including  a new and precisely characterised observable phenomenon, ``holographic noise''.  In the following section, we  adopt an interpretation of Matrix Theory in terms of macroscopic position observables. We  use wave mechanics to derive the low-frequency, large-longitudinal-separation limit of the theory and identify the quantum eigenstates of macroscopic spacetime, which describe   the observable relative positions of macroscopic bodies.  These eigenfunctions show a new relation between longitudinal and transverse position observables, different from the classical limit derived from field theory.  Position wavefunctions are shown to display the same surprisingly large transverse indeterminacy that was previously derived in \\cite{Hogan:2007pk,Hogan:2008zw} using a model based on classical wave correlations. This result supports the conjecture that holographic quantum geometry, characterised by large transverse indeterminacy, is a generic prediction of holographic unified quantum theories. In section III we support the assertions used to construct this wave model--- in particular, the interpretation of matrix traces as transverse position operators---  by evaluating the properties of a scattered wavefunction in Matrix theory.  The measurable phenomenon of holographic noise can be used to either rule out this interpretation of the physical meaning of Matrix theory,  or validate it.  In the latter case it will be possible to pursue an experimental program that studies properties of quantum geometry directly, including precise measurement of the fundamental frequency or minimum time. ",
        "conclusions": "The holographic noise derived here has previously been shown to approximately agree with the spectrum of currently unexplained continuum noise in the best operating interferometer, GEO600, in its most sensitive band, above about 500 Hz. Therefore, either the effect has already been detected,  or it is within reach to rule it out using existing technology. In the latter case, it  should be possible to exclude a broad class of theories that encompass the idea of holography; the position noise is directly related to the holographic information bounds around which these theories are constructed. The properties of holographic noise, such as its exact spectrum and its  shear character, are more specific than more generalized phenomenological descriptions of quantum spacetime noise (e.g., \\cite{AmelinoCamelia:2001dy,Smolin:2006pa}).  If the effect exists,  a diverse and detailed experimental program is motivated.  The effective theory for the degrees of freedom fixes both the  spectrum of  noise and the character of its geometrical correlations.  The spectrum of the noise depends only on $R$ and therefore provides a direct experimental window on the Planck scale. The expected signal correlations between nearby, but not coincident interferometers will be important to predict particular signatures that can distinguish holographic noise from other more prosaic sources: in-common signals occur, due to quantum wavefunction collapse, even between machines with no physical connection, and even on timescales too short to mimic other kinds of in-common disturbances, such as transmitted acoustic vibrations.  Eventually, other technologies such as atom interferometers will approach the sensitivity required to study holographic noise at smaller scales.   Interactions in that case occur with bodies with individual masses less than a Planck mass, in which case  a correct calculation of the collective behavior of the wavefunction may require a fuller account of the matrix degrees of freedom."
    },
    "0812/0812.3588_arXiv.txt": {
        "abstract": "{} {We present a sample of candidate quasars selected using the $KX$-technique. The data cover $0.68 \\deg^2$ of the X-ray Multi-Mirror (XMM) Large-Scale Structure (LSS) survey area where overlapping multi-wavelength imaging data permits an investigation of the physical nature of selected sources.} {The $KX$ method identifies quasars on the basis of their optical ($R$ and $z'$) to near-infrared (\\ks) photometry and point-like morphology. We combine these data with optical ($u^*,g'r',i',z'$) and mid-infrared ($3.6-24\\, \\mu m$) wavebands to reconstruct the spectral energy distributions (SEDs) of candidate quasars.} {Of 93 sources selected as candidate quasars by the $KX$ method, 25 are classified as quasars by the subsequent SED analysis. Spectroscopic observations are available for 12/25 of these sources and confirm the quasar hypothesis in each case.  Even more, 90\\% of the SED-classified quasars show X-ray emission, a property not shared by any of the false candidates in the KX-selected sample. Applying a photometric redshift analysis to the sources without spectroscopy indicates that the 25 sources classified as quasars occupy the interval $0.7 \\le z \\le 2.5$. The remaining 68/93 sources are classified as stars and unresolved galaxies.  } {} ",
        "introduction": "Despite the enormous progress that wide-field imaging surveys have provided in understanding quasar populations over the past decade, concerns remain over the extent to which observational biases contribute to a partial view of the quasar phenomenon.  Consequently, multi-wavelength observations are important not only to reproduce the spectral energy distribution (SED) of the objects under study, but also to provide as complete a census as is practical of the quasar population.  Nevertheless, selecting source populations using different energy bands encompasses a risk of confusion between different types of objects or different physical processes. Thus, studying the selection effects is an essential step to understanding the properties of the parent population.  A classic method to identify type-1 quasars in optical imaging surveys is known as the $UV$-excess ($UVX$; \\citealt{sandage65}, \\citealt{schmidt83}). The $UVX$ technique exploits the fact that quasars display an emission excess in blue wavebands compared to main sequence, late type Galactic stars, and therefore occupy a distinct (i.e. bluer) locus in a color-color diagram with respect to stars.  Based on a study of 323 radio-loud sources selected from the Parkes catalog~\\citep{parks90}, approximately 70\\% of which were identified spectroscopically as quasars, \\cite{webster95} claimed that up to 80\\% of quasars could be missed in optical surveys. Their suggestion was based on the $B_J - K$ color index of these radio-selected sources, which showed a much larger scatter ($1 < B_J - K < 8$) compared to the typical values for optically selected quasars.  Various interpretations were suggested, such as (a) intrinsic light absorption, (b) absorption from dust along the line of sight or (c) differences in the internal processes in the active nucleus.  Cases (a) and (b) would result in ``reddened\" quasars, whilst objects in category (c) would be considered as intrinsically ``red\" (in contrast to the ``classical\", blue type-1 quasars).  However, a number of studies have questioned the dusty quasar hypothesis.  \\cite{masci98} rejected the dust hypothesis and~\\cite{whiting01} suggesting that the $B_J-K$ spread could be accounted for by synchrotron emission. \\cite{benn98} claimed that the $B_J - K$ scatter could be attributed to observational effects such as variability, underestimated photometric uncertainties, etc. Although~\\cite{francis00} support the conclusions of~\\cite{webster95}, \\cite{richards03} demonstrate that, for the majority of the SDSS quasars, redder colors could be explained by intrinsic processes in the AGN and that the number of missed red and reddened quasars in surveys such as the SDSS (different from typical optical surveys because of the filter selection) comes down to ``only'' 15\\%. \\cite{brown06} found similar results, suggesting that red \\mbox{type-1} quasars correspond to about 20\\% of the type-1 population.  With the advent of large field, near infrared (NIR) surveys, \\cite{warren00} suggested a new colour excess criterion, designed to be less prone to the problem of dust reddening and extinction. Their method is the NIR analog of $UVX$: as the quasar SED follows a power-law, it displays a flux excess at NIR wavelengths compared to the Rayleigh-Jeans tail observed in stars. As a result quasars should also occupy a distinct region of the NIR color-color diagram, as they do in the $UBV$ color-color plane. The so-called $K$-excess criterion, hereafter $KX$, would serve as an alternative to the $UVX$ for selecting quasars at redshifts $z > 2.2$. At these redshifts, the UV/optical colors of quasars become virtually indistinguishable from those of late type stars.  Despite being a prominsing technique, relatively few $KX$-selected quasar samples have appeared in the literature \\---\\ mainly due to the considerable effort required to complete a large area NIR imaging survey. The conclusions drawn from the $KX$-selected quasar populations published to date are rather limited as various factors such as the small area of the survey~\\citep{croom01}, the relative shallow magnitudes sampled~\\citep{barkhouse01} or the inhomogeneity and incompleteness of the survey~\\citep{sharp02} have hampered an assessment of the potential of the method. The first detailed study of the $KX$ method was given by~\\cite{jurek07} while more recently~\\cite{smail08} also published a work on $KX$-selected quasars, raising the problem of contamination by foreground compact galaxies and their impact on spectroscopic follow-up observations.   In this paper we present an analysis of a sample of 93 $KX$-selected quasar candidates selected from 0.68 deg.$^2$ of \\rzk\\ imaging data located in the X-ray Multi-Mirror (XMM) Large Scale Structure (LSS) survey field \\citep{pierre04}.  The XMM LSS survey is an X-ray imaging survey covering approximately $10 \\deg^2$ to an approximate point-source flux limit of $8\\times 10^{-15} \\rm erg s^{-1} cm^{-2}$ in the [0.5-2]~keV energy band.  The XMM LSS is associated with a number of imaging data sets at different wavebands: the XMM LSS contains the XMM Medium Deep Survey (XMDS), a deeper X-ray imaging component over 2 deg.$^2$ in addition to the optical Canada-France-Hawaii Telescope Legacy Survey (CFHTLS) and the Spitzer Wide area IR Extragalactic (SWIRE) survey.  These additional data provide a panchromatic view of the quasar candidates and permit a detailed discussion of their physical nature.  The paper is organized as follows: Section 2 presents the optical and NIR data used in the $KX$ technique. Section 3 describes the multi-wavelength data available in the field. Section 4 describes the $KX$ selection criteria and discusses the physical nature of the selected sources using their multi-wavelength properties.  In Section 5 we draw some conclusions on the color properties of the confirmed quasars selected using $KX$ and comment on the overall efficiency of the technique. We also consider some of the challenges to be overcome in applying the $KX$ technique to large data sets. ",
        "conclusions": ""
    },
    "0812/0812.3541_arXiv.txt": {
        "abstract": "The form of the primordial power spectrum has the potential to differentiate strongly between competing models of perturbation generation in the early universe and so is of considerable importance. The recent release of five years of WMAP observations have confirmed the general picture of the primordial power spectrum as deviating slightly from scale invariance with a spectral tilt parameter of $n_{\\rm s} \\sim 0.96$. Nonetheless, many attempts have been made to isolate further features such as breaks and cutoffs using a variety of methods, some employing more than $\\sim 10$ varying parameters. In this paper we apply the robust technique of Bayesian model selection to reconstruct the \\emph{optimal} degree of structure in the spectrum. We model the spectrum simply and generically as piecewise linear in $\\ln$ $k$ between `nodes' in $k$-space whose amplitudes are allowed to vary. The number of nodes and their $k$-space positions are chosen by the Bayesian evidence so that we can identify both the complexity and location of any detected features. Our optimal reconstruction contains, perhaps, surprisingly few features, the data preferring just three nodes. This reconstruction allows for a degree of scale dependence of the tilt with the `turn-over' scale occuring around $k \\sim 0.016$ Mpc$^{-1}$. More structure is penalised by the evidence as over-fitting the data, so there is currently little point in attempting reconstructions that are more complex. ",
        "introduction": "The recent release by the Wilkinson Microwave Anisotropy Probe (WMAP) of five years of observations have confirmed that the primordial spectrum of density perturbations is consistent with being purely adiabatic and \\emph{close} to scale invariant, in perfect harmony with the simplest inflationary scenarios. This agreement appears remarkably robust when extended to independent datasets such as measures of the matter power spectrum from galaxy redshift surveys \\citep{TegmarkII}. Alternative models of the spectrum containing various features have been considered.  These include an exponential large scale cutoff \\citep{efstathioub} to explain the quadrupole power decrement, and theoretically motivated spectra to model the inflationary potential \\citep{Contaldi2007} or account for discontinuities from early universe phase transitions \\citep{Barriga}. Reconstructions of the spectrum, limiting \\emph{a priori} assumptions about its structure, have typically involved fitting some basis functions, such as wavelets \\citep{Wang}, some deconvolution method (\\citealt{Souradeep}; \\citealt{Silk}) or directly `binning' the spectrum into an arbitrary number of band powers \\citep{Bridle}. However, most previous methods fail to account for Occam's razor since they assume that more complexity, and typically more `detected' features, are necessarily important in explaining the data. Recently \\citet{Verde07} reconstructed the spectrum, while minimising the level of complexity needed via a cross-validation with a `hold-out' portion of the data. This approach is a timely progression, but in this paper we attempt a more statistically robust procedure with an optimal reconstruction using the Bayesian evidence to decide how much detail one should fit and where it is located in $k$-space, based solely on the data. ",
        "conclusions": "\\label{section:conclusions} In this paper we have attempted to fit an optimal degree of structure to the primordial power spectrum using Bayesian model selection tools as our discriminant criteria. We find that a scale invariant spectrum is significantly ruled out, the data instead favouring a tilted spectrum, with perhaps some slight scale dependence of $n_{\\rm s}$ located close to $k \\sim 0.01$ Mpc$^{-1}$. We fail to find any support in the data for further features beyond this simple scenario, the optimal reconstruction fitting between just two and three parameters. Previous authors (including ourselves) have regularly used many more degrees of freedom, finding a number of `interesting' features in the process. In this analysis, by accounting for Occams' razor we have found no statistically significant structure, much beyond a simple tilt, and there is, we feel, limited point in attempting more complex models at present, as the data simply cannot support them."
    },
    "0812/0812.1896_arXiv.txt": {
        "abstract": "{The symbiotic system \\object{HM\\,Sagittae} consists of a Mira star and a secondary White Dwarf component. The dust content of the system was severely affected by the nova outburst in 1975, which is still ongoing. The capabilities of optical interferometry operating in the mid-IR allow us to investigate the current geometry of the dust envelope.} {We test our previous spectro-interferometric study of this system with new interferometric configurations, increasing the $uv$ coverage and allowing us to ascertain the appearance of the source between 8 and 13~$\\mu$m.} {We used the MIDI instrument of the VLTI with the unit telescopes (UTs) and auxiliary telescopes (ATs) providing baselines oriented from PA=42$^{\\circ}$ to 127$^{\\circ}$. The data are interpreted by means of an elliptical Gaussian model and the spherical radiative transfer code {\\tt DUSTY}.} {\\rm We demonstrate that the data can be reproduced well by an optically thick dust shell of amorphous silicate, typical of those encountered around Mira stars, whose measured dimension increases from 8 to 13~$\\mu$m. We confirm that the envelope is more extended in a direction perpendicular to the binary axis. The level of elongation increases with wavelength in contrast to our claim in a previous study.} {The wider $uv$ coverage allows us to deepen our previous investigations of the close circumstellar structure of this object.} ",
        "introduction": "\\label{intro}  \\object{HM\\,Sge} is a D-type (dust dominated IR) symbiotic system that erupted as a symbiotic nova in 1975 \\citep{dokuchaeva76}, evolving from a 17$^{\\rm th}$ to a 11$^{\\rm th}$ magnitude star with a rich emission-line spectrum. The cool component is a Mira star and the hot component is a White Dwarf (WD), which had escaped detection being enshrouded in the dense Mira envelope before this dramatic event. \\\\ Previous works by \\citet{schild01} and \\citet{bogdanov01} were published on \\object{HM\\,Sge} the same year, evaluating the parameters of the dust in the system by fitting the ISO/SWS data using the {\\tt DUSTY} radiative transfer code \\citep{ivezic99}. Both groups used single and double component models but did not reach identical conclusions. \\citet{schild01} favored a 2-shell model, which accounted for the effect of the White Dwarf on the dust, whereas \\citet{bogdanov01} proposed a single-shell model. Since the extensions of the dusty shell differ from one model to another, mid-infrared interferometry data enabled \\citet{sacuto07} to determine the most suitable of the proposed models. The authors showed that the morphology of the dust was accounted for well by the single shell model of \\citet{schild01}, which included a compact and optically thick circumstellar environment, typical of a Mira star. However, because of its high optical depth, the model is inconsistent with the presence of non-absorbed emission of the Mira in the near-IR, and \\citet{sacuto07} indicated that this issue could be solved by assuming a more complex spatial distribution of dust. We attempt to confirm the morphology of the dusty envelope using observations acquired with five new interferometric baselines secured in 2007. ",
        "conclusions": "\\label{concl}  The wider $uv$ coverage data has allowed us to confirm that the dust envelope around the Mira is described well by an optically thick dust shell that consists predominantly of amorphous silicates. We also confirmed the longer extension of the circumstellar envelope in a direction perpendicular to the binary axis, and the increase in the size of the structure between 8 and 12.5 $\\mu$m. It has also been demonstrated that the amount of elongation is higher at 12.5 $\\mu$m and steadily decreases toward 8 $\\mu$m, but the ratio between the two axis remains limited at most to 0.85, implying only a slight departure from spherical symmetry. This behavior can naturally be interpreted. The layers emitting most of the 8 $\\mu$m flux, located closer to the photosphere, are the least affected by the White Dwarf, situated far (60 AU; \\citealt{eyres01}) from the Mira. At 13 $\\mu$m, the interferometer probes farther regions that are more affected by the energetic but diluted wind originating in the erupting White Dwarf. We also note that the new baselines, recorded in 2007, range from 60 to 89m whereas the previous ones ranged from 32 to 59m and probed outer regions of the dusty envelope, which were potentially more affected by the White Dwarf. Therefore, we expect the amount of elongation to be higher, as reported in \\citet{sacuto07}, although the poor $uv$ coverage of the 2005-2006 observations does not allow us to estimate it accurately.  These MIDI/VLTI interferometric observations probe a compact region that includes the giant star and its vicinity to a radius of about 30 mas, set by the smallest baselines ($\\sim$32m). There is a clear requirement for near and mid-IR high dynamics, imaging at the diffraction limit of a 8-10m class telescopes, to probe the inner nebula (0.8$\\arcsec$ wide) already studied with the HST \\citep{eyres01} and in the radio \\citep{richards99}.  We are close to the limits of the imaging capabilities of the VLTI used with the two telescopes recombiner MIDI for studying such a complex source. It is also uncertain whether the three telescopes recombiner AMBER, operating in the near-IR, can provide a substantial increase in information about the morphology of the system, because the (spherical) Mira and hot dust in the inner wind still dominate. A significant improvement in the quality of image reconstruction must await the advent of the second generation instrument MATISSE \\citep{lopez06}."
    },
    "0812/0812.1402_arXiv.txt": {
        "abstract": "We study four scenarios for the SCP~06F6 transient event that was announced recently. Some of these were previously briefly discussed as plausible models for SCP~06F6, in particular with the claimed detection of a $z=0.143$ cosmological redshift of a Swan spectrum of a carbon rich envelope. We adopt this value of $z$ for extragalactic scenarios. We cannot rule out any of these models, but can rank them from most to least preferred. Our favorite model is a tidal disruption of a CO white dwarf (WD) by an intermediate-mass black hole (IMBH). To account for the properties of the SCP~06F6 event, we have to assume the presence of a strong disk wind that was not included in previous numerical simulations. If the IMBH is the central BH of a galaxy, this explains the non detection of a bright galaxy in the direction of SCP~06F6. Our second favorite scenario is a type Ia-like SN that exploded inside the dense wind of a carbon star. The carbon star is the donor star of the exploded WD. Our third favorite model is a Galactic source of an asteroid that collided with a WD. Such a scenario was discussed in the past as the source of dusty disks around WDs, but no predictions exist regarding the appearance of such an event. Our least favorite model is of a core collapse SN. The only way we can account for the properties of SCP~06F6 with a core collapse SN is if we assume the occurrence of a rare type of binary interaction. ",
        "introduction": "SCP~06F6 is an optical transient discovered by Barbary et al. (2008). It raised to an optical peak flux of $F_{\\rm peak} \\simeq 2.5 \\times 10^{-14} \\erg \\s^{-1} \\cm^{-2}$ in $\\sim 100~$days, and then declined in another $\\sim 100~$days. Barbary et al. (2008) ruled out a microlensing event, and noted that if SCP~06F6 is associated with the galaxy cluster CL 1432.5+3332.8 at a redshift of $z=1.112$, then its peak luminosity is like that of the most bright SNe observed to date. In this case, this object projected distance from the center of the cluster would be $290 \\kpc$, but (1) there is no obvious host galaxy, making the association questionable (a faint host is possible at a projected distance of 12kpc for this redshift; Barbary et al. 2008). (2) The spectrum of this object is not like any other SN.  Gaensicke et al. (2008a) presented a fit of the spectrum of SCP~06F6 with a spectrum of a carbon-rich star at a redshift of $z=0.143$. They also argue that the lightcurve looks somewhat like that of a supernova type II. Based on that they suggest SCP~06F6 represents a new class of core collapse SNe. However, this explanation shares the weak points (1) and (2) above, and in addition: (3) The X-ray luminosity reported by Gaensicke et al. (2008a) is two orders of magnitude above that of a typical SN. (4) The spectral fit is not perfect.  We discuss alternative possible scenarios for the SCP~06F6 transient. In section 2 we describe the critical observations that set constraints on any model to explain the transient. In section 3 we discuss a Galactic WD-asteroid merger model. Barbary et al. (2008) considered as well the possibility that the progenitor was a cool WD in our galaxy. In that case the WD progenitor should be quite cold, and its distance should be $D>1.2 \\kpc$, assuming the WD is not colder than $3000 \\K$. In section 4 we discuss three possible extragalactic models. In addition to the core collapse SN model (Gaensicke et al. 2008a) we discuss a type Ia-like SN scenario, where the exploding WD is enshrouded in a carbon-rich nebula formed by the wind from the WD's giant companion. We also discuss a possible disruption of a WD passing close to an intermediate-mass black hole (IMBH; Rosswog et al. 2008b,c). We compare the four models with observations and summarize in section 5. ",
        "conclusions": "We compared four scenarios for the SCP~06F6 transient event (Barbary et al. 2008), one galactic and three extragalactic (with a redshift of $z=0.143$ as suggested by Gaensicke et al. 2008a):  1. \\emph{A Galactic event of an asteroid colliding with a WD},  2. \\emph{A new class of core collapse SN}, as suggested by Gaensicke et al. (2008a),  3. \\emph{Type Ia SN explosion}, resembling the hybrid Ia/IIn type,  4. \\emph{Tidal disruption of a CO WD by an intermediate mass black hole (IMBH)}, as recently studied in detail by Rosswog et al. (2008b,c) and Lodato et al. (2008).  In principle, all four scenarios can lead to a transient event, and we expect that such events will eventually be discovered. However, as we presently attempt to explain SCP~06F6, we compare the models to the observations of this transient in Table 1 \\begin{table}  Table 1: Comparison of the models  \\footnotesize \\bigskip \\begin{tabular}{|l||l|l|l|l|} \\hline Model      & Extragalactic    & Extragalactic   & Extragalactic       & Galactic WD-  \\\\ & core collapse SN & SN Ia-like      & IMBH-WD collision   & asteroid merger \\\\ \\hline Reference  &  Gaensicke       & This paper      & Rosswog             & This paper\\\\ & et al. (2008a)   &                 & et al. (2008c)      & \\\\ \\hline \\hline (1) Spectrum  & \\multicolumn{3}{|l|}{Like a redshifted carbon star with absorption Swan bands} & Absorption \\\\ & \\multicolumn{3}{|l|}{}                               & Swan bands \\\\ & \\multicolumn{3}{|l|}{}                               & distorted by  \\\\ & \\multicolumn{3}{|l|}{}                               & magnetic fields \\\\ \\hline (2) Optical & \\multicolumn{2}{|l|}{Within the range of SNe}  & Fits the Eddington      & Determines       \\\\ Luminosity  & \\multicolumn{2}{|l|}{}                         & luminosity of an IMBH   & the colliding    \\\\ & \\multicolumn{2}{|l|}{}                         & tidally disrupting a WD  & asteroid mass    \\\\ \\hline (3) Timescale & Typical of SNe,  & Typical of SNe, & Typical accretion        & KH timescale   \\\\ & but the event    & some SNe Ia     & time, but the decline    & of merged      \\\\ & declines faster  & have similar    & is faster than predicted & product        \\\\ & than core        & decline rates   &                          &                \\\\ & collapse SNe     &                 &                          &                \\\\ \\hline (4) Very bright&\\multicolumn{2}{|l|}{$\\times 100$ more than in typical SNe.} &\\multicolumn{2}{|l|}{X-rays from accretion disk} \\\\ in X-rays      &\\multicolumn{2}{|l|}{Requires X-rays to come from}           &\\multicolumn{2}{|l|}{and shocked expelled material} \\\\ &\\multicolumn{2}{|l|}{shocked polar ejecta}                   &\\multicolumn{2}{|l|}{} \\\\ \\hline (5) Undetected &\\multicolumn{2}{|l|}{Very faint host galaxy}  & Allowed BH mass     & Very cool WD \\\\ progenitor     &\\multicolumn{2}{|l|}{}                        & range accounts for  & progenitor \\\\ &\\multicolumn{2}{|l|}{}                        & a very faint galaxy &     \\\\ \\hline Strong points & a) Peak        & a) Spectrum:    & a) Spectrum: origin    & a) Spectrum:      \\\\ in explaining & optical        & carbon from     & of carbon is explained & origin of carbon  \\\\ SCP~06F6      & luminosity     & companion's     & b) Luminosity          & is explained      \\\\ & b) Timescale   & dense wind      & c) Time scale          & b) Luminosity     \\\\ &                & b) Peak         & d) Strong X-rays       & fits an asteorid  \\\\ &                & optical         & e) If BH is at         & c) Timescale      \\\\ &                & luminosity      & the nucleus the model  & d) Strong X-rays  \\\\ &                & c) Timescale    & accounts for a very    &                   \\\\ &                & d) Faint galaxy & faint galaxy           &                   \\\\ &                & is possible     & & \\\\ \\hline Weak  points  & a) Spectrum:        & a) Mass loss    & a) Observed decline is                 & a) Spectrum:    \\\\ in explaining & C-rich envelope     & rate extreme    & steeper than predicted                 & very strong     \\\\ SCP~06F6      & hard to explain     & for an AGB      & b) Difficult to account                & magnetic fields \\\\ & b) Strong X-rays    & star            & for a $60 \\AU$ ($\\sim 10^5 R_{\\rm g}$) & are required    \\\\ & c) No star-         & b) Special      & photosphere                            & to explain      \\\\ & forming galaxy      & geometry        & {\\it Note: An addition of}             & the location    \\\\ & is observed         & is required     & {\\it a disk wind solves}               & of Swan bands   \\\\ & d) Special          &                 & {\\it these problems}                   &       \\\\ & geometry            &                 &                                        &       \\\\ & is required         &                 &                                        &       \\\\ \\hline \\end{tabular}  \\footnotesize \\bigskip  \\normalsize \\end{table} with the following comments (numbers as in the first column of the table): \\newline (1) Redshifted carbon star-spectrum does better than what we can find for a Galactic source. A Galactic source requires absorption Swan bands distorted by strong magnetic fields, possibly with addition of some emission Swan bands, which also puts its distance to be $\\sim 1.5 \\kpc$. Despite its complexity, we cannot rule out such a Galactic source. The IMBH-WD disruption model with no nuclear burning, i.e., a CO WD seems to do better in accounting for the spectrum. A SN embedded in C-rich matter is also fine, but the origin of carbon needs explanation. The Ia-like SN scenario (SN Ia/IIn) seems more natural than core collapse, which would require a very unusual binary interaction. \\newline (2) For the Galactic source, total emitted energy is used to constrain the mass of the asteroid. For the extragalactic sources, the luminosity is within the range of SNe. For the IMBH-WD disruption model there is as yet no constraint from theory or observation on the luminosity. \\newline (3) As far as timescales, the SN models have the closest match, although the observed decline is faster than those of most core collapse SNe. Some SNe type IIn have equally fast decline in the IR bands and the SN Ia-like model (Ia/IIn) does good as well. The WD-asteroid model (where the event timescale is the Kelvin-Helmholtz timescale), were not observed in the past, but from theoretical point of view can explain the timescale. As for the IMBH-WD tidal disruption model, we must incorporate a new ingredient to existing calculations of the disruption process, which we suggest is a strong disk wind. \\newline (4) In the Galactic source the X-rays originate from close to the WD, as in dwarf novae. This can be the X-ray source at late times, when the space above and below the disk has been cleared from dense gas. Another possibility is the ejection of jets (or a collimated fast wind), which when collide with the circum-system material become the X-ray source. The same two mechanisms can operate in the IMBH-WD tidal interaction model. It seems that the SNe models should also employ a non-spherical geometry that allows fast polar ejecta to be shocked and emit X-rays. This further disfavors a core collapse SN based on a single star model; we can safely conclude that if SCP~06F6 is a core collapse SN, its peculiarity comes from a strong binary interaction. We recall that the peculiarity in the SN Ia-like model comes from a carbon star with a high mass loss rate, that also transfers mass to the pre-exploding WD. Namely, in the SNe models of SCP~06F6 we attribute the peculiarity to a rare binary companion. \\newline (5) The non detection of a galaxy at the location of the event is more of a problem for core collapse SN models because the galaxy is fainter than known star-forming galaxies. Less of a problem to SN Ia-like, and even less of a problem for the IMBH-WD disruption model, as a galaxy with no recent star formation can do with these two models. Moreover, if we assume that the IMBH is the BH at the nucleus of the galaxy, then a very faint dwarf galaxy is indeed expected. The non detection is not a problem for the WD-asteroid merger model, as it is used to constrain the progenitor distance to $D >1.2  \\kpc$ (Barbary et al. 2008).  We summarize the strong and weak points of the four models in the last two rows of Table 1. The most important thing to note is that IMBH-WD tidal disruption model has weak points if we use the existing numerical simulations for this process as they are. However, we can solve these problems if we add a very strong disk outflow at velocities of ${\\rm several} \\times 10^3 \\km \\s^{-1}$. (i) \\emph{Photosphere}: With a mass of $0.1-0.3 M_\\odot$ being lost at $3000-10^4 \\kms$, we can account for the photosphere (defined at $\\tau=2/3$) at tens of AU, and in particular of $60 \\AU$ at maximum light. (ii) \\emph{Lightcurve}: With the rapid removal of mass from the disk, and the interaction of this wind with inflowing gas, we can account for the lightcurve declining rate that is much more rapid than $t^{-5/3}$. This outflow is powered by accretion close to the Eddington limit, and possibly by extreme amplification of magnetic fields in the accretion disk.  For that, our favorite model is the IMBH-WD model. Second we rank the type Ia-like SN model, with a carbon star mass donor companion. Third we rank the WD-asteroid Galactic model. The core collapse model comes last, as it seems it must be very peculiar, and involve some kind of strong binary interaction.  The predictions of the different models are straightforward, allowing for an observational test. The Galactic scenario implies that a faint WD, with a possible IR excess, resides at the event location. The extra-galactic scenarios predict a very faint galaxy in that direction. The type II SN model implies the presence of a recently active star forming region."
    },
    "0812/0812.0318_arXiv.txt": {
        "abstract": "The extremely bright optical flash that accompanied GRB 080319B suggested, at first glance, that  the prompt $\\gamma$-rays in this burst were  produced by Synchrotron self Compton (SSC).   We analyze here the observed optical and $\\gamma$ spectrum. We find that the very strong optical emission poses, due to self absorption, very strong constraints on the emission processes  and put  the origin of the optical emission at a very large radius, almost inconsistent with internal shock. Alternatively it requires a very large random Lorentz factor for the electrons. We find that  SSC could not have produced the prompt $\\gamma$-rays. We also show that the optical emission and the $\\gamma$ rays could not have been produced by synchrotron emission from two populations of electron within the same emitting region. Thus we must conclude that the optical and the  $\\gamma$-rays were produced in different physical regions.   A possible interpretation of the observations is that the $\\gamma$-rays arose from internal shocks but the optical flash  resulted from external shock emission. This would have been consistent with the few seconds  delay observed between the optical and $\\gamma$-rays signals. ",
        "introduction": "The radiation mechanism for the $\\gamma$-ray burst (GRB) afterglow is widely accepted as synchrotron emission \\citep[see][for a review]{P04}. However, the mechanism for the prompt emission is still uncertain. Most recently, \\citet{km08} showed the inconsistency with the overall synchrotron model. On the other hand \\citet{psz08} found that  Synchrotron self-Compton (SSC) cannot explain the prompt emission unless the prompt optical emission  is  very high. With a naked eye (5th magnitude) optical flash \\citep{Cwiok08, Karpov08a}  GRB 080319B \\citep{Racusin08b}, was a natural candidate for SSC \\citep{kp08}.  GRB 080319B  was located at redshift $z=0.937$ \\citep{Vreeswijk08}. Its duration $T_{90}$ was $\\sim 57$s. The peak flux is $F_p \\sim 2.26 \\pm 0.21 \\times 10^{-5} {\\rm erg \\,cm^{-2}s^{-1}}$ and the peak of the $\\nu F_{\\nu}$ spectrum $E_p \\simeq 675{\\pm 22}$keV (i.e., $\\nu_p \\sim 1.6\\times 10^{20}$Hz, and consequently $F_{\\nu,p} \\sim 1.4 \\times 10^{-25} {\\rm erg\\, cm^{-2} Hz^{-1} s^{-1}}$).  The photon indexes below and above  $E_p$ are $-0.855^{+0.014}_{-0.013}$ and $-3.59^{+0.32}_{-0.62}$ respectively \\citep{r08}. Choosing standard cosmological parameters $H_0=70 {\\rm km}\\,{\\rm s}^{-1}\\,{\\rm Mpc}^{-1}, \\Omega_m=0.3, \\Omega_\\lambda=0.7$, GRB 080319B had  a peak luminosity $L_p \\sim 9.67\\times 10^{52} {\\rm erg}\\,{s}^{-1}$ and an isotropic equivalent energy $E_{\\rm \\gamma, iso} \\simeq 1.32 \\times 10^{54}$ erg \\citep{ga08}.  The optical observations were going on even before the onset of the GRB because TORTORA was monitoring the same region of sky at that moment. \\citet{Karpov08a} reported the optical V-band light curve in the prompt phase (from $\\sim$ -10 s to $\\sim$ 100 s). Variability was evident and there were at least 3 or 4 pulses in the optical light curve. The peak V-band magnitude reached 5.3, corresponding to a flux density $\\sim 28.7$ Jy.   Models for the optical emission accompanying the prompt $\\gamma$-rays have been extensively discussed since the early work of \\citet{Katz94} who used  a low energy spectrum $F_\\nu \\propto \\nu^{1/3}$ and found that the prompt optical emission could be bright to 18th mag. The observations of a prompt strong optical flash from GRB 990123 \\cite{Akerlof99} lead to a wave of interest in this phenomenon.  The $t^{-2}$ decline of the early optical flash favored a reverse shock model \\citep{sp99,mr99} that was suggested just few month earlier \\citep{SP99a}. Optical emission from later internal shocks and residual internal shocks have been discussed by \\citet{wyf06} and \\citet{lw08} respectively, and the late internal shocks model has been used to interpret the optical flares detected in GRB 990123, GRB 041219a, GRB 050904 and GRB 060111B \\citep{Akerlof99,Blake05,Boer06,Klotz06,zdx06}. However, none of the previous bursts had the data as plentiful as that of GRB 080319B and the constraints on the models were not very tight.  The very strong optical flash of  GRB 080319B lead naturally to the suggestion \\citet{kp08,r08} that  internal shocks  synchrotron  produced the prompt optical emission while  SSC  produced the prompt  $\\gamma$-rays.  We show here that self absorption of the optical poses major constraints on the source of the optical emission and we consider its implications on general models (including SSC) for the emission of this burst. The paper is structured as follows: In section 2 we consider the general constraints that arise from the optical emission. In section 4 we consider SSC model and in section 5  we make general remarks on Inverse Compton (IC).  In section 6 we examine the possibility that the optical and $\\gamma$-rays arose from two synchrotron emitting populations of electrons but in the same physical region. We find that the optical and the $\\gamma$-rays were originated in physically different regimes. We discuss the implications of these results in section 6. ",
        "conclusions": "\\citet{psz08} have shown recently that typical GRBs  with normal (weaker than 12th magnitude) prompt  optical emission cannot be produced by SSC of a softer  component.  The unique burst GRB 080319B had a very luminous prompt optical emission and one could expect, at first glance that the prompt $\\gamma$-rays were produced in SSC of the optical prompt emission. This was the accepted interpretation in the discovery paper \\citep{r08} as well as in several others \\citep{kp08,fp08}.   The numerous detailed  observations of this burst led to the hope that one can determine the physical parameters within the emitting regions here. However, a  careful analysis reveals a drastically different picture. There is no reasonable SSC solution. One can make an even stronger statement and argue that it is unlikely that the prompt $\\gamma$-rays are produced by inverse Compton scattering of seed photons produced by a source with no relativistic bulk motion relative to the IC electrons (regardless of the origin of the seed photons).  We also considered a situation with two populations of electrons that co-exist in the same emitting region. Both populations emit synchrotron radiation with the less energetic electrons producing the optical while the more energetic ones produce the $\\gamma$-rays. We have shown that even though the total energy released in the optical is orders of magnitude lower than the energy released in $\\gamma$-rays the lower energy electrons population should carry more energy - making, once more, the model energetically inefficient.  Combining these two results we conclude that   the optical emission and the soft $\\gamma$-rays do  not come from a single origin. This is the main conclusion of this letter. Once we relax the condition that both modes of prompt emission are produced in the same region, there are  many possibilities. However, even here the  limits obtained in section \\ref{sec:opt} indicate that the optical emission is produced at a very large radius which is most likely incompatible with internal shocks (Note that typical internal shocks will have a rather modest values of $\\gamma_e$.). This raises the possibility that the optical emission is produced in this burst by the early external shock (possibly by the reverse shock). While there are several problems with this model (in particular the fast rise of the optical emission that is faster than expected in this case \\cite{np04}) they are less severe than those  basic issues that arise with the internal shocks.  The fact that the optical emission follows to some extend the $\\gamma$-rays but with a time delay of several seconds \\citep{Guidorzi08} is consistent with this model."
    },
    "0812/0812.0632_arXiv.txt": {
        "abstract": "Measuring the 3D distribution of mass on galaxy cluster scales is a crucial test of the $\\Lambda$CDM model, providing constraints on the nature of dark matter.  Recent work investigating mass distributions of individual galaxy clusters (e.g. Abell 1689) using weak and strong gravitational lensing has revealed potential inconsistencies between the predictions of structure formation models relating halo mass to concentration and those relationships as measured in massive clusters.  However, such analyses employ simple spherical halo models while a growing body of work indicates that triaxial 3D halo structure is both common and important in parameter estimates.  We recently introduced a Markov Chain Monte Carlo (MCMC) method to fit fully triaxial models to weak lensing data that gives parameter and error estimates that fully incorporate the true shape uncertainty present in nature. In this paper we apply that method to weak lensing data obtained with the ESO/MPG Wide-Field Imager for galaxy clusters A1689, A1835, and A2204, under a range of Bayesian priors derived from theory and from independent X-ray and strong lensing observations. For Abell 1689, using a simple strong lensing prior we find marginalized mean parameter values $M_{200} = (0.83 \\pm 0.16)\\times10^{15}$ $h^{-1}$M$_{\\odot}$ and $C=12.2\\pm 6.7$, which are marginally consistent with the mass-concentration relation predicted in $\\Lambda CDM$.  With the same strong lensing prior we find for Abell 1835 $M_{200} =(0.67 \\pm 0.22)\\times10^{15}$ $h^{-1}$M$_{\\odot}$ and $C=7.1^{+ 10.6}_{-7.1}$, and using weak lensing information alone find for Abell 2204 $M_{200} =(0.50 \\pm 0.19)\\times10^{15}$ $h^{-1}$M$_{\\odot}$ and $C=7.1\\pm 6.2$.  The large error contours that accompany our triaxial parameter estimates more accurately represent the true extent of our limited knowledge of the structure of galaxy cluster lenses, and make clear the importance of combining many constraints from other theoretical, lensing (strong, flexion), or other observational (X-ray, SZ, dynamical) data to confidently measure cluster mass profiles.  If we assume CDM to be correct and apply a mass-concentration prior derived from CDM structure formation simulations, we find \\{$M_{200} =(0.99 \\pm 0.18)\\times10^{15}$ $h^{-1}$M$_{\\odot}$; $C=7.7\\pm 2.1$\\}, \\{$M_{200} =(0.68 \\pm 0.19)\\times10^{15}$ $h^{-1}$M$_{\\odot}$; $C=4.4\\pm 1.6$\\}, and \\{$M_{200} =(0.45 \\pm 0.13)\\times10^{15}$ $h^{-1}$M$_{\\odot}$; $C=5.0\\pm 1.7$\\} for A1689, A1835, and A2204 respectively. ",
        "introduction": "\\begin{figure*} \\centering \\rotatebox{0}{ \\resizebox{16cm}{!} {\\includegraphics{f1a.eps}\\includegraphics{f1b.eps}\\includegraphics{f1c.eps}} } \\caption{34' $\\times$ 34' $R$-band wide-field images of Abell 1689, 1835, and 2204, from the ESO/MPG WFI at La Silla.} \\label{fig:plot1} \\end{figure*} Galaxy clusters are ideal laboratories in which to study dark matter, being the most massive bound structures in the universe and dominated by their dark matter component ($\\sim90\\%$). Constraining the clustering properties of dark matter is crucial for refining structure formation models that predict both the shapes of dark matter halos and their mass function (e.g. \\cite{navarro97}; \\cite{evrard02}; \\cite{bahcall03}; \\cite{dahle06}).  Several methods are used to measure galaxy cluster dark matter profile shapes and halo masses on a range of scales, including X-ray and Sunyaev-Zeldovich (SZ) studies, dynamical analyses, and gravitational lensing.  However, all of these methods typically require simplifying assumptions to be made regarding the shape and/or dynamical state of the cluster in order to derive meaningful constraints from available data.  Crucially, while most parametric methods typically assume spherical symmetry of the halo, observed galaxy clusters often exhibit significant projected ellipticity and halos in CDM structure formation simulations (e.g. \\cite{bett07} (using the Millennium simulation); \\cite{shaw06}) show significant triaxiality in cluster-scale halos, with axis ratios between minor and major axes as small as 0.3.  Understanding and accurately incorporating the impact of this on cluster mass and parameter estimates is crucial for accurate comparisons between measured cluster properties and model predictions (e.g. the $\\Lambda$CDM mass-concentration relation).  Efforts to understand the impact of triaxiality in gravitational lensing and its potential role in explaining apparent discrepancies with CDM began when \\cite{oguri05} applied a fully triaxial NFW model to the shear map of Abell 1689 to find that it is consistent with $6\\%$ of cluster-scale halos. These continued as \\cite{gavazzi05} showed that a triaxial NFW can reconcile parameter values derived from observations of the cluster MS2137-23 to predictions from N-body simulations.  More generally, \\cite{corless07} demonstrated in the weak lensing regime that neglecting halo triaxiality in parameterised fits of NFW models to dark matter halos with axis ratios significantly less than one can lead to over- and underestimates of up to 50$\\%$ and a factor of 2 in halo mass and concentration, respectively.  While extreme cases of triaxiality are rare, such halos can be much more efficient lenses than their more spherical counterparts, especially when in configurations that hide most of the triaxial shape along the line of sight.  \\cite{oguri08} recently quantified this expectation, showing that the strongest lenses in the universe are expected to be a highly biased population, strongly favouring orientations along the line of sight, high levels of triaxiality, and apparent sphericity in projection; these biases correspond to a population of very triaxial prolate (``cigar''-shaped) halos oriented with their long axes approximately aligned along the line-of-sight.  Thus, for strong lensing clusters -- even those that are circular in projection on the sky such as Abell 1689 -- spherical symmetry is an unjustified and error-inducing assumption.  Further, even halos with less extreme axis ratios are inaccurately fit by spherical models.  We expect the vast majority of galaxy cluster scale dark matter halos to be triaxial, and so it is important to include this expectation in model fitting.  Moreover, \\cite{corless08} (herein CK08) showed that triaxiality leaves no quantifiable signature in lensing data (i.e. in the  maximum-likelihood values obtained fitting models), and so it is generally not possible to determine {\\it a priori} the expected impact of triaxiality on mass and concentration estimates for a given lensing halo.   In the past, triaxial models have not generally been fit to weak lensing data because they cannot easily be well-constrained, largely due to the intrinsic limitations of lensing.  Because lensing is determined only by the projected mass density and shear of the underlying mass distribution, it is inherently impossible to fully constrain a 3D structure without imposing strong priors on the shape of the halo or supplementing lensing data with other data types more sensitive to line-of-sight halo structure, such as dynamical information.  \\begin{table} \\centering \\caption{Properties of the three lensing clusters} \\label{tab:table1} \\begin{tabular}[t!]{ccccccccccccc} \\hline Cluster&RA&Dec&z&n [arcmin$^{-2}$]&$\\sigma_{\\epsilon}$\\\\ \\hline A1689&13 11 29.5&-01 20 17&0.1832&7&0.34\\\\ A1835&14 01 02.0&+02 51 32&0.2532&12&0.42\\\\ A2204&15 32 45.7&+05 34 43&0.1522&15&0.44\\\\ \\hline \\end{tabular} \\end{table}  Despite these difficulties, that triaxiality has been convincingly demonstrated to be an important factor in model parameter estimation in lensing analyses make it important to directly include it in NFW fits to galaxy clusters. For studies of individual clusters doing so is crucial if claims regarding the validity of the $\\Lambda$CDM paradigm based on NFW parameter estimates are to be meaningfully evaluated.  It is also of importance across populations to obtain more accurate distributions of galaxy cluster parameters, and especially in measuring the galaxy cluster mass function and constraining the scatter in mass-observable relations.  If the mass function is significantly sloped as expected (e.g. \\cite{evrard02}), even the best-case scenario of a symmetric scatter of mass estimates due to neglected triaxiality would lead to an asymmetric shift in the calculated mass function, because there are more low mass halos to shift up in mass than there are high mass halos to shift down (see e.g. \\cite{corless08b}).  In CK08 we introduced a Bayesian Markov Chain Monte Carlo (MCMC) method for fitting triaxial NFW models to weak gravitational lensing data.  We demonstrated that it accurately recovers mean parameter values and errors across a population of triaxial dark matter halos similar to those seen in CDM structure formation simulations.  In this paper we make the first application of that triaxial fitting method to the weak lensing signals of three well-known galaxy clusters: Abell 1689, Abell 1835, and Abell 2204.  We reanalyze preexisting lensing data obtained with 34x34 arcmin observations of the three cluster fields with the ESO/MPG Wide Field Imager, first presented and modelled with spherical mass models in \\cite{clowe01}, \\cite{clowe02}, and \\cite{king02a}.  The MCMC method allows for the simple and statistically robust imposition of prior probability function on the parameters of the triaxial model.  Though some would rather avoid priors altogether, simple models include hidden priors stronger than those we employ in this work; for example, a ``simple'' spherical model includes implicit $\\delta$-function priors on the triaxial axis ratios.  Priors are of such importance in this problem because models of 3D structures fit using only 2D lensing data are intrinsically underconstrained.  CK08 showed that carefully chosen priors are crucial to obtaining statistically accurate results across a population, but that slight offsets between the true distribution of parameters in nature and the distributions employed as priors lead only to a very small decline in the accurate performance of the triaxial fitting method.  We therefore continue in this paper the investigation of the behaviour and best use of various prior probability functions in fitting the weak lensing signals of galaxy clusters, in specific application to A1689, A1835, and A2204.   An additional advantage of the MCMC method is that it allows for the straightforward and statistically-robust inclusion of additional constraints from other independent and complementary observations such as strong lensing, SZ, X-ray and spectroscopic studies.  To demonstrate that capability, we investigate the inclusion of constraints from strong lensing and X-ray observations of the three clusters, imposing the external constraints as prior probability functions on the cluster model.  Section 2 briefly describes the observational data, reviews lensing by triaxial dark matter halos and the MCMC triaxial fitting method, and discusses the various priors employed in this paper.  Section 3 presents our results, and Section 4 concludes with some discussion. Throughout we assume a concordance cosmology with $\\Omega_m=0.3$, $h=0.7$, and a cosmological constant $\\Omega_{\\Lambda} = 0.7$.  \\begin{figure*} \\centering \\rotatebox{0}{ \\resizebox{13cm}{!} {\\includegraphics{f2.eps}} } \\caption{Prior probability distributions of axis ratios (Shaw), virial mass (Mass), mass-concentration (C), and triaxial line-of-sight orientation angle ($\\theta$), all derived from predictions of $\\Lambda$CDM structure formation or lensing simulations.} \\label{fig:plot2} \\end{figure*}  \\begin{table*} \\centering \\caption{The marginalized mean parameter values obtained fitting single triaxial NFWs to A1689, A1835, and A2204 under a range of prior probability distributions on the halo mass, concentration, shape and orientation derived from CDM structure formation simulations and external observational constraints, described in Sec. \\ref{sec:priors}.  $1\\sigma$ errors on the estimate of the mean parameter values are given; all values are for a concordance cosmology of $\\Omega_m = 0.3$, $\\Omega_{\\Lambda}=0.7$, $h=0.7$.} \\label{tab:table2} \\begin{tabular}[t!]{ccccccccccccc} \\hline Prior&$M_{200} [10^{15}M_{\\odot}]$&$C$&$a$&$b$&$M_{ES}$&$C_{ES}$\\\\ \\hline A1689&&&&\\\\ Flat&$1.31 \\pm 0.28$&$15.4\\pm 9.0$&$0.55\\pm 0.19$&$0.82\\pm0.14$&$1.28^{+0.30}_{-0.35}$&$15.3\\pm9.1$\\\\ Shaw&$1.27 \\pm 0.25$&$13.7\\pm 7.7$&$0.74\\pm 0.09$&$0.85\\pm0.08$&$1.27^{+0.26}_{-0.26}$&$13.7\\pm7.7$\\\\ Mass&$1.15 \\pm 0.24$&$19.5\\pm 12.3$&$0.56\\pm 0.19$&$0.81\\pm0.15$&$1.13^{+0.26}_{-0.30}$&$19.4^{+12.3}_{-12.4}$\\\\ C&$1.72 \\pm 0.37$&$4.0\\pm 0.5$&$0.53\\pm 0.20$&$0.71\\pm0.20$&$1.65^{+0.41}_{-0.48}$&$4.0^{+0.6}_{-0.7}$\\\\ MC&$1.41 \\pm 0.25$&$7.7\\pm 2.1$&$0.59\\pm 0.19$&$0.78\\pm0.17$&$1.38^{+0.27}_{-0.31}$&$7.7^{+2.1}_{-2.2}$\\\\ Sphere&$1.27\\pm0.24$&$13.3\\pm7.6$&&&&\\\\ $\\theta$&$1.18 \\pm 0.23$&$12.2\\pm 6.7$&$0.55\\pm 0.20$&$0.74\\pm0.18$&$1.16^{+0.25}_{-0.30}$&$12.1^{+6.8}_{-6.8}$\\\\ $\\theta_E$&$1.18 \\pm 0.20$&$18.8\\pm 6.5$&$0.55\\pm 0.20$&$0.81\\pm0.14$&$1.16^{+0.22}_{-0.25}$&$18.7^{+6.6}_{-6.7}$\\\\ $\\theta + \\theta_E$&$1.07 \\pm 0.18$&$15.3\\pm 5.2$&$0.55\\pm 0.20$&$0.74\\pm0.18$&$1.05^{+0.19}_{-0.22}$&$15.2^{+5.2}_{-5.3}$\\\\ $\\theta_E$ + MC&$1.19 \\pm 0.18$&$12.7\\pm 2.6$&$0.54\\pm 0.20$&$0.71\\pm0.20$&$1.17^{+0.20}_{-0.24}$&$12.6^{+2.7}_{-2.8}$\\\\ \\hline A1835&&&&\\\\ Flat&$1.10 \\pm 0.43$&$8.0^{+10.9}_{-8.0}$&$0.51\\pm 0.19$&$0.77\\pm0.16$&$1.07^{+0.45}_{-0.47}$&$7.9^{+11.0}_{-7.9}$\\\\ Shaw&$1.01 \\pm 0.32$&$7.3^{+10.9}_{-7.3}$&$0.73\\pm 0.09$&$0.84\\pm0.09$&$1.01^{+0.32}_{-0.32}$&$7.2^{+10.9}_{-7.2}$\\\\ Mass&$0.83 \\pm 0.28$&$13.3^{+17.6}_{-13.3}$&$0.52\\pm 0.19$&$0.76\\pm0.17$&$0.81^{+0.29}_{-0.30}$&$13.2^{+17.7}_{-13.2}$\\\\ C&$1.27 \\pm 0.41$&$3.5\\pm 0.6$&$0.53\\pm 0.19$&$0.78\\pm0.16$&$1.22^{+0.44}_{-0.47}$&$3.5^{+0.6}_{-0.7}$\\\\ MC&$0.97 \\pm 0.27$&$4.4\\pm 1.6$&$0.55\\pm 0.18$&$0.78\\pm0.16$&$0.94^{+0.29}_{-0.32}$&$4.4^{+1.6}_{-1.7}$\\\\ Sphere&$1.01\\pm0.30$&$6.8^{+9.5}_{-6.8}$&&&&\\\\ $\\theta$&$0.95 \\pm 0.31$&$7.1^{+10.6}_{-7.1}$&$0.51\\pm 0.20$&$0.74\\pm0.18$&$0.92^{+0.34}_{-0.35}$&$7.0^{+10.6}_{-7.0}$\\\\ X-ray&$1.17 \\pm 0.36$&$3.9\\pm 1.0$&$0.54\\pm 0.19$&$0.78\\pm0.16$&$1.13^{+0.39}_{-0.43}$&$3.8^{+1.0}_{-1.1}$\\\\ X-ray $+\\theta$&$0.98 \\pm 0.25$&$3.7\\pm 1.0$&$0.53\\pm 0.20$&$0.74\\pm0.18$&$0.95^{+0.28}_{-0.32}$&$3.7^{+1.0}_{-1.1}$\\\\ \\hline A2204&&&&\\\\ Flat&$0.72 \\pm 0.27$&$7.1\\pm 6.2$&$0.52\\pm 0.20$&$0.81\\pm0.15$&$0.70^{+0.29}_{-0.35}$&$7.1\\pm6.2$\\\\ Shaw&$0.68 \\pm 0.21$&$6.8\\pm 5.8$&$0.74\\pm 0.09$&$0.85\\pm0.08$&$0.67^{+0.22}_{-0.22}$&$6.8\\pm5.8$\\\\ Mass&$0.56 \\pm 0.20$&$10.4\\pm 9.7$&$0.53\\pm 0.20$&$0.79\\pm0.16$&$0.54^{+0.21}_{-0.26}$&$10.4^{+9.8}_{-9.7}$\\\\ C&$0.80 \\pm 0.25$&$4.1\\pm 0.8$&$0.54\\pm 0.20$&$0.82\\pm0.14$&$0.78^{+0.27}_{-0.29}$&$4.1^{+0.8}_{-0.9}$\\\\ MC&$0.64 \\pm 0.19$&$5.0\\pm 1.7$&$0.54\\pm 0.20$&$0.81\\pm0.15$&$0.62^{+0.20}_{-0.23}$&$5.0\\pm1.8$\\\\ Sphere&$0.69\\pm0.21$&$6.6^{+6.7}_{-6.6}$&&&&\\\\ \\hline \\end{tabular} \\end{table*}  \\begin{table*} \\centering \\caption{The approximate maximum-likelihood parameter values obtained fitting single triaxial NFWs to A1689, A1835, and A2204 under a range of prior probability distributions on the halo mass, concentration, shape and orientation derived from CDM simulations and external observational constraints, described in Sec. \\ref{sec:priors}.  $1\\sigma$ (68\\%) errors on the maximum-likelihood parameter values are given; all values are for a concordance cosmology of $\\Omega_m = 0.3$, $\\Omega_{\\Lambda}=0.7$, $h=0.7$.} \\label{tab:table3} \\begin{tabular}[t!]{ccccccccccccc} \\hline Prior&$M_{200} [10^{15}M_{\\odot}]$&$C$&$a$&$b$&$M_{ES}$&$C_{ES}$&$\\chi^2$&Reduced $\\chi^2$\\\\ \\hline A1689&&&&\\\\ Flat&$1.52^{+0.96}_{-0.59}$&$12.4^{+24.2}_{-7.7}$&$0.31^{+0.68}_{-0.20}$&$1.00^{+0.00}_{-0.70}$&$1.38^{+1.09}_{-0.79}$&$12.0^{+24.6}_{-8.0}$&5953.59&1.04084\\\\ Shaw&$1.38^{+0.63}_{-0.48}$&$10.6^{+18.7}_{-5.3}$&$0.81^{+0.15}_{-0.29}$&$0.89^{+0.11}_{-0.25}$&$1.37^{+0.63}_{-0.52}$&$10.6^{+18.7}_{-5.4}$&5953.6&1.04084\\\\ Mass&$1.22^{+0.83}_{-0.39}$&$11.8^{+31.7}_{-6.7}$&$0.38^{+0.62}_{-0.26}$&$0.86^{+0.14}_{-0.56}$&$1.14^{+0.90}_{-0.60}$&$11.6^{+31.9}_{-7.1}$&5953.85&1.04088\\\\ C&$1.17^{+1.49}_{-0.32}$&$3.6^{+2.4}_{-0.9}$&$0.19^{+0.81}_{-0.09}$&$0.31^{+0.69}_{-0.01}$&$0.88^{+1.79}_{-0.42}$&$3.3^{+2.7}_{-1.1}$&5958.32&1.04167\\\\ MC&$1.11^{+1.01}_{-0.28}$&$6.0^{+6.0}_{-2.8}$&$0.16^{+0.84}_{-0.06}$&$0.34^{+0.66}_{-0.04}$&$0.85^{+1.27}_{-0.40}$&$5.5^{+6.5}_{-2.8}$&5955.05&1.04109\\\\ Sphere&$1.33^{+0.44}_{-0.34}$&$10.0^{+9.3}_{-4.0}$&&&&&5954.12&1.0402\\\\ $\\theta$&$1.27^{+0.67}_{-0.55}$&$7.6^{+16.7}_{-4.4}$&$0.49^{+0.50}_{-0.39}$&$0.84^{+0.16}_{-0.54}$&$1.23^{+0.72}_{-0.83}$&$7.5^{+16.8}_{-4.9}$&5953.97&1.0409\\\\ $\\theta_E$&$1.31^{+0.67}_{-0.44}$&$16.9^{+18.8}_{-8.5}$&$0.32^{+0.68}_{-0.20}$&$0.96^{+0.04}_{-0.64}$&$1.21^{+0.77}_{-0.59}$&$16.5^{+19.3}_{-8.9}$&5952.58&1.04066\\\\ $\\theta + \\theta_E$&$1.10^{+0.51}_{-0.37}$&$11.8^{+12.8}_{-5.9}$&$0.49^{+0.50}_{-0.39}$&$0.78^{+0.22}_{-0.48}$&$1.07^{+0.54}_{-0.61}$&$11.7^{+12.9}_{-6.6}$&5952.58&1.04066\\\\ $\\theta_E$ + MC&$0.98^{+0.66}_{-0.21}$&$8.6^{+9.5}_{-1.5}$&$0.19^{+0.80}_{-0.09}$&$0.35^{+0.65}_{-0.05}$&$0.81^{+0.83}_{-0.29}$&$8.0^{+10.0}_{-1.9}$&5952.41&1.04063\\\\ \\hline A1835&&&&\\\\ Flat&$1.32^{+2.01}_{-0.73}$&$3.9^{+13.2}_{-2.8}$&$0.32^{+0.67}_{-0.22}$&$0.96^{+0.04}_{-0.65}$&$1.14^{+2.19}_{-0.91}$&$3.7^{+13.4}_{-2.9}$&8604.79&1.02134\\\\ Shaw&$1.19^{+0.70}_{-0.59}$&$3.6^{+10.8}_{-2.5}$&$0.73^{+0.23}_{-0.22}$&$0.83^{+0.17}_{-0.22}$&$1.18^{+0.72}_{-0.63}$&$3.5^{+10.8}_{-2.5}$&8605.71&1.02145\\\\ Mass&$0.86^{+1.05}_{-0.40}$&$4.0^{+22.0}_{-3.2}$&$0.50^{+0.49}_{-0.41}$&$0.96^{+0.04}_{-0.66}$&$0.82^{+1.09}_{-0.66}$&$3.9^{+22.0}_{-3.3}$&8607.37&1.02165\\\\ C&$1.20^{+2.12}_{-0.50}$&$3.4^{+1.5}_{-0.9}$&$0.35^{+0.65}_{-0.24}$&$0.94^{+0.06}_{-0.62}$&$1.05^{+2.28}_{-0.68}$&$3.3^{+1.7}_{-1.2}$&8604.99&1.02136\\\\ MC&$1.02^{+0.80}_{-0.49}$&$3.6^{+4.0}_{-1.8}$&$0.44^{+0.56}_{-0.34}$&$0.97^{+0.03}_{-0.67}$&$0.94^{+0.88}_{-0.70}$&$3.5^{+4.1}_{-2.1}$&8610.09&1.02197\\\\ Sphere&$1.07^{+0.55}_{-0.39}$&$3.6^{+5.4}_{-2.0}$&&&&&8605.82&1.02098\\\\ $\\theta$&$0.96^{+1.26}_{-0.53}$&$2.5^{+10.9}_{-1.8}$&$0.45^{+0.55}_{-0.34}$&$0.75^{+0.25}_{-0.45}$&$0.89^{+1.33}_{-0.74}$&$2.5^{+10.9}_{-2.0}$&8604.31&1.02128\\\\ X-ray&$1.06^{+1.68}_{-0.43}$&$4.3^{+1.7}_{-2.3}$&$0.41^{+0.58}_{-0.30}$&$0.93^{+0.07}_{-0.61}$&$0.98^{+1.76}_{-0.67}$&$4.2^{+1.8}_{-2.6}$&8605.83&1.02146\\\\ X-ray $+\\theta$&$1.04^{+0.98}_{-0.59}$&$3.9^{+2.0}_{-2.5}$&$0.14^{+0.85}_{-0.03}$&$0.44^{+0.56}_{-0.14}$&$0.73^{+1.29}_{-0.53}$&$3.5^{+2.4}_{-2.4}$&8605.55&1.02143\\\\ \\hline A2204&&&&\\\\ Flat&$0.79^{+1.07}_{-0.44}$&$3.6^{+13.9}_{-2.9}$&$0.33^{+0.66}_{-0.24}$&$0.98^{+0.02}_{-0.68}$&$0.69^{+1.18}_{-0.57}$&$3.4^{+14.0}_{-3.0}$&12470.7&1.05594\\\\ Shaw&$0.72^{+0.54}_{-0.35}$&$4.6^{+11.0}_{-3.1}$&$0.78^{+0.18}_{-0.27}$&$0.87^{+0.13}_{-0.24}$&$0.72^{+0.55}_{-0.37}$&$4.6^{+11.0}_{-3.2}$&12474.3&1.05625\\\\ Mass&$0.64^{+0.84}_{-0.35}$&$4.2^{+20.0}_{-3.6}$&$0.29^{+0.71}_{-0.19}$&$0.93^{+0.07}_{-0.63}$&$0.54^{+0.94}_{-0.45}$&$4.0^{+20.2}_{-3.6}$&12481.5&1.05686\\\\ C&$0.79^{+1.07}_{-0.44}$&$3.6^{+2.8}_{-1.2}$&$0.33^{+0.65}_{-0.24}$&$0.98^{+0.02}_{-0.68}$&$0.68^{+1.18}_{-0.51}$&$3.4^{+2.9}_{-1.5}$&12470.7&1.05594\\\\ MC&$0.64^{+0.76}_{-0.32}$&$4.2^{+5.4}_{-2.5}$&$0.29^{+0.71}_{-0.19}$&$0.93^{+0.07}_{-0.62}$&$0.54^{+0.86}_{-0.40}$&$4.0^{+5.6}_{-2.7}$&12481.5&1.05686\\\\ Sphere&$0.71^{+0.38}_{-0.26}$&$4.5^{+5.4}_{-2.4}$&&&&&12473.7&1.05584\\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "We first compare our fits under the Spherical prior to those of previous work on the three clusters -- that used weak lensing, strong lensing, and X-ray data -- to ascertain the agreement of our lensing catalogue and analysis with other methods and observations.  Once the accuracy of our method and catalogues is so established, we then move on to discuss the impact of triaxiality on the measure of the mass and concentration of A1689, A1835, and A2204. \\begin{figure*} \\centering \\rotatebox{0}{ \\resizebox{16cm}{!} {\\includegraphics{f9.eps}} } \\caption{The black contours give the 68 and 95 percent confidence contours for fits under a Spherical prior to the three clusters from this work, compared with the results of other authors fitting spherical models. Circles plot weak lensing-only results, diamonds weak lensing IR results, squares combined weak \\& strong lensing analyses, stars strong lensing-only measurements, and crosses X-ray results. The results and authors are compiled in Table \\ref{tab:table4}.  The Spherical contours do not represent the best estimates from this work, but are given for more direct comparison with past results to establish the consistency of the MCMC method with previous work.} \\label{fig:plot9} \\end{figure*}  \\subsection{Previous work} Table \\ref{tab:table4} collects the NFW $M_{200}$ and C values published in past work on the three clusters, all converted to a concordance cosmology with $h=0.7$; \\cite{limousin07} and \\cite{umetsu08} have previously summarised the existing work on Abell 1689 in a similar way.  Figure \\ref{fig:plot9} plots most of those results on top of the confidence contours for the fits under a Spherical prior in this work.  Though the Spherical prior fits are NOT our best estimates of the cluster parameters, they are used here in order to more directly compare our work to existing work in order to establish the consistency of the method.  Once this comparison establishes our results as consistent and reliable, we will move on to discuss the fully triaxial results.  \\subsubsection{Abell 1689} The concentration of Abell 1689 has been a source of significant controversy; \\cite{limousin07} offers an excellent and thorough discussion of the systematic issues that may affect concentration estimates in the case of this remarkable cluster.  Crucially, contamination of lensing catalogues with very faint cluster members is expected to lead to underestimates of the concentration when magnitude cuts alone are used to distinguish background from foreground and cluster galaxies.  Such was the method used in the early weak lensing studies of Abell 1689 in CS01 and \\cite{king02a}.  A version of that same catalogue is employed here, similar to that used in C03, improved with two bands of additional colour data and small improvements to the KSB PSF corrections.  The mean lensing redshift of the background galaxy population is changed in this analysis from $<z_s>=1$ to $<z_s>=0.8$, to account for the more stringent colour cuts in the galaxy population and make use of the improved photo-z information now available from the COSMOS field.  Applying the same fitting technique as in CS01 and C03 to the updated catalogue, we find maximum-likelihood parameter values of $M_{200} = 1.47 \\times 10^{15}$ M$_{\\odot}$ and $C=11.1$.  The higher mass value compared to the C03 fit can be attributed primarily to the change in the mean background source redshift assumed in the analysis: overestimating the mean background redshift as was done in C03 shifts the $M_{200}-C$ contours to lower masses.  In our MCMC maximum-likelihood fit under a spherical prior, we obtain marginalized mean parameter values $M_{200} = 1.27\\pm0.24\\times 10^{15}$ M$_{\\odot}$ and $C=13.3\\pm7.6$, and approximate maximum-likelihood parameters $M_{200} = 1.33^{+0.44}_{-0.34} \\times 10^{15}$ M$_{\\odot}$ and  $C=10.0^{+9.3}_{-4.0}$ which are consistent with this most recent result at $1\\sigma$.  The concentration values have increased with every new iteration of the catalogue, as expected as the level of contamination by faint cluster members decreases.  \\cite{king02b} performed a weak lensing analysis employing infrared data.  While the infrared provides an order of magnitude fewer background galaxies, it allows for a more secure determination of the cluster red sequence and removal of foreground galaxies.  The mass and concentration yielded in that study are both lower than those we find here; however, that is expected because that analysis did not include the $\\sim 1.1$ shear correction of \\cite{bacon01}.  Our marginalized spherical mass agrees well with those of the weak lensing results of \\cite{bardeau05} and \\cite{limousin07}, both using CFHT observations of the cluster, and is generally lower than those using the Subaru observation.  However, our mean concentration is generally higher than those inferred from the CFHT observations, and more consistent with those of the combined weak and strong analyses of \\cite{umetsu08} and \\cite{broadhurst05b}, though not so high as the recent weak-lensing only \\cite{medezinski07} result, which all employ the Subaru data.  Overall, our results fit well into the body of weak lensing and combined strong \\& weak lensing work that exists on the cluster.  Compared to the strong lensing-only results, our spherical mass is low compared to those found using ACS data, as with all of the weak lensing results.  Compared to the X-ray result, our mass and concentration are both high, though marginally consistent at $1\\sigma$.  This is unsurprising as \\cite{andersson04} find evidence for an ongoing or recent merger near the cluster centre, a finding supported by strong lensing and flexion (\\cite{leonard07}) measurements as well, and argue that X-ray analyses will underestimate the total mass in such cases.   \\subsubsection{Abell 1835} The catalogue employed for Abell 1835 is improved from that published in CS02 and C03, as the addition of two colour bands from CFHT observations allowed for a much more reliable discrimination between cluster, foreground, and background galaxies.  Using the catalogue constructed using a magnitude cut CS02 reported NFW parameter values $M_{200} = 1.44 \\times 10^{15}$ M$_{\\odot}$ and $C=2.96$.  Using the new catalogue, and a lower mean source redshift $<z_s>=0.8$, compared to the value of $<z_s>=1.0$ used in the CS02 and C03 analyses, we find for a spherical NFW marginalized mean parameters $M_{200} =1.01\\pm0.30\\times 10^{15}$ M$_{\\odot}$ and $C=6.8^{+9.5}_{-6.8}$ and approximate maximum-likelihood parameters $M_{200} =1.07^{+0.55}_{-0.39}\\times 10^{15}$ M$_{\\odot}$ and $C = 3.6^{+5.4}_{-2.0}$.  Our recovered concentration is higher, as expected given the significant improvement in the removal of faint galaxy contamination with the addition of two colour bands.  The recovered mass, however, is lower, though still comfortably consistent with our new results at $1\\sigma$.  Looking to other weak lensing studies, our spherical mass is lower than that of \\cite{bardeau07}, and our concentration higher.  However, our mass and concentration are in agreement with the \\cite{dahle06} result at $1\\sigma$.  Further, \\cite{smith05} find a total projected mass within 500 kpc of the cluster centre of $M=(2.9\\pm0.6) \\times 10^{14}$ $h^{-1}$M$_{\\odot}$, for an Einstein-De Sitter cosmology (they argue converting to a concordance cosmology would induce a change of no more than $10\\%$, less than the error bars).  For the same angular radius, our model in the concordance cosmology gives a projected mass of $M=3.0\\times 10^{14}$ $h^{-1}$M$_{\\odot}$ under a Flat prior, and for more direct comparison, $M=2.7\\times 10^{14}$ $h^{-1}$M$_{\\odot}$ under a Spherical prior, both in good agreement with the Smith result.  Our spherical mass and concentration estimates also agree very well with the parameter estimates derived from Chandra observations under assumptions of spherical symmetry (\\cite{schmidt01}). \\cite{zhang08} studied Abell 1835 with XMM and found a mass of $M_{500} = (5.90\\pm1.72) \\times 10^{14}$ $h^{-1}$M$_{\\odot}$ (where $M_{500}$ is the mass contained within a sphere with mean density 500 times the critical density); our marginalized mean model under a Flat prior gives $M_{500} = 5.8 \\times 10^{14}$ $h^{-1}$M$_{\\odot}$, also in good agreement. \\begin{table*} \\centering \\caption{Mass and Concentration estimates from previous work.  All lensing- and X-ray- derived values assumed spherical symmetry, and all X-ray analyses also assumed hydrostatic equilibrium.  Values in parentheses were fixed during fitting.} \\label{tab:table4} \\begin{tabular}[t!]{ccccl} \\hline $M_{200}$ $[10^{15}$ M$_{\\odot}]$&$C$&Method&Instrument&Author\\\\ \\hline A1689&&&&\\\\ 1.47&11.10&WL&ESO/MPG WFI&CS01 method with current catalogue\\\\ $1.76\\pm0.20$&$10.7^{+4.5+}_{-2.7}$&WL&Subaru&\\cite{umetsu08}\\\\ --&$22.1^{+2.9}_{-4.7}$&WL&Subaru&\\cite{medezinski07}\\\\ $1.55^{+0.31}_{-0.27}$&$4.28\\pm0.82$&WL&CFHT&\\cite{bardeau07}\\\\ $1.48\\pm0.22$&$7.6\\pm1.6$&WL&CFHT&\\cite{limousin07}\\\\ $1.07^{+0.46}_{-0.36}$&$3.5^{+0.5}_{-0.3}$&WL&CFHT&\\cite{bardeau05}\\\\ 1.13&7.9&WL&ESO/MPG WFI&C03\\\\ 1.07&5.7&WL (Infrared)&NTT SOFI&\\cite{king02b}\\\\ 0.85&4.8&WL&ESO/MPG WFI&\\cite{king02a}\\\\ $1.86\\pm0.16$&$10.1^{+0.8}_{-0.7}$&WL + SL&Subaru + ACS&\\cite{umetsu08}\\\\ $2.25\\pm0.20$&$7.6\\pm0.5$&WL + SL&Subaru + ACS&\\cite{halkola06}\\\\ $1.72\\pm0.19$&$10.9^{+1.1}_{-0.9}$&WL + SL&Subaru + ACS&\\cite{broadhurst05b}\\\\ --&$6.0\\pm0.6$ ($3\\sigma$)&SL&ACS&\\cite{limousin07}\\\\ $3.05\\pm0.30$&$6.0\\pm0.5$&SL&ACS&\\cite{halkola06}\\\\ 3.19&$6.5^{+1.9}_{-1.6}$&SL&ACS&\\cite{broadhurst05a}\\\\ $0.89^{+0.42}_{-0.32}$&$7.7^{+1.7}_{-2.6}$&X-ray&XMM&\\cite{andersson04}\\\\ \\hline A1835&&&&\\\\ $2.01^{+0.37}_{-0.33}$&$2.58\\pm0.48$&WL&CFHT&\\cite{bardeau07}\\\\ $1.33\\pm0.71$&(4.63)&WL&NOD/UHT&\\cite{dahle06}\\\\ 1.44&2.96&WL&ESO/MPG WFI&CS02\\\\ 0.83&--&X-ray&XMM&\\cite{majerowicz02}\\\\ $1.04^{+1.52}_{-0.59}$&$3.90^{+0.53}_{-0.63}$&X-ray&Chandra&\\cite{schmidt01}\\\\ \\hline A2204&&&&\\\\ $1.33\\pm0.82$&(5.03)&WL&NOD/UHT&\\cite{dahle06}\\\\ 0.86&6.3&WL&ESO/MPG WFI&CS02\\\\ \\hline \\end{tabular} \\end{table*}  \\subsubsection{Abell 2204} Though the colour selection in the catalogue used in this work is the same as that in CS02 and C03, as with the other two catalogues the PSF correction is improved and a lower mean background source redshift is employed to better reflect the expected mean redshift of the population after the colour selection.  Our marginalized mean parameters ($M_{200} =0.69\\pm0.21\\times 10^{15}$ M$_{\\odot}$; $C=6.6^{+6.7}_{-6.6}$) and maximum-likelihood parameters ($M_{200} =0.71^{+0.38}_{-0.26}\\times 10^{15}$ M$_{\\odot}$; $C=4.5^{+5.4}_{-2.4}$) under a Spherical prior agree with those of CS02 at $1\\sigma$.  Our mass is lower than that reported by \\cite{dahle06}, but within their (large) $1\\sigma$ errors.  \\cite{zhang08} studied Abell 2204 with XMM and found an X-ray mass $M_{500} = (4.12\\pm1.12) \\times 10^{14}$ $h^{-1}$M$_{\\odot}$; our marginalized model under a Flat prior gives $M_{500} = 3.7 \\times 10^{14}$ $h^{-1}$M$_{\\odot}$, fully consistent with the X-ray result.  Thus, though there continues to be a significant scatter in the parameter values returned for all three clusters, our parameter estimates under a spherical prior or in projection (where 3D structure does not matter) are all comfortably situated within the errors of previous work utilizing weak lensing, strong lensing, and X-ray observations.  The improvements in background galaxy discrimination are apparent in the increased concentrations of Abell 1689 and Abell 1835.  We are therefore confident our catalogues and MCMC sampler are returning reasonably accurate estimates of the mass profiles of these clusters, and understand how those profiles compare to those derived in previous work.  Now we are equipped to examine the mass profiles of each of the clusters freed from the unphysical assumption of spherical symmetry.   \\subsection{Triaxial parameter estimates} The parameter estimates for halo mass and concentration using the triaxial NFW model should provide a better estimate of the true parameters and errors of lensing galaxy clusters, as the model more closely reflects our true beliefs regarding the shape and structure of clusters.  Generally, we find that triaxial models under the most general Flat prior return higher mass and concentration estimates than their spherical counterparts, but that otherwise the impact of triaxial fitting under other more stringent priors, whether drawn from theory or independent observational constraints, can shift parameters estimates in either direction from the typical spherical values.  The question of which prior is most appropriate for a given question remains open, and crucially important for this underconstrained problem.  Generally, in the case where only weak lensing information is available, the most general Flat prior seems most useful, as it makes minimal assumptions about the halo shapes and parameter relations.  Until the question of overconcentration is resolved, it seems wise to avoid any prior that ties the concentration to the mass, as this may hide real discrepancies.  Given that, the application even of a less controversial Mass prior is unwise, because the shape of the mass-concentration degeneracy automatically requires higher concentration values as mass values are pushed to their lower limits.  This artificially exacerbates the mass-concentration relation discrepancy.  A mixed $MC$ prior may be of great use in the future if the relation is better understood and constrained, as it can provide simple physical constraints on this underconstrained problem without requiring independent observational data.  However, it is only valid once it represents our true prior expectation for galaxy cluster structure, a status which current observational data does not yet confer.  Similarly, employing a spherical prior imposes a prior constraint that does not represent our true best predictions for cluster structures.  The $M_{200}-C$ contours show that generally the contours under the Spherical prior are enclosed by those of the Flat triaxial models, though they occupy different subregions of the triaxial parameter space depending on the structure of the lensing cluster.  Thus, they not only give errors that are too small and unrepresentative of the true uncertainties in the problem, but unpredictably bias the mass and concentration estimates either high or low.  Though these biases were shown to cancel with large-scale averaging in CK08, they are a large problem when studying individual clusters.  Similarly, the overly-optimistic error estimates make proper comparisons of parameter values between different works inaccurate, as fully consistent results may appear excluded by inaccurately tight error bars; this is even more important because fitting spherical models to non-spherical data makes the details of the fitting method and its order of operations in measuring, averaging, and weighting shear far more important and likely to affect the final parameter estimate.  The inaccurate description of errors is also important when attempting to characterize the scatter in mass-observable relations, crucial knowledge for constraining the cluster mass function and with it cosmological questions such as the normalization of the matter power spectrum and the nature of dark energy. \\begin{figure} \\centering \\rotatebox{0}{ \\resizebox{8cm}{!} {\\includegraphics{f10.eps}} } \\caption{The red stars plot the maximum-likelihood $M_{200}$ and $C$ parameter values obtained fitting a spherical NFW to 500 randomly chosen halves of the Abell 2204 catalogue.  The green circle gives the maximum-likelihood parameter values, and the shaded black region gives the $2\\sigma$ confidence region, obtained fitting the same spherical NFW to the whole catalogue.  The black contours give the $2\\sigma$ confidence contours for 5 of the 500 half-catalogue realizations; the mass sheet degeneracy becomes very strong when the signal is weak, and causes the posterior probability distribution to develop very long tails to high concentrations and masses.} \\label{fig:plot10} \\end{figure}  As \\cite{oguri08} convincingly demonstrated over a large population of triaxial halos, and as we argued in CK08, the appearance of circular symmetry in a projected cluster DOES NOT justify the assumption of sphericity, most especially in strong lensing clusters.  \\cite{oguri08} found that the strongest lensing clusters in the universe are expected to be more triaxial, but look more circular in projection, than their weaker counterparts, as line-of-sight projection boosts the lensing strength; CK08 showed that such halos leave no signature in their lensing signal to differentiate them from truly spherical halos, making it impossible to determine beforehand the degree of triaxiality of a given system.  Thus, for all systems, and especially powerful strong lens systems such as Abell 1689, assuming spherical symmetry is never justified theoretically.  When external information is available, such as strong lensing, X-ray, or dynamical data, this MCMC method allows for the straightforward combination of those external constraints with the weak lensing data.  However, as emphasized in Sec. \\ref{sec:extpriors}, care must be taken that those external constraints do not reintroduce the very assumptions the triaxial model intends to avoid.  In the case of strong lensing, this is most easily done through the $\\theta$ orientation prior, which comes from a robust theoretical prediction favouring line-of-sight oriented halos to have the highest cluster lensing signals in the universe.  The Einstein radius prior is in some ways preferable in that it requires no prediction from CDM theory, but is problematic in that, applied to a dark matter only profile such as the NFW, it neglects the impact of baryons in the very centre of the cluster that are expected to significantly boost the strong lensing signal. This leads to a strong upward shift in the concentration value, which is unlikely to accurately represent the structure of the dark matter component of the cluster.  X-ray constraints may be applied when a halo appears to have a low degree of triaxiality, making the assumption of spherical symmetry of the (always more spherical) 3D gravitational potential less problematic, and when the cluster appears relaxed and undisturbed.  More detailed constraints may be applied by their direct inclusion in the likelihood function constraining the triaxial model.  For example, strong lensing may be used to directly constrain the mass and concentration if the strongly lensed features, or the 2D convergence map constrained by them (constructed with no assumptions about the 3D structure of the halo), are used to directly constrain the 3D triaxial model.  Doing so requires careful calibration of the weighting between the weak and strong lensing data, an optimization that has not yet been fully examined but has been begun by various groups (\\cite{bradac05}; \\cite{bradac06}; \\cite{diego07}), and should be a focus of the next generation of galaxy cluster studies.  \\subsection{Remaining systematics}\\label{sec:system} Though the parameters and the errors presented in this work now include the crucial contribution of halo triaxiality, there remain several outstanding systematics.  One of these is uncertainty in the mean redshift of the background galaxy population, changes in which would lead to changes in the mass estimates of several percent.  Another is residual contamination by cluster galaxies; though the colour cuts made in all three catalogues are very stringent, there is still the possibility of contamination by a few cluster members.  As with the uncertainties in the mean redshift of the background population, the effects should be negligible.  More importantly, stringent colour cuts to remove contamination often result in a significant decrease in the number density of background sources (for example, the catalogue from Abell 1689 used in this work had a source density of only 7/arcmin$^2$).  Low number density affects maximum-likelihood values and marginalized parameters significantly because of the large and non-Gaussian mass-sheet degeneracy between the NFW mass and concentration: mass and concentration constraints from low S/N data are poor because the tails of the highly asymmetric $M_{200}-C$ posterior probability distribution elongate dramatically.  Figure \\ref{fig:plot10} shows this effect, taking 500 randomly chosen halves of the Abell 2204 catalogue and treating them each as individual realizations to which a spherical NFW is fit.  Red stars plot the distribution of recovered maximum-likelihood parameters, compared to the $2\\sigma$ confidence contour for the spherical model fit to the complete catalogue.  The mean marginalized parameter values are even worse affected, with recovered values of concentration up to six times greater than the mean from the complete catalogue, because they are very sensitive to the long tails of the posterior when the signal is very weak, illustrated for five realizations in the Figure.  The mean marginalized parameters would be expected to be {\\it more} robust than their maximum-likelihood counterparts were the errors Gaussian; their sensitivity to the random removal of half the lensing signal reflects the very non-Gaussian errors induced by the mass-sheet degeneracy.  Because marginalized parameters are more meaningful in underconstrained problems such as that of fitting 3D mass models to lensing data, this is an important issue.  More efficient ways of removing foreground contamination are therefore very valuable, such as the Bayesian method of \\cite{limousin07}, or that of \\cite{medezinski07} that increases the efficiency of background selection by studying the radial tangential shear profiles of various colour populations, or improved methods of colour discrimination based on detailed photometric studies of the background galaxy population.  Future lensing work, including studies using the MCMC method employed here, can only be improved by increasing the background number density reliably available for inclusion in the lensing analysis.   Our triaxial model includes no substructure, while we know galaxy cluster scale halos often contain subclumps, and certainly contain mass concentrations at the positions of the cluster galaxies and their accompanying dark matter halos.  The weak lensing data does not have the resolution or signal-to-noise to constrain these; for example, \\cite{marshall06} found the Bayesian evidence supported fitting usually at most two components to cluster-scale substructures.   In the future, in combination with strong lensing or flexion, where detailed models including substructure are frequently constructed (e.g. \\cite{limousin07}; \\cite{jullo07}; \\cite{leonard07}), it should be possible to constrain more of the substructure of clusters.  However, \\cite{clowe04} showed that the effects of substructure on parameter estimate are small compared to those of triaxiality, and so we expect the inclusion of multiple components to be a minor perturbation to the results we present in this paper.  Finally, our quoted errors do not include a contribution from uncorrelated line-of-sight structure, shown by \\cite{hoekstra03} to induce an unbiased scatter in parameter estimates that can increase the errors on concentration and mass by up to a factor of 2.  \\cite{dodelson04} demonstrated that this can be somewhat mitigated by including a noise term for large scale structure in the lensing analysis; building on the greater coherence scale of large scale structure noise compared with the noise associated with the intrinsic ellipticity dispersion of galaxies, Dodelson notes that the errors on cluster mass can be reduced by $\\sim 50\\%$ for wide-field data, though this has yet to be implemented in any weak lensing analysis.  \\subsection{Bayesian Evidence} We choose not to use the Bayesian evidence to discriminate between our fits under different priors, because we are not attempting to choose between them as models of the universe.  The limitations of the evidence are clear when the Flat vs Spherical cases are considered: in many cases the evidence would likely favour a Spherical over a Flat prior, due to a significant decrease in available parameter space coupled with a relatively small decrease in likelihood values under the Spherical prior, as seen in the very similar reduced $\\chi^2$ values between all priors. However, unlike in cases of fitting multiple halos to account for substructure, or in another context, adding additional parameters to cosmological models, we know a priori from physical observations that a triaxial model is a better model than a spherical model -- we see clearly in non-parametric lensing mass maps that galaxy clusters have significant levels of ellipticity and are not spherical!  Thus, the prior on the spherical {\\it model} is close to zero, but difficult to quantify.  The various other theoretical priors studied in this paper are all derived from a single CDM model, and so treating them as separate models to choose between is inappropriate.  We study them here independently to understand how various aspects of the CDM model interact with the lensing data; in application to a large survey, all priors that represent the true prior expectation for the problem without interfering with the scientific questions asked by the study should be employed simultaneously.  Similarly, external constraints from strong lensing or X-ray do not represent different models to be selected between, only appropriate or inappropriate measures to constrain the triaxial models.  The evidence is a powerful tool for model discrimination, but is not the appropriate method for choosing priors on a single physical model, well-founded in theory and observations.  \\subsection{Summary} We fit a fully triaxial NFW model to weak lensing data from massive galaxy clusters Abell 1689, Abell 1835, and Abell 2204, deriving parameter estimates and errors under a range of theoretical and observational priors.  Under a relatively weak prior drawn from the strong lensing orientation bias for Abell 1689, that same strong lensing prior combined with an X-ray derived mass and concentration constraint for Abell 1835, and under the very general Flat prior for Abell 2204, our best mean parameter estimates in a concordance cosmology with $h=0.7$ are \\begin{itemize} \\item Abell 1689: $M_{200}=(1.18 \\pm .23)\\times10^{15}$ M$_{\\odot}$; $C=12.2\\pm 6.7$; \\item Abell 1835: $M_{200}=(0.98 \\pm .25)\\times10^{15}$ M$_{\\odot}$; $C=3.7\\pm 1.0$; \\item Abell 2204: $M_{200}=(0.72 \\pm .27)\\times10^{15}$ M$_{\\odot}$; $C=7.1\\pm 6.2$; \\end{itemize} All results are consistent with previous work, but importantly with larger error contours that better reflect the true uncertainty in mass profile estimates.  Triaxiality does not easily return Abell 1689 to the fold of average CDM clusters; the predicted CDM mass-concentration relation is enclosed in its $2\\sigma$ $M_{200}-C$ confidence contour under most priors, but is excluded under some observational priors, making Abell 1689 at the moment weakly consistent with the predictions of CDM at $2\\sigma$.  However, consistency does not mean truth: the primary result of this work is to demonstrate the necessity for improved methods of combining diverse observational constraints to counter the very large uncertainties that accompany fits of triaxial NFW profiles to galaxy cluster lenses.  That way lies the future of galaxy cluster studies, whether of individual clusters or large populations and mass-observable relations: without such optimized and robust combination methods, most results will continue to be consistent with one another and with theory simply through the size of their errors, rather than through physical concordance."
    },
    "0812/0812.0297_arXiv.txt": {
        "abstract": "{ High resolution synthetic stellar libraries are of fundamental importance for the preparation of the Gaia Mission. We present new sets of spectral stellar libraries  covering two spectral ranges: 300 --1100 nm at 0.1 nm resolution, and 840 -- 890 nm at 0.001 nm resolution. These libraries span a large range in atmospheric parameters, from super-metal-rich to very metal-poor (-5.0 $<$[Fe/H]$<$+1.0), from  cool to hot (\\teff=3000--50000 K) stars, including peculiar abundance variations. The spectral resolution, spectral type coverage and number of models represent a substantial improvement over previous libraries used in population synthesis models and in atmospheric analysis. ",
        "introduction": "\\begin{table*}[th!] \\scriptsize \\caption{Synthetic stellar libraries for Gaia. References: $^1$\\cite{Bouret08},$^2$\\cite{Hind}, $^3$\\cite{Koch05}, $^4$\\cite{Gust08}, $^5$\\cite{Alva98}, $^6$\\cite{Brott05}, $^7$\\cite{WD}} \\label{sinopt} \\begin{center} \\vspace{-0.3cm} \\begin{tabular}{l l l l l} \\hline \\\\ [-3pt] Name &  T$_{eff}$  (K)& $\\log g$ & [Fe/H]  &Notes \\\\[3pt] \\hline \\\\ O,B,A stars $^1$  & 8000 -- 50000 & 1.0 -- 5.0&-5.0 -- +1.0&TLUSTY code, NLTE, wind, mass loss\t\t\\\\[2pt] sdO   $^2$             &26000 -- 100000 & 4.8-6.4  & +0.0    0&TMAP, NLTE, $\\Delta$\\teff= 500 K, $\\Delta \\log g$= 0.2 \\\\[2pt] A stars\t$^3$      & 6000 -- 16000 & 1.0 -- 5.0&+0.0        &Magnetic field\t\\\\[2pt] MARCS\t$^4$      & 4000 -- 8000  &-0.5 -- 5.0&-5.0 -- +1.0& Variations in individual $\\alpha$-elements abundances\t\\\\[2pt] C stars\t$^{4,5}$  & 4000 -- 8000  & 0.0 -- 5.0&-5.0 -- +0.0& $\\Delta$\\teff= 500 K; [C/Fe]=0,1,2,3; [$\\alpha$/Fe]=+0.0, +0.4\t\t\\\\[2pt] PHOENIX\t$^6$      & 3000 -- 10000 &-0.5 -- 5.5&-3.5 -- +0.5& $\\Delta$\\teff= 100 K; [$\\alpha$/Fe]=-0.2--+0.8, $\\Delta$[$\\alpha$/Fe]=0.2\\\\[2pt] Emission lines    &$\\geq$ 15000 & 2.8 -- 4.0  &+0.0        &Be, WR models                                  \\\\[2pt] WD$^7$            & 6000 -- 90000 &7.0--9.0     &            &WDA \\& WDB, LTE\t\t\t\t\t\\\\[2pt] \\hline \\end{tabular} \\end{center} \\vspace{-0.4cm} \\end{table*} \\normalsize  ESA's Gaia mission, to be launched in 2011, is meant to obtain accurate position, parallax and  proper motion for 10$^9$ object all over the sky, up to magnitude $G$=20 ($V$=20--22) with an astrometric accuracy at the $\\mu$ arcsec level, providing low-dispersion spectroscopy for each object and radial velocity up to $G$=16 (for more detailss see Cacciari 2008, this meeting). The final catalogue is expected to provide the discrete classification of sources (single stars, galaxies, QSOs, asteroids) and the astrophysical parameters (APs) for single stars (i.e. T$_{\\rm eff}$, $\\log g$,...) and possibly the parametrization of special sources (galaxies). Gaia will produce an unprecedented amount of data (40--50~GB/day) leading to 100~TB of compressed data in 5 years. This huge number of observed objects must be classified in an automated way. In the Gaia project, all the classification algorithms  but one are based on supervised methods  \\citep[see for example][]{Coryn07}, classifying a source or estimating its APs by comparison with a set of templates. In order to do this,  simulations of Gaia observations covering the whole AP space are required: high resolution and high quality synthetic libraries are of fundamental importance. Here we focus on new stellar libraries, while galaxies, QSOs, asteroids are discussed elsewere \\citep{Vivi, Claeskens, Warell}. ",
        "conclusions": "In this paper we summarize the work recently done by a large scientific comunity in the calculation of  a large database of state-of-the-art synthetic stellar spectra, to be used in preparation of the ESA's Gaia mission. They cover the whole optical range with 0.1 resolution and the Ca II triplet region at 0.001 nm resolution, for a large set of astrophysical parameters. Generally, these libraries are of great interest in stellar population synthesis and abundance analysis works. These large grids offer excellent possibilities to explore the variation of the spectrum properties with stellar parameters.  Even if  theoretical stellar spectra modelling has greatly improved in the paste decades, still the effects  and the inconsistency  introduced by the underlying simplifying assumptions need to be analysed in detail. Here   a few comparison are shown, using broad band colors and direct spectra comparison. While (V-I) and (V-R) color are a good reproduction of the empirical calibration for MARCS, Basel, A stars, (B-V) colors give a poorer result, being greatly affected by the well known problem of the  line list. Direct spectra comparison shows that hot star spectra are in reasonable agreement. The differences between MARCS and PHOENIX are of the order of 10-20\\%, in particular for cool stars and at the bluest end of the spectrum."
    },
    "0812/0812.0859_arXiv.txt": {
        "abstract": "With use of CompHEP package we've made the detailed estimate of the influence of double $e^+e^-$ pair production (DPP) by photons on the propagation of ultra high energy electromagnetic (EM) cascade. We show that in the models in which cosmic ray photons energy reaches few$\\times10^3$~EeV refined DPP analysis may lead to substantial difference in predicted photon spectrum compared to previous rough estimates. ",
        "introduction": "Ultra-high energy (UHE) photons have not been recognized so far by any of present generation experiments~\\cite{AGASA_eest,Yakutsk_spec,HiRes_spec,Auger_spec}, although their existence is predicted by Greisen--Zatsepin--Kuzmin effect~\\cite{g,zk} as well as by most of hypothetical top-down models of UHE cosmic rays origin. There are several bounds on fraction and flux of ultra-high energy photons above $10 - 100\\EeV$ obtained by independent experiments~\\cite{A+Y,Ylim,Auger_sdlim}. Photon limits are used to constrain the parameters of top-down models (see for example~\\cite{Gelmini:2005wu}). Future bounds may also limit considerable part of parameter space of astrophysical models, in which photons are produced as secondaries from interactions of primary protons or nuclei with cosmic microwave background (CMB). Understanding interactions of UHE photons with universal backgrounds is a crucial point for building such constraints.  In the wide energy range the spectra of electron and photon components of cosmic rays follow each other due to relatively rapid processes transferring $\\gamma$-rays to electrons and backwards. Pair production (PP) and inverse Compton scattering (ICS) are the main processes that drive the EM cascade. In the Klein--Nishina limit where $s \\gg m_e^2$, either electron or positron produced in a pair production event typically carries almost all of the initial total energy. The produced electron (positron) then undergoes ICS losing more than 90~\\% of energy and finally the background photon carries away almost all of the initial energy of the UHE photon. Due to this cycle the energy loss rate of the leading particle in the EM cascade is more than one order of magnitude less than interaction rate. However in presence of a random extragalactic magnetic field (EGMF) the electrons may lose substantial part of their energy by emitting synchrotron radiation. In this case, starting from certain energy the synchrotron loss rate for the electrons becomes to dominate over ICS rate, which leads to suppression of the EM cascade development. Its penetration depth is then defined by the photon mean free path. Depending on the value of EGMF this transition may occur between $\\sim 1\\EeV$ and $\\sim 10^6\\EeV$.  In this article we consider higher order process, double $e^+e^-$~pair production (DPP) by photons. The DPP cross section grows rapidly with $s$ near the threshold and quickly approaches the asymptotic value $\\sigma (\\infty) \\simeq 6.45~\\mb$ \\cite{Cheng:1970ef,Cheng:1970zb,Lipatov:1971ax}. The explicit energy dependence of the DPP cross section was estimated in Ref.~\\cite{brown} by calculating the dominant contribution from two $e^+e^-$ pairs to the absorptive part of gamma-gamma forward scattering amplitude. Since PP cross section decreases with the increase of $\\sqrt{s}$ the DPP rate starts to dominate over PP rate above certain energy. For interactions with CMB the transition occurs above $\\sim 1000\\EeV$. In presence of the radio background this energy goes up somewhat. If the EGMF is less than $10^{-11}\\G$ the EM cascade still exists at these energies and one should accurately count the secondary electrons from DPP. So far the EM cascade simulations such as~\\cite{propag, Lee:1996fp} roughly estimated DPP effect by utilizing the total cross section and assuming that one $e^+e^-$~pair of the two carries all the initial energy while two particles in the pair are produced with the same energy. By making use of CompHEP package~\\cite{Boos:2004kh,Pukhov:1999gg,cite_comphep} we numerically calculate differential cross section for DPP and compare the influence of DPP on propagation of ultra high energy EM cascade with previous estimates.  The paper is organized as follows. In Sec.~\\ref{CS} we present the results of the calculation of DPP cross section. In Sec.~\\ref{TE} we write transport equations for EM cascade and calculate the coefficients for transport equations for photons, and secondary $e^+$,$e^-$ related to DPP. In Sec.~\\ref{sec_models} we illustrate the influence of DPP in model example. In Sec.~\\ref{Conclusion} we summarize our results. ",
        "conclusions": "\\label{Conclusion}  In this work we have considered in detail the DPP process. We have estimated the distribution of secondary electrons and positrons and made the improved cosmic rays simulation based on the new estimate. We have shown that in certain cases the DPP process may modify the photon component of the spectrum substantially. However this modification can only be seen if radio background is close to the minimal model~\\cite{clark} and EGMF is lower than $10^{-11}{\\rm G}$. In this case there is an energy range where DPP is the main attenuation mechanism for $\\gamma$~rays and therefore differences in DPP estimates can clearly be seen. Although in the vast majority of the models which do not contradict to the present experimental bounds on the photon fraction in UHE cosmic rays DPP process doesn't make a substantial contribution to the attenuation and therefore can be treated simplistically or even disregarded.  We are indebted to D.~Gorbunov for inspiring us to prepare this work. We also thank D.~Semikoz for valuable discussions. This work was supported in part by the Russian Foundation of Basic Research grant 07-02-00820, by the grant of the President of the Russian Federation NS-1616.2008.2, MK-1966.2008.2 (O.K.). Work of S.D. was supported in part by the grant of the Russian Science Support Foundation, Russian Foundation of Basic Research grants 08-02-00473-a and 08-02-00768-a. Numerical part of the work was performed on the Computational cluster of the Theoretical Division of INR RAS."
    },
    "0812/0812.3966_arXiv.txt": {
        "abstract": "The large optical reflector ($\\sim 100\\,$m$^2$) of a H.E.S.S. Cherenkov telescope was used to search for very fast optical transients of astrophysical origin. 43~hours of observations targeting stellar-mass black holes and neutron stars were obtained using a dedicated photometer with microsecond time resolution. The photometer consists of seven photomultiplier tube pixels: a central one to monitor the target and a surrounding ring of six pixels to veto background events. The light curves of all pixels were recorded continuously and were searched offline with a matched-filtering technique for flares with a duration of 2\\,$\\mu$s\\,--\\,100\\,ms. As expected, many unresolved ($<3\\,\\mu$s) and many long ($>500\\,\\mu$s) background events originating in the earth's atmosphere were detected. In the time range 3\\,--\\,500\\,$\\mu$s the measurement is essentially background-free, with only eight events detected in 43\\,h; five from lightning and three presumably from a piece of space debris. The detection of flashes of brightness $\\sim 0.1$\\,Jy and only 20\\,$\\mu$s duration from the space debris shows the potential of this setup to find rare optical flares on timescales of tens of microseconds. This timescale corresponds to the light crossing time of stellar-mass black holes and neutron stars. ",
        "introduction": "\\label{sec:introduction}  At optical wavelengths, the sky has been surveyed in deep exposures typically lasting minutes to hours. Much faster optical flares may have escaped detection, since the dynamic sky at sub-second timescales has rarely been explored in the past \\cite{eichler00}. The availability of high quantum-efficiency, fast photodetectors (photomultiplier tubes (PMTs), avalanche photodiodes or charge-coupled devices (CCDs)) has lead a few groups to construct high-time-resolution photometers, imagers or spectrographs and operate them on medium to large optical telescopes \\cite{shvartsman97, straubmeier01, ryan06, dhillon07, hormuth08}.  The very large light collecting areas ($>100\\,$m$^2$) of reflectors of imaging Cherenkov telescopes such as H.E.S.S. \\cite{hinton04}, MAGIC \\cite{lorenz04} and VERITAS \\cite{krennrich04} provide an interesting alternative to normal optical telescopes \\cite{eichler00}. Up to now, the only optical observations performed using the reflector of a Cherenkov telescope concerned the measurement of the optical pulsed emission from the Crab pulsar \\cite{hinton06, lucarelli08}.  Compact objects like neutron stars and stellar-mass black holes are prime candidates as sources of fast optical flares. The fastest expected time variability of the emission from these systems is given by the light travel time through the immediate vicinity of the compact object, 10\\,--\\,100\\,$\\mu$s for a typical size of tens of kilometers \\cite{eichler00}. Such short flares can be produced by the accretion of lumps of matter, for instance in X-ray binary systems from the stellar companion \\cite{dolan01}. Relativistic outflows can further shorten the observed timescales by Doppler boosting (for exceptional examples in the case of super-massive black holes see \\cite{gaidos96, aharonian07, albert07}). Optical variability in the millisecond range has been observed for pulsars \\cite{shearer03} and X-ray binaries \\cite{bartolini94, beskin94, kanbach01} and has been searched for in objects known to exhibit flares in other wavelength regions, for instance at radio frequencies from rotating radio transients (RRATs) \\cite{mclaughlin06, dhillon06}. No optical variability on timescales shorter than 100\\,$\\mu$s has been detected from any astronomical target so far.  When trying to detect astronomical flares, natural and artificial flaring events occurring in the earth's atmosphere or orbit have to be considered. Known producers of such terrestial optical flares on sub-second timescales are airplanes, satellites, lightning and shooting stars \\cite{schaefer87}. Additionally, cosmic-ray induced air showers produce light flashes lasting a few nanoseconds \\cite{galbraight53}.  This paper reports on 43\\,h of high-time-resolution optical photometry of X-ray binaries and one RRAT using the reflector of a H.E.S.S. Cherenkov telescope during moonshine. A comparison of the performance of H.E.S.S. with regular optical telescopes for such fast photometry observations is presented in Sec.~\\ref{sec:comparison}. The custom-built detector that was used to record continuous light curves at microsecond time resolution is described in Sec.~\\ref{sec:camera}. An overview of the observations and a description of the flare-finding algorithm that was used to search for flares on timescales of 2\\,$\\mu$s\\,--\\,100\\,ms is given in Sec.~\\ref{sec:observations}. The events that were detected are presented in Sec.~\\ref{sec:results}, and Sec.~\\ref{sec:outlook} gives an outlook on how this new technique might be improved in the future. ",
        "conclusions": ""
    },
    "0812/0812.4541_arXiv.txt": {
        "abstract": "The ACS Survey of Galactic Globular Clusters is a $Hubble$ $Space$ $Telescope$ $(HST)$ Treasury program designed to provide a new large, deep and homogeneous photometric database. Based on observations from this program, we have measured precise relative ages for a sample of 64 Galactic globular clusters by comparing the relative position of the clusters' main sequence turn offs, using main$-$sequence fitting to cross$-$compare clusters within the sample. This method provides relative ages to a formal precision of 2-7\\%. We demonstrate that the calculated relative ages are independent of the choice of theoretical model. We find that the Galactic globular cluster sample can be divided into two groups -- a population of old clusters with an age dispersion of $\\sim$5\\% and no age$-$metallicity relation, and a group of younger clusters with an age$-$metallicity relation similar to that of the globular clusters associated with the Sagittarius dwarf galaxy.  These results are consistent with the Milky Way halo having formed in two phases or processes. The first one would be compatible with a rapid ($<$0.8 Gyr) assembling process of the halo, in which the clusters in the old group were formed. Also these clusters could have been formed before reionization in dwarf galaxies that would later merge to build the Milky Way halo as predicted by $\\Lambda$CDM cosmology. However, the galactocentric metallicity gradient shown by these clusters seems difficult to reconcile with the latter. As for the younger clusters, it is very tempting to argue that their origin is related to their formation within Milky Way satellite galaxies that were later accreted, but the origin of the age$-$metallicity relation remains unclear. ",
        "introduction": "One of the keys to understanding the structure and evolution of the Milky Way is the ability to divide its stars and clusters into separate Galactic Populations as first hinted at by \\citet{O26} and formally introduced by \\citet{B44}. In the early 1950s, Sandage published the first deep color$-$magnitude diagrams (CMDs) of the northern Galactic globular clusters (GGCs) M3 and M92 \\citep{S53, A53} and discovered the faint sequence of stars similar to the main sequence of Population I. It was almost immediately realized that GGCs are ancient compared to the stars in the Solar Neighborhood.  Over the last half century, it has become clear that GGCs are astronomical fossils providing important information on the formation and evolution of the Milky Way. One characteristic of GGCs that makes them so valuable is that they are the oldest Galactic objects for which reliable ages can be measured.  GGC ages are of great interest to a variety of cosmological and comogonical issues since their {\\it absolute} ages place a lower limit on the age of the Universe. Unfortunately, absolute age measurements are still compromised by a number of uncertainties, particularly GGC distances and foreground reddenings as well as metallicities, bolometric corrections and some poorly known aspects of the stellar evolution. Relative ages are almost (although not totally) free of these effects and can still reveal fundamental information about Galaxy formation mechanisms and time scales. For this reason they have become of paramount importance in the context of the study of GGC.  The study of GGC ages has evolved apace with improvements in instrumentation, in our theoretical understanding of stellar evolution and with the advance of computational facilities. Reliable age determination requires deep photometry that reaches at least the main$-$sequence turn off (MSTO), together with a good understanding of the stellar evolution theoretical modeling. During the first half of the 20$^{\\rm th}$ century, the available instrumentation did not allow observations deeper than the horizontal branch of most GGCs. Over the past three decades, however, low quantum efficiency non$-$linear photographic imaging has given way to digital imaging using high quantum efficiency linear CCDs.  This has resulted in ever deeper and more precise cluster CMDs, that reached the MSTO and allowed the application of age dating techniques. First relative age studies were carried out by \\citet{G85}, \\citet{P87}, \\citet{SK89} and \\citet{SD90}, as representative examples. However, uncertainties in both the input physics \\citep{Ch98} and the color$-$T$_{\\rm eff}$ transformation of the theoretical models into the observational plane \\citep{B98} constitute a significant source of errors in the final relative age determination, particularly when ages of clusters with different metallicities are compared. In the last decade, many efforts have been made in GGC relative ages, improving our understanding with each technological or theoretical advance \\citep[e. g. ][]{Ch96, R96, Stet96, SW98, B98, R99, Van00, DeA05}.  Relative ages can be estimated to high precision by measuring the position of the MSTO relative to CMD features that have little or no age dependence \\citep{Stet96, SarCha97}. However, a key aspect of any relative age experiment is having a large homogenous database of globular cluster photometry.  Such databases have only recently appeared in the literature but have significantly advanced our understanding of Milky Way formation and evolution. The largest dedicated efforts are those of \\citet{R00a, R00b}, a ground-based survey of 35 nearby GGCs \\citep{R99}, and \\citet{DeA05}, an analysis of 55 clusters from \\citet{P02}'s HST snapshot catalog.  These and other recent relative age studies \\citep{B98, SW98, R99, Van00, DeA05} consistently point towards a scenario in which metal$-$poor clusters are largely coeval but metal$-$rich clusters show some age dispersion, $\\sim$1.0 Gyr (rms), and a possible age$-$metallicity relation. No correlation between age and Galactocentric distance has been conclusively demonstrated.  In this paper, we present an analysis of a uniform set of HST/ACS data for 64 GGCs\\footnote{\\citet{Sar07} presented observations of 65 GCs, but Pal 2 has been left out of this study due to its high differential reddening}.  This dataset, previously introduced by \\citet{Sar07}, has produced well-defined CMDs for all of the target clusters from the tip of the red giant branch (RGB) to at least $\\sim$6.5 magnitudes below the MSTO \\citep{A08}. The determination of GGC ages was a principle goal of the HST/ACS survey and exposures times were defined to maximize the S/N ratio at the MSTO.  This makes the survey an excellent resource for the evaluation of relative ages.  Due to the difficulty in determining the position of the MSTO\\footnote{ The MSTO point is the current position in the CMD where stars in the cluster start burning hydrogen in a shell and are leaving the main$-$sequence. The empirical MSTO is then the bluest point on the main$-$sequence$-$to$-$subgiant$-$branch portion of the cluster's CMD.} to the high precision required for age studies, many different methods have been developed in previous relative ages works \\citep[see][for a complete discussion]{MW06}. The two most used relative age indicators are those based on the color difference between the MSTO and the RGB at a given magnitude level (horizontal method), and the magnitude difference between the MSTO and the zero$-$age HB\\footnote{ The zero$-$age HB is the position in the CMD where stars in the cluster start burning helium into carbon and oxygen in a helium rich core and hydrogen into helium in a shell.} (vertical method). \\citet{VBS90} were pioneers in measuring GGC relative ages making use of the horizontal method. They essentially derived relative ages comparing the color difference between the MSTO and the location of a well defined point in the lower RGB, located 2.5 magnitudes brighter than a point in the upper main$-$sequence (MS) that is 0.05 magnitudes redder than the MSTO. However, horizontal method results are sensitive to the assumed value of the mixing length parameter and to the color$-$T$_{\\rm eff}$ transformation of the theoretical models into the observational plane. As a result, relative ages determined using the horizontal method are model dependent.  Potentially, more reliable age indicators are those related to the brightness of the MSTO, in particular the vertical method. First surveys of GGC relative ages using the vertical method were carried out by \\citet{G85} and \\citet{P87}. They measured relative ages from a sample of 26 and 41 GGCs, respectively, collected from the available literature. Later on, \\citet{SK89} used the brightness difference between the HB and the MSTO to estimate the ages of 31 GGCs. Nevertheless, accurate determination of the zero$-$age HB is delicate, in particular for the clusters at the extreme boundaries of the metallicity distribution, which generally have a red (or blue) HB and no RR Lyrae stars. Moreover, the dependence of the zero$-$age HB luminosity on the cluster metallicity is still controversial.  In this paper, we use an alternative approach. Our data reaches at least $\\sim$6.5 magnitudes below the MSTO, providing a very well defined MS. This allows, for the first time, the use of relative MS fitting for a large number of clusters. As it will be shown later, the relative MS fitting procedure used in this study does essentially the following. The MS of two clusters, with similar metallicities, are superimposed in the CMD by shifting one of the cluster's mean ridge line in both color and magnitude  (explained in detail in section~\\ref{MRLandMSF}). This way, any differences between the two cluster's distance and reddening are compensated, and an intrinsic MSTO magnitude difference is obtained. This relative MSTO brightness is then used to derive relative ages. This allows us to use what is, in essence, an improved version of the traditional vertical method that substitutes the well-defined MS for the contentious zero$-$age HB.  We also take the additional step of comparing our observational measures to updated theoretical models. We have primarily used the most recent state$-$of$-$art models of \\citet{D07} (hereafter D07), transformed to the observational plane with an empirical Color$-$T$_{\\rm eff}$ transformation.  However, we also compare our results to the models of \\citet{P04} (hereafter P04), \\citet{B94} (hereafter B94) and \\citet{G00} (hereafter G00).  We find that the various theoretical models agree closely on the MSTO absolute magnitude and color, resulting in negligible model dependence.  This paper is organized as follows. Section~\\ref{obs} describes the observations and data reduction pipeline, and Section~\\ref{metScale} presents the globular cluster database. The mean ridge line fitting and MS$-$fitting procedures used in this study are described in Section~\\ref{MRLandMSF}. The analysis of the uncertainties is presented in Section~\\ref{errorDet}. Finally, relative age results and a discussion are presented in Sections~\\ref{results} and \\ref{discussion}, respectively. The paper ends with a summary of conclusions (Section~\\ref{conclusions}). ",
        "conclusions": "Normalized ages have been derived for a sample of 64 GGCs using the stellar evolution models of D07 and both ZW and CG metallicity scales.  We have also performed an analysis using the stellar models of P04, B94 and G00.  The result is the most extensive and precise database of normalized ages so far produced . Our results are:  \\begin{itemize}  \\item We find that we are able to measure relative ages to a formal precision of 2-7\\%. The relative ages are independent of the choice of theoretical model.  Four independent sets of stellar evolution libraries (D07, P04, B94, G00) produce essentially identical results.  \\item We find that the GGCs fall into two well-defined groups.  The first one represents a population of old clusters that have the same age, a dispersion of 5\\% in relative age and no age$-$metallicity relation. If we assume an absolute age of 12.8 Gyr, the absolute age dispersion of the old clusters is $\\sim$0.6 Gyr. Accounting for the measurement uncertainties, we obtain an intrinsic age dispersion of $\\sim$0.4 Gyr. The second group of clusters shows a clear age$-$metallicity relation, with young clusters being more metal rich than older ones. Also in this case, the intrinsic age dispersion with respect to its age$-$metallicity relation is $\\sim$0.4 Gyr. Besides, there is no age dispersion$-$metallicity relation if we look at the two samples (old and young) separately.  \\item These results are consistent with a scenario in which the formation of the Milky Way CCG system took place in two phases. The first one produced clusters with ages within a range of $\\sim$0.8 Gyr. This age range is compatible with the timescale for the collapse of a protogalaxy of the same mass and scalelength as that of the Milky Way dark matter halo. In other words, the age dispersion of the old group of globular clusters is not in contradiction with the formation from the colapse of a single protosystem, resembling the model proposed by \\citet{ELS62}. The standard $\\Lambda$CDM scenario would also predict a significant star formation in protogalactic building blocks before reionization, which, after merging, would produce an old, coeval population of globular clusters in a large galaxy like the Milky Way. However, to account for the galactocentric metallicity gradient of this group seems difficult in the context of this mechanism. The second phase spanned a time interval as long as $\\sim$6 Gyr and resulted in a group of GGCs with a clear age$-$metallicity relation. It is very tempting to argue that the origin of this second group of clusters is related to their formation within Milky Way satellite galaxies that were later accreted. However, this would not account for the fact that the clusters of this group share the same age$-$metallicity relation.  \\item  We find that the increasing dispersion in age with galactocentric distance is explained by the age$-$metallicity relation of GGCs.  \\item These results are independent of the assumed metallicity scale.  \\item It is worth mentioning that our database is biased by distance. No outer halo clusters, with galactocentric distances larger than $\\sim$20 kpc, have been considered. The analysis of the galactic halo's age structure is therefore limited by the galactocentric distance range of the database.  \\end{itemize}"
    },
    "0812/0812.3157_arXiv.txt": {
        "abstract": "The Diffuse Supernova Neutrino Background (DSNB) provides an immediate opportunity to study the emission of MeV thermal neutrinos from core-collapse supernovae. The DSNB is a powerful probe of stellar and neutrino physics, provided that the core-collapse rate is large enough and that its uncertainty is small enough. To assess the important physics enabled by the DSNB, we start with the cosmic star formation history of Hopkins \\& Beacom (2006) and confirm its normalization and evolution by cross-checks with the supernova rate, extragalactic background light, and stellar mass density. We find a sufficient core-collapse rate with small uncertainties that translate into a variation of $\\pm 40$\\% in the DSNB event spectrum. Considering thermal neutrino spectra with effective temperatures between 4--6 MeV, the predicted DSNB is within a factor \\mbox{4--2} below the upper limit obtained by Super-Kamiokande in 2003. Furthermore, detection prospects would be dramatically improved with a gadolinium-enhanced Super-Kamiokande: the backgrounds would be significantly reduced, the fluxes and uncertainties converge at the lower threshold energy, and the predicted event rate is 1.2--5.6 events yr$^{-1}$ in the energy range 10--26 MeV. These results demonstrate the imminent detection of the DSNB by Super-Kamiokande and its exciting prospects for studying stellar and neutrino physics. ",
        "introduction": "Core-collapse supernovae (CCSNe) occur at a rate of several per second in the Universe, each releasing a prolific $\\sim$ $10^{58}$ neutrinos and antineutrinos. Their detection provides a rich bounty for stellar and neutrino studies, shown by the detection of the neutrinos from SN 1987A \\cite{Hirata:1987hu,Bionta:1987qt,Hirata:1988ad,Bratton:1988ww,Arnett:1987iz,Bahcall:1987nx,Raffelt:1987yt,Barbieri:1988nh,Lattimer:1988mf,Arnett:1990au,Jegerlehner:1996kx,Lunardini:2004bj}. While a core-collapse supernova in the Milky Way would easily be detected in neutrinos, the occurrence rate is only $\\lesssim$ 3 per century \\cite{VanDenBergh:1991ke,Tammann:1994ev,Diehl:2006cf}. Proposed neutrino detectors should be able to detect supernovae up to 10 Mpc away with an occurrence rate of $\\sim$ 1 per year \\cite{Ando:2005ka,Kistler:2008us}. The vast majority of supernovae are therefore undetectable. However, the cumulative emission from all past core-collapse supernovae, which forms the Diffuse Supernova Neutrino Background (DSNB), has promising detection prospects \\cite{BisnovatyiKogan:1984,Krauss:1983zn,Dar:1984aj,Woosley:1986aa,Totani:1995rg,Totani:1995dw,Malaney:1996ar,Hartmann:1997qe,Kaplinghat:1999xi,Fukugita:2002qw,Ando:2002zj,Ando:2002ky,Strigari:2005hu,Lunardini:2005jf,Yuksel:2007mn,Chakraborty:2008zp}. The Super-Kamiokande (SK) limit of $\\phi(E_{\\bar{\\nu}_e} > 19.3 \\,\\mathrm{MeV}) < 1.2$ cm$^{-2}$ s$^{-1}$ \\cite{Malek:2002ns} on the DSNB flux is already close to theoretical predictions \\cite{Strigari:2005hu,Yuksel:2007mn}.  Predicting the DSNB for a given supernova neutrino emission model requires knowledge of the rate of core-collapse supernovae. In the past, this was not well known, and various studies provided insights on the supernova rate \\emph{from} the DSNB \\cite{Totani:1995dw,Fukugita:2002qw,Ando:2004sb}. Indeed, the SK limit on the DSNB flux is strong enough to rule out some supernova rate evolution models, assuming a fiducial neutrino emission model. However, our understanding of the cosmic star formation history (CSFH) has been greatly augmented by improved direct measurements in different wavebands and redshifts (see, e.g., Refs.~\\cite{Hopkins:2004ma,Hopkins:2006bw} and references therein). Cross-checks with other well-measured observables are now constraining, so that the CSFH is well determined by methods other than the DSNB. Thus it is both timely and important to study the prospects of the DSNB for probing stellar and neutrino physics.  In this paper we start with the CSFH of Hopkins \\& Beacom (hereafter HB06 \\cite{Hopkins:2006bw}) which is based on a compilation of recent data, and assess it by cross-checks, with the aim of evaluating the uncertainties that are carried forward into DSNB predictions through the core-collapse supernova rate. We henceforth refer to these as the \\underline{\\it astrophysical} inputs and uncertainties on the DSNB. On the other hand, we refer to the supernova neutrino emission and neutrino properties as the \\underline{\\it emission} inputs and uncertainties on the DSNB.  As cross-check material we consider measurements of the rate of core-collapse supernovae, which have been significantly updated \\cite{Cappellaro:2004ti,Dahlen:2004km,Botticella:2007er,Smartt:2008zd,Dahlen:2008aa}, the extragalactic background light, which records the total stellar emission over all time (for a recent review, see, e.g., Ref.~\\cite{Hauser:2001xs}), and, finally, the stellar mass density (see e.g., Ref.~\\cite{Wilkins:2008be} and references therein). While our approach is similar to previous studies such as Ref.~\\cite{Strigari:2005hu}, we perform novel checks and with higher precision. In particular, our analysis delivers fiducial inputs with unprecedentedly small uncertainties.  Using our constrained astrophysical inputs, we find that the DSNB uncertainty is dominated by the emission inputs, demonstrating the potential to study stellar and neutrino physics using the DSNB. Furthermore, taking into account both astrophysical and emission inputs, we find that the predicted DSNB is within at most a factor \\mbox{$\\sim$ 4} of the SK limit set in 2003. For a 6 MeV thermal spectrum, typical of scenarios with neutrino mixing, this factor reduces to $\\sim$ 2. At the lower detection threshold energy for a gadolinium-enhanced SK \\mbox{($\\sim 10$ MeV} \\cite{Beacom:2003nk,Yuksel:2005ae}), the combined uncertainty on the predicted DSNB is remarkably only a factor 2, with a event rate of 1.2--5.6 events yr$^{-1}$ in the energy range 10--26 MeV. Thus, the core-collapse rate and the neutrino emission per supernova are large enough to allow the imminent detection of the DSNB in SK. The successful detection will confirm the ubiquitous emission of neutrinos from core-collapse supernovae and initiate the much-anticipated study of stellar and neutrino physics, while a non-detection would require new stellar or neutrino physics.  The paper is organized as follows. In Sections \\ref{sec:astrophysics} and \\ref{sec:dsnb} we discuss the astrophysical and emission inputs for the DSNB. We discuss the DSNB detectability and make future predictions for a gadolinium-enhanced SK in Section \\ref{sec:sk}, and finish with conclusions in Section \\ref{sec:conclusion}. We adopt the standard $\\Lambda$CDM cosmology with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "} Neutrinos are the only probes (possibly apart from gravitational waves) of the central regions of core-collapse events, and their study and detection are strongly motivated by many areas of physics. In this paper we assess the uncertainties for DSNB searches and implications for stellar and neutrino physics. The results can largely be divided into three categories.  \\subsection{On astrophysical inputs} We start with an up-to-date compilation of the CSFH and cross-check it with other observables. While our aim is to constrain the astrophysical inputs for the DSNB (i.e., the CCSN rate), these new results are of value in their own right.  \\begin{itemize} \\item \\emph{Consistency with CCSN rates}: Using a collection of recently observed CCSN rates, we show that they are in agreement with predicted CCSN rates from the CSFH, up to $z \\sim 1$. Importantly, this check is almost independent of the IMF. \\item \\emph{Consistency with EBL}: Using new EBL measurements, including TeV gamma-ray constraints, we show that predictions from the CSFH are in agreement with observations. The observed total EBL is $73^{+21}_{-26} \\, \\mathrm{nW \\, m^{-2} \\, sr^{-1}}$, while we predict $78^{+31}_{-24} \\, \\mathrm{nW \\, m^{-2} \\, sr^{-1}}$ for the BG IMF. \\item \\emph{Consistency with stellar mass density}: The choice of the BG IMF is independently supported by studies of the stellar mass density, which show agreement in $z \\lesssim 0.7$. \\end{itemize}  In conclusion, our fiducial CSFH together with the BG IMF gives excellent agreement among all observations considered, preparing the astrophysical inputs for the DSNB. In particular, we stress how the astrophysical inputs cannot be too low, in order to maintain the self-consistent picture of the birth, life, and death of stars, as illustrated by the cross-checks.  \\subsection{On emission inputs} Next we assess the range of neutrino emission expected from supernovae. Coupled to our cross-checked astrophysical inputs, this allows us to discuss how constraining the present limit on the DSNB is. Our best-determined results are shown in Fig.~\\ref{fig:DSNBevent}.  \\begin{itemize} \\item \\emph{Neutrino emission}: Ultimately, this is a quantity to be studied using supernova neutrino detections. We thus discuss the generic emission of $\\bar{\\nu}_e$ from supernovae using simple arguments based on the energetics of core collapse, neutrino mixing, and the observed SN 1987A neutrino burst. We also show the neutrino temperature depends weakly on the remnant parameters. We conclude that the relevant high-energy $\\bar{\\nu}_e$ emission cannot be made too small. \\item \\emph{Present constraints}: The SK limit is already probing interesting parameter regions. The dominance of an effective $\\bar{\\nu}_e$ temperature above 8 MeV is constrained. Similarly, a high rate of dark core collapses is also prohibited by an amount that depends on the assumed temperature. \\item \\emph{Implications for stellar and neutrino physics}: If an effective temperature is ruled out, this rules out neutrino mixing and initial neutrino temperatures that would lead to that effective temperature in the DSNB energy range. The correspondance between the initial and effective neutrino emissions is dependent on the neutrino mixing scenario, which has been studied systematically by various authors. \\end{itemize}  To conclude, from general and SN 1987A considerations, the neutrino emission cannot be too low. The current SK limit is already constraining interesting neutrino effective temperatures. Noting the SK limit was placed in 2003, more data and improved cuts will allow better sensitivity.  \\subsection{On detection in SK} With the DSNB inputs and their uncertainties checked, we discuss implications for DSNB detection in the future.  \\begin{itemize} \\item \\emph{SK prospects}: The DSNB is near detection in SK. The 6 MeV spectrum, typical of scenarios with neutrino mixing, is within a factor $\\sim$ 2 of the current SK limit. The factor increases to $\\sim$ 4 for the lower 4 MeV and SN 1987A reconstructed spectra. \\item \\emph{SK with gadolinium}: Intriguingly, the fluxes and uncertainties converge at the improved detection threshold (\\mbox{$\\sim$ 10 MeV}), so that predictions span an uncertainty of a factor $\\sim$ 2. The predicted event rate between 10--26 MeV is 1.2--5.6 events yr$^{-1}$, leaving no room for the DSNB to escape detection. \\item \\emph{Future physics with SK}: The effective temperature $T_{\\bar{\\nu}_e}$ contains physics concerning neutrino emission from the collapsed core and on neutrino mixing. Stellar and neutrino physics can be extracted by future comparisons between the measured and theoretical $T_{\\bar{\\nu}_e}$. Future SK analysises should report results directly in terms of the time-integrated luminosity and effective temperature \\cite{Yuksel:2005ae}, using the astrophysically-measured supernova rate. \\end{itemize}  To conclude, while the current SK will continue probing interesting physics, a gadolinium-enhanced SK is almost guaranteed to detect DSNB events. A non-detection would require novel stellar or neutrino physics. Together with the decreasing astrophysical uncertainties, these results strongly support the imminent detection of the DSNB and solidify its important role in understanding supernova and neutrino physics."
    },
    "0812/0812.3682_arXiv.txt": {
        "abstract": "Deep submillimeter (submm) continuum imaging observations of the starburst galaxy M\\,82 are presented at 350, 450, 750 and 850\\,$\\mu$m wavelengths, that were undertaken with the Submillimetre Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope in Hawaii. The presented maps include a co-addition of submm data mined from the SCUBA Data Archive.  The co-added data produce the deepest submm continuum maps yet of M\\,82, in which low-level 850\\um\\ continuum has been detected out to 1.5\\,kpc, at least 10\\% farther in radius than any previously published submm  detections of this galaxy.  The overall submm morphology and spatial spectral energy distribution of M\\,82 have a general north-south asymmetry consistent with H$_\\alpha$ and X-ray winds, supporting the association of the extended continuum with outflows of dust grains from the disk into the halo.  The new data raise interesting points about the origin and structure of the submm emission in the inner disk of M\\,82.  In particular, SCUBA short wavelength evidence of submm continuum peaks that are asymmetrically distributed along the galactic disk suggests the inner-disk emission is re-radiation from dust concentrations along a bar (or perhaps a spiral) rather than edges of a dust torus, as is commonly assumed.  Higher resolution submm interferometery data from the Smithsonian Submillimeter Array and later Atacama Large Millimeter Array should spatially resolve and further constrain the reported dust emission structures in M\\,82. ",
        "introduction": "\\label{sec:intro_m82}  Massive ejections of gas and dust originating in galactic nuclei have been observed at scales of a few kpc in optical emission lines, submillimeter (submm) molecular lines, mid-infrared (MIR), ultraviolet (UV), submm and radio continuum emission, and soft X-rays \\citep[e.g.,][]{wat84, hec90, dev99, hoo05, eng06}.  One explanation for the outflows is that a high supernova rate in the galactic nucleus heats up the surrounding gas to high temperatures with sound speeds exceeding the escape velocity of the galaxy, creating a $wind$ that expands outward from the galaxy \\citep{che85}.  The wind entrains cosmic rays, warm and cool gas, as well as cool dust, making the outflow directly visible in many wavebands.  The outflows are usually oriented along the minor axes of the galaxies and are thus most easily observed in edge-on galaxies. In local starbursts and high-z Lyman Break galaxies with high enough global star-formation rate per unit area, the superwinds are common and responsible for expelling metals from these galaxies and enriching the inter-galactic medium (IGM) and therefore have implications in the evolution of galaxies and the IGM \\citep[see, e.g,][for recent reviews]{hec03, vei05}.   M\\,82 (NGC\\,3034) is a nearby and popular object in which to investigate the physical association between galactic nuclei and large-scale outflows.  The galaxy is edge-on with an inclination of about $10^{\\circ}$ at position angle $72^{\\circ}$ and is classified as IrrII.  At an estimated distance of 3.63\\,Mpc (as determined for M\\,81 by \\citet{fre94}), M\\,82 has optical dimensions of $11'.2 \\times 4'.3$, i.e. $\\sim 11.8 \\times 4.5$\\,kpc (image scale is $\\sim$17.6 arcsec$^{-1}$).  The nuclear region, within $4' \\times 2'$ about the major axis of the galactic disk, has numerous point sources or emission concentrations, some of which originate from supernovae and massive star clusters and have been detected from the X-ray to radio wavebands.  Layers of dust filaments laden these inner regions producing severe optical extinction and copious infrared to submm re-radiated emission.   New Hubble Heritage Team optical images obtained with a deep 6-point mosaic in B (0.45\\,$\\mu$m), V (0.55\\,$\\mu$m), I (0.81\\,$\\mu$m), and H$_\\alpha$ (0.65\\,$\\mu$m) filters of the Advance Camera for Surveys (ACS) on board the {\\it Hubble Space Telescope} ($HST$) exhibit detailed, filamentary outflows of the M\\,82 \\citep[][]{mut07}, especially in H$_\\alpha$. As described above, it is thought that this outflow is being driven by the copious  formation of massive stars (or a starburst) and subsequent explosions of supernovae. The starburst outflow not only provides the ejection mechanism for the material from the galactic nucleus, but also heats the gas and ionizes the hydrogen, causing it to glow with the red light of H$_\\alpha$ emission line.  These new optical images as well as earlier detailed natural-color composite images of M\\,82 obtained with $HST$ \\citep[see, e.g.,] []{gri01} and the Subaru Telescope \\citep[][]{ohy02} show more than 100 compact groupings of about $10^{5}$ stars in very bright star clusters sprinkled throughout the galaxy's central region, prominent dust lanes that crisscross the disk, knotty filaments of ionized gas that have rich nebular spectra that are not especially enriched in nitrogen, hydrogen gas in a strong galactic wind that is clearly below the galactic center and to the right of the central region, along with many other regions of varying star-formation environments in the nuclear parts of this galaxy.  The huge clusters of massive stars, numerous X-ray and radio detected supernovae, gas concentrations, optically-dramatic dust filaments, galactic winds and other active nuclear features have been attributed to a large burst of star formation $10^7$ to $10^8$ years ago, that was probably triggered by a tidal interaction with the nearby spiral galaxy M\\,81 and dwarf starburst galaxy NGC\\,3077 (see, e.g., \\citet[][]{yun94} for evidence of H\\,I tails linking the three and \\citet []{sch00} for a recent review).   The proximity of this galaxy makes it possible to observe the region of interaction between the star-formation regions and the halo, including expanding shells or bubbles and ``chimneys'' that are producing a clearer picture of the localized driving mechanisms for the outflows \\citep[e.g.,][]{hec90, wil99, wes07}.  The proximity also allows the detections of {\\em low-level}\\, emission in the halo and consequently the determination of the amounts and possible origin of material that results in this emission \\citep[e.g.,][]{sea01,eng06}. Studying the contents and interactions between the star-formation regions and the halo is important for understanding their role in the evolution of M\\,82 and may provide clues to general galaxy evolution as well as details of the composition of the intergalactic material.  This paper focuses on the Submillimetre Common-User Bolometer Array (SCUBA) maps of the copious submm re-radiated emission that results from the dust-laden, star-forming disk, as well as large-scale, low-level emission that is associated with the outflows in the halo of M\\,82. Submm continuum observations of the dusty central regions in M\\,82 were previously obtained with the submm continuum receiver UKT\\,14 on the James Clerk Maxwell Telescope (JCMT) by \\citet{hug90,hug94} and later with SCUBA, also on the JCMT, by \\citet{lee98_pro} and \\citet{alt99}. The current study was intended to extend these previous imaging submm observations in spatial extent, sensitivity and to all the four submm wavelengths available with the SCUBA array  (see Section~\\ref{sec:obs_m82}).  The study includes a co-addition of data mined from the SCUBA Data Archive.  The co-added data produce the deepest submm continuum maps yet of M\\,82, in which low-level emission is detected out to 1.5\\,kpc for the first time in the submm continuum of this galaxy.  The deep maps are used in a detailed morphological study of the nuclear and large-scale detections (see Sections~\\ref{sec:general} and ~\\ref{sec:m82_dis}), including a focused comparative analysis to optical (see Section~\\ref{sec:m82_submm_vs_opt}) and high-resolution CO\\,(1-0) (see Section~\\ref{sec:m82_submm_vs_co}) morphology. The maps are also used in the computation the first submm spatial spectral energy distribution (SED) of separate locations in the nuclear star-forming region of M\\,82 (see Section~\\ref{sec:m82_dis}). These observational results are used in the discussion of the origin and structure of submm continuum morphology and spatial SED of M\\,82 and reviewed in the context of relevant interpretations by other researchers, including those who use data from other wavelengths (see Section~\\ref{sec:m82_dis}).  In particular, (a) the commonly assumed interpretation that the double emission peaks that are seen in the mm to infrared continuum are due to emission from the edges of an inclined, dusty-molecular torus is challenged (see Section~\\ref{sec:m82_origin}), (b) an analytical review of CO results is undertaken to assess if CO emission may significantly contaminate the continuum observed in SCUBA filters (see Section~\\ref{sec:m82_contamination}), and (c) a morphological comparison is conducted to check whether the localized outflows that are reported in high resolution radio, CO, and SiO maps respectively by \\citet[][]{wil99}, \\citet[][]{wei99}, and \\citet[][]{gar01} can be seen in the SCUBA maps (see Section~\\ref{sec:m82_implications}). Finally, the overall implication of the results are discussed and possible future work is outlined (see Section~\\ref{sec:m82_sum}). ",
        "conclusions": "\\label{sec:m82_sum}  SCUBA 350, 450, 750 and 850\\,$\\mu$m imaging observations have been presented of the dust-laden, star-forming inner disk and large-scale, low-level emission that is associated with the outflows in the halo of M\\,82.  The displayed maps include co-added data that were mined from the SCUBA  Data Archive, resulting in the deepest submm continuum maps of M\\,82.  The 850\\,\\um\\ morphology has a single emission peak that is centered about $9''$ west of the galactic nucleus, while the 450 and 350\\,\\um\\ maps have two emission peaks centered about $10''$ and $6''$ respectively east and west  of the nucleus along the galactic disk, similar to previous continuum observations at mm, submm and mid-infrared wavelengths, as well as to CO line transitions and H$_{\\alpha}$ observations of comparable resolution (see Section~\\ref{sec:general}).  In the 750\\um\\ image, the single emission peak seen in the 850\\micron\\ image (see Figure~\\ref{fig:m82_850450} is predictably beginning to be resolved-out into the double peaks seen in the 450 and 350\\micron\\ images.  Low-level emission is detected out to 1.5\\,kpc for the first time in the 850\\,\\um\\ continuum of this galaxy, i.e. at least $\\sim 160$\\,pc radius farther-out than other recent studies.  The deep maps were used in a detailed morphological study of the disk and large-scale detections, including a comparative analysis of submm to optical morphology (see Section~\\ref{sec:m82_submm_vs_opt}). The overall, extended submm morphology of M\\,82 generally resembles the optical picture in that the disk emission has an apparently elliptical shape whose major axis is clearly aligned with that of the galactic disk at position angle $\\sim 72^{\\circ}$.  However, the submm morphology is much smoother than the optical picture, and some prominent dust cloud and filamentary lanes that are seen in the optical are not obvious in the submm continuum. One simple explanation for the submm versus optical correspondence (or lack of it) is the different resolutions and sensitivies of the presented observations. If the optical features emit submm emission, it is possible the emission is at a level lower than the current mapped SCUBA sensitivities or smeared and thus erased by the relatively larger SCUBA beam. This can be varified by future higher resolution and more sensitive submm observations that should become possible with ALMA.  A comparative analysis was further conducted of submm to high-resolution CO\\,(1-0) morphology (see Section~\\ref{sec:m82_submm_vs_co}). Some resolved peaks in the CO maps could be associated with unresolved features in the submm maps. However, there are differences that show the CO and dust emission do not trace each other in a very simple way and suggest that although the dust and CO generally  appear mixed in the central star-forming region, in fact there are differences in their spatial distributions. The different concentrations are most probably locations of varying gas-to-dust densities or star-formation environments, that have been reported in M\\,82 \\citep[e.g.,][]{gri01,pet00}. This can be varified by future higher resolution observations of dust in M\\,82 that should be possible with the SMA and later ALMA.   The SCUBA maps were also used in a computation of the first submm spatial SED analysis of locations within and outside the central star-forming region of M\\,82 (see Section~\\ref{sec:m82_fluxes}) and in the discussion of the origin and structure of submm maps (see Section~\\ref{sec:m82_dis}).  In particular, (a) the commonly assumed interpretation that the double emission peaks that were seen in the mm to infrared continuum are due to emission from the edges of an inclined, dusty-molecular torus was challenged (see Section~\\ref{sec:m82_origin}), (b) an analytical review of CO results was undertaken to assess if CO emission might significantly contaminate the continuum observed in SCUBA filters (see Section~\\ref{sec:m82_contamination}), and (c) a morphological comparison was conducted to check whether the localized outflows that were reported in radio and SiO maps respectively by \\citet[][]{wil99} and \\citet[][]{gar01} could be seen in the SCUBA maps (see Section~\\ref{sec:m82_implications}).  Evidence in this paper in terms of clearly asymmetric inner-disk submm emission seen in the presented intensity seems to disfavor the commonly assumed interpretation of a dusty torus in M\\,82. Arguments were presented to explane the inner-disk submm maps of M\\,82 in the context of emission from a rather complex distribution of dust concentrations that are in regions of different star-formation environments, as has been reported from various studies using data at other wavelengths \\citep[e.g.,][]{ach95,sch00,gri01}.  It is not obvious if the CO contribution to the higher frequencies of SCUBA will be less than or higher than the 47\\% estimated to the SCUBA-850\\,\\um\\ band by \\citet{sea01}.  However, it is clear that the CO contamination to higher-frequency continuum may also be significant and warrants detailed investigation.  Future work on this galaxy will attempt to acquire the data of the CO lines in the 450 and 350\\,\\um\\ bands and make a quantitative comparison of these data in order to determine the possible contributions of CO to the high-frequency-SCUBA data.  The overall submm low-level morphology has a general north-south asymmetry that is similar to the H$_\\alpha$ winds and CO and X-ray outflows that have been detected in M\\,82 \\citep[e.g.,][]{sho98}.  The submm spatial SED analysis also shows thermal properties that change along the minor axis of the galaxy disk, similar to the radio spectral index gradient by \\citet[][]{sea91} and \\citet[][]{wil99} and consistent with the north-south asymmetric, large-scale X-ray and H$_{\\alpha}$ winds.  Therefore, the current results support the simple interpretation \\citep[e.g.,][]{lee98_pro,alt99} that the asymmetric morphology in the submm maps is a manifestation of corresponding outflows of dust grains from the galactic disk into the halo. As noted above, this work has presented low-level 850\\,\\um\\ continuum emission out to 1.5\\,kpc for the first time in this galaxy, i.e. at least $\\sim 160$\\,pc radius farther-out than other recent studies."
    },
    "0812/0812.1708_arXiv.txt": {
        "abstract": "We present the results of a targeted 3-mm spectral line survey towards the eighty-three 6.67\\,GHz methanol maser selected star forming clumps observed by \\citet{Purcell2006}. In addition to the previously reported measurements of \\hcop\\,(1\\,--\\,0), \\hthirteencop\\,(1\\,--\\,0), and \\chthreecn\\,(5\\,--\\,4) \\& (6\\,--\\,5), we used the Mopra antenna to detect emission lines of \\ntwohp\\,(1\\,--\\,0), HCN\\,(1\\,--\\,0) and HNC\\,(1\\,--\\,0) towards 82/83 clumps (99 per cent), and \\chthreeoh\\,(2\\,--\\,1) towards 78/83 clumps (94 per cent).  The molecular line data have been used to derive virial and LTE masses, rotational temperatures and chemical abundances in the clumps, and these properties have been compared between sub-samples associated with different indicators of evolution. The greatest differences are found between clumps associated with 8.6\\,GHz radio emission, indicating the presence of an Ultra-Compact \\hii~region, and `isolated' masers (without associated radio emission), and between clumps exhibiting \\chthreecn~emission and those without. In particular, thermal \\chthreeoh~is found to be brighter and more abundant in Ultra-Compact \\hii~(\\uchii) regions and in sources with detected \\chthreecn, and may constitute a crude molecular clock in single dish observations.  Clumps associated with 8.6\\,GHz radio emission tend to be more massive {\\it and} more luminous than clumps without radio emission. This is likely because the most massive clumps evolve so rapidly that a Huper-Compact\\hii~or \\uchii~region is the first visible tracer of star-formation.  The gas-mass to sub-mm/IR luminosity relation for the combined sample was found to be  L$\\propto$\\,M$^{0.68}$, considerably shallower than expected for massive main-sequence stars. This implies that the mass of the clumps is comparable to, or greater than the mass of the stellar content.  We find also that the mass of the hot core is correlated with the mass of the clump in which it is embedded. ",
        "introduction": "Molecular emission is a powerful tool when used to investigate the physical and chemical conditions in hot cores. Transitions requiring different temperatures and densities for excitation constitute an excellent probe of physical structure. Because the chemical properties of hot cores vary with time, the relative molecular abundances can also be used as indicators of evolution.  In the preceding decades, representative line surveys of a limited sample of cores have begun to identify the molecules most suited to investigating the process of massive star formation. To date these have been restricted to a few objects: Orion-KL (\\citealt{Blake1986,Turner1989,Ziurys1993,Schilke1997}), Sgr-B2 (\\citealt{Cummins1986,Turner1989,Sutton1991}), G34.3+0.15 (\\citealt{MacDonald1996,Kim2000,Kim2001}), IRAS 17470$-$2853 (also known as G5.89$-$0.39, \\citealt{Kim2002}), G305.2+0.2, \\citep{Walsh2006}.  Chemical models, including both gas-phase and grain-surface reactions have been developed for specific objects whose physical structure and conditions are well known, e.g. G34.3+0.15 \\citep{Millar1997b,Thompson1999}, as well as for general cases, e.g. \\citet{Rodgers2003}. These and earlier models are limited by the computational power available at the time and consider either time-dependent chemistry  at a test position, or a `snapshot' of spatial abundances at a single time. Advances in computational power have recently allowed the development of models which consider both time- {\\it and} space-dependent chemistry simultaneously, e.g., the model of AFGL\\,2591 by \\citet{Doty2002}. Today it is the lack of observational constraints which restrict our understanding of chemistry in young massive stellar objects. The priority in the coming years must be to assemble a large sample of molecular abundance measurements towards a broad range of massive star forming regions, at different stages of evolution and with luminosities and masses from B3 to O5 type massive stars.  In this paper we present the remaining results of the Mopra `Hot Molecular Cores' (HMC) survey towards 6.67\\,GHz methanol (\\chthreeoh) maser selected targets. The initial results were presented in \\citet{Purcell2006} (hereafter Paper 1), wherein we reported the detection of the hot-core species methyl-cyanide (\\chthreecn) and the observations of the formyl ion (\\hcop), and its isotopomer (\\hthirteencop). Here we expand the analysis to cover thermal transitions of methanol (\\chthreeoh), carbon-monoxide (\\thirteenco), diazenylium (\\ntwohp), hydrocyanic acid (HCN) and hydroisocyanic acid (HNC). We analyse the full complement of molecules for evidence of evolution.  Note: Tables~3-9 are available in electronic form only along with additional supporting text and figures. ",
        "conclusions": "In continuation of the `Hot Molecular Cores' survey we used the Mopra telescope to search for 3-mm transitions of \\chthreeoh\\,(2\\,--\\,1), \\ntwohp\\,(1\\,--\\,0), \\thirteenco\\,(1\\,--\\,0), HCN\\,(1\\,--\\,0) and HNC\\,(1,--\\,0). This is in addition to the previously reported measurements of \\hcop\\,(1\\,--\\,0), \\hthirteencop\\,(1\\,--\\,0), and \\chthreecn\\,(5\\,--\\,4) \\& (6\\,--\\,5), in Paper 1. Molecular emission was detected in all but one source (G10.10$-$0.72), with \\ntwohp, \\thirteenco, HNC and HCN detected towards all of the remaining 82 targets. Thermal \\chthreeoh~emission was detected in 78 sources (94 per cent). The following conclusions have been drawn from the data:  \\begin{enumerate} \\item Virial masses estimated from \\ntwohp~are comparable to or lower than the gas masses of the sample. Values for M$_{\\rm gas}$ have been calculated from the 1.2-mm dust emission and range from $30$\\,\\msun~to $5\\times\\,10^3$\\,\\msun, while  values for M$_{\\rm virial}$ range from $\\sim$\\,40 to $2.5\\times\\,10^3$\\,\\msun. We note that \\ntwohp~exhibits particularly narrow linewidths with an average of 3.0\\,\\kms.  \\item H$_2$ volume densities have also been estimated from the surface brightness and extent of 1.2-mm emission. The mean value is $5\\times\\,10^4$\\,\\cmmthree, however, we note that the assumptions inherent in the calculation introduce large unknown uncertainties in values for individual sources.  \\item Rotational temperatures derived from \\chthreeoh~range from 3.0\\,K to 14.0\\,K, with a mean value of 6.67\\,K. These anomalously low temperatures strongly suggest that the \\chthreeoh~is sub-thermally excited. Assuming a common excitation temperature, we find the abundance ratio of A to E-type \\chthreeoh~ranges from 0.5 to 2.8, with a mean of 1.5, consistent with previously published values.  \\item \\ntwohp~is found to be optically thin towards most of the sample. Likewise the \\chthreeoh~rotational diagrams are consistent with low optical depths in the J\\,=\\,2\\,--\\,1 transitions. \\thirteenco, HCN and HNC all exhibit line profile asymmetries and may be significantly optically thick.  \\item The {\\it luminosity/gas-mass} relationship for the sample is found to follow a rough power law of the form L$\\propto$\\,M$^{0.68}$, meaning that the lower mass clumps are over-luminous and the higher mass clumps are under-luminous. We find that the gas-mass of the clumps is comparable to or greater than the mass of the stellar content (implied by the luminosity). We interpret this as an indicator of young ages, as the newborn massive stars still retain their natal molecular clouds.  \\item The radio loud sub-sample (\\uchii~regions) is on average more massive {\\it and} more luminous than the radio-quiet sub-sample (isolated masers). We argue that the the most likely explanation for this disparity is due to the comparatively rapid timescale over which the most massive objects evolve into \\hii~regions. The first significant tracer of star-formation in these objects may be radio-continuum emission from a HC\\hii~or \\uchii~region. It is also possible that some of the isolated maser sample contains hyper-compact \\hii~regions which are too optically thick to detect at 8\\,GHz. In either case there is a clear age difference between the subsamples.  \\item We separated the source list into sub-samples associated with different tracers of evolution and compared their median physical and chemical properties. Our findings are as follows:  \\begin{itemize} \\item {\\it All} spectral lines are brighter and more luminous towards the radio loud sample, extending the result found in Paper 1.  \\item The thermal \\chthreeoh~brightness and abundance is found to be enhanced (compared to other species) in radio-loud and \\chthreecn-bright sources. We suggest that in single-dish observations thermal \\chthreeoh~may be used as a crude molecular clock to indicate the evolutionary state of a clump.  \\item Given the small sample size, conclusions drawn from a comparison of MSX-bright and MSX-dark sources are not significant. It is worth noting, however, that linewidths and line-luminosities are generally greater in the infrared dark clouds, indicating perhaps that when selected using \\chthreeoh~masers, dark-clouds conceal objects at later stages of evolution.  \\item We divided the source list into sub-samples with and without high velocity linewings and \\hcop-infall-profiles, but found no significant differences in chemical or physical properties. \\end{itemize}  \\item We find that the beam-averaged \\chthreecn~column density is correlated with the gas-mass derived from 1.2-mm emission. This suggests that the mass of the hot-core is dependent on the mass of the clump in which it is embedded. \\end{enumerate}"
    },
    "0812/0812.2221_arXiv.txt": {
        "abstract": " ",
        "introduction": "Approximately half of the neutron(\\emph{n})-capture elements (atomic number $Z>30$) in the Universe are created by slow \\emph{n}-capture nucleosynthesis (the ``\\emph{s}-process'').  The \\emph{s}-process can occur in low- and intermediate-mass stars (1--8~M$_{\\odot}$), the progenitors of planetary nebulae (PNe), during the thermally-pulsing asymptotic giant branch (AGB) phase.  Free neutrons are released by the reaction $^{13}$C($\\alpha$,$n$)$^{16}$O --- or $^{22}$Ne($\\alpha$,$n$)$^{25}$Mg in more massive AGB stars ($>3.5$~M$_{\\odot}$) --- in the intershell region between the H- and He-burning shells.  Fe-peak nuclei experience a series of \\emph{n}-captures interspersed with $\\beta$-decays to transform into isotopes of heavier elements.  The enriched material is transported to the stellar envelope via convective dredge-up, and is expelled into the ambient interstellar medium by stellar winds and PN ejection \\citep{bus99}.  Nebular spectroscopy uniquely reveals information about \\emph{n}-capture nucleosynthesis that cannot be obtained from stellar spectra.  For example, the lightest \\emph{n}-capture elements ($Z=30$--36) and noble gases are not well-studied in their sites of synthesis, due to the difficulty in detecting these species in evolved stars or supernova remnants \\citep[e.g.,][]{wally95}.  In contrast, these elements are readily detected in PNe.  Spectroscopy of PNe also enables investigations of \\emph{s}-process nucleosynthesis in intermediate-mass stars ($M=3.5$--8~M$_{\\odot}$).  These stars are obscured by dusty, optically thick circumstellar envelopes during the AGB and post-AGB phases \\citep{hab96, van03} that preclude optical or UV spectroscopy.  Hence, little is known about their contribution to \\emph{s}-process nuclei in the Universe \\citep{trav04}.  These stars produce Type~I PNe \\citep[characterized by He and N enrichments;][]{peim78} in which \\emph{n}-capture elements can be detected.  The low cosmic abundances of \\emph{n}-capture elements \\citep{asp05} prevented their identification in nebular spectra until \\citet[][hereafter PB94]{pb94} discovered emission lines of Se, Br, Kr, Xe, and possibly other trans-iron elements in the optical spectrum of the PN NGC~7027.  Many of these identifications were later confirmed with higher resolution spectroscopy performed by \\citet{zh05} and \\citet{sharp07}.  PB94's analysis led \\citet{din01} to identify two anonymous emission lines in the $K$~band spectra of PNe as fine-structure transitions of [Se~IV] and [Kr~III].  Shortly thereafter, \\citet{ster02} identified Ge~III~$\\lambda$1088.46 in absorption in the UV spectra of four PNe, and found that these objects exhibit Ge enrichments indicative of \\emph{s}-process nucleosynthesis in their progenitor stars.  Together, these studies determined abundances of \\emph{n}-capture elements in only 11 PNe, too small a sample to characterize \\emph{s}-process enrichments in PNe as a population.  \\citet[][hereafter SD08]{sd08} performed the first large-scale survey of \\emph{n}-capture elements in PNe by detecting the near-infrared emission lines [Kr~III]~2.199 and/or [Se~IV] 2.287~$\\mu$m in 81 of 120 Galactic PNe.  To correct for the abundances of unobserved Se and Kr ions, analytical ionization correction factors (ICFs) were derived from a large grid of photoionization models \\citep{ster07}.  We derived Se and Kr abundances (or upper limits) in each object, increasing the number of PNe with known \\emph{n}-capture element abundances by nearly an order of magnitude.  Se and Kr were found to be enriched by a factor of two or more in $\\sim$45\\% of the observed PNe, indicating efficient \\emph{s}-process nucleosynthesis and convective dredge-up in their progenitor stars.  Both exhibit a positive correlation with the C/O ratio, as has been found for other \\emph{n}-capture elements in AGB \\citep{smith90} and post-AGB stars \\citep{van03}.  Type~I PNe, which have intermediate-mass progenitors, display little or no enrichment of Se and Kr compared to other PN subclasses.  These trends have been shown to be consistent with theoretical predictions \\citep{kar09}.  The $K$~band Se and Kr lines were recently detected in three nearby extragalactic PNe by \\citet{din09}, providing insight to the nucleosynthesis of these elements in low-metallicity environments for the first time.  The PN Hen~2-436 in the Sagittarius dwarf spheroidal galaxy exhibits large enrichments of Se and Kr (Figure~1), comparable to the most highly enriched Galactic PNe observed by SD08, while two PNe observed in the Magellanic Clouds do not appear to be enriched.  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.43, angle=0]{figure1.eps} \\caption{$K$~band spectrum of the PN Hen~2-436 in the Sagittarius dwarf spheroidal galaxy, measured with the GNIRS spectrograph on Gemini South \\citep[from][]{din09}.  Both [Kr~III]~2.199 and [Se~IV]~2.287~$\\mu$m were detected, and are identified in the figure.}\\label{fig1} \\end{center} \\end{figure}  The Se and Kr abundances derived in these studies are typically accurate to within a factor of two or three, based on the error analyses of \\citet{ster07} and SD08, and similar uncertainties can be expected for other \\emph{n}-capture elements observed in ionized nebulae.  These uncertainties stem from (1) the detection of only one ion each of Se and Kr, which can lead to large and uncertain ICFs; and (2) the poorly known atomic data that controls the ionization balance of these elements, which precludes the \\linebreak derivation of accurate ICFs.  We present initial results from a multi-disciplinary investigation, comprised of new optical observations and atomic data determinations, designed to address these sources of uncertainty and enable much more accurate abundance analyses of \\emph{n}-capture elements in ionized nebulae. ",
        "conclusions": ""
    },
    "0812/0812.0012_arXiv.txt": {
        "abstract": "\\noindent Motivated by experimental probes of general relativity, we adopt methods from perturbative (quantum) field theory to compute, up to certain integrals, the effective lagrangian for its $n$-body problem. Perturbation theory is performed about a background Minkowski spacetime to $\\mathcal{O}[(v/c)^4]$ beyond Newtonian gravity, where $v$ is the typical speed of these $n$ particles in their center of energy frame. For the specific case of the 2 body problem, the major efforts underway to measure gravitational waves produced by in-spiraling compact astrophysical binaries require their gravitational interactions to be computed beyond the currently known $\\mathcal{O}[(v/c)^7]$. We argue that such higher order post-Newtonian calculations must be automated for these field theoretic methods to be applied successfully to achieve this goal. In view of this, we outline an algorithm that would in principle generate the relevant Feynman diagrams to an arbitrary order in $v/c$ and take steps to develop the necessary software. The Feynman diagrams contributing to the $n$-body effective action at $\\mathcal{O}[(v/c)^6]$ beyond Newton are derived. ",
        "introduction": "In this paper we are concerned with the problem of describing the gravitational dynamics of arbitrary $n \\geq 2$ compact non-rotating bodies moving in a background Minkowski spacetime. By assuming non-relativistic motion, this problem can be approached in a perturbative manner, by approximating these compact objects as point masses and calculating the effective lagrangian $L_{\\text{eff}}[\\{ \\vec{x}_a, \\vec{v}_a, \\dot{\\vec{v}}_a, \\dots \\}]$ for their coordinates $\\{\\vec{x}_a| a = 1,2,\\dots,n\\}$ and their time derivatives $\\{\\vec{v}_a, \\dot{\\vec{v}}_a, \\dots\\}$, up to some given order in the typical speed $v$ of these $n$ objects:\\footnote{We use units where all speeds or velocities are measured in multiples of the speed of light, i.e. $c = 1$.} Newtonian gravity starts at $\\mathcal{O}[v^0]$ and the Einstein-Infeld-Hoffman lagrangian \\cite{EinsteinInfeldHoffman}, that describes the precession of the perihelion of elliptical orbits, is of $\\mathcal{O}[v^2]$ (1 PN).\\footnote{The nomenclature is: $\\mathcal{O}[v^{2Q}] \\leftrightarrow Q$ PN.} The $n$ body problem at $\\mathcal{O}[v^4]$ was first tackled by Ohta et al. \\cite{Ohta}. Some computational and coordinate issues encountered there were clarified by Damour and Sch\\\"{a}fer \\cite{DamourSchaeferNBody}. In the latter, some integrals could not be evaluated. A portion of these were later performed by Sch\\\"{a}fer \\cite{3BodySchaefer}, so that currently, up to the $n=3$ case is known. But to know the effective lagrangian for arbitrary $n$ at this order, one needs to further calculate the integrals for the $n=4$ case. As we will see later, once $L_{\\text{eff}}$ is known up to $n=4$, the arbitrary $n$-body lagrangian will follow from a limited form of superposition.  We will examine this problem using perturbative field theory techniques introduced in \\cite{nrgr}. The motivations are two-fold, both of them stemming from experimental probes of gravitational physics: one requires the 2-body effective lagrangian to higher than $\\mathcal{O}[v^7]$, and the other may need the $n$-body counterpart at $\\mathcal{O}[v^4]$.  \\textsl{Gravitational Waves Detection} \\quad The recent years have seen an array of gravitational wave detectors such as GEO, LIGO, TAMA, and VIRGO coming online. These experiments seek to detect gravitational waves produced by binary black holes and/or neutron stars as they spiral towards each other. Within their frequency bandwidth, these detectors are able to track the frequency evolution of the gravitational waves from these binaries over $\\mathcal{O}[10^4]$ orbital cycles and hence make very accurate measurements. To be able to do so, however, theoretical templates need to be constructed so that the raw data can be integrated against them to determine if there is a significant correlation. Via a generalized Kepler's third law relating orbital frequency to the binary separation distance, these templates are based on energy balance: the rate of energy loss of these binaries is equal to the power in the gravitational radiation emitted. Both the notion of energy and expressions for the flux of gravitational radiation require the knowledge of the dynamics of these binaries, which in turn is encapsulated in their effective lagrangian. Due to the high accuracy to be attained, this effective lagrangian needs to be computed up to 3 PN and higher.\\footnote{Blanchet \\cite{Blanchet2BodyLADM} offers a review of the post-Newtonian framework and its relation to gravitational wave experimental observables.}  Currently, the dynamics of compact astrophysical binaries is known up to 3.5 PN.\\footnote{The half integer PN order lagrangians, scaling as odd powers of $v$ relative to Newtonian gravity, describe dissipation -- gravitational waves produced by and interacting with the $n$ compact objects. In this present paper, we shall focus only on the conservative part of their dynamics up to 2 PN.} (See \\S1.3 of Blanchet \\cite{Blanchet2BodyLADM} and the references therein.) To obtain the dynamics at 4 PN and beyond is a challenging task. Because of the need to regularize the divergences that arise from approximating compact objects as point particles, one may wish to engage field theoretic methods to handle them. Such a pursuit was initiated in \\cite{nrgr}, where it was shown how to carry out the field theory effective lagrangian calculation in a systematic manner by first doing some dimensional analysis. One of the main thrusts of this present work is to attempt to make as methodical as possible such a route in post-Newtonian calculations. In particular, we advocate using the computer to automate the process, so that at the end only those Feynman diagrams that truly require human intervention are left for manual evaluation. Given the computational effort required at 4 PN and beyond, we believe this is necessary not only to save time and energy, but also to reduce human errors. For example, at such a high PN order, even the derivation of the necessary diagrams will itself be non-trivial -- the reader not convinced of this fact is encouraged to look at appendix \\ref{3PNDiagramsSection} containing the 3 PN diagrams -- but an efficient implementation of the algorithm that we will sketch in the main body of this work will allow automatic generation of Feynman diagrams to arbitrary PN order, modulo computing power.  \\textsl{Solar System Gravity} \\quad Closer to Earth, the Einstein-Infeld-Hoffman lagrangian, at $\\mathcal{O}[v^2]$ beyond Newton, is routinely used to compute the solar system ephemerides, and to analyze spacecraft trajectories and space based gravitational experiments. A range of experiments, such as the new lunar ranging observatory APOLLO, proposals to land laser ranging missions on Mars and/or Mercury, and spacecraft laboratories -- GTDM, LATOR, BEACON, etc. -- will begin to probe the non-Euclidean nature of the solar system's spacetime geometry beyond 1 PN by measuring the timing and deflection of light propagation more precisely than before. (See, for instance, Turyshev \\cite{TuryshevTestGR} for a recent review.)  Within the point particle approximation, both the solar system dynamics and its geometry can be gotten at simultaneously by computing from general relativity the effective $n$-body lagrangian. Because general relativity is a non-linear field theory, knowledge of the 2 body lagrangian is not sufficient to deduce its $n$-body counterpart, as superposition is not obeyed. That the $n$-body $L_{\\text{eff}}$ encodes not only dynamics $\\{\\vec{x}_a[t]\\}$ but also the geometry $g_{\\mu\\nu}[t,\\vec{x}]$ can be seen by adding a test particle to the $n$-body system.\\footnote{This observation can be found, for example, in Damour and Esposito-Farese \\cite{Damour_Esposito-Farese_ReadgOffLeff}.} Denoting the latter's mass and coordinate vector as $M_\\epsilon$ and $y^\\mu \\equiv (t, \\vec{y})^\\mu$ respectively, in the limit as $M_\\epsilon$ tends to zero relative to the rest of the other masses in the system, we know its exact action has to be\\footnote{We work in cartesian coordinates and employ the $\\eta_{\\mu\\nu} = \\text{diag}[1,-1,\\dots,-1]$ sign convention. The Einstein summation convention is adopted. Greek letters run from 0 to $d-1$ while English alphabets run from $1$ to $d-1$.}  {\\allowdisplaybreaks \\begin{align*} - M_\\epsilon & \\int \\dx{t} \\sqrt{ \\bar{g}_{\\mu\\nu} \\frac{\\text{d} y^\\mu}{\\dx{t}} \\frac{\\text{d} y^\\nu}{\\dx{t}} } \\\\ &= - M_\\epsilon \\int \\dx{t} \\bigg( 1 - \\frac{1}{2} \\left( \\frac{\\dd \\vec{y}}{\\dx{t}} \\right)^2 + \\frac{1}{2} \\delta g_{00}[z] \\nonumber \\\\ &\\qquad + \\delta g_{0i}[z] \\frac{\\text{d} y^i}{\\dx{t}} + \\frac{1}{2} \\delta g_{ij}[z] \\frac{\\text{d} y^i}{\\dx{t}} \\frac{\\text{d} y^j}{\\dx{t}} + \\dots \\bigg), \\\\ \\bar{g}_{\\mu\\nu} &\\equiv \\eta_{\\mu\\nu} + \\delta g_{\\mu\\nu} \\\\ z &\\equiv \\{\\vec{x}_a,\\vec{v}_a,\\dot{\\vec{v}}_a,\\dots\\},t,\\vec{y}; \\quad a = 1,2,\\dots,n, \\end{align*}}  since it now moves along a geodesic on the spacetime metric generated by the rest of the $n$ masses. Therefore, if $L_{n+1}$ is the $(n+1)$-body lagrangian less the $M_\\epsilon(-1 + (1/2) \\left( \\dd \\vec{y}/\\dx{t} \\right)^2)$, the deviation of the spacetime metric from Minkowski $\\delta g_{\\mu\\nu}$ can be read off the action of the test particle using the prescription:  {\\allowdisplaybreaks \\begin{align*} &\\delta g_{00}[t, \\vec{x}] = -\\left. 2 \\frac{\\partial}{\\partial M_\\epsilon} L_{n+1}[\\vec{y} = \\vec{x}] \\right\\vert_{\\dot{\\vec{y}} = \\ddot{\\vec{y}} = \\dots = M_\\epsilon = 0} \\\\ &\\delta g_{0i}[t, \\vec{x}] = -\\left. \\frac{\\partial}{\\partial M_\\epsilon} \\frac{\\partial L_{n+1}}{\\partial (\\dd y^i/\\dx{t})} [\\vec{y} = \\vec{x}] \\right\\vert_{\\dot{\\vec{y}} = \\ddot{\\vec{y}} = \\dots = M_\\epsilon = 0} \\\\ &\\delta g_{ij}[t, \\vec{x}] \\\\ &= -\\left. \\frac{\\partial}{\\partial M_\\epsilon} \\frac{\\partial^2 L_{n+1}}{\\partial (\\dd y^i/\\dx{t}) \\partial (\\dd y^j/\\dx{t})}[\\vec{y} = \\vec{x}]\\right\\vert_{\\dot{\\vec{y}} = \\ddot{\\vec{y}} = \\dots = M_\\epsilon = 0} \\end{align*}}  We see that understanding and testing the dynamics -- the equations that govern the time evolution of the $\\{\\vec{x}_a\\}$ -- is intimately tied to understanding and testing the spacetime geometry $g_{\\mu\\nu}$ of the solar system.  The outline of the paper is as follows. In section \\ref{setup}, we set up a lagrangian description of the system of $n$ compact astrophysical objects by approximating them as $n$ point particles. Einstein's equations can then be solved perturbatively as a Born series, whose graphical representation are the Feynman diagrams containing no graviton loops; the result of summing the diagrams yield the $n$-body effective action we seek. (Our description will be brief because it will merely be an overview of the methods developed in \\cite{nrgr}.) We then sketch the algorithm that could be used to generate the necessary Feynman diagrams contributing to the effective action up to an arbitrary PN order. In section \\ref{diagramresults}, we calculate the individual diagrams that occur at the Newtonian thru 2 PN order and present the effective action of the $n$-body system up to certain integrals (\\ref{0PNLeff}, \\ref{1PNLeff}, \\ref{2PNnBody}). As a by-product, we reproduce the known 2 PN 2 body lagrangian. In the appendixes, we discuss the integrals encountered in the diagrams; the algorithm for generating the $N \\geq 2$ graviton Feynman rules on a computer; and also display the Feynman diagrams occurring at the 3 PN order. ",
        "conclusions": "In this paper we have, following \\cite{nrgr}, used a point mass approximation for the $n$-body in general relativity, allowing us to obtain a lagrangian description at the cost of introducing an infinite number of terms in the action. Because we are seeking the 2 PN effective lagrangian, however, only the minimal terms $\\{-M_a \\int\\dx{s}_a\\}$ are necessary. By examining the physical scales in the problem, we have described how to organize our action (\\ref{fullaction}) and outlined an algorithm that would allow us in principle to generate all the necessary Feynman diagrams up to an arbitrary PN order for a given set of point particle actions (minimal or not); as well as automate the computation process so that as few of the Feynman diagrams as possible are left for human evaluation. This way, the post-Newtonian program can be pursued in an efficient and systematic manner within the framework of perturbative field theory, and the necessary software may be developed to tackle the effective 2 body lagrangian calculation at 4 PN and beyond. In the bulk of this work, we obtained in closed form the conservative portion of the effective lagrangian $L_{\\text{eff}}[\\{\\vec{x}_a,\\vec{v}_a,\\dot{\\vec{v}}_a\\}]$ up to 1 PN for the general case of $n$ point masses in $d \\geq 4$ spacetime dimensions and up to 2 PN for 2 point masses in (3+1)-dimensions.  It is apparent that the primary bottleneck of higher post-Newtonian calculations is one of calculus. For the $n$-body problem, it is the analytic evaluation of integrals such as $I_{22}$ (\\ref{I22Def}) and the $I_{1\\dots N}$ (\\ref{npoint}) for $N \\geq 4$. At the same time, it is possible that choosing a different gauge from the one used in (\\ref{gaugefixing}) and/or a different parametrization of $h_{\\mu\\nu}$ may help reduce the number of diagrams and the amount of work needed in manipulating the momentum dot products from the tensor contraction of graviton vertices in fourier space. The reduction of diagrams at 2 PN was recently demonstrated in the 2 body case computed by Gilmore and Ross \\cite{GilmoreRoss}. They used the full de Donder gauge\\footnote{This will modify the $N$-graviton interaction in the Einstein-Hilbert action to all orders in $h_{\\mu\\nu}$.} $S_{\\text{gf}} = \\intx{d}{x} \\sqrt{g} g^{\\alpha\\beta} g^{\\mu\\nu} g^{\\sigma\\rho} \\Gamma_{\\sigma\\mu\\nu} \\Gamma_{\\rho\\alpha\\beta}$ and the Kaluza-Klein parametrization for $h_{\\mu\\nu}$, first introduced to the PN problem by Kol and Smolkin \\cite{KolSmolkin}.\\footnote{Comparing the calculus involved, however, when using the Kol-Smolkin parametrization for $h_{\\mu\\nu}$, one notes that Gilmore and Ross \\cite{GilmoreRoss}, for the 2 body problem, encountered the same master integrals as the ones used in this paper, (\\ref{integral_PropI}) and (\\ref{integral_OneLoopRd}). Moreover, for the $n$-body case, the most difficult integrals at 2 PN, arising from the contraction of two 3-graviton vertices, namely $I_{0000-0000}^{\\text{III}}[q,r,s,u]$ in (\\ref{masterintegralI0000-0000_III}), are identical in form to the ones in their equation (52), keeping all particles distinct.} Some other possible choices include the ADM variables normally associated with the (3+1)-decomposition of the spacetime metric. Yet another possibility is to employ the gravitational lagrangian constructed by Bern and Grant \\cite{BernGrant_EH=YMYM} using quantum chromodynamics gluon amplitudes, up to the 5-graviton interaction; this is sufficient, however, only to 3 PN.  We end with a cautionary remark against taking superposition too literally within the post-Newtonian framework. By using the 1 PN lagrangian (\\ref{1PNLeff}) in 3 spatial dimensions, and taking the continuum limit, the force $\\vec{F}$ experienced by a stationary point mass $M_x$ located at $\\vec{r}$ away from the center of a static, spherical, hollow shell of surface mass density $\\sigma$ and coordinate radius $R$ can be shown to be\\footnote{This computation arose out of discussions on Birkhoff's theorem in GR with Dai De-Chang and Glenn Starkman. The following formula can also be found in their recent paper with Matsuo \\cite{FDivergesAtShell}. Its derivation actually only requires the 2 body portion of the $L_{\\text{eff}}$, because the 3-body portion integrates to a constant.}  \\begin{align*} \\vec{F} &\\equiv M_x \\frac{\\dd^2 \\vec{r}}{\\dx{t}^2} \\nonumber \\\\ &= -M_x^2 \\frac{\\pi R G_{\\text{N}}^2 \\sigma}{1 + 12 \\pi R G_{\\text{N}} \\sigma} \\frac{\\partial}{\\partial \\vec{r}} \\left( \\frac{1}{|\\vec{r}|} \\ln\\left\\vert \\frac{|\\vec{r}|+R}{|\\vec{r}|-R} \\right\\vert \\right), \\end{align*}  which evidently diverges as the point mass approaches the surface of the shell. Because the force would vanish if gravity were purely Newtonian, such a result for a first order calculation most likely indicates the breakdown of perturbation theory in this regime, since the post-Newtonian lagrangian was derived with an implicit assumption that the point masses involved were well separated, i.e. $r_{\\text{s}} \\ll r$."
    },
    "0812/0812.0362_arXiv.txt": {
        "abstract": "We study predictions for dark matter phase-space structure near the Sun based on high-resolution simulations of six galaxy halos taken from the Aquarius Project. The local DM density distribution is predicted to be remarkably smooth; the density at the Sun differs from the mean over a best-fit ellipsoidal equidensity contour by less than $15\\%$ at the $99.9\\%$ confidence level. The local velocity distribution is also very smooth, but it differs systematically from a (multivariate) Gaussian distribution. This is not due to the presence of individual clumps or streams, but to broad features in the velocity modulus and energy distributions that are stable both in space and time and reflect the detailed assembly history of each halo. These features have a significant impact on the signals predicted for WIMP and axion searches. For example, WIMP recoil rates can deviate by $\\sim 10\\%$ from those expected from the best-fit multivariate Gaussian models. The axion spectra in our simulations typically peak at lower frequencies than in the case of multivariate Gaussian velocity distributions. Also in this case, the spectra show significant imprints of the formation of the halo. This implies that once direct DM detection has become routine, features in the detector signal will allow us to study the dark matter assembly history of the Milky Way.  A new field, ``dark matter astronomy'', will then emerge. ",
        "introduction": "In the 75 years since \\cite{1933AcHPh...6..110Z} first pointed out the need for substantial amounts of unseen material in the Coma cluster, the case for a gravitationally dominant component of non-baryonic dark matter has become overwhelmingly strong. It seemed a long shot when \\cite{1982ApJ...263L...1P} first suggested that the dark matter might be an entirely new, weakly interacting, neutral particle with very low thermal velocities in the early universe, but such Cold Dark Matter (CDM) is now generally regarded as the most plausible and consistent identification for the dark matter.  Particle physics has suggested many possible CDM particles beyond the standard model. Two promising candidates are WIMPs \\citep[weakly interacting massive particles, see][]{1977PhRvL..39..165L, 1978ApJ...223.1015G,1984STIN...8526469E} and axions \\citep{1977PhRvL..38.1440P, 1977PhRvD..16.1791P, 1978PhRvL..40..223W, 1978PhRvL..40..279W}. Among the WIMPs, the lightest supersymmetric particle, the neutralino, is currently favoured as the most likely CDM particle, and the case will be enormously strengthened if the LHC confirms supersymmetry. However, ultimate confirmation of the CDM paradigm can only come through the direct or indirect detection of the CDM particles themselves.  Neutralinos, for example, are their own antiparticles and can annihilate to produce $\\gamma$-rays and other particles. One goal of the recently launched Fermi Gamma-ray space telescope is to detect this radiation \\citep{1999APh....11..277G,2008Natur.456...73S}.  Direct detection experiments, on the other hand, search for the interaction of CDM particles with laboratory apparatus.  For WIMPs, detection is based on nuclear recoil events in massive, cryogenically cooled bolometers in underground laboratories \\citep{1996PhR...267..195J}; for axions, resonant microwave cavities in strong magnetic fields exploit the axion-photon conversion process \\citep{1985PhRvD..32.2988S}.  Despite intensive searches, the only experiment which has so far reported a signal is DAMA \\citep[][]{2007hepi.conf..214B} which has clear evidence for an annual modulation of their event rate of the kind expected from the Earth's motion around the Sun. The interpretation of this result is controversial, since it appears to require dark matter properties which are in conflict with upper limits established by other experiments \\citep[see][for a discussion and possible solutions]{2004PhRvD..70l3513S,2005PhRvD..71l3520G,2006JPhCS..39..166G}. Regardless of this, recent improvements in detector technology may enable a detection of ``standard model'' WIMPS or axions within a few years.  Event rates in all direct detection experiments are determined by the local DM phase-space distribution at the Earth's position. The relevant scales are those of the apparatus and so are extremely small from an astronomical point of view. As a result, interpreting null results as excluding specific regions of candidate parameter space must rely on (strong) assumptions about the fine-scale structure of phase-space in the inner Galaxy. In most analyses the dark matter has been assumed to be smoothly and spherically distributed about the Galactic Centre with an isotropic Maxwellian velocity distribution \\citep[e.g.][]{1988PhRvD..37.3388F} or a multivariate Gaussian distribution \\citep[e.g.][]{2001JHEP...03..049U, 2001PhRvD..63d3005G, 2002PhRvD..66f3502H}.  The theoretical justification for these assumptions is weak, and when numerical simulations of halo formation reached sufficiently high resolution, it became clear that the phase-space of CDM halos contains considerable substructure, both gravitationally bound subhalos and unbound streams. As numerical resolution has improved it has become possible to see structure closer and closer to the centre, and this has led some investigators to suggest that the CDM distribution near the Sun could, in fact, be almost fractal, with large density variations over short length-scales \\citep[e.g.][]{2008PhRvD..77j3509K}. This would have substantial consequences for the ability of direct detection experiments to constrain particle properties.   Until very recently, simulation studies were unable to resolve any substructure in regions as close to the Galactic Centre as the Sun \\citep[see][for example]{2001PhRvD..64f3508M,2002PhRvD..66f3502H, 2003MNRAS.339..834H}. This prevented realistic evaluation of the likelihood that massive streams, clumps or holes in the dark matter distribution could affect event rates in Earth-bound detectors and so weaken the particle physics conclusions that can be drawn from null detections \\citep[see][for recent discussions]{2006PhRvD..74d3531S,2008PhRvD..77j3509K}.  As we shall show in this paper, a new age has dawned. As part of its Aquarius Project \\citep{2008arXiv0809.0898S} the Virgo Consortium has carried out a suite of ultra-high resolution simulations of a series of Milky Way-sized CDM halos. Simulations of individual Milky Way halos of similar scale have been carried out by \\cite{2008Natur.454..735D} and \\cite{2008arXiv0808.2981S}. Here we use the Aquarius simulations to provide the first reliable characterisations of the local dark matter phase-space distribution and of the detector signals which should be anticipated in WIMP and axion searches.  \\begin{figure} \\center{ \\includegraphics[width=0.4\\textwidth]{DPDF_A.eps} \\includegraphics[width=0.4\\textwidth]{DPDF_2.eps} } \\caption{Top panel: Density probability distribution function (DPDF) for all resimulations of halo Aq-A measured within a thick ellipsoidal shell between equidensity surfaces with major axes of $6$ and $12$~kpc. The local dark matter density at the position of each particle, estimated using an SPH smoothing technique, is divided by the density of the best-fit, ellipsoidally stratified, power-law model. The DPDF gives the distribution of the local density in units of that predicted by the smooth model at random points within the ellipsoidal shell. At these radii only resolution levels 1 and 2 are sufficient to follow substructure.  As a result, the characteristic power-law tail due to subhalos is not visible at lower resolution. The fluctuation distribution of the smooth component is dominated by noise in our 64-particle SPH density estimates. The density distribution measured for a {\\it uniform} (Poisson) particle distribution is indicated by the black dashed line.  Bottom panel: As above, but for all level-2 halos after rescaling to $V_{\\rm max}=208.49$~km/s. In all cases the core of the DPDF is dominated by measurement noise and the fraction of points in the power law tail due to subhalos is very small. The chance that the Sun lies within a subhalo is $\\sim 10^{-4}$. With high probability the local density is close to the mean value averaged over the Sun's ellipsoidal shell. } \\label{fig:DPDF} \\end{figure} ",
        "conclusions": "We have characterised the local phase-space distribution of dark matter using the recently published ultra-high resolution simulations of the Aquarius Project. Our study provides new insights relevant to searches for the elusive CDM particles. This results from the unprecedented resolution and convergence (in a dynamical sense) of our simulations, as well as from the fact that they provide a sample of six Milky Way-like dark matter halos.  We have measured the probability distribution function of the DM mass density between $6$ and $12$~kpc from the centre of the halo, finding it to be made up of two components: a truly smooth distribution which scatters around the mean on ellipsoidal shells by less than $5\\%$ in all the halos of our sample, and a high-density tail associated with subhalos. The smooth DM component dominates the local DM distribution.  With $99.9\\%$ confidence we can say that the Sun lies in a region where the density departs from the mean on ellipsoidal shells by less then $15\\%$.  Experimentalists can safely adopt smooth models to estimate the DM density near the Sun.  We find that the local velocity distribution is also expected to be very smooth, with no sign of massive streams or subhalo contributions.  The standard assumption of a Maxwellian velocity distribution is not correct for our halos, because the velocity distribution is clearly anisotropic. The velocity ellipsoid at each point aligns very well with the shape of the halo. A better fit to the simulations is given by a multivariate Gaussian. Even this description does not reproduce the exact shape of the distributions perfectly. In particular the modulus of the velocity vector shows marked deviations from such model predictions. Velocity distributions in our six different halos share common features with respect to the multivariate Gaussian model: the low-velocity region is more populated in the simulation; the peak of the simulation distribution is depressed compared to the Gaussian; at high velocities there is typically an excess in the simulation distribution compared to the best-fit multivariate Gaussian.  Furthermore the velocity distribution shows features which are stable in time, are reproduced from place to place within a given halo, but differ between different halos. These are related to the formation history of each individual halo.  The imprints in the modulus of the velocity vector reflect features in the energy distribution. We explicitly show that the phase-space distribution function as a function of energy contains characteristic wiggles. The amplitude of these wiggles with respect to the average distribution function of our sample of six halos rises from high to low-binding energies. After appropriate scaling, the most bound part of the distribution function looks very similar in all halos, suggesting a (nearly) universal shape.  The weakly bound part of the distribution, on the other hand, can deviate in any given halo by an order of magnitude from the mean.   We have used our simulations to predict detector signals for WIMP and axion searches. We find that WIMP recoil spectra can deviate about $10\\%$ from the recoil rate expected from the best-fit multivariate Gaussian model. The energy dependence of these deviations looks similar in all six halos; especially at higher binding energies.  We find that the annual modulation signal peaks around the same day as expected from a multivariate Gaussian model with no clear trend over our halo sample for varying recoil velocity thresholds.  The maximum recoil modulation amplitude, on the other hand, shows a clear threshold-dependent difference between the signal expected for a multivariate Gaussian model and that estimated from the simulation. We have also explored the expected signal for direct detection of axions. We find the axion spectra to be smooth without any sign of massive streams. The spectra show characteristic deviations from those predicted by a multivariate Gaussian model; the power at low and high frequencies is higher than expected. The most pronounced effect is that the spectra peak at lower frequencies than predicted.  Since the frequencies in the axion detector are directly proportional to the kinetic energy of the axion particles, the bumps in the DM velocity and energy distributions are also clearly visible in the axion spectra. All the effects on the various detector signals are driven by differences in the velocity distribution. Individual subhalos or streams do not influence the detector signals however, since they are sub-dominant by a large factor in all six halos.  Our study shows that, once direct dark matter detection has become routine, the characterisation of the DM energy distribution will provide unique insights into the assembly history of the Milky Way halo. In the next decade, a new field may emerge, that of ``dark matter astronomy''."
    },
    "0812/0812.0648_arXiv.txt": {
        "abstract": "The temperature in the optically thick interior of protoplanetary discs is essential for the interpretation of millimeter observations of the discs, for the vertical structure of the discs, for models of the disc evolution and the planet formation, and for the chemistry in the discs. Since large icy grains have a large albedo even in the infrared, the effect of scattering of the diffuse radiation in the discs on the interior temperature should be examined. We have performed a series of numerical radiation transfer simulations including isotropic scattering by grains with various typical sizes for the diffuse radiation as well as for the incident stellar radiation. We also have developed an analytic model including isotropic scattering to understand the physics concealed in the numerical results. With the analytic model, we have shown that the standard two-layer approach is valid only for grey opacity (i.e. grain size $\\ga10$ \\micron) even without scattering. A three-layer interpretation is required for grain size $\\la10$ \\micron. When the grain size is 0.1--10 \\micron, the numerical simulations show that isotropic scattering reduces the temperature of the disc interior. This reduction is nicely explained by the analytic three-layer model as a result of the energy loss by scatterings of the incident stellar radiation and of the warm diffuse radiation in the disc atmosphere. For grain size $\\ga10$ \\micron\\ (i.e. grey scattering), the numerical simulations show that isotropic scattering does not affect the interior temperature. This is nicely explained by the analytic two-layer model; the energy loss by scattering in the disc atmosphere is exactly offset by the ``green-house effect'' due to scattering of the cold diffuse radiation in the interior. ",
        "introduction": "Protoplanetary discs are planet formation sites. We observe the electro-magnetic radiation from the discs to understand their physical conditions, and then, to know the planet formation, especially, its beginning. Property of the radiation from a disc is essentially determined by the structure of the disc. The structure follows the response of the disc to the radiation from the central star. To interpret the radiation from protoplanetary discs, therefore, we should have a robust link between the radiation from the central star and the disc structure, in particular, the temperature structure.  The temperature structure of the discs is also essential for chemical reactions in the discs. The evolution of various molecules in the discs and the exchange of these molecules between the gas phase and the solid phase on grains will be discussed in detail with the ALMA in near future \\citep[e.g.,][]{nom08}. The condensation front of icy molecules, so-called 'snow line', is also determined by the temperature structure of the discs \\citep[e.g.,][]{sas00,oka09}. The location of the snow line is very important because it significantly enhances the amount of solid materials to make planetary cores and affects the supply of water to rocky planets. The temperature structure also affects the evolution of the discs themselves. The accretion activity in the discs is supposed to be driven by the magnetorotational instability \\citep[e.g.,][]{san00}. This requires a certain degree of the ionization of disc materials which depends on the structure of the discs.  A milestone in the research of the vertical temperature structure of protoplanetary discs is the work by \\cite{cg97} (hereafter CG97). They proposed a two-layer model consisting of the super-heated layer directly exposed by the stellar radiation and the interior warmed by the super-heated layer. This model explained the shape of the spectral energy distribution of the discs very well. The computational cheapness of the analytic approach makes the CG97 model very useful to compare with a large sample of the discs observed. Therefore, some attempts to refine the simple model of CG97 were performed \\citep{chi01,dul01,dul03a,raf06,gar07}.  Despite a great success of the CG97 model, numerical simulations of the radiation transfer in the discs showed a significant decrement of the equatorial temperature relative to the prediction by the CG97 model \\citep{dul03a}. The internal energy loss by the radiation at a long wavelength where the optical depth of the disc is relatively small is suggested as a cause of the decrement \\citep{dul02,dul03a}. In the CG97 model and refined ones, the wavelength dependence of dust opacity was taken into account in terms of mean opacity. \\cite{dul02} argued the importance of using full wavelength dependent opacity. However, we propose an alternative approach in this paper: a three-layer model with mean opacity which reproduces the temperature reduction quite well. In addition, we show that the two-layer approximation in the CG97 model is valid only when the dust opacity is ``grey'' which is expected if the size of dust grains is larger than about 10 \\micron.  Effect of scattering on the vertical temperature structure was not considered in the literature very much. Scattering of the stellar radiation was taken into account analytically by \\cite{cal91} and numerically by \\cite{dul03} who showed that the temperature of the disc interior is slightly reduced by the scattering. How about scattering of the diffuse disc radiation? As the size of grains increases, the scattering albedo increases. In particular, the albedo of large icy grains is close to unity even for the infrared wavelength. This may affect the disc structure significantly. Nevertheless, it has not been examined so far.  This paper discusses the effect of scattering of the diffuse radiation as well as that of the stellar radiation on the vertical temperature structure of protoplanetary discs. Since the albedo depends on the grain size, we examine the scattering effect as a function of the grain size. Although there are 2-D/3-D numerical radiation transfer codes available publicly, most of them have a serious difficulty in solving the radiation equilibrium in very high optical depth ($\\tau\\sim10^6$) found in protoplanetary discs \\citep{pas04,ste06}. The RADICAL developed by \\cite{dul00} can solve such a problem without any difficulty thanks to a variable Eddington tensor method. However, it treats scattering of only the stellar radiation. We present a variable Eddington factor code with both scatterings of the stellar radiation and of the diffuse radiation but in a 1-D geometry. We also present an analytic model to interpret the numerical results. This simple model would be very useful to understand the physics determining the temperature structure of protoplanetary discs.  The rest of this paper consists of three sections; in section 2, we develop a numerical radiation transfer code taking into account both of the scatterings but only for isotropic case and show the obtained numerical solutions. In section 3, we construct an analytic model to interpret the numerical solutions and discuss the physical mechanism determining the temperature structure in protoplanetary discs. In the final section, we summarise our findings. ",
        "conclusions": "We have examined the effect of scattering of the diffuse radiation on the vertical temperature structure of protoplanetary discs. This is motivated by the fact that scattering albedo increases as the size of dust grains grows in the discs. In particular, large icy grains have a significant albedo even in the infrared wavelength. For this aim, we have developed a 1D plane-parallel numerical radiation transfer code including isotropic scattering of the diffuse radiation as well as that of the incident radiation. We have also developed an analytic model with isotropic scattering both of the diffuse and the incident radiations in order to interpret the solutions obtained from the numerical simulation. All results of the numerical simulations has been nicely reproduced by the analytic model.  The analytic model presented in this paper is an extension of the seminal two-layer model by \\cite{cg97}; we have introduced a new layer between the super-heated surface layer and the disc interior. This middle layer (or disc atmosphere) is required when the absorption of the radiation from the super-heated layer occurs well above the photosphere of the opaque isothermal interior. This situation is realised if the dust opacity is negatively proportional to the wavelength. Thus, we should consider three layers rather than two layers if the grain size is smaller than about 10 \\micron. On the other hand, for grey opacity, which is realised if the grain size is $\\ga10$ \\micron, the standard treatment with two layers is justified.  We have found from the numerical simulation that the dust temperature of the disc surface is almost not affected by scattering. This is because the temperature is determined by the radiation equilibrium of grains in the incident radiation field locally. The grain size has an effect on the surface temperature via the wavelength dependence of the dust opacity. There are two asymptotic temperatures: the higher temperature is realised by small grains ($\\la0.1$ \\micron) and the lower temperature is realised by large grains ($\\ga1$--10 \\micron). This is because the emission efficiency relative to absorption efficiency of the small grains is smaller than that of the large grains which have grey opacity. This trend of the numerical solutions is excellently reproduced by the analytic model of the super-heated layer.  The numerical simulations without scattering show that the dust temperature of the optically thick interior also has two asymptotic values: the lower temperature for small grains ($\\la0.1$ \\micron) and the higher temperature for large grains ($\\ga10$--100 \\micron). Thus, the trend is opposite from the surface temperature. The higher asymptotic temperature for large grains is well matched with the prediction of the two-layer model. On the other hand, the lower asymptotic temperature for small grains is a factor of $(1/2)^{1/4}$ lower than the prediction of the two-layer model. In fact, the flux input from the super-heated layer is always same although the interior temperature is different as a function of the grain size. This phenomenon has been already found by \\cite{dul03a} who attributed it to the energy loss from the interior by the radiation at a long wavelength \\citep[see also][]{dul02}. We have proposed another interpretation by the middle layer between the super-heated layer and the interior; the super-heated layer gives a half of the absorbed energy of the incident radiation to the middle layer which gives a half of the obtained energy to the interior. This three-layer model exactly predicts a factor of $(1/2)^{1/4}$ reduction of the interior temperature for the small grain case.  The numerical simulations with isotropic scattering also show two asymptotic temperatures of the interior. Interestingly, these asymptotic temperatures of no scattering and isotropic scattering cases are the same. For small grains ($\\la0.1$ \\micron), since the scattering albedo is negligible for all wavelengths interest, the same temperature is trivial. The same temperature for large grains ($\\ga10$--100 \\micron) has been nicely explained by the two-layer model with grey opacity. The physical mechanism is the exact offset between the reduction of the flux input into the interior by scattering in the super-heated layer and the reduction of the flux output from the interior by scattering in itself (i.e. green-house effect).  For grain sizes of 0.1--10 \\micron, which are expected by a moderate growth of the size in the discs, we have found a further reduction of the interior temperature in isotropic scattering cases relative to that without scattering from the numerical simulations. This reduction has been well explained by the three-layer analytic model. The physical mechanism is the wavelength dependence of albedo; the flux input into the interior is reduced by scattering in the super-heated and middle layers, whereas the flux output from the interior is not reduced because of negligible (or weak) scattering. Note that the interior flux has a longer wavelength typically, thus, the albedo for the radiation is smaller.  In conclusion, the scattering of the diffuse radiation can affect the vertical temperature structure of protoplanetary discs significantly when the grain size grows to be about 1--10 \\micron. We need to investigate the effect in a global disc model in future. The analytic model presented in this paper could be useful to understand the physics determining the temperature structure in the discs."
    },
    "0812/0812.2684_arXiv.txt": {
        "abstract": "{We report the discovery of the first variable extreme horizontal branch star in a globular cluster ($\\omega$ Cen). The oscillation uncovered has a period of 114 s and an amplitude of 32 mmags. A comparison between horizontal branch models and observed optical colours indicates an effective temperature of 31,500$\\pm$6,300 K for this star, placing it within the instability strip for rapidly oscillating B subdwarfs.  The time scale and amplitude of the pulsation detected are also in line with what is expected for this type of variable, thus strengthening the case for the discovery of a new subdwarf B pulsator. } ",
        "introduction": "Subdwarf B (sdB) stars are hot (20,000 K $\\lesssim T_{\\rm eff} \\lesssim$ 40,000 K), compact (5.2 $\\lesssim \\log{g} \\lesssim$ 6.2), evolved objects found on the extreme horizontal branch (EHB), a high temperature extension of the horizontal branch \\citep{heber2008}. Over the last decade, a small subset ($\\sim$ 5\\%) of them have been found to exhibit rapid luminosity variations with typical periods of 100$-$200 s and amplitudes on the mmag scale (see e.g. \\citealt{fontaine2006} for a recent review). Also known as EC 14026 stars after the prototype \\citep{kilkenny1997}, the rapid pulsators are found among the hotter sdB stars in a well-defined instability strip from 29,000$-$36,000 K. Pulsation driving is believed to occur through the action of a classical $\\kappa$-mechanism associated with a local overabundance of iron peak elements in the driving region, which in turn relies on radiative levitation processes \\citep{charp1996,charp1997}. The modelling of these stars has been very successful, to the point where their non-adiabatic pulsation properties are understood not only qualitatively, but also quantitatively. Indeed, full asteroseismological analyses leading to a precise determination of the targets' fundamental parameters have now been carried out for 12 out of at least 35 known EC 14026 stars (see e.g. \\citealt{val2008, charp2008} for recent results).  It is generally accepted that sdB stars are the progeny of red giant branch stars that underwent significant mass loss around the time of the core helium flash, leaving them with very thin hydrogen-rich envelopes. These stellar structures do not experience the asymptotic giant branch phase, instead evolving as so-called {\\em AGB-Manqu\\'e} stars \\citep{greggio1990}, however we still lack firm constraints on the evolutionary channel(s) followed. A number of different scenarios involving e.g. a common envelope phase, stable and unstable Roche lobe overflow, or the merger of two He white dwarfs have been proposed \\citep{han2002,han2003}, but remain to be tested observationally. More recently, it has been suggested \\citep{han2008} that the low fraction of close binaries identified among cluster EHB stars ($\\approx$ 4\\% from \\citealt{monibidin2008}) compared to the field population ($\\approx$ 60\\% according to \\citealt{maxted2001}; $\\approx$ 40\\% from \\citealt{napi2006}) can be explained naturally in terms of an age effect. In the case of an older population, the white dwarf merger scenario forming a single sdB star becomes more important than common envelope evolution producing an sdB star in a close binary. It has also been suggested  that cluster EHB stars might be the progeny of the ``hot helium-flasher\" scenario (\\citealt{castellani1993, dcruz1996,castellani2006, miller2008}). In this scenario, red giant stars approaching the tip of the RGB lose a significant fraction of their envelope due to a violent mass-loss event. The ensuing red giants have a higher mass than the limit for core helium ignition, and undergo a He-core flash either when approaching, or along the white dwarf cooling sequence \\citep{brown2001}. Another working hypothesis is that cluster EHB stars are the progeny of the helium-enhanced sub-populations \\citep{dantona2002,lee2005,dantona2007} recently detected in a few globular clusters \\citep{bedin2004,piotto2008}.  This is where asteroseismology becomes important: since the different evolutionary scenarios leave an imprint on the fundamental parameters of the resulting sdB population, contrasting the predictions from population synthesis models with the corresponding asteroseismic parameter distributions can be used to distinguish between the proposed formation channels. First attempts in this direction have been made \\citep{fontaine2008}, however the number of EC 14026 stars analysed up to now is too small for quantitative conclusions.  Moreover, the sample of known sdB pulsators has so far been confined to the Galactic field population. While several attempts have been made to detect variable EHB stars in globular clusters (e.g. \\citealt{catelan2008,kaluzny2008,reed2006}), no convincing EC 14026 star candidates were uncovered.  Rapidly pulsating subdwarf B stars in globular clusters are of high interest in the context of an asteroseismic characterisation of the EHB population in these systems. The first necessary step in this direction is obviously the detection of pulsating subdwarf B stars in globular clusters. This is why, when given a couple of free hours at the end of a night of observations dedicated to a different programme, we decided to monitor a field in $\\omega$ Cen for rapid variability among the EHB population. We present the outcome of this mini-survey in what follows. ",
        "conclusions": "On the basis of 2 hours of SUSI2 time-series data, as well as the availability of WFI and ACS multi-colour photometry, we discovered what appears to be the first rapidly pulsating subdwarf B star in a globular cluster. Both the pulsation properties and the effective temperature estimate are in line with what is observed for this type of variable in the field population; however, a longer time series of photometry and high S/N spectroscopy are needed to verify the results presented here. It will be particularly interesting to see whether more frequencies can be uncovered, since all known sdB pulsators show multi-periodic luminosity variations. However, given the magnitude of the pulsator discovered and the shortness of the time series obtained, it is not surprising that only one periodicity was detected. Indeed, the period spectra of fast sdB pulsators are quite often dominated by one high-amplitude mode, and individual frequencies can be too closely spaced to be distinguishable in just one two-hour run. The amplitude of the oscillation measured is among the highest found for any sdB pulsator (the highest amplitude pulsation is around 50  mmags for Balloon 09010001, see \\citealt{oreiro2004}; typical amplitudes are on the 1-5 mmag scale), and lower amplitude luminosity variations will be challenging to uncover given the faintness of the star and the crowded field. Nevertheless, we hope that the results presented here will constitute a first step towards the asteroseismology of globular cluster subdwarf B stars, and create the unique opportunity to provide quantitative constraints on their evolutionary history."
    },
    "0812/0812.2960_arXiv.txt": {
        "abstract": "The Vela pulsar is the brightest persistent source in the GeV sky and thus is the traditional first target for new $\\gamma$-ray observatories. We report here on initial {\\it Fermi} Large Area Telescope observations during verification phase pointed exposure and early sky survey scanning. We have used the Vela signal to verify {\\it Fermi} timing and angular resolution. The high quality pulse profile, with some 32,400 pulsed photons at $E\\ge 0.03$\\,GeV, shows new features, including pulse structure as fine as 0.3\\,ms and a distinct third peak, which shifts in phase with energy. We examine the high energy behavior of the pulsed emission; initial spectra suggest a phase-averaged power law index of $\\Gamma=1.51^{+0.05}_{-0.04}$ with an exponential cut-off at $E_c=2.9\\pm 0.1$\\,GeV. Spectral fits with generalized cut-offs of the form $e^{-(E/E_c)^b}$ require $b\\le 1$, which is inconsistent with magnetic pair attenuation, and thus favor outer magnetosphere emission models. Finally, we report on upper limits to any unpulsed component, as might be associated with a surrounding synchrotron wind nebula (PWN). ",
        "introduction": "Radio pulsations at $P$=89\\,ms from PSR B0833$-$45 (=PSR J0835$-$4510) in the Vela supernova remnant were discovered by Large, Vaughan \\& Mills (1968). This pulsar is bright ($S_{1.4}\\,{\\rm GHz} \\approx 1.5$\\,Jy), young (characteristic age $\\tau_c = P/2{\\dot P} = 11$\\,kyr), and energetic (${\\dot E}_{SD}=6.9\\times 10^{36}I_{45}\\, {\\rm ergs/s}$, for a neutron star moment of inertia $10^{45}I_{45}\\,{\\rm g\\,cm^2}$). It is embedded in a flat spectrum radio synchrotron nebula (see Frail et al. 1997) and is surrounded by a bright X-ray wind nebula displaying remarkable toroidal symmetry (Helfand et al. 2001; Pavlov et al. 2003). A VLBI parallax measurement provides a well-determined distance of $D=287_{-17}^{+19}$pc (Dodson et al. 2003), improving on earlier optical HST measurements (Caraveo et al. 2001). This proximity means that the spindown energy flux at Earth, ${\\dot E}_{SD}/4\\pi d^2$, is second only to that of the Crab pulsar and ensures that Vela is among the most intensely studied neutron stars, particularly at high energies.  Pulsed $\\gamma$-ray emission from Vela was first detected during the {\\it SAS-2} mission (Thompson et al. 1975); it is, in fact, the brightest persistent source of celestial $\\gamma$-rays and has been the proving ground of GeV observatories ever since.  The basic source properties discovered by {\\it SAS-2} were elaborated using observations by {\\it COS-B} (Kanbach et al. 1980), and {\\it CGRO}/EGRET (Kanbach et al. 1994): the source is $\\sim$100\\% pulsed with no convincing evidence for year-scale variations, the two $\\gamma$-ray peaks are separated by 0.42 in pulsar phase and the first peak lags the radio peak by 0.12 in phase. The source spectrum is hard with an average photon index $\\Gamma= 1.7$ with evidence of a cut-off above 2--4\\,GeV. The most detailed study, to date, of the $\\gamma$-ray pulsations is that of Fierro et al. (1998), which produced a high signal-to-noise pulse profile and evidence for phase-resolved spectral variations. As the present work was being prepared for submission, the AGILE team reported their initial results on $\\gamma$-ray pulsars, including Vela (Pellizoni et al. 2008); the results, given the more limited count rates and energy resolution, are broadly consistent with our conclusions described below.  With the successful launch of the {\\it Fermi} Gamma-ray Space Telescope, formerly GLAST, observatory on June 11, 2008, we have a new opportunity to examine the high energy behavior of the Vela pulsar and to study this archetype of the young, energetic pulsars in detail. During Launch and Early Operations phase (L\\&EO) the {\\it Fermi} mission targeted the Vela pulsar for a number of pointed observations, in addition to coverage during initial tests of the sky survey mode. In the latter mode the instrument axis is offset North and South of the zenith during alternate orbits to provide near-uniform sky coverage every three hours.  One main purpose of these early observations was to tune the Large Area Telescope (LAT) performance on celestial $\\gamma$-ray sources. However, the initial results on Vela itself are of interest, including new high energy features in the pulse profile, an improved measurement of the high energy cut-off in the pulsar spectrum, and a search for associated pulsar wind nebula (PWN) emission at GeV energies.  \\placetable{1} \\begin{table*} \\centering \\caption{\\label{1} Early LAT observations of PSR B0833$-$45} \\begin{tabular}{lcccc} \\hline\\hline Date       & MJD &Primary mode& $N$(E$>$0.03GeV)$^a$ & Notes\\\\ \\hline 2008 Jun 30--Jul 4&54647.40--54651.36& Sky survey &1859&LAT First light.\\\\ 2008 Jul 4--Jul 15&54651.38--54662.10& Survey+Pointings &5172&LAT Calibrations.\\\\ 2008 Jul 15--Jul 19&54662.12--54666.08& Pointings+Survey+limb Following &1170&Pointed Obs. Tuning\\\\ 2008 Jul 19--Jul 22*$^b$ &54666.10--54669.14& Pointed Vela+ 2$^{nd}$ Target &7751&Pointed Obs. Tuning\\\\ 2008 Jul 22--Jul 24&54669.16--54671.41& Pointed Vela+ 2$^{nd}$ Target &3212&Pointed Obs. Tuning\\\\ 2008 Jul 24--Jul 30*&54671.36--54677.45& Pointed Vela+ 2$^{nd}$ Target &7607&Pointed Obs. Tuning\\\\ 2008 Jul 30--Aug 3&54677.45--54681.66& Polar study + Sky Survey Tuning&886&Nominal Ops.\\\\ 2008 Aug 3--Sep 14*&54687.68--54723.91&Sky Survey&26680&Nominal Ops.$^c$ \\\\ \\hline \\multicolumn{5}{l}{$^a$All photons (pulsed and background) within 5$^\\circ$ of PSR B0833-45, zenith angle $z<$105$^\\circ$, event class \\emph{Pass6-Diffuse}.}\\\\ \\multicolumn{5}{l}{$^b$Observations marked by asterisk are those used for spectral studies.}\\\\ \\multicolumn{5}{l}{$^c$L$\\&$EO ended on Aug.11.} \\end{tabular} \\end{table*} \\begin{figure*} \\includegraphics[scale=0.90]{f1.eps}% \\caption{\\label{fig:Lightcurves} Vela broad band ($E=0.1-10$\\,GeV) pulse profile for all photons from an energy dependent ROI. Two pulse periods are shown. The peak of the radio pulse is at phase $\\phi=0$. The count rate is shown in variable width phase bins with a constant number of counts. The dashed line shows the background level, as estimated from a surrounding annulus during the off-pulse phase. Insets show the pulse shape near the peaks and in the off-pulse region.} \\end{figure*} ",
        "conclusions": "\\subsection{Light Curve}  The early LAT data have already greatly clarified the $\\gamma$-ray emission of the Vela pulsar and, even before detailed phase-resolved spectral measurements and model comparisons, significantly constrain the origin of high energy pulsar emission. Pulsar particle acceleration is believed to occur in the open zone, on the field lines above each magnetic pole which extend through the light cylinder. These subtend $\\sim 0.5$ of pulsar phase at high altitude, so with a total phase in the narrow pulses of $\\sim 0.06$ one would expect a comparable fraction of the open zone, and hence fraction of the spin-down power, to contribute to the $\\gamma$-ray emission.  The narrow peaks are often interpreted as caustics (Morini 1983; Romani \\& Yadigaroglu 1995, Dyks et al.\\ 2004), where the combination of field retardation, aberration and time delay causes emission from the boundary of an acceleration zone (gap) to pile up in pulsar phase. A simple pulse from a 1-D fold caustic would increase to a maximum as $I(\\phi) \\propto (\\phi-\\phi_0)^{-1/2}\\theta[\\phi-\\phi_0]$ as one approaches the destruction of an image pair at phase $\\phi_0$ or fall off in mirror form after image pair creation. Vela's P2 structure with a slow rise and abrupt fall looks like a trailing (image pair destruction) caustic, albeit with a steeper increase to the maximum, suggesting intensity variation along the emission surface or a higher-order catastrophe. The shape of P1 is less clear, but the apparent fast rise and slow fall would indicate a leading (image pair creation) caustic (Figure 1).  The broad structure of the P3 interpulse region suggests a beam of emission from a bundle of field lines bracketed by the two caustic-forming zones. The similarity between the $\\gamma$-ray P3 phase structure and that of the lower frequency optical/UV emission at $\\phi \\approx 0.25$ suggests a common origin -- they both cover similar phase intervals and the peak shifts to later phase at higher energy.  One appealing possibility is to associate the optical/UV P3 with synchrotron emission from relatively low altitude pair cascades produced by the few-GeV curvature radiation dominating the $\\nu F_\\nu$ spectral energy distribution (SED). These $\\gamma\\gamma\\longrightarrow$\\,e$^\\pm$ pairs, with electron energy $\\gamma_e \\sim 2\\,{\\rm GeV}/m_e c^2 \\sim 10^{3.5}$, will produce synchrotron emission with characteristic peak energy $\\epsilon \\sim 12 B_{12}\\gamma_e{\\rm sin}\\Psi_{-1}$\\,keV as they radiate away their pitch angle $\\sim 0.1\\Psi_{-1}$ radians in a magnetic field of $10^{12}$\\,G. For Vela, we have $B_{12} \\sim 3 (r/R_\\ast)^{-3}$, so to produce $\\epsilon \\sim 5$\\, eV photons the pairs would reside at $r\\sim 80 \\Psi_{-1}^{-1/3}R_\\ast$, where $R_\\ast$ is the neutron star radius. If the synchrotron emission is outward beamed, along with the pairs, then the Compton scattered flux would appear at $\\sim \\Psi \\gamma_e^2 \\epsilon \\sim 5$\\,GeV, matching the LAT P3 component. This suggests that the pulse of seed photons at IR to UV energies should have energy-dependent phase shifts tracking their Compton-up-scattered P3 progeny.  \\subsection{Luminosity and Beaming}  While the VLBI parallax provides a well-determined distance of $D=287^{+19}_{-17}$\\,pc, the pattern of the $\\gamma$-ray beam and hence the true $\\gamma$-ray luminosity are uncertain. To compute the true luminosity \\begin{equation} L_{\\gamma} = 4\\pi f_\\Omega F_{\\rm E,obs}  D^2 \\end{equation} we need a correction factor $f_\\Omega$ from the observed phase-averaged energy flux F$_{\\rm E,obs}$ at the Earth line of sight (at angle $\\zeta_E$ to the rotation axis) to an average over the sky. For a given model, the correction depends on the geometry, so we have \\begin{equation} f_\\Omega (\\alpha,\\zeta_E)=\\int F_\\gamma(\\alpha,\\zeta,\\phi){\\rm d}\\Omega/(F_{exp}4\\pi), \\end{equation} where the model for a given magnetic inclination $\\alpha$ from the rotation axis gives $F_\\gamma(\\alpha,\\zeta,\\phi)$, the pulse profile as a function of phase $\\phi$ seen at viewing angle $\\zeta$. $F_{exp}$ is the expected phase-averaged flux for the actual Earth line-of-sight $\\zeta_E$.  In general highly-aligned models, such as the polar cap scenario, have small sky coverage and $f_\\Omega<0.1$, while outer-magnetosphere models have much larger $f_\\Omega$ (Romani \\& Yadigaroglu 1995). For Vela, the viewing angle is estimated to be $\\zeta \\approx 63^\\circ$ from the geometry of the {\\em Chandra X-ray Observatory}-measured X-ray torus (Ng \\& Romani 2008). The magnetic impact angle $\\beta=\\zeta-\\alpha$ is constrained by the radio polarization data. Together, these allow us to compute the high energy beam shape (assuming uniform emissivity along gap surfaces). For the outer gap (OG) picture we compute $f_\\Omega=1.0$ for $\\zeta=64^\\circ$. For the two-pole caustic (TPC) model appropriate to the slot gap picture, we find $f_\\Omega = 1.1$ (Watters et al.\\ 2008).  The observed energy flux obtained by integrating $E_\\gamma$ times the photon number flux in Equation (1) over the range 0.1$-$10\\,GeV is (7.87$\\pm0.33\\pm1.57$)$\\times$10$^{-9}$\\,ergs/cm$^2$/s. For a distance of 0.287 kpc, the $\\gamma$-ray luminosity of Vela is then 7.8$\\times$ 10$^{34}f_\\Omega$\\,ergs/s. Thus the observed $\\gamma$-ray flux implies an efficiency of $\\eta_\\gamma = 0.011 f_\\Omega/I_{45}$ with $f_\\Omega \\ge 1$ for both the OG and TPC models. Note that this is a substantial fraction of the geometrical bound on the efficiency from the peak widths noted above.  \\subsection{Gamma-Ray Pulsar Models}  Three general classes of models have been discussed for $\\gamma$-ray pulsars. In polar cap models (Daugherty \\& Harding 1994, 1996; Sturner \\& Dermer 1995; Harding \\& Muslimov 2004), the particle acceleration and $\\gamma$-ray production takes place in the open field line region within one stellar radius of the magnetic pole of the neutron star.  In outer gap models (Cheng, Ho \\& Ruderman 1986; Romani 1996; Hirotani 2005), the interaction region lies in the outer magnetosphere in vacuum gaps associated with the last open field line. Other recent models hark back to the `slot-gap' picture of Arons (1983), with the polar cap rim acceleration extending to many stellar radii (Harding \\& Muslimov 2004).  The appreciable offset from the radio peak, coupled with the presence of only one radio pulse, suggests that the $\\gamma$-rays arise at high altitude. This is also supported by the lack of hyper-exponential absorption in the spectrum, which is a signature of $\\gamma$-B pair attenuation of low altitude emission. The remaining two geometrical models can produce the general double peak profile of Vela (and many other $\\gamma$-ray pulsars). In `Outer gap' models emission starts near the `null charge' $\\Omega \\cdot B=0$ surface at radius $r_{\\rm NC}$ and extends toward the light cylinder $r_{\\rm LC} = cP/2\\pi$. One magnetic pole dominates the emission in each hemisphere and the two peaks represent leading and trailing edges of the hollow cone of emission from this pole. If on the other hand emission extends well below the null charge surface toward the neutron star, both magnetic poles can contribute toward emission in a given hemisphere. This is the two-pole caustic model of Dyks \\& Rudak (2003) which might be realized in `slot gap' acceleration models. In this case the leading `image pair creation' caustic from high altitudes should not be visible, and the first $\\gamma$-ray pulse represents the trailing caustic from emission at $r<r_{NC}$ below the null-charge surface (the radio magnetic pole) while the second pulse represents a trailing caustic at higher altitudes from the opposite pole; both pulses should generally show a slow rise and steeper fall. Detailed $\\gamma$-ray pulse profiles and, especially, phase resolved spectra can help distinguish these scenarios.  Another geometric difference in the two scenarios is the origin of the emission outside the main pulse. Low altitude field lines in the two-pole model tend to produce emission along all lines of sight. In the outer gap picture high altitude field lines can contribute faint emission extending through the off pulse region.  Further, as noted above, the energy spectrum also presents challenges for near-surface emission.  In the polar cap models, a sharp turnover is expected in the few to 10\\,GeV energy range due to attenuation of the $\\gamma$-ray flux in the magnetic field (Daugherty \\& Harding 1996). The spectral change at $\\sim$2.9\\,GeV does not appear to fit this model. We conclude that low altitude radiation subject to $\\gamma$-B pair production cannot account for the bulk of the Vela $\\gamma$-ray emission. Indeed we can use the observed cut-off to estimate a minimum emission height as $r\\approx (\\epsilon_{max} B_{12}/1.76{\\rm GeV})^{2/7}P^{-1/7} R_\\ast$, where $\\epsilon_{max}$ is the unabsorbed photon energy, $P$ is the spin period and the surface field is $10^{12}B_{12}$\\,G (Baring 2004). Using Vela parameters ($P=0.089$\\,s, $B_{12}=3.4$) we see that our cut-off implies that the bulk of the emission arises from $>2.2R_\\ast$. Since we see pulsed photons up to $\\epsilon_{max} =17$\\,GeV, this emission must arise at $r>3.8R_\\ast$. A similar conclusion has been recently made for the Crab pulsar by the MAGIC team, who observe pulsed $E>25$\\,GeV emission, using the imaging air Cerenkov technique (Aliu et al.\\ 2008).    % Finally, it should be noted that while typical radio pulsar emission is inferred to arise at relatively low altitudes, for Vela-type pulsars recent radio models infer emission heights of as much as $100R_\\ast$ (Karastergiou \\& Johnston 2007). All of these factors point to a high-altitude origin of the Vela $\\gamma$-ray emission. \\bigskip \\bigskip \\bigskip \\bigskip"
    },
    "0812/0812.2198_arXiv.txt": {
        "abstract": "We present a model that describes stellar infrared excesses due to heating of the interstellar (IS) dust by a hot star passing through a diffuse IS cloud. This model is applied to six $\\lambda$ Bootis stars with infrared excesses. Plausible values for the IS medium (ISM) density and relative velocity between the cloud and the star yield fits to the excess emission. This result is consistent with the diffusion/accretion hypothesis that $\\lambda$ Bootis stars (A- to F-type stars with large underabundances of Fe-peak elements) owe their characteristics to interactions with the ISM. This proposal invokes radiation pressure from the star to repel the IS dust and excavate a paraboloidal dust cavity in the IS cloud, while the metal-poor gas is accreted onto the stellar photosphere. However, the measurements of the infrared excesses can also be fit by planetary debris disk models. A more detailed consideration of the conditions to produce $\\lambda$ Bootis characteristics indicates that the majority of infrared-excess stars within the Local Bubble probably have debris disks. Nevertheless, more distant stars may often have excesses due to heating of interstellar material such as in our model. ",
        "introduction": "\\label{sec:introduction}  The vast majority of stars follow simple trends in mass, luminosity, age, metallicity, and other parameters as captured in the simplest form through the Russel-Vogt Theorem. The exceptions to these trends therefore must represent interesting events in stellar evolution. $\\lambda$ Bootis stars are an outstanding example. They are late B- to early F-type Population I stars with peculiar surface abundances: while light elements like C, N, O and S are roughly solar, heavy elements (Fe-peak) are depleted up to a factor of 100 \\citep{Paunzen02}. The ages of $\\lambda$ Bootis stars span the entire main sequence lifetimes for their types, with a distribution that has a peak at a rather evolved stage ($\\sim$ 1 Gyr). Less than 2\\% of all objects in the relevant spectral domain are believed to belong to the $\\lambda$ Bootis group, suggesting either a very short time scale (10$^6$ years) for the phenomenon, or uncommon conditions for its occurrence.  Understanding $\\lambda$ Bootis stars will provide important insights to more general problems. Due to their extremely shallow convection zones, A-type stars in general and chemically peculiar ones in particular provide excellent natural laboratories for the study of fundamental astrophysical processes such as diffusion, meridional circulation, stellar winds, and accretion of circumstellar and interstellar material.  \\cite{Michaud86} suggested that the peculiar chemical abundances on the surfaces of $\\lambda$ Bootis stars are due to accretion of circumstellar material that is mixed in the shallow convection zone of the star by the joint action of gravitational settling and radiative acceleration. \\cite{Charbonneau93}, \\cite{Venn90}, and \\cite{Turcotte93} developed this hypothesis further. It naturally explains why the anomalous abundance pattern is similar to that found in the gas-phase of the interstellar medium (ISM), where metallic elements like iron and silicon have condensed into dust grains. \\cite{Andrievsky00} suggested a specific mechanism: that $\\lambda$ Bootis behavior results from condensation of dust grains in the circumstellar environment, their removal through radiation pressure and the accretion of the remaining metal-poor gas. However, their diffusion/accretion model has many free parameters and fails to explain some observations such as the small proportion of stars with $\\lambda$ Bootis characteristics. Alternative theories invoke contact binary systems merging into a single star \\citep{Andrievsky97} and unresolved spectroscopic binaries \\citep{Faraggiana99}.  More recently, \\cite{Kamp02a} have suggested that the $\\lambda$ Bootis phenomenon results from interaction along the lines discussed by \\cite{Andrievsky00} but between the star and the local ISM (Figure \\ref{interaction}). This new hypothesis removes the ad-hoc assumptions and some of the degeneracy of previous models. Different levels of underabundance are produced by different amounts of accreted material relative to the photospheric mass. The small fraction of $\\lambda$ Bootis stars is explained by the low probability of a star-cloud interaction and by the effects of meridional circulation, which washes out any accretion pattern a few million years after the accretion has stopped \\citep{Kamp02a}. The behavior is suppressed in low temperature stars because of their more massive, difficult-to-contaminate convection zones. Strong stellar winds prevent the accretion of material for very hot stars.  \\begin{figure}[h] \\epsscale{0.85}  \\plotone{f1.eps}  \\caption{\\label{interaction}Geometry of the star-cloud interaction. Two different avoidance radii, corresponding to two different grain sizes are shown.} \\end{figure}  In addition to their peculiar abundance pattern, some $\\lambda$ Bootis stars show infrared excesses, suggestive of the presence of dust. Since any ongoing interaction with the ISM will heat interstellar dust, these cases can probe the model proposed by \\cite{Kamp02a}. Using IRAS photometry, \\cite{King94} found that 5 out of a sample of 56 $\\lambda$ Bootis stars have infrared excess in at least one of the satellite's bands. More recently, \\cite{Paunzen03} have compiled data from the ISO and IRAS satellites to identify $\\lambda$ Bootis stars with excess infrared emission. According to their work, 23\\% of 26 well established members of this group show infrared excesses with derived fractional dust luminosities (L$_{\\mathrm{IR}}$/L$_*$) ranging from $2.2 \\times 10^{-5}$ to $4.3 \\times 10^{-4}$. Both the incidence of excesses and the fractional luminosities are similar to the characteristics of planetary debris disks for stars of similar spectral type \\citep{Rieke05, Su06}.  \\cite{Kamp02b} used a sample of 21 A-type stars including debris-disks, dust-free stars and six $\\lambda$ Bootis stars, and found no fundamental difference between the atmospheric compositions of the stars with and without debris disks. This result suggests that circumstellar material is {\\it unlikely} to be linked to the anomalous $\\lambda$ Bootis abundances. Thus, the IR excesses of the $\\lambda$ Bootis-type stars might be due to a different mechanism, e.g., accretion of interstellar material distributed in a diffuse cloud. This idea is also supported by coronagraphic observations by \\cite{Kalas02} that showed some stars previously catalogued as debris disk stars are surrounded by an extended reflection nebulosity.  In this paper we develop a thermal emission model to explore the \\cite{Kamp02a} ISM interaction hypothesis. We emphasize the basic physics behind the model rather than detailed simulations of individual stars. Our most important conclusion is that simple infrared excess SEDs can often be fitted equally well by either a debris disk or interstellar interaction model. As a result, further probing of ISM-related excesses requires detailed modelling of stars where other lines of evidence positively identify such interactions \\citep[e.g., for $\\delta$ Velorum, ][]{Gaspar08}.  In \\S 2 we provide a description of the observations used in this paper and discuss the physical parameters of our sample of stars. We introduce our thermal emission model in \\S 3, where we describe the basic physical parameters involved in the calculation of the spectral energy distributions. \\S 4 explores the effect of varying different model parameters on the output SEDs. We discuss the implications of our results for the identification of debris disks and the nature of $\\lambda$ Bootis stars in \\S 5 and conclude the paper in \\S 6.  \\medskip \\medskip \\medskip \\medskip ",
        "conclusions": "We have constructed a geometrical and thermal emission model for the interaction of hot stars with diffuse interstellar clouds. We find that the predicted flux emitted by dust in the clouds is consistent with {\\it Spitzer} photometry of $\\lambda$ Bootis stars. However, we also find that simple SED fitting is not enough to distinguish between infrared excesses generated in this manner and those due to planetary debris disks. Within the Local Bubble, we estimate that the contamination of debris disk samples will be low, but further away it may be significant.  Our model provides order of magnitude estimates for the hypothetical cloud parameters. Small ($\\sim$ 0.02 pc) and diffuse ($\\sim$ 0.5 - 5 cm$^{-3}$) clouds with masses of about 10\\% of a lunar mass in solid material are sufficient to produce the observed infrared excess around these stars. It may be possible to identify examples of such interactions by observing bow shocks like the one found by \\cite{Gaspar08} for $\\delta$ Velorum. For these cloud parameters, we estimate accretion rates of gas onto the stellar surface between 10$^{-16}$ and 10$^{-14}$ M$_{\\odot}$ yr$^{-1}$, probably too low to produce $\\lambda$ Bootis characteristics. However, our results suggest that ISM interactions under certain favorable conditions (low relative velocity between cloud and star, higher than average cloud density, larger clouds) may yield both the infrared excesses and abundance anomalies of such stars. These conditions may be relatively common outside the Local Bubble."
    },
    "0812/0812.4389_arXiv.txt": {
        "abstract": "The inner crust of neutron stars, formed of a crystal lattice of nuclear clusters immersed in a sea of unbound neutrons, may be the unique example of periodic nuclear systems. We have calculated the neutron specific heat in the shallow part of the crust using the band theory of solids with Skyrme nucleon-nucleon interactions. We have also tested the validity of various approximations. We have found that the neutron specific heat is well described by that of a Fermi gas, while the motion of the unbound neutrons is strongly affected by the nuclear lattice. These apparently contradictory results are explained by the particular properties of the neutron Fermi surface. ",
        "introduction": "\\label{sec:bandtheory}  Following the standard assumptions, we consider a body centered cubic crystal with only one nuclear cluster per lattice site~\\cite{haen07}. We have taken the composition $(N,Z)$ of the clusters calculated by Negele and Vautherin~\\cite{ne73}. But the nucleon distributions have been recalculated at each temperature $T$ by solving the finite temperature Hartree-Fock equations with the Skyrme SLy4 effective nucleon-nucleon interaction~\\cite{chabanat98}. Typical temperatures of interest for cooling isolated neutron stars and for accreting neutron stars lie in the range between 10$^7$ to few 10$^9$~K~\\cite{gne01}.   Fully self-consistent calculations within the band theory are computationally very expensive. We have thus solved the Hartree-Fock equations in two steps~\\cite{cha07}.  First, we have determined self-consistently the nucleon distributions and mean fields in the Wigner-Seitz approximation with mixed Neumann-Dirichlet boundary conditions. We have removed the spurious fluctuations of the neutron density by averaging in the interstitial region~\\cite{ne73}. In a second stage, we have solved the Schr\\\"odinger equation with Bloch boundary conditions using these self-consistent mean fields. We have fixed the lattice spacing so that the volume of the Wigner-Seitz polyhedron is equal to the volume of the spherical Wigner-Seitz cell of radius $R_{\\rm cell}$.  We have employed the Linearized Augmented Plane Wave method described in details in Ref.~\\cite{cha05}. As shown in Ref.~\\cite{cha07}, the one-body spin-orbit potential has a negligible effect on the neutron level ordering and was therefore neglected.  \\begin{table}[b] \\centering \\setlength{\\tabcolsep}{.05in} \\renewcommand{\\arraystretch}{1.8} \\caption{For the three shallow layers considered in this paper, are given: the total density ($\\rho$), the number of protons and neutrons ($Z$, $N$ respectively), the nuclear mass number $A$ (at $T=0$), the radius of the cell ($R_\\mathrm{cell}$) and the neutron gas density at $T$=0 ($\\rho_n^G$).} \\label{tab.struc} \\vspace{.5cm} \\begin{tabular}{cccccc} \\hline \\hline $\\rho$ [g.cm$^{-3}$] & $Z$ & $N$ & $A$ & $R_{\\rm cell}$ [fm] & $\\rho_n^G$ [fm$^{-3}$] \\\\ \\hline $6.69 \\times 10^{11}$ & 40 & 160 & 133 & 49.24 & 1.3 10$^{-4}$ \\\\ $1.00 \\times 10^{12}$ & 40 & 210 & 141 & 46.33 & 2.6 10$^{-4}$ \\\\ $1.47 \\times 10^{12}$ & 40 & 280 & 140 & 44.30 & 4.9 10$^{-4}$ \\\\ \\hline \\hline \\end{tabular} \\end{table}  The specific heat (per unit volume) is defined as \\begin{equation} \\label{eq.cv} c_V(T)=\\frac{\\partial U}{\\partial T} \\biggr\\vert_V= T \\frac{\\partial S}{\\partial T}\\biggr\\vert_V \\, , \\end{equation} (taking the Boltzmann's constant $k_{\\rm B}=1$) where $U$ is the total internal energy density, $S$ the entropy density and the partial derivatives are evaluated at constant volume $V$. We have found that the latter expression is numerically more convenient. The entropy density of the unbound neutrons is given by \\begin{equation} S=-\\sum_\\alpha\\hspace{-0.1cm}\\int\\hspace{-0.1cm}\\frac{{\\rm d}^3\\pmb{k}}{(2\\pi^3)} \\biggl[ f_{\\alpha\\pmb{k}} \\ln  f_{\\alpha\\pmb{k}}  + (1- f_{\\alpha\\pmb{k}} )\\ln (1-f_{\\alpha\\pmb{k}} )\\biggr] \\label{eq:s} \\end{equation} where $\\alpha$ is the band index, $\\pmb{k}$ is the Bloch wave vector and $f_{\\alpha\\pmb{k}}$ is the Fermi-Dirac distribution. The latter depends on the temperature and the neutron chemical potential $\\mu_n$ which is determined from the total density $\\rho^G_n$ of unbound neutrons by \\begin{equation} \\rho_n^G=\\sum_\\alpha \\int \\frac{{\\rm d}^3\\pmb{k}}{(2\\pi^3)} f_{\\alpha\\pmb{k}} \\, . \\end{equation} Integrations have been carried out using the special point method (see \\cite{cha05}). Note that the number of unbound neutrons per cluster is determined by the shape of the mean fields, which vary with temperature. Consequently $\\rho^G_n$ can also depend on the temperature. In particular, with increasing temperature the most loosely bound neutrons may be excited into the delocalized states. ",
        "conclusions": "Modelling the recently observed thermal relaxation of neutron star crusts requires the knowledge of their thermal properties. In this paper, we have computed the specific heat of the neutron ocean permeating the inner crust at densities below $1.5\\times 10^{12}$ g.cm$^{-3}$, by applying the band theory of solids with the Skyrme SLy4 nucleon-nucleon interaction. We have compared the results obtained using different approximations. We have found that for temperatures $T=10^7-10^9$K relevant to neutron stars, the neutron specific heat is essentially given by that of a Fermi gas. The Thomas-Fermi expression~(\\ref{eq.cv-tf}) yields nearly undistinguishable results. The specific heat calculated in the Wigner-Seitz approximation agrees reasonably well, while the de Gennes approximation~(\\ref{eq.cv-dg}) leads to specific heats a factor of 2-3 smaller than those obtained in the band theory.  The results on the specific heat might suggest that the unbound neutrons are not affected by the presence of the nuclear clusters. It is however not true since the dynamical effective mass of the same unbound neutrons is very different from the bare one unvealing strong interactions with the periodic lattice. This apparent paradox can be explained by the fact that the neutron Fermi surface area and the neutron group velocity are both reduced by the same factor in the presence of clusters, leading to a nearly unchanged density of states and specific heat."
    },
    "0812/0812.1565_arXiv.txt": {
        "abstract": "{ We present our discovery of pulsational radial-velocity variations in the cool Ap star \\hd, an object spectroscopically similar to the bright, rapidly-oscillating Ap (roAp) star \\equ. Based on high-resolution time-series spectroscopy obtained with the HARPS spectrometer at the European Southern Observatory 3.6-m telescope, we detected oscillations in \\ndt\\ and \\nd\\ lines with a period close to 9~min and amplitudes of 20--30~\\ms. Substantial variation in the pulsational amplitude during our 3.8~h observing run reveals the presence of at least three excited non-radial modes. The detection of extremely low amplitude pulsations in \\hd\\ indicates that the roAp excitation mechanism produces variability in the radial velocity amplitude of between a few tens \\ms\\ and several \\kms. This supports the idea that many, if not all, cool Ap stars occupying the roAp instability strip may harbour non-radial pulsations, which currently remain undetected due to their small photometric and radial-velocity amplitudes. } ",
        "introduction": "\\label{intro}  Rapidly-oscillating Ap (roAp) stars are magnetic chemically-peculiar stars pulsating in non-radial, magnetoacoustic {\\it p-}modes of periods close to 10~min. Pulsations in roAp stars are believed to be driven by the opacity mechanism operating in the hydrogen-ionization zone \\citep{balmforth:2001}. The presence of strong magnetic fields in roAp stars enhances the driving of the high-overtone oscillations by the suppression of convection, influences pulsation frequencies, and determines the global geometry of the pulsational perturbations \\citep{dziembowski:1996,saio:2005}.  Strongly inhomogeneous vertical distributions of chemical elements combined with the rapid transformation of the outwardly-propagating pulsation waves are responsible for the unique spectroscopic pulsation signature of roAp stars \\citep{ryabchikova:2002}. In particular, the lines of rare-earth elements (REEs) and the cores of the hydrogen lines pulsate with a factor of 10--100 higher amplitudes than the remaining spectral features \\citep{kochukhov:2001b,mkrtichian:2003,kurtz:2006b}.  The observed roAp instability strip is \\textbf{limited} to the \\teff\\ range 6400--8100~K, although theoretical stability calculations \\citep{cunha:2002} predict pulsations in stars as hot as 9500~K and are unable to account for the unstable modes observed in stars with \\teff\\,$\\le$\\,7400~K. At the same time, the coexistence of pulsating and apparently constant Ap stars in the same region of the H-R diagram has been a long-standing puzzle \\citep[e.g.,][]{martinez:1994}. It is now understood that the time-resolved photometric techniques employed to detect and study variability in most of the 37 known roAp stars are insensitive to the low-amplitude pulsations observed in spectroscopic time-series analyses \\citep{hatzes:2004,elkin:2005}. This leads to the suggestion that all magnetic Ap stars in a certain temperature range may oscillate, but some have amplitudes below the photometric detection threshold \\citep{kochukhov:2002}.  We test this hypothesis by completing a high-precision survey of a sample of bright cool magnetic Ap stars using the  High Accuracy Radial velocity Planet Searcher (HARPS) spectrograph at the European Southern Observatory (ESO). The first result of our observations -- the discovery of 10.9-minute oscillations in the Ap star HD\\,115226 -- was reported by \\citet{kochukhov:2008b}. Here we present the discovery of a new rapidly oscillating Ap star, \\object{\\hd}, which pulsates with one of the lowest radial-velocity (RV) amplitudes measured for roAp stars. ",
        "conclusions": "\\label{discus}  We have established the presence of multiperiodic pulsations in the cool magnetic Ap star \\hd\\ using combined RV measurements of the lines belonging to \\ndt\\ and \\nd. The star exhibits oscillations with three frequencies, which have different amplitude ratios for the two Nd ions. The phase lag between RV curves of \\ndt\\ and \\nd\\ can be interpreted in the framework of the outwardly propagating pulsational perturbation, which first reaches the layer where \\ndt\\ lines form and, after some delay, is seen in the higher atmospheric layer probed by stronger \\nd\\ lines \\citep{ryabchikova:2007b,mashonkina:2005}. The difference in the amplitude ratios of the frequency components of the \\ndt\\ and \\nd\\ RV curves can be ascribed to different vertical cross-sections of the three pulsation modes.  \\hd\\ is the 38th known roAp star. Its discovery is significant because the star's pulsation amplitude is noticeably lower than for other roAp stars discovered to date using time-resolved spectroscopy. For example, HD\\,218994 \\citep{gonzalez:2007} and HD\\,115226 \\citep{kochukhov:2008b} pulsate with an amplitude $\\ge$\\,500~\\ms, while for HD\\,116114 \\citep{elkin:2005} and HD\\,154708 \\citep{kurtz:2006b} pulsations with an amplitude of 50--100~\\ms\\ were reported. Only for $\\beta$\\,CrB (HD\\,137909) comparable RV amplitudes of 20--30~\\ms\\ were found in individual lines of singly ionized REEs \\citep{kurtz:2007a}. However, $\\beta$\\,CrB is in many respects different from other roAp stars. It is an evolved star with a long pulsation period and a chemical composition deviating from that of a typical roAp star \\citep{ryabchikova:2004a}. In contrast, \\hd\\ appears to have average roAp characteristics and is, in fact, a spectroscopic twin of the well-known roAp star \\equ. Nevertheless, it pulsates with an unusually low RV amplitude. This shows that although the atmospheric chemical composition, in particular the REE ionization anomaly, is helpful in selecting roAp candidates, it has no direct connection with the amplitude of oscillations in the line-forming region.  \\citet{kurtz:2006b} noted the tendency for weaker roAp oscillations to be found in stars with stronger fields. However, \\hd\\ has a mean field modulus that is significantly weaker than many roAp stars but also shows an exceptionally low pulsation amplitude. We therefore conclude that a low-amplitude roAp pulsation can be present in cool Ap stars of any field strength. Although there are theoretical reasons to believe that the magnetic field alters the amplitude of the photospheric oscillations in few strong-field stars, a parameter other than the field intensity defines pulsation amplitude for other roAp stars.  The detection of \\textbf{very} low amplitude pulsations in \\hd\\ suggests that the roAp excitation mechanism produces oscillations with no apparent lower amplitude threshold. Thus, many cool Ap stars may possess pulsations with RV amplitudes $\\ll$\\,100~\\ms, which can be currently detected by time-resolved spectroscopy only in bright sharp-line stars such as \\hd."
    },
    "0812/0812.0339_arXiv.txt": {
        "abstract": "The X-ray structure of Kepler's supernova remnant shows a rounded shape delineated by forward shocks.  We measure proper-motions of the forward shocks on overall rims of the remnant, by using archival {\\it Chandra} data taken in two epochs with time difference of 6.09 yr. The proper-motions of the forward shocks on the northern rim are measured to be 0$^{\\prime\\prime}$.076 ($\\pm$0$^{\\prime\\prime}$.032 $\\pm$0$^{\\prime\\prime}$.016) -- 0$^{\\prime\\prime}$.11 ($\\pm$0$^{\\prime\\prime}$.014 $\\pm$0$^{\\prime\\prime}$.016) yr$^{-1}$, while those on the rest of the rims are measured to be 0$^{\\prime\\prime}$.15 ($\\pm$0$^{\\prime\\prime}$.017 $\\pm$0$^{\\prime\\prime}$.016) -- 0$^{\\prime\\prime}$.30 ($\\pm$0$^{\\prime\\prime}$.048 $\\pm$0$^{\\prime\\prime}$.016) yr$^{-1}$, here the first-term errors are statistical uncertainties and the second-term errors are systematic uncertainties.  Combining the best-estimated shock velocity of 1660$\\pm$120\\,km\\,sec$^{-1}$ measured for Balmer-dominated filaments in the northern and central portions of the remnant (Sankrit et al.\\ 2005) with the proper-motions derived for the forward shocks on the northern rim, we estimate the distance of 3.3$^{+1.6}_{-0.4}$\\,kpc to the remnant.  We measure the expansion indices, $m$, (defined as $R \\varpropto t^m$) to be 0.47--0.82 for most of the rims.  These values are consistent with those expected in Type-{\\scshape I}a SN explosion models, in which the ejecta and the circumstellar medium have power-law density profiles whose indices are 5--7 and 0--2, respectively.  Also, we shuold note the slower expansion on the northern rim than that on the southern rim.  This is likely caused by the inhomogeneous circumstellar medium; the density of the circumstellar medium is higher in the north than that in the south of the remnant.  The newly estimated geometric center, around which we believe the explosion point exists, is located at $\\sim$5$^{\\prime\\prime}$ offset in the north of the radio center. ",
        "introduction": "Kepler's supernova remnant (SNR; SN 1604) is one of the historical SNRs.  The X-ray structure of the remnant shows a clear circular shape with a radius of about 100$^{\\prime\\prime}$.  The distance estimations to the remnant have been scattered in a range of 3--5\\,kpc (e.g., Green 1984; Schaefer 1994).  Although the far side of 5\\,kpc supported by H {\\scshape I} observations (Reynoso \\& Goss 1999) has been adopted in recent literature, optical studies combining proper-motions with shock velocities of Balmer-dominated filaments have preferred a near side of 2.9$\\pm$0.4\\,kpc (Blair et al.\\ 1991) or 3.9$^{+1.4}_{-0.9}$\\,kpc (Sankrit et al.\\ 2005).  The SN type of Kepler's SNR is very interesting.  Based on the light curve of the Kepler's SN, Baade (1943) classified this remnant as a result of Type-{\\scshape I}a SN.  {\\it Ginga} spectrum from this remnant showed a strong Fe K line, supporting that the Kepler's SN was a Type-{\\scshape I}a SN (Hatsukade et al.\\ 1990).  {\\it ASCA} spectrum of this remnant showed strong K-shell lines from Si, S, and Fe, and the measured relative abundances supported Type-{\\scshape I}a origin (Kinugasa \\& Tsunemi 1999).  Recent deep {\\it Chandra} observations detected strong emission lines from heavy elements almost everywhere in the remnant, again supporting Type-{\\scshape I}a origin (Reynolds et al.\\ 2007).  On the other hand, the density of the ambient medium around the remnant was estimated to be so high (at least 1\\,cm$^{-3}$; e.g., Hughes \\& Helfand 1985) at a large height of $\\sim$470\\,pc (assuming the distance of 4\\,kpc to the remnant) from the Galactic plane.  The high density of the ambient medium was considered to be a signature of circumstellar material (CSM) which was blown off from a progenitor star as a stellar wind.  In addition to the existence of the dense surroundings, Nitrogen-rich materials as a result of CNO-processing in a massive progenitor star were detected in some optical knots (van den Bergh \\& Kamper 1977; Dennefeld 1982).  These facts suggested that the Kepler's remnant was core-collapse in origin. Also, Bandiera (1987) suggested a massive runaway star as a possible progenitor of Kepler's SNR, since it was able to naturally account for the presence of the CSM and the asymmetric structure of the remnant in optical wavelength.  Recently, Blair et al.\\ (2007) proposed that the Kepler's SN was categorized as a Type-{\\scshape I}a explosion in a region with significant CSM, which was a small but growing class of Type-{\\scshape  I}a SNe named as Type-{\\scshape I}a/IIn by Kotak et al.\\ (2004).  Due to the large expansion of Kepler's SNR, it is the best target from which we can investigate the detailed preexisting structures of the CSM for this rare kind of SNe.  Here, we report proper-motions of the forward shocks on the overall rims of Kepler's SNR, by using archival {\\it Chandra} data.  We derive the distance to the remnant, combining the proper-motions we measure with the optically determined forward shock velocity (Sankrit et al.\\ 2005).  Also, we present evolutional states in various portions of the remnant, which gives us critical information on structures of the surrounding CSM. ",
        "conclusions": "We have measured proper-motions of the forward shocks on the overall rims of Kepler's SNR, using the archival {\\it Chandra} data.  The expansion indices measured at various parts of the rim supports the Type-{\\scshape I}a SN which has the situation that the ejecta and the CSM, respectively, have power-law density profiles with indices of 5--7 and 0--2 rather than the core-collapse SN.  We find that the shock velocities are asymmetric: the shock velocities on the northern rim are 1.5--3 times slower than those on the rest of the remnant.  We attribute this asymmetry to the density inhomogeneities of the CSM surrounding the remnant.  The shape of the X-ray boundary of the remnant as well as the inhomogeneous CSM structures lead us to consider that the expansion center is located at $\\sim$5$^{\\prime\\prime}$ offset in the north of the radio center."
    },
    "0812/0812.2799_arXiv.txt": {
        "abstract": "We present a new multi--fluid, grid MHD code PIERNIK, which is based on the Relaxing TVD scheme~\\cite{Jin95}. The original scheme (see Trac \\& Pen~\\cite*{2003PASP..115..303T} and Pen~\\etal~\\cite*{2003ApJS..149..447P}) has been extended by an addition of dynamically independent, but interacting fluids: dust and a diffusive cosmic ray gas, described within the fluid approximation, with an option to add other fluids in an easy way.  The code has been equipped with shearing--box boundary conditions, and a selfgravity module, Ohmic resistivity module, as well as other facilities which are useful in astrophysical fluid--dynamical simulations. The code is parallelized by means of the MPI library. In this paper we introduce the multifluid extension of Relaxing TVD scheme and present a test case of dust migration in a two--fluid disk composed of gas and dust. We demonstrate that due to the difference in azimuthal velocities of gas and dust and the drag force acting on both components dust drifts towards maxima of gas pressure distribution. ",
        "introduction": "The basic set of conservative MHD equations (see Paper I, \\cite{2008arXiv0812.2161H}, this volume) describes a single fluid. The Relaxing TVD scheme by Pen, Arras \\& Wong~\\cite*{2003ApJS..149..447P} can be easily extended for multiple fluids by concatenation of the vectors of conservative variables for different fluids \\begin{equation} \\mathbf{u} =\\big( \\underbrace{\\rho^{i}, m_x^{i}, m_y^{i}, m_z^{i}, e^{i}}_{\\textrm{\\scriptsize ionized gas}}, \\underbrace{\\rho^{n}, m_x^{n}, m_y^{n}, m_z^{n}, e^{n}}_{\\textrm{\\scriptsize neutral gas}}, \\underbrace{\\rho^{d}, m_x^{d}, m_y^{d}, m_z^{d}}_{\\textrm{\\scriptsize dust}}  \\big), \\end{equation} representing ionized gas, neutral gas, as well as dust treated as a pressureless fluid.  In a short notation this can be written as \\begin{equation} \\mathbf{u} = (\\mathbf{u}^i, \\mathbf{u}^n, \\mathbf{u}^d). \\end{equation} The corresponding fluxes for these fluids are combined in a similar way \\begin{equation} \\mathbf{F}(\\mathbf{u},\\mathbf{B}) = \\left(\\mathbf{F}^i(\\mathbf{u}^i,\\mathbf{B}), \\mathbf{F}^n(\\mathbf{u}^n), \\mathbf{F}^d(\\mathbf{u}^d)\\right), \\end{equation} where the elementary flux vectors like $\\mathbf{F}^i(\\mathbf{u}^i,\\mathbf{B})$, $\\mathbf{F}^n(\\mathbf{u}^n)$ and $\\mathbf{F}^d(\\mathbf{u}^d)$ are considered as fluxes computed independently for each fluid in the single fluid description. In multidimensional computations the fluxes $\\mathbf{G}(\\mathbf{u},\\mathbf{B})$  and $\\mathbf{H}(\\mathbf{u},\\mathbf{B})$, corresponding to the transport of conservative quantities in $y$ and $z$--directions, are constructed in a similar manner. The full multifluid system of MHD equations, including source terms \\begin{equation} \\partial_t \\mathbf{u} + \\partial_x \\mathbf{F} = \\mathbf{S}(\\mathbf{u}), \\end{equation} is subsequently solved by means of the Relaxing TVD scheme described in Paper~I, together with the induction equation coupling magnetic filed $\\mathbf{B}$ to the ionized gas component. The term $\\mathbf{S}(\\mathbf{u})$ includes any source terms (like gravity)  corresponding to individual fluids as well as the terms coupling the dynamics of different fluids.  As an example of mutual fluid interaction we consider the environment of protoplanetary discs, where  aerodynamic interaction of gas and dust particles takes place. The interaction of the neutral gas and dust components induces the effects of the drag force in the momentum and energy equations for these fluids \\begin{equation} \\vec{S}_m^d = \\alpha^{dn} \\rho^d \\rho^n \\left(\\vec{v}^n-\\vec{v}^d\\right),\\qquad \\vec{S}_m^n = - \\vec{S}_m^d, \\label{eqn:drag-force} \\end{equation} \\begin{equation} S_{e}^n = \\alpha^{dn} \\rho^d \\rho^n \\vec{v}_n \\cdot \\left( \\vec{v}_x^d-\\vec{v}_x^n\\right), \\end{equation} where % \\begin{equation} \\alpha^{dn} = \\frac{1}{\\rho^n t_{\\textrm{\\scriptsize stop}}} \\end{equation} is the inverse product of gas density and particle stopping time. In general $\\alpha^{dn}$ depends on properties of dust grains (theirs sizes and bulk densities) and gas, and their relative velocities. For simplicity in the calculations presented we assumed its constant value of $10$. Friction forces can be incorporated, when needed, between any pair of fluids. ",
        "conclusions": ""
    },
    "0812/0812.3340_arXiv.txt": {
        "abstract": "The two brightest and so far the best studied gamma ray bursts (GRBs), 990123 and 080319B, were ordinary, highly collimated GRBs produced in a core collapse supernova explosion within a high-density wind environment and observed from a very near-axis viewing angle. Inverse Compton scattering (ICS) and synchrotron radiation (SR), the two dominant radiation mechanisms in the cannonball (CB) model of GRBs, together with the burst environment, provide a very simple and sufficiently accurate description of the multiwavelength lightcurves of their prompt and afterglow emissions. ",
        "introduction": "Two models have been used extensively to analyze Gamma ray bursts (GRBs) and their afterglows (AGs), the fireball (FB) model (for recent reviews see, e.g., M\\'{e}sz\\'{a}ros~2002, 2006; Zhang~2007) and the cannonball (CB) model (Dar \\& De R\\'ujula~2004, hereafter DD2004; Dado \\& Dar 2008a,b, hereafter DD2008a,b, and references therein). Despite their similar names, the two models are (or were) entirely different (e.g., Piran~1999), hence only one of them, if either, may provide a faithful physical description of GRBs. Until recently, the FB model has been widely accepted as that model. However, the rich data on GRBs accumulated from space based observations, in particular from the Swift satellite, complemented by early time data from ground based robotic telescopes and late-time follow-up observations with larger telescopes, have challenged this prevailing view. Kumar et al.~(2007) concluded that the prompt $\\gamma$-ray emission cannot be produced in shocks, internal or external. Zhang, Liang and Zhang~(2007) found that the fast decay and rapid spectral softening ending the prompt gamma-ray and X-ray emission cannot be explained simultaneously by high latitude emission. The X-ray and optical afterglows (AGs) of the Swift GRBs were found to be chromatic at early time (Covino et al.~2006) and to have chromatic breaks (Panaitescu et al.~2006) which differ significantly from the jet breaks expected in the collimated fireball model of AGs. Burrows and Racusin~(2007) examined the XRT light curves of the first $\\sim\\! 150$ Swift GRBs and reported that the expected jet breaks are extremely rare. In particular, Liang et al.~(2008) have analyzed the Swift X-ray data for the 179 GRBs detected between January 2005 and January 2007 and the optical AGs of 57 pre- and post-Swift GRBs. They found that not a single burst satisfies all the criteria of a jet break.  In spite of the above, not all authors are so critical and they believe that the GRB data require only some modifications of the standard FB model in order to accommodate the results (e.g.,~Panaitescu~2006; Dai et al.~2007; Sato et al.~2007; Kumar and Panaitescu~2008). Other authors seem to ignore the failures of the FB model and continue to interpret the observations with the FB model hypotheses (`internal and external shocks', `colliding conical shells', `forward and reverse shocks', `continuous energy injection') and parametrize the data with freely adopted formulae (e.g., `segmented power laws') which were never derived explicitly from underlying physical assumptions (for recent examples, see, e.g., Bloom et al.~(21 authors)~2008; Racusin et al.~(93 authors)~2008).   The situation concerning the CB model is different. The predictions of the model were derived in fair approximations from its underlying assumptions. They were shown to describe correctly the main observed properties of GRBs and reproduce successfully the diverse broad band light curves of large representative sets of both long GRBs and XRFs (e.g., DD2004; DD2008a and references therein) and short hard bursts (DD2008b). In this paper, we demonstrate these for the GRBs 990123 and 080319B, the brightest GRBs observed so far that have the best sampled and the most accurately measured multiwavelegth light curves of the prompt and afterglow emission. We show that the two underlying radiation mechanisms of the CB model, inverse Compton scattering (ICS) and synchrotron radiation (SR), and the burst environment suffice within the CB model provide a very simple and successful description of the observed light curves of their prompt and afterglow emissions. ",
        "conclusions": "Accurate data that have been accumulated in recent years from space based and ground based observations have challenged the prevailing views on GRBs. This is true in particular for the brightest and best studied GRBs, 990123 (e.g., Maiorano et al.~2005; Corsi et al.~2005) and 080319B (Bloom et al.~2008; Racusin et al.~2008):  Their prompt multiwavelength emission cannot be explained by a single radiation mechanism. Their hard X-ray and $\\gamma$-ray emission cannot be explained as SR from internal shocks generated by collisions between conical shells. Their prompt optical emission is not correlated with their hard X-ray and $\\gamma$-ray emission. Their afterglows are chromatic, and roughly decay like a single power-law without a jet break. Their spectral and temporal power-law indices do not satisfy the closure relations for conical blast waves.  To rescue the FB model, SR as the source of the hard X-ray and $\\gamma$-ray emission was replaced by inverse Compton scattering of the self produced SR. In this so called `synchrotron self Compton' (SSC) mechanism, the prompt SR is produced either by internal shocks (e.g., Kumar \\& Panaitescu~2008) or by relativistic magnetic turbulences without internal shocks (e.g. Kumar \\& Narayan~2008). The SSC mechanism implies that the SR emission begins before the X-ray and $\\gamma$-ray emission, and both are correlated. But, contrary to these expectations, the observed prompt optical emission lags consideraby behind the hard X-ray and $\\gamma$-ray emission and no temporal correlation is observed between the hard X-ray and $\\gamma$-ray peaks and the prompt optical peaks. Moreover, the SSC model predicts a vanishing polarization of the prompt hard X-ray and $\\gamma$-ray emission. The measured polarization, so far in 4 ordinary GRBs, suggests a large polarization (Cobb et al.~2004; Willis et al.~2005; Kalemci et al.~2007; McGlynn et al.~2007).   The situation concerning the FB model interpretation of the observed afterglow of GRBs 990123 and 080319B is similar. Although some aspects of the data have been explained by invoking structured jets, multiple blast waves propagating into the ISM, backward shocks and evolving microphysical parameters, no unique falsifiable predictions to test these suggestions were made and it is not clear whether the agreement between theory and some aspects of the data which was claimed is significant or results only from a flexible parametrization and sufficient adjustable parameters.   The situation concerning the CB model is entirely different. The predictions of the model were derived in fair approximations from its original underlying assumptions long ago. They were shown to predict correctly the main observed properties of GRBs and reproduce successfully the diverse broad band light curves of large representative sets of both long GRBs and XRFs (e.g., DD2004; DD2008a and references therein) and short hard bursts (DD2008b). In this paper, we have demonstrated these for GRBs 990123 and 080319B, the two brightest and so far the best studied gamma ray bursts, without invoking any new physics. They imply that GRBs 990123 and 080319B were ordinary, highly collimated GRBs produced in a core collapse supernova explosion within a high-density wind environment and observed from a very near-axis viewing angle. Massive wind/ejecta blown from the progenitor star before the supernova explosion probably created a matter-free bubble around the progenitor star which was filled with radiation. Inverse Compton scattering of this light by the jet of CBs produced the hard X-rays and $\\gamma$-rays which preceded the broadband SR emission when the jet crossed the wind.  As in all other GRBS, inverse Compton scattering and synchrotron radiation, the two dominant radiation mechanisms in the cannonball model of GRBs, together with the burst environment -wind blown into a constant density ISM- provide a very simple and sufficiently accurate description of the multiwavelength lightcurves of their prompt and afterglow emissions. The dependence of the bend/break frequency on density variations along the jet trajectory can explain the observed spectral variations but obviously cannot predict it. The predicted  general trend $\\beta_O(0)\\!\\sim\\! 0.5\\!\\rightarrow$ $\\beta_O(t)\\!\\sim\\! 1.1$ at late time is well satisfiesd.  Two predictions which can be tested in future observations of extremely bright GRBs are as follows:  Because of their small viewing angles, the polarization of the prompt X-rays and $\\gamma$-rays in very bright GRBs, such as 990123 and 080319B, is predicted to be small (Shaviv \\& Dar~1995; DD2004), unlike that of ordinary GRBs with $\\gamma\\, \\theta\\approx 1$ where it is predicted to be very large.  Collisions of the CBs with the dense wind and glory, which are needed to produce very luminous GRBs, are expected to produce also detectable fluxes of high energy photons (Dado \\& Dar~ 2005: Dado \\& Dar, in preparation): Electrons from the wind and/or ISM which are swept into, or scattered by, the CBs should produce sub-GeV photons by ICS of glory light. Hadronic collisions of the thermal nuclei in the CBs' plasma with wind nuclei are expected to produce detectable fluxes of sub-TeV photons and marginal fluxes of neutrinos through $\\pi$ production, while wind/ISM nuclei swept into or scattered by the CBs are expected to produce fluxes of sub-PeV photons which are detectable only from very nearby bursts because of the opacity of the infrared background to high energy photons. (Dar \\& De R\\'ujula~2006).    {\\bf Acknowledgment:} We thank Elisabetta Maiorano for making available to us tabulated data of the BeppoSAX measurements of the X-ray light curves of GRB 990123."
    },
    "0812/0812.1942_arXiv.txt": {
        "abstract": "We have calculated  optical spectra of hydrogen-rich (DA) white dwarfs with magnetic field strengths between 1 MG and 1000 MG for temperatures between 7000 K and 50000 K. Through a least-squares minimization scheme with an evolutionary algorithm, we have analyzed the spectra of 114 magnetic DAs from the SDSS (95 previously published plus 14 newly discovered within SDSS, and five discovered by SEGUE). Since we were limited to a single spectrum for each object we used only centered magnetic dipoles or dipoles which were shifted along the magnetic dipole axis. We also statistically  investigated the distribution of magnetic-field strengths and geometries of our sample. ",
        "introduction": "\\label{Introduction}  The Sloan Digital Sky Survey (SDSS) project covering $10^4$ ${\\rm deg}^2$ of the sky has discovered thousands of new white dwarfs, among them 105 with magnetic fields (MWDs) \\citep{Gaeansickeetal02, Schmidtetal03, Vanlandinghametal05}.White dwarfs with magnetic fields between $10^4$ and $ 10^9$ G are thought to represent more than 10\\% of the total population of white dwarfs \\citep{Liebertetal03}. The SDSS Data Release 3 (DR3, $\\approx 4200\\, {\\rm deg}^2$) increased the number of known magnetic white dwarfs from 65 \\citep{WickramasingheFerrario00,Jordan01} to 170 \\citep{Kawkaetal07}.  In this work we present the re-analysis of the 95 DA MWDs discovered by \\citet{Schmidtetal03} and \\citet{Vanlandinghametal05},  plus the analysis of 19 additional objects from SDSS up to DR6\\footnote{http://www.sdss.org/dr6/} ($9583\\,{\\rm deg}^2$).  \\citet{Schmidtetal03} and \\citet{Vanlandinghametal05} determined the field strengths  and the  inclinations of magnetic dipoles by comparing visually the observed spectra  with model spectra calculated using a simplified radiation transfer code \\citep{LatterSchmidt87}. Their analyses resulted in dipolar field strengths for the  SDSS MWDs  between $1.5$\\,MG and $\\sim1000$\\,MG. From their sample of 105 MWDs, 97 were classified as DAs. ",
        "conclusions": "Overall, our results are consistent with the former analyses of MWDs (see Fig.\\,\\ref{fig:comparison}), which shows that simple atmosphere models with preassumed dipole magnetic values are good approximations for these objects, if only single-phase spectroscopy without polarization information is considered. Our method is able to account for a diversity in low magnetic field strengths ($B\\,<\\,20$\\,MG). This is apparent in the deviation of our best-fit magnetic field strength values from literature values in the horizontal $\\sim$2\\,MG axis of Figure\\,\\ref{fig:comparison}. In many cases offset dipole models resulted in significantly better fits than the models with centered dipoles, however, at smaller field strength this can be in part due to the simplifications in the treatment of Stark broadening.  The distribution of the magnetic field strengths of the MWDs from the \\citet{Schmidtetal03} sample had  a maximum around $\\sim$5\\,-\\,30 MG. We have updated the values of DA white dwarfs and created a histogram of all known magnetic DA WDs  (Fig.\\,\\ref{fig:histogram}) and added results for the new objects. The magnetic field strengths of non-SDSS MWDs are from the compilation by \\citet{Kawkaetal07}. In spite of the SDSS biases, the same $\\sim$5\\,-\\,30 MG peak as in \\citet{Schmidtetal03} is apparent in Fig.\\,\\ref{fig:histogram}.  The dipole magnetic field ohmic decay timescale is $10^{10}$ yr. Even the higher multipoles can live for such a long period of time (see, e.g. \\cite{Muslimovetal95}). Therefore, no significant correlation between temperature and magnetic field strength is expected if temperature is assumed as an indicator of age (Fig.\\,\\ref{fig:bvst}). This lack of correlation supports the fossil ancestry of these fields inherited from earlier stages of stellar evolution.  High field MWDs are thought to be remnants of magnetic Ap and Bp stars. If flux conservation is assumed, the distribution of the polar field strengths of high field MWDs should be largest in the interval \\mbox{50--500\\,MG}. In our sample, objects with magnetic field strengths lower than 50 MG are more numerous than the objects with higher magnetic field strengths (see Fig.\\,\\ref{fig:histogram}), part of this effect is due to our biases (see Section \\ref{SDSS data}).  Nevertheless it is consistent with previous results and supports the hypothesis that magnetic fossil fields from Ap/Bp stars only are not sufficient to produce high field MWDs \\citep{WickramasingheFerrario05MN}. \\citet{Auriereetal07} argued that dipole magnetic field strengths of magnetic Ap/Bp stars have a ``magnetic threshold'' due to large scale stability conditions, and this results in a steep decrease in the number of magnetic Ap/Bp stars below polar magnetic fields of 300\\,G.  A possible progenitor population for MWDs with dipolar field strengths below 50 MG is the currently unobserved population of A and B stars with magnetic field strengths of 10\\,-\\,100\\,G. \\citet{WickramasingheFerrario05MN} suggested that if $\\sim40\\%$ of A/B stars have magnetism, this would be sufficient to explain the observed distribution of MWDs. However, the existence of this population seems to be highly unlikely since investigations of \\citet{Shorlinetal02} and \\citet{Bagnuloetal06} for magnetism in this population yielded null results, for median errors of 15\\,-\\,50\\,G and 80\\,G respectively. Another candidate group is the yet undetected magnetic F stars for these MWDs with lower field strengths \\citep{Schmidtetal03}. But this conclusion is strongly affected by SDSS MWD discovery biases.   \\ack{ Part of this work is supported by the DLR under grant 50 OR 0802. Funding for the creation and distribution of the SDSS archive has been provided by the Alfred P. Sloan Foundation, the SDSS member institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, Monbusho, and the Max Planck Society. The SDSS World Wide Web site is http://www.sdss.org/.}"
    },
    "0812/0812.3947_arXiv.txt": {
        "abstract": "We present images from a long term program (MOJAVE: Monitoring of Jets in AGN with VLBA Experiments) to survey the structure and evolution of parsec-scale jet phenomena associated with bright radio-loud active galaxies in the northern sky. The observations consist of 2424 15 GHz VLBA images of a complete flux-density limited sample of 135 AGN above declination $-$20 degrees, spanning the period 1994 August to 2007 September. These data were acquired as part of the MOJAVE and 2~cm Survey programs, and from the VLBA archive.  The sample selection criteria are based on multi-epoch parsec-scale (VLBA) flux density, and heavily favor highly variable and compact blazars. The sample includes nearly all the most prominent blazars in the northern sky, and is well-suited for statistical analysis and comparison with studies at other wavelengths. Our multi-epoch and stacked-epoch images show 94\\% of the sample to have apparent one-sided jet morphologies, most likely due to the effects of relativistic beaming. Of the remaining sources, five have two-sided parsec-scale jets, and three are effectively unresolved by the VLBA at 15 GHz, with essentially all of the flux density contained within a few tenths of a milliarcsecond. ",
        "introduction": "High angular resolution studies of the kinematics of AGN jets over many years have led to a better understanding of the process of acceleration and collimation of the relativistic jets associated with the ejection of relativistic plasma from a central supermassive black hole.  Early observations were made with ad hoc arrays comprised of relatively few antennas and focused on a few relatively strong and fast superluminal sources (e.g., \\citealt*{C75, S75, SCM75}).  In these studies it was difficult to trace details of the jet outflow due to limited temporal and interferometric coverage of the arrays, which consisted of existing antennas with differing instrumental characteristics. Moreover, by concentrating on only the fastest or most interesting sources, (e.g., \\citealt*{C77,ZCU95}), the kinematic results did not reflect the full range of speeds seen in the overall population. With the start of VLBA operations in the mid 1990s, we were able to obtain a reasonably uniform set of multi-epoch 15 GHz (2 cm) observations on 132 AGN \\citep{K98,2002AJ....124..662Z} which covered the period 1994 to 2002. From these data we were able to discuss the statistical properties of AGN jets \\citep {KL04,KKL05,H06}, including the intrinsic properties of the parent population \\citep{CLH07} and the detailed nature of a few sources of particular interest \\citep{H03,VRK03,LKV03,K08,KLH07}.  Although these studies included most known ``flat spectrum'' ($\\alpha > -0.5$, where $S_\\nu \\propto \\nu^\\alpha$) radio sources with total flux density above well defined limits, they included or excluded some sources because of the contribution of extended emission, errors in the measured spectral index and/or spectral curvature.  As described in paper I of this series \\citep{LH05}, our new MOJAVE (Monitoring of Jets in Active Galactic Nuclei with VLBA Experiments) sample includes all sources with measured total VLBA flux density at 15 GHz greater than a defined flux density at any epoch during the period 1994.0 to 2003.0.  For the 96 sources in common with the studies of \\cite{K98} and \\cite{2002AJ....124..662Z}, the time scale for determining the speed and possible accelerations or decelerations is now extended from 7 to 13 years. \\cite{LH05} give more details of the sample definition and also discuss the single-epoch linear polarization properties of the MOJAVE sample, while \\cite{HL06} discuss the circular polarization properties (Paper II).  The MOJAVE/2 cm Survey program is among the largest multi-epoch AGN VLBI surveys carried out to date, and is complemented by several programs at other wavelengths. These include the 5 GHz Caltech-Jodrell Flat-Spectrum Survey \\citep{Britzen07,2008AA...484..119B}, the 2 \\& 8 GHz Radio Reference Frame program \\citep{PMF07}, and the 43 GHz Boston University AGN monitoring program \\citep{J07}. Each of these programs offers various trade-offs in terms of sample size, angular resolution, and image sensitivity. Shorter wavelength studies, such as the Boston University program, provide better angular resolution than at 15 GHz, but the sensitivity is poorer, jet features fade faster, and the impact of tropospheric conditions  is more severe.  Thus, short wavelength observations need to be more frequent, and consequently must focus on much smaller samples of extremely bright sources.  Longer wavelength studies have better surface brightness sensitivity and so can trace the jet motions out to greater distances, but at the expense of degraded angular resolution.  This paper is the fifth in a series describing the results of the MOJAVE program, with the most recent ones discussing the kiloparsec-scale jet properties \\citep{CLK07}(Paper III) and parent luminosity function \\citep{CL08}(Paper IV) of the sample. Here we present a complete set of 15 GHz VLBA images of the flux-density limited MOJAVE sample of 135 compact extragalactic radio jets obtained during the period 1994 August to 2007 September. The overall layout of the paper is as follows: in \\S~\\ref{SD} we discuss the sample definition and selection criteria, in \\S~\\ref{obs} we discuss our observational program and data reduction method, and in \\S~\\ref{morph} we summarize the overall parsec-scale morphological properties of the AGN jets and describe our current monitoring program.  In subsequent papers we will discuss their kinematic properties, overall demographics, and polarization properties. ",
        "conclusions": "}  The MOJAVE program represents the largest full-polarization high resolution monitoring program of powerful AGN jets carried out to date. The complete flux-density limited sample of 135 sources contains the brightest and most highly beamed AGN visible in the northern sky. Composed primarily of blazars, it is especially useful for comparisons with studies at other wavelengths, especially at very high energies, where blazars tend to dominate the all-sky catalogs. In 2006, we expanded the monitoring sample to include an additional 65 AGN on the basis of their known gamma-ray emission, low luminosity, and/or unusual jet kinematics. This list will be further supplemented with new AGN detections from the Fermi gamma-ray observatory.  By combining archival data from the 2 cm Survey and other VLBA programs, we have obtained 2424 individual 15 GHz VLBA images on the flux-density limited MOJAVE sample, covering the period 1994 August to 2007 September. The full set of images is publicly available on the MOJAVE website. These data indicate that 94\\% of the jets are one-sided, likely due to relativistic beaming effects. Of the remainder, three AGN (0235+164, 1741$-$038, and 1751+288) are essentially unresolved on VLBA scales, having essentially all of their flux density contained within a few tenths of a milliarcsecond. Five other AGN: 0238-084(NGC~1052), 0316+413 (3C~84), 1228+126 (M87), 1413+135, 1957+405 (Cygnus A), have two-sided jet morphologies. The statistical and kinematic analysis of the full dataset will be presented in subsequent papers in this series."
    },
    "0812/0812.1459_arXiv.txt": {
        "abstract": "We study effects of possible tachyonic perturbations of dark energy on the CMB temperature anisotropy. Motivated by some models of phantom energy, we consider both Lorentz-invariant and Lorentz-violating dispersion relations for tachyonic perturbations. We show that in the Lorentz-violating case, the shape of the CMB anisotropy spectrum generated by the tachyonic perturbations is very different from that due to adiabatic scalar perturbations and, if sizeable, it would be straightforwardly distinguished from the latter. The tachyonic contribution improves slightly the agreement between the theory and data; however, this improvement is not statistically significant, so our analysis results in limits on the time scale of the tachyonic instability. In the Lorentz-invariant case, tachyonic contribution is a rapidly decaying function of the multipole number $l$, so that the entire observed dipole may be generated without conflicting the data at higher multipoles. On the conservative side, our comparison with the data places limit on the absolute value of the (imaginary) tachyon mass in the Lorentz-invariant case. ",
        "introduction": "Recently, a number of suggestions have been put forward for explaining the observed accelerated expansion of the Universe. Among them there are the presence of the cosmological constant, modification of gravity at ultra-large scales, existence of new light fields (for reviews, see, e.g., Refs.~\\cite{Padmanabhan:2002ji, Sahni:2004ai, Copeland:2006wr, Sahni:2006pa, Frieman:2008sn}). In the latter case, dark energy can be characterized by equation of state $p=w\\rho$, where the parameter $w$ is different from $-1$ and generically depends on time. In a simple version, dark energy is due to a scalar field (quintessence), and the parameter $w$ obeys the constraint $w>-1$, while $w=-1$ for the cosmological constant. However, the value of $w$  may be strongly negative, $w<-1$; this is the case for phantom dark energy. Existing data do not exclude also a possibility ~\\cite{Komatsu:2008hk, Sahni:2008xx, Xia:2008ex} that at relatively large redshift  the dark energy is of quintessence type $w>-1$, and later it becomes phantom with $w<-1$.  Phantom energy violates null energy condition, which is usually a signal for instabilities. As an example, the simplest model of  scalar field with the wrong sign of kinetic term~\\cite{Caldwell:2002} suffers from the presence of ghost (negative energy state) at arbitrarily high spatial momenta. This implies catastrophic vacuum instability. However, in phantom models that break Lorentz-invariance, violation of the null energy condition at cosmological scales (related to the property that $w<-1$ for spatially homogeneous phantom field) does not necessarily imply unacceptable instabilities at shorter scales. This suggests that Lorentz-violating phantom theories may be viable. Indeed, models of this sort have been recently constructed~\\cite{Senatore, Crem:2006,Rubakov:2006, Libanov:2007}. A property of one class of these models~\\cite{Rubakov:2006, Libanov:2007} is that there is a tachyonic mode in the perturbation spectrum about the homogeneous phantom background. This mode  occurs at sufficiently small spatial momenta only, so that the time scale of the tachyonic instability may be roughly comparable to the age of the Universe. This is not particularly dangerous.  It is conceivable that the existence of  tachyonic modes at low spatial momenta is a fairly generic property of a class of phantom models: the violation of the null energy condition may show up precisely in this way. Therefore, it is of interest to study observable consequences of such models. In this paper we consider one of these consequences, namely, the effect of tachyonic modes on the anisotropy of CMB temperature. We adopt the phenomenological approach, and instead of using results obtained within a concrete model, we parametrize the tachyonic instability by a few parameters. We will consider models  with a Lorentz-violating dispersion relation and Lorentz-invariant one.  In the Lorentz-violating case, we parametrize the dispersion relation as follows: \\begin{equation} \\omega^2 = \\alpha |{\\bf p}| ( M - |{\\bf p}|), \\label{dispers} \\end{equation} where $\\alpha$  and $M$ are constant parameters, and ${\\bf p}$ is the physical spatial momentum. Our convention is that positive values of $\\omega ^{2}$ correspond to exponential growth of perturbations, while the usual oscillatory behaviour occurs at negative $\\omega ^{2}$. Thus, the parameter $M$ equals to the momentum below which the mode is tachyonic. For given $\\alpha$, the parameter $M$ determines also the time scale of instability. The parametrization (\\ref{dispers}) is chosen in accord with Refs.~\\cite{Rubakov:2006, Libanov:2007} where similar dispersion relation has been found in a concrete model of phantom energy. The analysis presented below can be straightforwardly generalized to other forms of dispersion relation.  In the Lorentz-invariant case the dispersion relation has the form \\begin{equation} \\omega ^{2}=M^{2}-\\mathbf{p}^{2} \\label{Eq/Pg3/1:PhanTach} \\end{equation} In this case too, the tachyon mass $M$ equals to the spatial momentum below which the mode is unstable, and $M^{-1}$ is the time scale of instability. From our perspective, the important difference between the dispersion relations (\\ref{dispers}) and (\\ref{Eq/Pg3/1:PhanTach}) is that in the latter case the ``frequency'' is nonzero at ${\\bf p}=0$ and monotonously decreases as $|{\\bf p}|$ increases, whereas in the former, the ``frequency'' vanishes as  ${\\bf p}=0$ and has a maximum at finite $|{\\bf p}|$. This will lead to qualitatively different shapes of  the contributions to the CMB anisotropy spectra. It is worth noting in this regard that the Lorentz-invariant model may be viewed as a representative of a class of theories with tachyonic perturbations: analogous results would hold for Lorentz-violating models with dispersion relations similar to (\\ref{Eq/Pg3/1:PhanTach}).   In the cosmological context, the dispersion relations (\\ref{dispers}) and (\\ref{Eq/Pg3/1:PhanTach}) are written as follows: \\begin{eqnarray} \\omega^2(t) &=& \\alpha \\dfrac{k}{a(t)}\\left( M - \\dfrac{k}{a(t)}\\right) \\; , \\nonumber \\\\ \\omega^2 (t) &=&  M^2 - \\frac{k^2}{a^2(t)} \\; , \\nonumber \\end{eqnarray} respectively, where $k$ is the time-independent conformal momentum and $a(t)$ is the scale factor. In the expanding Universe, a mode of given $k$ is first normal, and after the physical momentum gets reshifted down to $k/a=M$, it becomes unstable.   For $ M  >  H_{0}$, where $H_{0}$ is the present value of the Hubble parameter, the tachyonic modes in both models start to grow exponentially at times preceding the present cosmological epoch. The growth of the tachyonic perturbations gives rise to the growth of the gravitational potential $\\Phi$ generated by these perturbations\\footnote{Refs.~\\cite{Libanov:2007, rub08} consider the Lorentz-violating case, but that analysis is straightforwardly repeated in the Lorentz-invariant situation.}~\\cite{rub08} \\begin{eqnarray} \\Phi (t,\\mathbf{x})&=&\\frac{1}{(2\\pi )^{3/2}}\\int \\limits_{}^{}\\!d\\mathbf{k}\\Phi (t,\\mathbf{k})\\mathrm{e}^{i\\mathbf{kx}} + \\mbox{h.c.}, \\nonumber\\\\ \\Phi(t, {\\bf k}) &=&{A({\\bf k} )} \\exp\\left(\\int\\limits^t_{t_k}\\omega(t^{\\prime})\\;dt^{\\prime}\\right), \\label{gravpot} \\end{eqnarray} where $A( {\\bf k} )$ is the amplitude of primordial fluctuations in the mode with conformal momentum ${\\bf k}$ at the time $t_k$ when this mode becomes unstable. We will discuss the range of  the amplitudes $A(\\mathbf{k})$ later on. We stress that the gravitational potential (\\ref{gravpot}) is generated by tachyonic perturbations rather than inhomogeneities in the ordinary matter.  In this paper we calculate CMB multipoles generated by the gravitational potential (\\ref{gravpot}). In the Lorentz-violating case, the mechanism we discuss gives rise to the contribution to the CMB anisotropy spectrum which is quite different from the standard spectrum coming from adiabatic scalar perturbations generated, e.g., at inflationary stage (for the latter see, e.g., Refs.~\\cite{book1, book2, book3}). As we will see below, the tachyonic contribution leads to potentially observable features in the CMB spectrum at relatively low multipoles. On the other hand, there are hints towards the existence of deviations in the observed spectrum~\\cite{Komatsu:2008hk, Nolta:2008ih} from the predictions based on flat (Harrison-Zeldovich) or almost flat spectrum of primordial adiabatic perturbations. So, one is tempted to employ tachyonic instabilities to explain these deviations. We add the contribution due to the tachyonic perturbations to the standard contribution of the adiabatic modes and compare the result with the observed spectrum. We show that the tachyonic contributions slightly improve the agreement between the theory and data, but this improvement is statistically insignificant. So, our analysis only leads to  limits on the parameters of the tachyonic perturbations.  In the Lorentz-invariant case, tachyonic perturbations contribute to the lowest multipoles only. An interesting possibility here is that they may generate  the entire observed dipole without getting in conflict with measurements of higher multipoles.  This paper is organized as follows. We begin with preliminaries on the cosmological model in section~\\ref{sec:background}, and then discuss in detail the growth of the gravitational potential due to the tachyonic modes in section~\\ref{sec:growth}, first in the Lorentz-violating model and then in the Lorentz-invariant one. In section~\\ref{sec:multipoles} we calculate the contributions to the CMB multipoles, again distinguishing Lorentz-violating and Lorentz-invariant cases. We compare our results with the data in section~\\ref{sec:data}, and conclude in section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we have considered the effects on the anisotropy of CMB temperature due to possible tachyonic perturbations of dark energy. Because of the exponential growth, these perturbations may generate large gravitational potential $\\Phi$ at the recent cosmological epoch, and only at that epoch. This results in a sizeable Sachs-Wolfe effect. Note that the tachyonic perturbations we have discussed are unrelated to perturbations in baryons or dark matter, so their contribution to the CMB anisotropy does not correlate with the distribution of structure in the Universe.  Our analysis was mostly motivated by the Lorentz-violating models of phantom energy. Hence, we have studied in detail the tachyonic perturbations with the dispersion relation (\\ref{dispers}). We have seen that their effect on the CMB angular spectrum has a pronounced maximum whose position depends on one of the parameters, $\\alpha$, and is practically insensitive to  other parameters. It is expected that similar shape of the angular spectrum is characteristic to a wide class of models with Lorentz-violating tachyonic perturbations, as it is closely related to the fact that these perturbations become sizeable at late times only.  We have also considered tachyonic perturbations with Lorentz-invariant dispersion relation (\\ref{Eq/Pg3/1:PhanTach}). In that case, an interesting possibility is that the largest contribution to the CMB ansotropy is received by the dipole component, and the angular spectrum rapidly decays with the increase of $l$. We have seen that even if the entire observed dipole anisotropy is attributed to the tachyonic perturbations, the quadrupole component generated by them is still consistent with the observational data. It is worth noting that this result should be inherent not only in the Lorentz-invariant model, but also in Lorentz-violating models with tachyonic dispersion relations, provided that the ``frequency'' does not vanish at zero momentum and decreases as momentum increases.  Our main conclusion is that even if perturbations of the tachyonic type exist in the Universe, their contribution to the CMB anisotropy is small. Nevertheless, we do not exclude a possibility that growing precision of observations, and especially elaborate analysis of correlations between CMB anisotropy and structures in the Universe, may lead to hints toward the possible exotic property of dark energy, the tachyonic behaviour of its perturbations."
    },
    "0812/0812.2036_arXiv.txt": {
        "abstract": "Based on deep imaging from the Advanced Camera for Surveys aboard the Hubble Space Telescope, we present new evidence that stellar feedback created a $\\sim$ 1 kpc supergiant \\hi\\ shell (SGS) and triggered star formation (SF) around its rim in the M81 Group dwarf irregular galaxy IC~2574.  Using photometry of the resolved stars from the HST images, we measure the star formation history of a region including the SGS, focusing on the past 500 Myr, and employ the unique properties of blue helium burning stars to create a movie of SF in the SGS.  We find two significant episodes of SF inside the SGS from 200 $-$ 300 Myr and $\\sim$ 25 Myr ago.  Comparing the timing of the SF events to the dynamic age of the SGS and the energetics from the \\hi\\ and SF, we find compelling evidence that stellar feedback is responsible for creating the SGS and triggering secondary SF around its rim. ",
        "introduction": "Vigorous (SF) is one mechanism thought to be responsible for the creation of the holes and shells observed in the interstellar medium (ISM) of a wide variety of galaxies.  Following a strong episode of SF, energy input into the ISM from young massive stars (e.g., stellar winds, Type~{\\sc II}\\ supernovae) can provide sufficient mechanical energy to create these holes and shells and may even trigger secondary SF around the edges \\citep[e.g.,][]{ttb88, puc92}.  Although observations of young stellar clusters inside these ISM features confirm that this is one possible scenario, a majority of holes and shells do not contain bright associated young clusters \\citep[e.g.,][]{vdh96, rho99, pas08}. One explanation is that only a select few ISM holes and shells have been observed with sufficient photometric depth to detect the associated clusters.  Alternatively, a number of competing theories suggest the origins of holes and shells in the ISM are from mechanisms other than stellar feedback, e.g., high velocity cloud collisions \\citep[e.g.,][]{kam91}, interactions \\citep[e.g.,][]{bur04}, gamma ray bursts \\citep[e.g.,][]{loeb98}, stochastic SF propagation \\citep{har08}.  Dwarf irregular galaxies (dIs) provide excellent environments in which we can study the creation mechanisms of these ISM features.  Larger galaxies exhibit differential rotation and sheer, which can destroy the holes and shells in the ISM on timescales as short as 10$^{7}$ year.  However, dIs are typically solid body rotators \\citep[e.g.,][]{ski88,ski96}, which allows features in the ISM to grow to large sizes ($\\sim$ 1 kpc) and stay coherent over longer time scales \\citep[e.g.,][]{wal99}.  Nearby dIs are particularly convenient probes of the interplay between SF and the ISM, because we can directly compare the resolved stellar populations to the \\hi\\ component to better understand mechanisms of feedback.  At a distance of $\\sim$ 4 Mpc \\citep{kar02}, IC~2574, an M81 Group dI, contains a ``supergiant shell'' (SGS) \\citep[denoted by the red ellipse in panels (a) and (b) of Figure \\ref{f1};][]{wal98, wal99} and is an excellent candidate for specifically studying both the formation of the SGS and subsequent triggered SF around its rim.  \\citet{wal99} classify the SGS as an ellipse ($\\sim$ 1000 $\\times$ 500 pc), however note that many of the holes and shells in IC~2574 are elongated due to projection effects (the de-projected diameter is $\\sim$ 700 pc). They further note there is some evidence that the SGS is a superposition of two merged shells \\citep{wal99}, which could account for its elliptical shape. In addition to the shape, the expansion of the SGS is somewhat uncertain.  \\citet{wal99} measure an \\hi\\ expansion velocity of 25 \\kms\\ from a break in the PV-diagram, which is only an indirect measure of expansion, and in fact is consistent with stalled expansion as well.  We discuss the implications of these scenarios in \\S 4.  UV observations show a young stellar cluster \\citep[$\\sim$ 11 Myr old,][]{ste00} inside the SGS, but no \\halpha\\ emission is observed from the interior.  Around the rim, strong \\halpha\\ (and mid-IR) emission is observed \\citep[see panel (c) of Figure \\ref{f1}, e.g.,][]{wal98, can05} indicating SF at these locations within the past $\\sim$ 5 Myr.  In this paper, we use deep HST/ACS imaging of the resolved stellar populations of the SGS to measure its recent SFH ($<$ 500 Myr).  We use the unique properties of resolved blue helium burning stars (BHeBs) to trace the spatial locations of the SF episodes, resulting in a movie that shows how SF propagated in the vicinity of the SGS over the past 500 Myr.  From these lines of evidence, we are able to directly explore the SF event(s) responsible for the creation of the SGS and triggering of secondary SF. ",
        "conclusions": "\\begin{figure}[t] \\epsscale{1.2} \\plotone{f3.eps} \\caption{ (\\textit{top}) The SFH of the SGS region over the past 500 Myr.  The red dashed line indicates the SFR of the SGS averaged over the lifetime of the galaxy; (\\textit{bottom}) The cumulative energy from SF in the SGS over the past 500 Myr calculated with STARBURST99.} \\label{f3} \\end{figure}      To understand the current state of the SGS, we consider the energetics of both the \\hi\\ and SF, the role of triggered SF, and the timing of the events that led to the evacuation of the \\hi\\ and secondary SF around the rim.  To do this we examine the efficiency of energy, i.e. the feedback efficiency, transfered from stellar feedback into moving the ISM in the context of both an expanding and stalled shell.  \\begin{figure*}[!ht] \\epsscale{1.1} \\plotone{f4.eps} \\caption{\\footnotesize{Selected still frames from the spatially resolved recent SFH of the SGS region.  The white ellipse corresponds to the elliptical outline of the SGS itself shown in Figure \\ref{f1}.   The spatial resolution of the images is $\\sim$ 8\\arcsec, similar to that of the \\hi\\ observations.}  The movie can be seen at http://www.astro.umn.edu/$\\sim$dweisz/2574/sgs.mov .} \\label{f4} \\end{figure*}  Following the model of \\citet{che74}, one can compute the energy necessary to evacuate the \\hi\\ mass interior to the SGS to be $\\sim$ 2$\\times$10$^{53}$ erg \\citep[e.g.,][]{wal98,wal99,can05}.  Using STARURST99 \\citep{lei99}, we can independently calculate the energy associated with SF in the SGS region by inputting our SFH and assuming the same IMF as was used to measure the SFH (see \\citealt{wei08}).  We computed the total cumulative kinetic energy (i.e., stellar winds and SNe) of the SGS region over the past 500 Myr to be $\\sim$ 3$\\times$10$^{55}$ erg (Figure \\ref{f3}).   Because of our relatively coarse time resolution, we need to consider that our elevated SF periods, may be shorter bursts averaged out over a longer time.  Thus, we used STARBURST99 to model the SFH as both a series of constant episodes of SF, as well instantaneous bursts for the elevated SF episodes and found the energetics to be consistent.   At first pass, it may seem that SF generates an unreasonably large amount of energy compared to the \\hi\\ energy requirement.  However, by considering the feedback efficiency we can reconcile the apparent discrepancy.  Assuming that each SN imparts 10$^{51}$ erg in into moving the ISM we find a $\\sim$ 1\\%\\  stellar feedback efficiency,  which is on the low end of the spectrum that ranges from $\\sim$ 1 $-$ 20\\%\\ \\cite[e.g.,][]{the92, cole94, pad97, thor98}.  However, we have really computed a lower limit on the efficiency, as the \\hi\\ energy is a lower bound, and the SF energy is an upper bound.  The \\hi\\ energy was computed using an assumed smooth uniform \\hi\\ distribution inside the SGS. Given that the central SF event in the SGS was extremely powerful, it is likely that there was a higher density of gas interior to the SGS in the past \\citep[e.g.,][]{oey02}.  For example, the model of \\citet{che74} predicts that an increase in a factor of 10 in the gas density would change our feedback efficiency to $\\sim$ 10\\%.   When we consider the energy input into the ISM due to SF, only a fraction of this energy is used to move the ISM, with the rest lost to heating and radition.  Simulations suggest $\\sim$ 10$^{50}$ erg per SN is a more realistic amount of SN energy imparted into moving the ISM \\citep{cole94, thor98}, which also increases our feedback efficiency.  Finally, if the SGS has stalled expansion and broken out of the plane of the galaxy, the energy from  SF simply escapes and does not go into moving the ISM.  Despite the large uncertainties in the calculations of both energy quantities, we find that our feedback efficiency falls in the range of the values typically calculated from simulations, indicating the SF generated an appropriate amount of energy for creation of the SGS.   The SGS presents a clear case of not only sequential SF, but also triggered (i.e., causally connected) SF.  From the spatially resolved SFH movie (Figure \\ref{f4}), we see clearly that the central burst predates the SF on the rim, which is in excellent agreement with age estimates in the literature \\citep[e.g.,][]{wal98, can05, pas08}.  Establishing a causal relationship between SF events is generally non-trivial and only a handful of cases of triggered SF have been established (see \\citealt{oey05} and references therein).  In the case of the SGS, we have a recent SF episode interior to the SGS, which peaked $\\sim$ 25 Myr ago and the bulk of the recent SF on the rim is $<$ 10 Myr old, indicated by both the SF movie and the \\halpha\\ observations.  Thus, it appears that the central SF event initiated the expansion of the shell, which swept up the ISM, and when the gas on the rim becomes sufficiently dense,  secondary SF began.  As a slight twist to this scenario, we can speculate about the role of the older (200 $-$ 300 Myr), central SF event.  Because of the location of the older event and energy input into the ISM, it is possible that this older event cleared out some or all of the SGS, and the more recent event simply served to accelerate the expansion and trigger the SF around the rim.  Or perhaps the older event created one spherical shell, which merged with a younger shell to account for the ellipticity of the SGS.  To better our understanding of the role of both young and old SF events in reshaping the ISM, we will need to employ similar analysis techniques on a statistically significant sample of holes and shells in the ISM of nearby dIs."
    },
    "0812/0812.2346_arXiv.txt": {
        "abstract": "{We measure $F814W$ Surface Brightness Fluctuations (SBFs) for a sample of distant shell galaxies observed with the Advanced Camera for Survey (ACS) on board of HST.  To evaluate the distance at galaxies, theoretical SBF magnitudes for the ACS@HST filters are computed for single burst stellar populations covering a wide range of ages (t=1.5-14 Gyr) and metallicities (Z=0.008-0.04).  Using these stellar population models we provide the first $\\bar{M}_{F814W}$ versus $(F475W-F814W)_0$ calibration.  The results suggest that \\emph{optical} SBFs can be measured at $d \\geq 100 Mpc$ using high resolution spatial optical data. ",
        "introduction": "The SBF method is a powerful technique to derive distance to galaxies as far as 130 Mpc with uncertainties lower then 10\\%.  The SBFs can be evaluated for elliptical, spiral galaxies with prominent bulge and dwarf galaxies.  An interesting case is represented by shell elliptical galaxies.  Shell structures are considered robust indicator of past interactions events and the stellar population in the shell depends on the galaxy with which the merger has taken place (e.g. \\citet{Malin&Carter83}) and on the time spent since the shell structure has formed. Then, it is reasonable to expect that the presence of shells might influence the SBF signal of the galaxy. In this respect, the high quality of ACS images is crucial to allow the measurement of SBF of the galaxy, even at high distances. ",
        "conclusions": ""
    },
    "0812/0812.0482_arXiv.txt": {
        "abstract": "{ We describe the spectral classification of white dwarfs and some of the physical processes important for their understanding. In the major part of this paper we discuss the input physics and computational methods for one of the most widely used stellar atmosphere codes for white dwarfs. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.2508_arXiv.txt": {
        "abstract": "{We examine the dynamics of a self--gravitating magnetized neutron gas as a source of a  Bianchi I spacetime described by the Kasner metric. The set of Einstein-Maxwell field equations can be expressed as a dynamical system in a 4-dimensional phase space. Numerical solutions of this system reveal the emergence of a point--like singularity as the final evolution state for a large class of physically motivated initial conditions. Besides the theoretical interest of studying this source in a fully general relativistic context, the resulting idealized model could be helpful in understanding the collapse of local volume elements of a neutron gas in the critical conditions that would prevail in the center of a compact object.} ",
        "introduction": "\\label{intro}  Critical stellar configurations, such as white dwarfs, neutron stars, supermassive stars, relativistic star clusters and black holes are important astrophysical systems in which relativistic effects cannot be ignored. These astrophysical systems, denoted generically by the term ``compact objects'', can be thought of as natural laboratories to understand a wealth of phenomena relevant to theoretical and experimental physics under critical conditions. These conditions are ideally suited to propose and test theoretical models of strong magnetic fields associated with self--gravitating sources under critical conditions. In such conditions, gravitation couples with other fundamental interactions (strong, weak and electromagnetic) and from this interplay important clues of their unification could emerge.  The presence and effects of strong magnetic fields in compact objects has been studied  in the literature (see \\cite{shapiro,bocquet,Peng:2007uu,Reisenegger:2008et,Suh:2000ni,Cardall}, and references quoted therein). At a very basic level, the simplest approach is to consider various types of self--gravitating and magnetized plasmas of neutron matter under equations of state that are appropriate for the critical conditions of a compact object \\cite{salgado1,salgado2}. Since astrophysical systems exhibit, in general, angular momentum and pressure anisotropies, it is interesting to examine how these effects could influence their stability.  Following the basic known methodology \\cite{shapiro}, an of equation of state based on a nuclear ferromagnetic model (not related to electric currents) was examined in \\cite{Aurora1}, with the purpose of studying the interplay of pressure anisotropy and the magnetic field within a Newtonian framework. Hence, in the present work we consider a relativistic generalization that will allows us to examine (under specific restriction of the geometry) the evolution of this type of magnetized neutron gas under strong gravity. In particular, our aim is to provide the minimally basic qualitative and quantitative elements to address the question of whether the magnetic field can ``slow down'', or even reverse or stop, gravitational collapse, as well as the issue of the evolution to a stable configuration.   It is evident that a proper and comprehensive study of a magnetized fermion gas, as a source of a compact object, requires a spacetime with axial symmetry (or without symmetries), leading to dynamical equations that must be solved by means of hydrodynamic codes of high complexity \\cite{bocquet,Hartle1,Hartle2,Bednarek}. In previous articles \\cite{Alain2,Ulacia Rey:2006hx} we studied a self--gravitating gas of magnetized electrons, described by an appropriate equation of state, as the source of an anisotropic Bianchi-I spacetime described by the Kasner metric. While a Bianchi I model is obviously inadequate as the metric of a compact object, it is among the less complicated geometries compatible with a magnetized source.  Thus, the evolution of such a source under a much simplified Bianchi geometry is basically a toy model, but as such it can still be useful to obtain qualitative features of the sources under local critical conditions, specially in conditions near the center of the object where angular momentum plays a minor dynamical role. These qualitative results could provide a better understanding, and/or a useful approximated description, of the behavior of local internal volume elements near the center of a more realistic configuration. In this article we extend previous work on the electron gas to the case of a neutron gas. Hence, we follow a similar methodology based on re--writing the field and conservation equations for the magnetized neutron gas in the Bianchi I geometry as a dynamical system, evolving in a 4--dimensional phase space. This system is then analyzed qualitatively and numerically.  It is worthwhile remarking the basic differences between the magnetized electron gas examined in \\cite{shapiro,Alain2,Ulacia Rey:2006hx} and the neutron gas that we consider in this article. Electrons interact with a magnetic field through their charge, leading to the Landau diamagnetism characterized by a quantization effect associated with the Landau energy levels. However, neutrons interact with the magnetic field by their anomalous magnetic moment (AMM), in the context of Pauli's paramagnetism and the equations of Pauli--Dirac. Consequently, one expects the magnetic neutron interaction to be weaker, though conditions of degenerated neutron gases in compact objects are also expected to be more critical (because of the higher densities) than those of degenerated electron gases. Therefore, relativistic effects of gravity are more likely to play a dominant dynamical role in neutron gases. Still, it is important to mention that a self--gravitating and magnetized neutron gas is a simplified model of a source for a compact object, as protons and chemical equilibrium potentials should also be considered to evaluate local effects~\\cite{shapiro,salgado1,salgado2}. Nevertheless, the magnetized neutron gas already exhibits important qualitative and quantitative differences in comparison with electron gases previously examined.  The paper is organized as follows:  we derive in section II the equation of state for a magnetized neutron gas. The Einstein--Maxwell field equations for the Kasner metric and this source are derived in section III. In section IV we examine the limit of a weak magnetic field. The set of ordinary non--linear differential equations that follows from the field equations is transformed into a set of evolution equations in sections V and VI. This dynamical system is analyzed qualitatively and numerically in section VII. Our conclusions are presented in sections VIII. ",
        "conclusions": "We have used a Bianchi I model to study the evolution of a magnetized neutron gas characterized by a physically motivated and fully relativistic equation of state.  As far as we are aware, this equation of state has not been considered previously in a general relativistic context. Besides the general theoretical interest in undertaking such a study, we argue that the simplified Bianchi geometry roughly approximates a grand canonical subsystem of a magnetized neutron source in the conditions prevailing near the center of a compact object, hence our treatment can be conceived as a toy model that can be useful in understanding the local evolution of volume elements of this source under these conditions. However, a proper examination of the specific limitations of the dynamics of this toy model and/or its connection with concrete astrophysical studies of actual compact objects lies beyond the scope of the present article. As we comment further ahead, we will consider these important tasks in forthcoming articles by resorting to perturbation techniques, more elaborate numerical methods and less idealized sources.  The Einstein-Maxwell field equations for a magnetized neutron gas in the Bianchi I geometry were transformed into a set of non--linear evolution equations, which were solved numerically for generic collapsing initial conditions and analyzed qualitatively as a proper dynamical system. The results that we found are: \\begin{itemize} \\item The final state in the collapsing evolution of local volume elements is an isotropic point--like singularity. This final state occurred for a wide range of initial conditions associated with parameter values that would be typical in a compact object. \\item The magnetic field increases, but its values are always below the critical field. This result is consistent with numerical values of maximal field intensities compatible with stability in numerical studies of magnetized rotating configurations (see \\cite{bocquet}). \\item The study of the phase space associated with the dynamical equations shows that the system evolves, for $H_0<0$, to an equilibrium point, ({\\it i.e.} into a stable configuration). It is possible to introduce a temperature dependence in the equation of state. Buy doing so, the of evolution of the neutron gas could be associated with high temperature neutron sources in the context of early universe conditions in cosmological models dealing with primordial magnetic field \\cite{King:2006cy}. \\end{itemize}  It is important to remark that, unlike the dynamical study of a magnetized electron gas~\\cite{Alain2}, anisotropic ``cigar--like'' singularities did not occur for a wide range of initial conditions. Since a dynamical effect of the magnetic field under critical conditions is a final state anisotropic singularity aligned in the direction of the field, the exclusive occurrence of point--like isotropic in a magnetized neutron gas suggests that the final stage of the evolution of this gas is more intensely dominated by the focusing effect associated with strong gravity than the electron gas. This is consistent with the fact that electrons are strongly coupled to the magnetic field through their electric charge, whereas neutrons have a weaker coupling associated with their anomalous magnetic moment.  It is important to stress that the geometry of Bianchi I models in a comoving frame has stringent limitations in dealing adequately with the dynamical effects associated with a magnetic field. This is important when considering neutrons as a source in which electric charge vanishes but not the magnetic moment. By being spatially homogeneous with a 4--velocity orthogonal to flat 3--dimensional hypersurfaces of maximal symmetry, the Lorentz force is necessarily zero: $f^{b}=qu_{a}F^{a b}=0$. Also, the fact that the Bianchi I model is spatially flat makes it inadequate to examine (even as a toy model) the interplay of local collapse and the magnetic field in localized objects, as such interplay is necessarily associated with strong positive spatial curvature. However, in our case these inadequacy can be overcome (at least partially) by considering a general perturbation scheme on a Bianchi I model, in which spatial curvature and 4-acceleration are perturbative but not zero. In this case it is possible to examine the effects of the magneto--curvature coupling associated with a non--trivial Lorentz force and a nonzero deceleration parameter in the Raychaudhuri equation (see \\cite{Tsagas} for general detail). The use of such a perturbed Bianchi I model for the description of the neutron gas source considered in this article is presently under consideration in a separate article.  Besides the introduction of a perturbation scheme in a Bianchi I model, another possible improvement on the dynamical description of the source under consideration would be to consider Bianchi models I, V, VII or IX with a tilted 4--velocity, which are endowed with more degrees of dynamical freedom, including even the possibility of nonzero rotation (see \\cite{Coley:2008gh} and references quoted therein). These models would allow for a less restrictive study of the dynamical effects, reported in \\cite{Tsagas}, in which magnetic tension and gravitational collapse may present non--trivial coupling with a nonzero and non--perturbative 4--acceleration and vorticity with the magnetic field.  Finally, as we mentioned in the introduction, a magnetized gas consisting only of neutrons can be theoretically interesting but it is too idealized as a potential source for a compact object. Thus, we will consider as an extension of this work, besides the extra degrees of freedom in the dynamics (mentioned above), a gas mixture of neutrons, electrons and protons, complying with suitable balance conditions and adequate chemical potentials, in comparison with other types of equations of state \\cite{bocquet,salgado1,salgado2}. These extension of the present work are also under consideration for future articles."
    },
    "0812/0812.2813_arXiv.txt": {
        "abstract": "The complex structure of the light curves of \\emph{Swift} GRBs has made the identification of breaks, and the interpretation of the blast wave caused by the burst, more difficult than in the pre-\\emph{Swift} era. We aim to identify breaks, which are possibly hidden, and to constrain the blast wave parameters; electron energy distribution, $p$, density profile of the circumburst medium, $k$, and the continued energy injection index, $q$. We do so by comparing the observed multi-wavelength light curves and X-ray spectra of our sample to the predictions of the blast wave model. We can successfully interpret all of the bursts in our sample of 10, except two, within this framework and we can estimate, with confidence, the electron energy distribution index for 6 of the sample. Furthermore  we identify jet breaks in a number of the bursts. A statistical analysis of the distribution of $p$ reveals that, even in the most conservative case of least scatter, the values are not consistent with a single, universal value.  The values of $k$ suggest that the circumburst density profiles are not drawn from only one of the constant density or wind-like media populations. ",
        "introduction": "\\label{breaks:introduction}  The afterglow emission of Gamma-Ray Bursts (GRBs) is generally well described by the blast wave model \\citep{rees1992:MNRAS258,meszaros1997:ApJ476,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, progenitor models, 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 approximately one 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} it has become clear that this model for GRBs cannot, in its current form, explain the full complexity of observed light curve features and the lack of observed achromatic temporal breaks. The unexpected features detected, such as steep decays, plateau phases (e.g., \\citealt{tagliaferri2005:Natur436,nousek2006:ApJ642,obrien2006:ApJ647}) and a large number of X-ray flares (e.g., \\citealt{burrows2007:RSPT365,chincarini2007:ApJ671,falcone2007:ApJ671}) have revealed the complexity of these sources up to about one day since the initial event, which is yet to be fully understood. These superimposed features also make it difficult to measure the underlying power-law features on which the blast wave model is based, and may lead to misinterpretations of the afterglows.  Achromatic temporal breaks are not observed in the majority of bursts, 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{liang2008:ApJ675}). Previous comprehensive studies of \\emph{Swift} light curves in multiple wavelengths \\citep{liang2007:ApJ670,liang2008:ApJ675} have attempted to identify achromatic temporal breaks due either to jet breaks, or the cessation of continued energy injection (i.e. end of the plateau phase), mainly as they concern the energy requirement of bursts.  Identification of jet breaks is also important as a test of the blast wave physics as it represents the time at which sideways spreading of the jet becomes important and the edges of the jet become visible. The post jet break phase also gives a direct measurement of the electron energy distribution, $p$, as the temporal decay, at all wavelengths, during this phase should have a value equal to that of $p$. However, due to differences between broken power-law fits and the underlying parameters \\citep{johannesson2006:ApJ640}, this method may systematically underestimate the true value of $p$. The value of $p$ may also be measured from the spectral information of the optical or X-ray bands; the accuracy of this method is then dependent on the quality of the spectra used. A study of \\emph{BeppoSAX} bursts used broadband optical to X-ray Spectral Energy Distributions (SEDs) to accurately constrain the value of $p$ via the spectral and temporal parameters of both bands, and a comparison to the blast wave model \\citep{starling2008:ApJ672}. The same authors also use this method to estimate the density profile of the circumburst medium, $k$, as the temporal slope may depend on this value as defined by $\\rho \\propto r^{-k}$; a constant density medium, $k=0$ causing a more shallow decay than a wind like medium, $k=2$. Similar interpretations of the temporal and spectral information may also be used to estimate the effect, if any, that continued energy injection has on the light curves. Continued energy injection is used to explain the shallow decay phase observed in many \\emph{Swift} GRBs and is usually assumed to take the form of $E \\propto t^{q}$ \\citep{nousek2006:ApJ642}. The additional energy can be due to slower shells of matter catching up with to blast wave, or a continued activity of the central engine. In either case, the addition of the extra energy to the blast wave causes the flux to drop off less rapidly than it would otherwise, and is observed as a shallower than expected decay. The cessation of the energy injection is marked by a temporal break to the normal decay phase as predicted by the blast wave model which implies, if it is due to continued activity of the central engine, that the engine must be active for up to about a day. Such late activity of the central engine has been observed in the form of episodic flares  \\citep{curran2008:A&A487} but such late, smooth and continuous activity has yet to be confirmed.    In this paper we interpret a sample of afterglows of \\emph{Swift} GRBs, to constrain the blast wave parameters: electron energy distribution, $p$, density profile of the circumburst medium, $k$, and the continued energy injection index, $q$. In \\S\\ref{breaks:observations} we introduce our sample and the method of reduction, while in \\S\\ref{breaks:results} we present the results of our temporal and spectral analysis. In \\S\\ref{breaks:discussion} we interpret the results in the context of the blast wave model and discuss their implications in the overall context of GRB observations. We summarise our findings in \\S\\ref{breaks:conclusion}. Throughout, we use the convention that a power-law flux is given as $F \\propto t^{-\\alpha} \\nu^{-\\beta}$ where $\\alpha$ is the temporal decay index and $\\beta$ is the spectral index.  All uncertainties are quoted at the $1\\sigma$ confidence level. ",
        "conclusions": "\\label{breaks:conclusion}    Throughout this paper we have applied the blast wave model \\citep{rees1992:MNRAS258,meszaros1998:ApJ499}, assuming on-axis viewing, a standard jet structure and no evolution of the microphysical parameters, to a selection of GRBs with well sampled temporal and spectral data. We attempted to constrain three parameters of interest and to test the validity of the blast wave model for light curves observed by \\emph{Swift}, the complexity of which has made the identification of breaks and the interpretation of the blast wave more difficult than in the pre-\\emph{Swift} era. We find that the majority of the afterglows are well described within the frame work of the blast wave model and that the parameters derived are consistent with those values found by previous authors.   After an inspection of all the \\emph{Swift} X-ray light curves and a comparison with the optical data available in the literature, we identified a sample of 10 bursts that had enough data to well constrain the optical and X-ray temporal indices, and the X-ray spectral indices, of the afterglows. We analysed these data, fitting power-law decays to the light curves in an attempt to identify possible breaks; taking into account the possibility of hidden breaks and the effect of host galaxy contribution to the optical, where possible.  In the case of well constrained X-ray spectra, we show that $p$ can be estimated from the X-ray spectra alone, with the caveat that it will give two possible values of $p$; these must be differentiated by either broadband SEDs or a comparison with the predictions of the blast wave model, or both. We compare the observed light curves and spectra of each burst in our sample to those values expected from the blast wave model and find that we can successfully interpret all bursts, except two, within this frame work; these interpretations require  `hidden' X-ray breaks in GRB\\,060206 and GRB\\,061126. Furthermore, we identify, reasonably unambiguously, jet breaks in 5 afterglows out of our sample of 10, including a jet break in GRB\\,060729 that was not identified in previous analysis of the light curve.  After interpretation within the blast wave model,  we are able to confidently estimate the electron energy distribution index, $p$, for 6 of the bursts in our sample. A statistical analysis of the distribution of $p$ reveals that, even in the most conservative case of least scatter, the values are not consistent with a single, universal value suggested by some studies; this has important implications for theoretical particle acceleration studies. In a number of cases, we are also able to obtain values for the circumburst density profile index, $k$, and the index of continued energy injection, $q$. The calculated values of $q$ are consistent with both suggested sources of continued energy injection and a much larger sample of afterglows will be required to constrain the sample properties of $q$. The values of $k$, consistent with previous works on the matter, suggest that the circumburst density profiles are not drawn from only one of the constant density or wind-like media populations."
    },
    "0812/0812.4574_arXiv.txt": {
        "abstract": "We investigate the effects of neutrino heating and $\\alpha$-particle recombination on the hydrodynamics of core-collapse supernovae.  Our focus is on the non-linear dynamics of the shock wave that forms in the collapse, and the assembly of positive energy material below it.  To this end, we perform time-dependent hydrodynamic simulations with FLASH2.5 in spherical and axial symmetry. These generalize our previous calculations by allowing for bulk neutrino heating and for nuclear statistical equilibrium between $n$, $p$ and $\\alpha$. The heating rate is freely tunable, as is the starting radius of the shock relative to the recombination radius of $\\alpha$-particles.  An explosion in spherical symmetry involves the excitation of an overstable mode, which may be viewed as the $\\ell = 0$ version of the `Standing Accretion Shock Instability'. In two-dimensional simulations, non-spherical deformations of the shock are driven by plumes of material with positive Bernoulli parameter, which are concentrated well outside the zone of strong neutrino heating.  The non-spherical modes of the shock reach a large amplitude only when the heating rate is also high enough to excite convection below the shock. The critical heating rate that causes an explosion depends sensitively on the initial position of the shock relative to the recombination radius.  Weaker heating is required to drive an explosion in two dimensions than in one, but the difference also depends on the size of the shock.  Forcing the infalling heavy nuclei to break up into $n$ and $p$ below the shock only causes a slight increase in the critical heating rate, except when the shock starts out at a large radius. This shows that heating by neutrinos (or some other mechanism) must play a significant role in pushing the shock far enough out that recombination heating takes over. ",
        "introduction": "Although tremendous progress has been made on the mechanism of core-collapse supernovae in recent years, we still do not have a clear picture of a robust path to an explosion in stars that form iron cores. Heating by the absorption of electron-type neutrinos significantly modifies the settling flow below the bounce shock but -- in spite of early positive results \\citep{bethe85} -- explosions are obtained in spherical collapse calculations only if the progenitor star is lighter than about 10-12 $M_\\sun$ \\citep{kitaura06}. More massive stars fail to explode in spherical symmetry \\citep{liebendoerfer01}.  Two-dimensional collapse calculations show strong deformations of the shock and convective motions below it \\citep{burrows95,janka96,buras06a,buras06b,burrows06a,burrows07,marek07}. It has long been recognized that convection increases the residency time of settling material in the zone of strong neutrino heating \\citep{herant92}.  It is also becoming clear that multidimensional explosions require the assembly of a smaller amount of material with positive energy, but the details of how this happens remain murky.  An early treatment of shock breakout by \\citet{bethe97} focused on the strong gradient in the ram pressure of the infalling material, but implicitly assumed that the shocked material had already gained positive energy. If large-scale density inhomogeneities are present below the shock, they will trigger a finite-amplitude, dipolar instability, thereby allowing accretion to continue simultaneously with the expansion of positive-energy fluid (Thompson 2000). The accretion shock is also capable of a dipolar oscillation which leads, above a critical amplitude, to a bifurcation between freshly infalling material and material shocked at earlier times \\citep{BM03}. Such an oscillation is easily excited in a spherical flow composed of a zero-energy, polytropic fluid (\\citealt{BM06}) via a linear feedback between ingoing vortex and entropy waves and outgoing sound waves \\citep{F07}.  It can also be excited indirectly by neutrino heating, which if sufficiently strong will drive large-scale buoyant motions below the shock (\\citealt{herant94}; \\citealt{foglizzo06}). For relatively weak heating, the dipolar oscillation can decrease the damping effect of advection and trigger convective motions that would otherwise be suppressed \\citep{foglizzo06,scheck08}.  A significant sink of thermal energy in the accretion flow arises from the dissociation of heavy nuclei. The heavy elements that flow through the shock are broken up into $\\alpha$-particles and nucleons when exposed to the high temperature ($>1$~MeV) of the postshock region. The Bernoulli parameter $b$ of the shocked fluid then becomes substantially negative. A significant fraction of this dissociation energy can be recovered if nucleons recombine into $\\alpha$-particles \\citep{bethe96}. But for this to occur, a decrease in the temperature is required and thus the shock must expand significantly beyond the radius at which it typically stalls ($\\sim 100-150$~km).  One of the primary goals here, and in a previous paper \\citep{FT08a} [hereafter Paper I], is to gauge the relative importance of these effects in setting the stage for a successful explosion.  The persistence and amplitude of a dipolar oscillation can only be reliably measured in fully three-dimensional simulations (there are preliminary indications that it is less prominent in three spatial dimensions than in axial symmetry; \\citealt{iwakami08}). On the other hand, the interplay between $\\alpha$-particle recombination and hydrodynamical instabilities has received little attention.  Although recombination is certainly present in previous numerical studies which employ finite-temperature equations of state (EOSs), it should be kept in mind that considerable uncertainties in the EOS remain at supranuclear densities.  A softening or hardening of the EOS feeds back on the position of the shock for a given pre-collapse stellar model \\citep{marek07}.  The parametric study of the critical neutrino luminosity by \\citet{murphy08} is based on a single EOS;  they obtain explosions in which the shock seems to break out from nearly the same radial position at $\\sim 250-300$~km. Variations in the density profile of the progenitor star will similarly modify the position of the shock, the concentration of $\\alpha$-particles below it, and the critical neutrino luminosity for an explosion.  In this paper, we study the interplay between non-spherical shock oscillations, neutrino heating, and $\\alpha$-particle recombination, when the heating rate is pushed high enough for an explosion to occur. Our focus is on the stalled shock phase, between $\\sim 100$ ms and 1 s after bounce. In a one-dimensional calculation, the accretion flow reaches a quasi-steady state during this interval, and the shock gradually recedes (e.g. \\citealt{liebendoerfer01,buras06a}).  Our approach is to introduce EOSs of increasing complexity into one- and two-dimensional, time-dependent hydrodynamic simulations. To this end, we use the code FLASH2.5 \\citep{fryxell00}, which is well tested in problems involving nuclear energy release in compressible fluids \\citep{calder02}. We adopt a steady state model as our initial condition, and a constant mass accretion rate, neutrino luminosity, and fixed inner boundary. The steady-state approximation to the stalled shock phase was first introduced by \\citet{burrows93}, and has recently been used by \\citet{ohnishi06} to study the non-linear development of the shocked flow with a semi-realistic equation of state and neutrino heating. In Paper I we modeled the accretion flow as a polytropic fluid, from which a fixed dissociation energy is removed immediately below the shock, and allowed for neutrino cooling but not heating.  Here we generalize this model to allow both for heating, and for nuclear statistical equilibrium (NSE) between neutrons, protons, and $\\alpha$-particles. Heating is introduced in a simple, parametrized way, without any attempt at simulating neutrino transport. The nuclear abundances are calculated as a function of pressure $p$ and density $\\rho$, using a complete finite-temperature, partially degenerate EOS. Our model for the shocked material retains one significant simplification from Paper I:  we do not allow the electron fraction $Y_e$ to vary with position below the accretion shock. Here there are two competing effects: electron captures tend to reduce $Y_e$, whereas absorption of $\\nu_e$ and $\\bar\\nu_e$ tends to drive high-entropy material below the shock toward $Y_e \\simeq 0.5$. Since we are interested especially in the dynamics of this high-entropy material, we set $Y_e = 0.5$ when evaluating the $\\alpha$-particle abundance. To obtain a realistic density profile, we continue to approximate the internal energy of the fluid as that of a polytropic fluid with a fixed index $\\gamma = {4\\over 3}$. In reality, the equation of state between the neutrinosphere and the shock depends in a complicated way on the degeneracy of the electrons and the effects of electron captures. The consequences of introducing these additional degrees of freedom will be examined in future work.  In spite of these simplifications, our results already show many similarities with more elaborate collapse calculations. Spherical explosions are due to a global instability resembling the one-dimensional Standing Accretion Shock Instability (SASI), but modified by heating. As in Paper I, we find that the period of the $\\ell = 0$ mode remains close to twice the post-shock advection time. Strong deformations of the shock in two-dimensional runs are driven by material with positive Bernoulli parameter, which generally resides outside the radius $r_\\alpha$ where the gravitational binding energy of an $\\alpha$-particle is equal to its nuclear binding energy.  The recombination of $\\alpha$-particles plays a major role in creating this positive-energy material, but for this to happen the shock must be pushed beyond $\\sim 200$ km from the neutronized core.  In this paper, we consider only neutrino heating as the impetus for the initial expansion of the shock, rather than more exotic effects such as rotation or magnetic fields. We find that the critical heating rate is a strong function of the initial position of the shock with respect to $r_\\alpha$, which implies that a much higher neutrino luminosity is needed to revive a shock that stalls well inside $r_\\alpha$. The difference in the critical heating rate between one- and two-dimensional simulations also depends on the size of the shock, and thence on the structure of the forming neutron star.  We find that the shock develops a dipolar oscillation with a large amplitude only when the heating rate is also high enough to trigger a strong convective instability. We therefore surmise that buoyant motions driven by neutrino heating play a major role in driving the dipolar oscillations that are seen in more complete simulations of core collapse. Some evidence is found that acoustic wave emission by the convective motions also can play a role. We investigate the possibility of a heat engine within the gain layer (the region with a net excess of neutrino heating over cooling). We find that most of the heat deposition by neutrinos is concentrated in lateral flows at the base of the prominent convective cells.  At the threshold for an explosion, neutrino heating plays a key role in pushing material to positive Bernoulli parameter, but only if the shock starts well inside $r_\\alpha$.  The plan of the paper is as follows.  Section \\ref{s:model} presents our numerical setup, treatment of heating, cooling, and nuclear dissociation, and outlines the sequences of models. Sections \\ref{s:1d} and \\ref{s:2d} show results from one- and two-dimensional simulations, respectively.  We focus on the relative effectiveness of neutrino heating and $\\alpha$-particle recombination in driving an explosion, and the relation between Bernoulli parameter and large-scale deformations of the shock.  The critical heating rate for an explosion is analyzed in \\S\\ref{s:thresholds}, and the competition between advective-acoustic feedback and convective instability is discussed in \\S\\ref{s:turbulence}.  We summarize our findings in \\S\\ref{s:summary}. The appendices contain details about our EOS and the numerical setup. ",
        "conclusions": "\\label{s:summary}  We have investigated the effects of $\\alpha$-particle recombination and neutrino heating on the hydrodynamics of core-collapse supernovae, when the heating rate is pushed high enough to reach the threshold for an explosion.  The effect of dimensionality has been probed by comparing one- and two-dimensional time-dependent hydrodynamic calculations. Our main results can be summarized as follows:  \\vskip .1in \\noindent 1. --  The critical heating parameter that yields an explosion depends sensitively on the starting position of the shock relative to $r_\\alpha$. This means that the critical neutrino luminosity depends sensitively on the stall radius of the shock and, in turn, on the core structure of the progenitor star and the density profile in the forming neutron star. Within the framework explored in this paper, we find two extreme types of explosion.  In the first, neutrino heating does most of the work, with a significant final boost from $\\alpha$-particle recombination. In the second, neutrino heating is generally less important at promoting material below the shock to positive energies.  \\noindent 2. --  During the final stages of an explosion, the heat released by $\\alpha$-particle recombination is comparable to the work done by adiabatic expansion.  This heat is concentrated in material that has previously been heated by neutrinos. Significantly more energy is lost through $\\alpha$-particle dissociation in fresh downflows, so that nuclear dissociation remains on balance an energy sink within the accretion flow.  \\noindent 3. -- The large-amplitude oscillations that are seen in one-dimensional runs near an explosion are the consequence of the $\\ell=0$ SASI as modified by heating.  In contrast with the $\\ell = 1,2$ modes of a laminar accretion flow, the period of these oscillations is close to twice the post-shock advection time. The critical heating rate for an explosion (assuming constant mass accretion rate, neutrino luminosity, and inner boundary) corresponds to neutral stability for the $\\ell = 0$ mode.  \\noindent 4. -- The critical heating parameter $H_{\\rm cr}$ for an explosion is generally lower in two dimensions than in one, but the difference becomes smaller as the starting radius of the shock approaches $r_\\alpha$. The difference depends somewhat on the ratio $r_*/\\rs0$ and thus on the cooling efficiency and equation of state.  \\noindent 5. -- Non-spherical deformations of the shock are tied to the formation of large-scale plumes of material with positive energy.  Our two-dimensional explosions with a super-critical heating rate involve a large-scale convective instability that relies on the accumulation of vorticity on the largest spatial scales.  Volume-filling convective cells are apparent in a time-averaged sense. Transient heating events create positive-energy material that accumulates in between the convective cells and the shock. A significant fraction of the heating occurs in horizontal flows at the base of the convective cells, which are fed by a dominant equatorial accretion plume. If the heating parameter is large enough, this results in an amplifying cycle and explosion.  \\noindent 6. -- The amplitude of the $\\ell = 1$ and 2 modes correlates strongly with the value of the heating parameter, and is coupled to the appearance of vigorous neutrino-driven convection below the shock. In agreement with the work of \\citet{foglizzo06} and \\citet{scheck08}, we find that $\\chi\\approx 3$ marks the transition from a strong linear instability in a nearly laminar flow below the shock, to a volume-filling convective instability.  In all of our simulations, the threshold for explosion lies well within the latter regime. This highlights a basic difference between one- and two-dimensional explosions: the mechanism is fundamentally non-linear in two dimensions.  \\noindent 7. -- We have explored essentially  one ratio of cooling radius to shock radius, namely $r_*/r_{\\rm s0} = 0.4$ at zero heating (corresponding to $r_*/r_{\\rm s} \\sim 0.2$ near the threshold for an explosion).  The growth of the $\\ell = 1$ SASI mode is strongest for this particular aspect ratio when $\\varepsilon = 0$ (see Figure 12 of Paper I). As dissociation is introduced into the flow, we found that the peak growth rate moves to larger values of $r_*/r_{\\rm s0}$. On the other hand, detailed collapse calculations indicate ratios of neutrinosphere radius to shock radius that are even smaller than $\\sim 0.2$ following $\\sim 100$ ms after collapse (e.g. \\citealt{marek07}). We conclude that the $l=1$ SASI mode is not being artificially suppressed by our choice of initial shock size.  \\noindent 8. -- Vortical motions with a Mach number $\\sim 0.3$-0.5 first appear at the onset of convective instability around the radius of maximal neutrino heating, but before the dipolar mode of the shock reaches its limiting amplitude.  These vortices are a source of acoustic waves, which have a similar period to the large-scale oscillation of the shock. Near the threshold for an explosion, the turbulence in the gain region becomes supersonic, as the existence of widespread secondary shocks attests. These shocks convert turbulent kinetic energy to internal energy, increasing the effective heating rate.  \\vskip .1in There are at least two reasons why explosions by the mechanism investigated here may be more difficult in fully three-dimensional simulations. First, the existence of more degrees of freedom for the low-order modes of the shock in three dimensions implies that the amplitude of individual shock oscillations is lower. As a result, it is more difficult for the shock to extend out to the radius where $\\alpha$-particle recombination gives it the final push.  Second, an axisymmetric explosion that is driven by neutrino heating involves the accumulation of vorticity on the largest spatial scales, an effect that is special to two dimensions.  A full resolution of these issues is possible only with high-resolution three-dimensional simulations."
    },
    "0812/0812.2328_arXiv.txt": {
        "abstract": "{At the Very Large Telescope Interferometer, the purpose of the fringe-tracker FINITO is to stabilize the optical path differences between the beams, allowing longer integration times on the scientific instruments AMBER and MIDI.} {Our goal is to demonstrate the potential of FINITO for providing $H$-band interferometric visibilities, simultaneously and in addition to its normal fringe-tracking role.} {We use data obtained during the commissioning of the Reflective Memory Network Recorder at the Paranal observatory. This device has permitted the first recording of all relevant real-time data needed for a proper data-reduction.} {We show that post-processing the FINITO data allows valuable scientific visibilities to be measured. Over the several hours of our engineering experiment, the intrinsic transfer function is stable at the level of $\\pm2$\\%. Such stability would lead to robust measurements of science stars even without the observation of a calibration star within a short period of time. We briefly discuss the current limitations and the potential improvements.} {} ",
        "introduction": "The purpose of the FINITO facility at the Very Large Telescope Interferometer \\citep[VLTI, see][]{Scholler-2006jul} is to compensate for the random optical delay (OPD) introduced by atmospheric turbulence. If not corrected, such perturbation would make the interferometric fringes jitter on the detector, preventing the use of long integration time on scientific instruments. FINITO was offered to the astronomical community in September 2007, and since then several AMBER \\citep{Petrov-2007mar} programs have already benefited from the enhanced data quality produced by fringe-tracking \\citep{LeBouquin-2008apr,Wittkowski-2008feb}.  Nevertheless, FINITO can potentially tackle another important issue at VLTI: measuring robust, absolutely calibrated, interferometric visibilities in the near-IR. Indeed, AMBER has been designed to be a highly rewarding spectro-interferometer \\citep{Petrov-2007mar,Weigelt-2007mar,Domiciano-2007mar,Meilland-2007mar}, but this has required technical trade-offs. It typically reads out more slowly than an atmospheric coherence time, and so is critically sensitive to fringe stabilization, and hence to seeing and mechanical vibrations. This issue is exacerbated by VLTI service-mode only providing a few calibration stars per night, which is insufficient for properly sampling the time variations of the instrumental response (called transfer function). All together, this makes the \\emph{absolute} calibration of AMBER fringe contrast difficult \\citep{Millour-2007may}.  We have recently carried out the necessary developments to make FINITO able to provide visibility measurements. No hardware modifications were needed as FINITO already interferometrically combined the light of 2 or 3 telescopes to stabilize the OPD. The main issue was the high frame rate of the system running generally at more than 1kHz. This was overcome by the commissioning of the so-called VLTI \\emph{Reflective Memory Network Recorder} (RMNrec), which listens to the communications between several systems of VLTI and is able to store the real-time data into proper FITS files.  In this paper, we investigate the possibility of post-processing these real-time data to measure interferometric visibilities and to properly calibrate them. In Sect.~\\ref{sec:observationsAndDataReduction} we summarize the FINITO and RMNrec instrumental setup, the test observations made, and the simple data reduction used for this demonstration. In Sect.~\\ref{sec:resultsAndDiscussion}, as a classical test, we compare the measured angular diameter of some well-known stars with previously published values. We discuss the results, the accuracy obtained, and the perspectives for reaching higher precision in the future. The paper ends with brief conclusions and perspectives. ",
        "conclusions": "We have demonstrated the possibility of post-processing the real-time fringe-tracking data of the Very Large Telescope Interferometer for astronomical purposes. We proposed a simple strategy for reducing the data, based on scan selection and power integration. We were able to recover the $H$-band diameters of four already-known giant stars. The good general agreement proved the capability of the FINITO/OPDC/RMNrec facility to provide meaningful interferometric visibility measurements.  Interestingly, the transfer function level is dominated by the integration time. It is stable over several hours, with possible significant changes from one day to the next. We have not been able to find any obvious correlation with the atmospheric seeing and fringe-tracking performances, even if we agree that our dataset does not allow us to draw definite conclusions. Post-processing the FINITO/OPDC data therefore provides a rough, but robust, estimation of the $H$-band absolute visibility even without the observation of a calibration star close in time. This complements the scientific observations very well, generally done in another spectral range (near-IR for AMBER, and $N$-band for MIDI) and with higher spectral resolution (up to $10\\,000$ for AMBER), but with only poor absolute calibration of the continuum visibility. We emphasize that this additional, crucial measure is obtained simultaneously and without any additional overhead time.  This demonstration is also promising in the context of precision interferometry at the VLTI. We were able to reach accuracy of $\\pm{}2\\%$ on the raw visibilities with a simple data reduction method. We showed that the overall accuracy is most probably limited by systematic errors. Statistical errors only about $\\pm0.1\\%$ for 30s integration on $\\mathrm{H}=2^{\\rm m}$ unresolved targets with the 1.8m Auxiliary Telescopes.  Based on the success of this demonstration, we now plan to routinely log and archive the VLTI fringe-tracking data. In this context, we emphasize the need for a dedicated public data reduction pipeline."
    },
    "0812/0812.3781_arXiv.txt": {
        "abstract": "$\\gamma$-ray astronomy has produced for several years now sky maps for low photon statistics, non-negligible background and comparatively poor angular resolution. Quantifying the significance of spatial features remains difficult. Besides, spectrum extraction requires regions with large statistics while maps in energy bands allow only qualitative interpretation.\\\\ The two main competing mechanisms in the VHE domain are the Inverse-Compton emission from accelerated electrons radiating through synchrotron in the X-ray domain and the interactions between accelerated hadrons and the surrounding medium, leading to the production and subsequent decay of $\\pi^{0}$ mesons. The spectrum of the VHE emission from leptons is predicted to steepen with increasing distance from the acceleration zone, owing to synchrotron losses (i.e. cooled population). It would remain approximately constant for hadrons.  Ideally, spectro-imaging analysis would have the same spatial scale in the TeV and X-ray domains, to distinguish the local emission mechanisms. More realistically, we investigate here the possibility of improving upon the currently published HESS results by using more sophisticated tools. ",
        "introduction": "The most simple procedure to produce hardness maps relies on bins of fixed size, defined independently of the source morphology. The bin size is a compromise between increasing the number of excess events per bin and keeping small-scale structures. This guarantees the absence of any internal artefact but is not adapted in case of large flux variations. \\newline {\\bf Core equations}\\\\ \\begin{eqnarray*} E \\, &= &N_{\\mathrm{on}} - \\frac{Exposure_{on}}{Exposure_{off}} \\times N_{\\mathrm{off}} \\quad \\mathrm{C}=\\frac{N_{\\mathrm{high}} - N_{\\mathrm{low}}}{N_{\\mathrm{high}} + N_{\\mathrm{low}}} \\\\ S_C &= &\\frac{S_{low} \\times S_{high}}{\\sqrt{S_{low}^2 + S_{high}^2}} \\end{eqnarray*} \\noindent For $\\gamma$-like ($\\mathrm{on}$) and background ($\\mathrm{off}$) events in two separate high and low energy bands, these formula define the excess $E$~\\cite{berge}, the contrast $C$ (used here instead of the immediate ratio of excesses) between the bands and its significance $S_C$, combining the significance in each band, given by the Li\\&Ma formula~\\cite{lima}. The selection of the energy bands, starting at the energy threshold for the observation, aims at balanced counts of excess events. \\newline {\\bf Adaptive binning}\\\\ In the Weighted Voronoi Tessellation (WVT) adaptive binning algorithm \\cite{wvtpaper}, the core procedure assigns pixels to non-overlapping bins, whose shape is maintained close to circular. From a given position $P_{start}$, neighbouring pixels are accreted until the requested significance or maximum size is reached. The assignment of pixels is repeated to reduce the spread of bin significance and shapes. Identified point sources are treated here as fixed bins, to avoid contamination by the background as well as smearing out of the source. To make the whole algorithm independent of the start position, we loop over random $P_{start}$, keeping the average of the re-binned maps. This procedure is illustrated in Fig.~\\ref{figure:methods2}. The target significance is 7 $\\sigma$, with an upper limit on bin radius of 0.12$^{\\circ}$. \\newline {\\bf Adaptive smoothing}\\\\ Adaptive smoothing methods, like ASMOOTH~\\cite{asmooth} algorithm (CSMOOTH in the CIAO package \\cite{ciao}), adjust the scale of a smoothing kernel to maintain a constant significance. For a given scale, the image is smoothed and pixels reaching the desired significance are added to the output image and removed from the input one. Remaining pixels are treated with progressively larger scales. To build the contrast map, the two energy bands are treated separately, see~\\cite{asmooth} for details, with a 9 $\\sigma$ target here. \\newline {\\bf Photon index}\\\\ From the contrast, we estimate the photon index and flux normalization, assuming a power-law spectrum. A hardness ratio derived from excess counts instead of fluxes is affected by the instrument response. To correct for this, the effective area of the instrument is integrated in the energy bands, accounting for zenith angles and offsets from the observations.  \\begin{figure}[ht!] \\begin{minipage}[c][4.6cm][c]{0.3\\textwidth} \\centering \\includegraphics[height=4.5cm,angle=0]{SHELLLEPTON_High9_slices.eps} \\end{minipage} \\begin{minipage}[c][4.6cm][c]{0.3\\textwidth} \\centering \\includegraphics[height=4.5cm,angle=0]{Reg-binning-Vignette.eps} \\end{minipage} \\begin{minipage}[c][4.6cm][c]{0.3\\textwidth} \\centering \\includegraphics[height=4.5cm,angle=0]{ASMOOTH-Vignette.eps} \\end{minipage} \\caption{Applying different methods to the same count map. From {\\it left} to {\\it right}, fixed-scale smoothing, re-binning, ASMOOTH. The contours are for significance, from 3 to 7 $\\sigma$ or 5 to 9$\\sigma$ for ASMOOTH). \\label{figure:methods1} } \\end{figure} \\begin{figure}[h!] \\begin{minipage}[c][4.6cm][c]{0.3\\textwidth} \\centering \\includegraphics[height=4.5cm,angle=0]{WVT-simple-Vignette.eps} \\end{minipage} \\begin{minipage}[c][4.6cm][c]{0.3\\textwidth} \\centering \\includegraphics[height=4.5cm,angle=0]{WVT-averaged-Vignette.eps} \\end{minipage} \\begin{minipage}[c][4.6cm][c]{0.3\\textwidth} \\centering \\includegraphics[height=4.5cm,angle=0]{WVT-size-Vignette.eps} \\end{minipage} \\caption{The WVT algorithm. From {\\it left} to {\\it right}, the original WVT, the adapted WVT and the associated average bin sizes. Significance contours from 3 to 7 $\\sigma$. \\label{figure:methods2} } \\end{figure}  \\begin{figure}[h!] \\includegraphics[width=16cm,angle=0]{Leptonic_scenario-Slices.eps} \\caption{Slices for model I, with the geometry of Figure~\\ref{figure:methods1}{\\it left}. In black is shown the model. Green, blue and pink are respectively the fixed binning (0.04$^{\\circ}$ integration radius), WVT and ASMOOTH algorithms, with 1 $\\sigma$ error bands. } \\label{figure:slicesL} \\end{figure}  The acceptance profiles are built for the same energy bands from regions in the sky devoid of $\\gamma$-ray source. For each position on the map, the spectral power-law index is varied until the hardness ratio agrees with the data. The flux is then scaled to the excess. \\newline {\\bf Disclaimer}\\\\ Several important issues are not investigated here, notably the variation of the PSF with energy (smaller at higher energy) and its effect on the actual spatial resolution, nor the inclusion of trial factors in the significance.  \\begin{figure}[ht!] \\includegraphics[width=16cm,angle=0]{Hadronic_scenario-Slices.eps} \\caption{Slices for model II, following the conventions of Fig.\\ref{figure:slicesL}. } \\label{figure:slicesH} \\end{figure} ",
        "conclusions": "The methods discussed here give results compatible with HESS publication. They may help exploring the morphology of extended VHE sources, although as illustrated in a simulation, the uncertainties might prove too large for detailed studies.    \\begin{theacknowledgments} We wish to thank the HESS collaboration for allowing us to use its data in this work. \\end{theacknowledgments}"
    },
    "0812/0812.3748_arXiv.txt": {
        "abstract": "We describe a method for incorporating ambipolar diffusion in the strong coupling approximation into a multidimensional magnetohydrodynamics code based on the total variation diminishing scheme. Contributions from ambipolar diffusion terms are included by explicit finite difference operators in a fully unsplit way, maintaining second order accuracy. The divergence-free condition of magnetic fields is exactly ensured at all times by a flux-interpolated constrained transport scheme. The super time stepping method is used to accelerate the timestep in high resolution calculations and/or in strong ambipolar diffusion. We perform two test problems, the steady-state oblique C-type shocks and the decay of Alfv\\'en waves, confirming the accuracy and robustness of our numerical approach. Results from the simulations of the compressible MHD turbulence with ambipolar diffusion show the flexibility of our method as well as its ability to follow complex MHD flows in the presence of ambipolar diffusion. These simulations show that the dissipation rate of MHD turbulence is strongly affected by the strength of ambipolar diffusion. ",
        "introduction": "In galactic molecular clouds, ambipolar diffusion, which arises in partially ionized plasmas, is a key ingredient of the mechanism of star formation \\citep[e.g.,][]{mes56,mou76,shu87}. In the central portions of molecular clouds, the molecular gas is dense enough that recombination is nearly total, so that very low fractions ($\\lesssim10^{-7}$) of the gas remains ionized while the rest of the gas is neutral \\citep{mye85}. This small residual ionization is usually attributed to cosmic rays, which can penetrate nearly all clouds. Ambipolar diffusion causes the relative drift of ions coupled to the magnetic field and neutrals in the molecular cloud cores and so it enable the cloud cores to collapse gravitationally.  Star formation assisted by ambipolar diffusion has been studied extensively in the context of magnetically subcritical or supercritical models \\citep[see e.g.,][]{mou99,des01,bas04,tas07}. In a current paradigm of star formation, magnetically supported molecular cloud cores must lose magnetic support through the action of ambipolar diffusion so that star formation can take place. Recent works have focused on the role of turbulence in the formation of protostellar cores \\citep[e.g.,][]{nak08}. Including the effect of turbulence on the mechanism of ambipolar diffusion can enhance the ambipolar diffusion rate \\citep{fat02,zwe02,hei04}, so that the ambipolar diffusion timescale is significantly shorter than that estimated for a similar, but quiescent, medium. Using three-dimensional numerical simulations, \\citet{ois06} and \\citet{li08} have studied the properties of turbulence with ambipolar diffusion in a two-fluid approximation, while \\citet{pad00} have investigated the heating through ambipolar diffusion in turbulent molecular clouds using a single-fluid approximation.  Shock waves in molecular clouds spread into steady-state continuous shocks, or C-type shocks, through ambipolar diffusion. If the shock speed is slower than the ion Alfv\\'en speed but faster than the neutral sound speed, the ions coupled to magnetic fields drag the neutrals into the postshock region, producing a continuous structure \\citep{dra80}. These steady-state C-type shocks can, however, be unstable on a short enough timescale to be of astrophysical interest. \\citet{war91} showed that if the magnetic field lines are perturbed slightly and ions collect in the magnetic valleys, the ion-neutral friction may overcome the magnetic forces in the shock front and derive an exponentially growing instability.  Numerical treatments of ambipolar diffusion have been commonly derived from ideal magnetohydrodynamic (MHD) models. Extensive numerical methods including ambipolar diffusion have been proposed in the study of the dynamics of partially ionized plasmas within the frame of single or two fluid models \\citep{tot94,mac95,mac97,smi97,sto97,li06,til08}. \\citet{mac95} have described an explicit method for one-fluid ambipolar diffusion in the strong coupling limit, while \\citet{tot94} has used a semi-implicit scheme for two-fluid ambipolar diffusion to investigate instability in C-type shocks. \\citet{til08} have also presented a semi-implicit method for ambipolar diffusion using a two-fluid approximation. Implicit schemes for the multifluid treatment of Hall diffusion and ambipolar diffusion have been suggested by \\citet{fal03} and \\citet{osu06,osu07}.  In this work we describe a fully explicit method for incorporating the single-fluid ambipolar diffusion into a multidimensional MHD code based on the total variation diminishing scheme. The divergence-free condition of the magnetic field is ensured by a flux-interpolated constrained transport scheme, and a super time stepping method is used in order to considerably accelerate the otherwise painfully short diffusion-driven time steps.  The organization of this paper is as follows. In \\S 2 the MHD equations are presented along with the approximations we have made and in \\S 3 our numerical methods are described in detail. Two test problems are presented in \\S 4, while MHD turbulence simulations with significant ambipolar diffusion follow in \\S 5. A summary is given in \\S 6. ",
        "conclusions": "In this paper we describe specific numerical methods for incorporating ambipolar diffusion into a multidimensional MHD code based on the total variation diminishing scheme. We assume the strong coupling approximation, that magnetic force and the neutral-ion drag force in weakly ionized plasmas are almost equal and so the plasma can be treated as a single fluid. Since our numerical methods described in this paper are fully explicit and maintain a second-order accuracy, it is straightforward to extend them to parallelized versions and other geometries. The divergence-free constraint on the magnetic field has been exactly enforced through the flux-interpolated constraint transport scheme at all times. By using the super time stepping method to accelerate the timestep for ambipolar diffusion, we remove the severe restriction on the stable timestep that would arise at high numerical resolution and/or in strong ambipolar diffusion.  Ambipolar diffusion has been tested through the direct comparison with analytic solutions of diffusion problems. We have computed test problems that include oblique C-type shocks and the decay of Alfv\\'en waves. For both of these test problems, comparisons of numerical results to analytic solutions are possible and they demonstrate the good accuracy and robustness of our methods. We have also performed simulations of the compressible MHD turbulence in the presence of ambipolar diffusion and they confirm the ability of our code to follow complex MHD flows. We have shown that the dissipation rate of MHD turbulence is strongly affected by the strength of ambipolar diffusion in both decaying turbulence and forced turbulence.  This multidimensional MHD code incorporating an explicit scheme for solving the ambipolar diffusion term allows us to study astrophysical systems such as molecular cloud cores and protostellar discs in which ambipolar diffusion is thought to be important. Currently this code is being used to study the evolution of compressible MHD turbulence with ambipolar diffusion, and results of three-dimensional, high resolution MHD simulations will be reported elsewhere."
    },
    "0812/0812.1860_arXiv.txt": {
        "abstract": "{ We report the discovery of 4 strong gravitational lensing systems by visual inspections of the Sloan Digital Sky Survey images of galaxy clusters in Data Release 6 (SDSS DR6). Two of the four systems show Einstein rings while the others show tangential giant arcs. These arcs or rings have large angular separations ($>8''$) from the bright central galaxies and show bluer color compared with the red cluster galaxies. In addition, we found 5 {\\it probable} and 4 {\\it possible} lenses by galaxy clusters. ",
        "introduction": "The strong gravitational lensing by galaxy clusters not only constrains their mass distributions \\citep[e.g.][]{tc01,szb+08} but also provides an independent way to study high-redshift faint background galaxies, which are otherwise unobservable \\citep[e.g.][]{mf93,mkm+03}. Furthermore, the statistics of gravitational arcs in lensing clusters can be used to constrain the paradigm of structure formation and cosmology \\citep[e.g.][]{bhc+98,lmj+05}.  Up to now, about 120 strong lensing clusters have been discovered \\citep[e.g.][]{lp89,sef+91,lgh+99,zg03,gky+03,ste+05,hgo+08}. Almost all previous searches for the lensing systems are based on the ground-based telescopes with seeing between $0.5''$ and $1.5''$. \\citet{ste+05} used the {\\it HST} WFPC2 data and found 104 giant arcs in 54 clusters. Recently, based on 240 rich clusters identified from the Sloan Digital Sky Survey \\citep[SDSS,][]{yaa+00} and the follow-up deep imaging observations from the Wisconsin-Indiana-Yale NOAO 3.5\\,m telescope and the University of Hawaii 88 inch telescope, \\citet{hgo+08} uncovered 16 new lensing clusters with giant arcs, with 12 likely lensing clusters and 9 possible candidates.  In this letter, we report the discovery of 4 strong gravitational lensing systems by clusters in the SDSS DR6. The Einstein rings or giant arcs we found from the SDSS images suggest that they are excellent strong lensing systems. In addition, we found 5 {\\it probable} and 4 {\\it possible} lensing systems by galaxy clusters. ",
        "conclusions": ""
    },
    "0812/0812.3865_arXiv.txt": {
        "abstract": "A set of cosmological models that takes into account the variation of the particle number is presented. In this context both dark matter and dark energy can be explained using a single component, without assuming any exotic equation of state, solving directly the cosmic coincidence problem. ",
        "introduction": "Currently the observational evidence coming from supernovae studies \\cite{SNIa}, cosmic background radiation fluctuations \\cite{cmbr}, and baryon acoustic oscillations \\cite{bao}, set a strong case for a cosmological (concordance) universe model composed by nearly $70$ percent of a mysterious component called dark energy, responsible for the current accelerated expansion, nearly $25$ percent of dark matter, which populate the galaxies halos and a small percent (around $4$) composed by baryonic matter. The nature of these two dark component remains so far obscure \\cite{DEDM}. In the case of dark matter, we have a set of candidates to be probed with observations and detections in particle accelerators, however the case of dark energy is more elusive. This component not only has to fill the universe at the largest scale homogeneously, and to have the appropriate order of magnitude to be comparable with that of dark matter, but also has to have a negative pressure equation of state, something that was assumed first to explain inflation in the early universe, inspired by high energy theories of particle physics, but which seems to be awkward to appeal for, at these low energy scales. The alternative way used to account for the cosmic acceleration is to consider a modification of general relativity at large scales \\cite{alterGR}. However it seems to be a difficult route to follow, considering the subtleties in the interpretation of the observations even in the simplest model based on general relativity.  To describe the universe at the largest scale, as is the case in cosmology, we have to make certain assumptions regarding the local distribution of matter and its characteristics. As is well known (and explained in several classic books \\cite{tolman}, \\cite{weinberg}, \\cite{peebles}) the energy-momentum tensor used in the Einstein's equations is imported from considerations made in special relativity, through a covariant generalization supported by the equivalence principle. The assumption about homogeneity and isotropy leads to consider the matter content of the universe as a fluid. In the case of a perfect fluid, defined as a mechanical medium incapable of exerting transverse stresses, it takes the standard diagonal form \\begin{equation}\\label{temunu} T^{\\mu \\nu} = (\\rho + p)U^{\\mu}U^{\\nu}-g^{\\mu \\nu}p, \\end{equation} where $\\rho$ is identified as the energy density of the fluid, $p$ is the proper hydrostatic pressure, $U^{\\mu}=dx^{\\mu}/ds$ are the components of the macroscopic velocity of the fluid with respect to the actual coordinate system and $g^{\\mu \\nu}$ is the metric of it. The expression (\\ref{temunu}) is what is measured by an observer at rest in the fluid, who examines an element of the fluid small enough so that gravitational curvature can be neglected. The conclusion of applying the energy conservation relation to this tensor \\begin{equation} \\label{conslaw} T^{\\mu\\nu}_{\\hspace{0.3cm};\\nu}=0, \\end{equation} is that perfect fluids behave adiabatically when examined by a local observer, satisfying the equation \\begin{equation}\\label{ds0} d(\\rho V)+ pdV=0, \\end{equation} where $V$ is the physical volume of the region considered. Is usual to take as the volume $V =d_f^{3}$, where $d_f$ is a physical distance of the order of the region we are interested in. Using that $d_f = d_c R(t)$, leads to the well know result \\begin{equation} \\dot{\\rho}+3\\frac{\\dot{R}}{R}(\\rho + p)=0, \\end{equation} To get this result, it was assumed also that particle number is conserved $N^{\\mu}_{;\\mu}=0$, where $N^{\\mu}=nU^{\\mu}$ with $n$ the particle number density.  A way to modify (\\ref{ds0}) without spoiling adiabatic evolution is by adding a work term due to the change of the particle number $N$. In \\cite{prigogine}, the authors studied this case as an alternative cosmological model, considering the change of the particle number, assuming that such a correction have to be considered during matter creation. This work actually suggested for the first time, a way to incorporate the particle creation process in the context of cosmology, in a self consistent way. In fact, the original claim made by Zeldovich \\cite{zeldov70}, that gravitational particle production can be described phenomenologically by a negative pressure, is realize here in a beautiful way. A covariant formulation of the model was presented in \\cite{CLW90}.  The context where these considerations would be important were mentioned to be; the steady-state cosmological model, the warm inflationary scenario and within the standard inflationary scenario, during the reheating phase. Based on this work, the studies of cosmological models with matter creation \\cite{lima96} were initiated, which were rapidly recognized as be potentially important to explain dark energy \\cite{harko}. In particular, in Ref.\\cite{zsbp01} the authors established a model where dark energy can be mimicked by self-interactions of the dark matter substratum. Within the same framework, models of interacting dark energy and dark matter were proposed \\cite{interact}. Actually, in these studies is possible to have consistently, a universe where matter creation proceeds within an adiabatic evolution. More recently, Lima et al.,\\cite{lima08} have presented a study of a flat cosmological model where a transition from decelerated to accelerated phase exist. They explicitly show that previous models considered \\cite{lima99}, does not exhibit the transition, and study the observational constraints on the model parameters.  In this work I consider non flat models where by modifying the first law, enable us to explain the current acceleration of the universe expansion through a fictitious pressure component coming from changes in the total particle number. In this way, the dark energy component is basically an effect due to the dark matter creation process in the universe, fact that enable us to explain easily the cosmic coincidence problem; it is not strange to have a similar contribution from these two component, because they are just one component; dark matter. I also demonstrate that a non constant matter creation rate can indeed lead to a transition from a decelerated to an accelerated regime even in the case of non flat geometries. In the next section I derive the equations of motion in the case of matter creation. Then, I open section III with a very simple model with a non constant creation rate, that resemble the $\\Lambda$CDM model. After this, I discussed the case of the transition from a decelerated to an accelerated expansion in non flat models. ",
        "conclusions": ""
    },
    "0812/0812.1117_arXiv.txt": {
        "abstract": "We perform a detailed phase-space analysis of various phantom cosmological models, where the dark energy sector interacts with the dark matter one. We examine whether there exist late-time scaling attractors, corresponding to an accelerating universe and possessing dark energy and dark matter densities of the same order. We find that all the examined models, although accepting stable late-time accelerated solutions, cannot alleviate the coincidence problem, unless one imposes a form of fine-tuning in the model parameters. It seems that interacting phantom cosmology cannot fulfill the basic requirement that led to its construction. ",
        "introduction": "Recent cosmological observations support that the universe is experiencing an accelerated expansion, and that the transition from the deceleration phase to the accelerated phase happened in the recent past \\cite{observ}. In order to explain this remarkable phenomena, one can modify the theory of gravity, such as $f(R)$ gravity \\cite{invr,fr}, Dvali-Gabadadze-Porrati model \\cite{dgp}, and string inspired models \\cite{brane}. An alternative direction is to introduce the concept of dark energy which provides the acceleration mechanism. The simplest candidate of dark energy which is consistent with current observations is the cosmological constant. Due to the lack of a good explanation of the small value of the cosmological constant, many dynamical dark energy models were explored, such as a canonical scalar field (quintessence) model \\cite{quint}, a phantom model that has equation of state parameter $w<-1$ \\cite{phant,Boisseau:2000pr,ArmendarizPicon:2000ah}, or the combination of quintessence and phantom in a unified model named quintom \\cite{quintom}. Moreover, many theoretical studies are devoted to shed light on dark energy within some quantum gravitational principles, such as the so-called ``holographic dark energy'' proposal \\cite{HOLO}.   The dynamical nature of dark energy introduces a new cosmological problem, namely why are the densities of vacuum energy and dark matter nearly equal today although they scale independently during the expansion history. To solve the problem, one would require that the matter density and dark energy density always approach their current values independent of the initial conditions.  The elaboration of this 'coincidence' problem led to the consideration of generalized versions of the aforementioned models with the inclusion of a coupling between dark energy and dark matter. Thus, various forms of ``interacting'' dark energy models \\cite{interacting,Guo:2004xx} have been constructed in order to fulfil the observational requirements.  One decisive test for dark energy models is the investigation of their phase-space analysis. In particular, to examine whether they possess attractor solutions corresponding to accelerating universes and ratio $\\Omega_{{\\text{dark energy}}}/\\Omega_{{\\text{dark matter}}}$ of the order 1. If these conditions are fulfilled, then the universe will result to that solution at late times, independently of the initial conditions, and the basic observational requirements will be satisfied. Although non-interacting quintessence \\cite{expon,Copeland:1997et}, non-interacting phantom \\cite{phannonin} and non-interacting quintom models \\cite{quinnonin} admit late-time accelerated attractors, they possess $\\Omega_{{\\text{dark energy}}}=1$ and thus they are unable to solve the coincidence problem.  When dark energy - dark matter interaction is switched on in quintessence, then one can find accelerated attractors which moreover give $\\Omega_{{\\text{dark energy}}}/\\Omega_{{\\text{dark matter}}}\\approx{\\cal{O}}(1)$ \\cite{Wetterich:1994bg,GarciaCompean:2007vh,Boehmer:2008av,Chen:2008pz}, but paying the price of introducing new problems such as justifying a non-trivial, almost tuned, sequence of cosmological epochs \\cite{Amendola:2006qi}. In cases of interacting phantom models \\cite{Guo:2004vg}, the existing literature remains in some special coupling forms with mainly numerical results, which suggest that the coincidence problem might be alleviated \\cite{Guo:2004xx,Curbelo}.  In the present work we are interested in performing a detailed phase-space analysis of the interacting phantom paradigm, considering new and general forms of the interaction term $Q$. In addition, we go beyond the simplified non-local forms of the literature, considering also a local $Q$-form proportional to a density \\cite{Boehmer:2008av,Cen:2000xv,Malik:2002jb,Ziaeepour:2003qs}. Doing so we find that although late-time accelerated attractors do exist in all the models, as it is expected for phantom cosmology, almost all of them correspond to complete dark energy domination and thus they are unable to solve the coincidence problem. Only for a narrow area of the parameter space of one particular model, can the corresponding solution alleviate the aforementioned problem.  The plan of the work is as follows: In section \\ref{phantcosm} we construct the interacting phantom cosmological scenario and we present the formalism for its transformation into an autonomous dynamical system, suitable for a phase-space stability analysis. In section \\ref{models} we perform the phase-space analysis for four different models, using various interaction terms, and in section \\ref{cosmimpl} we discuss the corresponding cosmological implications. Finally, section \\ref{conclusions} is devoted to the summary of the obtained results. ",
        "conclusions": "\\label{conclusions}  In this work we performed a detailed phase-space analysis of various phantom cosmological models, where the dark energy sector interacts with the dark matter one. Our basic goal was to examine whether there exist late-time scaling attractors, corresponding to accelerated universe and possessing $\\Omega_{{\\text{dark energy}}}/\\Omega_{{\\text{dark matter}}}\\approx{\\cal{O}}(1)$, thus satisfying the basic observational requirements. We investigated four different interaction models, including one with a local, i.e. $H$-independent, form of the interaction term (interacting model 4). We extracted the critical points, determined their stabilities, and calculated the basic cosmological observables, namely the total equation-of-state parameter $w_{tot}$ and $\\Omega_{{\\text{dark energy}}}$ (attributed to the phantom field). The key point of solving the coincidence problem is that the universe is in the attractor now and it has normal radiation dominated and matter dominated eras before it reached the attractors. Note that as long as the interaction term is not too strong, the standard cosmology can be always recovered. Once it is in the attractor, the results do not depend on the initial conditions. Thus, one can switch on the interaction and consider as initial conditions the end of the known epochs of standard Big Bang cosmology, in order to avoid disastrous interference.  In all the examined models we found that stable late-time solutions do exist, corresponding moreover  to an accelerating universe. This feature was expected since phantom cosmology has been constructed in order to always satisfy this condition. Indeed, for all the studied models, we did not find any non-accelerating stable solution. However, in almost all the cases the late-time solutions correspond to a complete dark energy domination and thus are unable to solve the coincidence problem. The only case in which this is possible is in interacting model 1, if we select the parameter values from a very narrow region of the $2D$ parameter space (see figure \\ref{Model1B}).  In conclusion, we see that the examined models of interacting phantom cosmology can produce acceleration (which is ``embedded'' in phantom cosmology in general) but cannot solve the coincidence problem, unless one imposes a form of fine-tuning in the model parameters. This result has been extracted by the negative-kinetic-energy  realization of phantom, which does not cover the whole class of phantom models, but since it is a qualitative statement it should intuitively be robust for general interacting phantom scenarios, too. An alternative direction would be to consider more complicated interaction terms, suitably constructed in order to solve the coincidence problem. But the interaction term was introduced in phantom cosmology in order to solve the coincidence problem in a simple and general way, avoiding the assumptions and fine-tunings of conventional cosmology. Although promising and with many advantages, interacting phantom cosmology needs further investigation."
    },
    "0812/0812.2052_arXiv.txt": {
        "abstract": "We present an investigation of the relationships between the radio properties of a giant radio galaxy MRC B0319-454 and the surrounding galaxy distribution with the aim of examining the influence of intergalactic gas and gravity associated with the large-scale structure on the evolution in the radio morphology.  Our new radio continuum observations of the radio source, with high surface brightness sensitivity, images the asymmetries in the megaparsec-scale radio structure in total intensity and polarization.  We compare these with the 3-D galaxy distribution derived from galaxy redshift surveys.  Galaxy density gradients are observed along and perpendicular to the radio axis: the large-scale structure is consistent with a model wherein the galaxies trace the ambient intergalactic gas and the evolution of the radio structures are ram-pressure limited by this associated gas.  Additionally, we have modeled the off-axis evolution of the south-west radio lobe as deflection of a buoyant jet backflow by a transverse gravitational field: the model is plausible if entrainment is small.  The case study presented here is a demonstration that giant radio galaxies may be useful probes of the warm-hot intergalactic medium believed to be associated with moderately over dense galaxy distributions. ",
        "introduction": "In a powerful radio galaxy, reaccelerated jet material inflates the synchrotron lobes after passage through the termination shocks at the jet ends \\citep{scheuer74,begelman84}. The synchrotron lobes interact with the surrounding thermal gas and their morphologies are a result of this interaction. In cluster environments, where XMM-Newton and Chandra have been able to detect and image the relatively dense intra-cluster medium (ICM) in X-rays, the total intensity radio and X-ray contours follow each other and X-ray holes are observed at the locations of the radio lobes \\citep[e.g.][]{boehringer93,mcnamara00,nulsen02,birzan04}. These holes are evidence that the expanding radio lobes interact with, and displace the X-ray emitting gas. Lobe interaction with intra-cluster gas is also apparent in the bending of radio plumes in wide-angle and narrow-angle tailed radio sources \\citep[e.g.][]{bliton98,hardcastle05,douglass08}. The morphologies of these sources suggest that the radio tails are deflected behind the moving host galaxy, due to the ram-pressure exerted by the ICM. \\par Giant radio galaxies, which have linear sizes of $>1$~Mpc, tend to reside outside of rich clusters \\citep{jamrozy04}. Their giant lobes, which extend well beyond the  interstellar media (ISM) and coronal halos of  their host galaxies, represent the interaction between the light synchrotron lobe plasma and the heavier intergalactic medium (IGM). If galaxies are a reliable tracer of the IGM gas, then we expect the distribution of gas around a giant radio galaxy to depend on the position of the host galaxy relative to large-scale structure in the IGM. Since the evolution of the radio lobes depends on their interaction with this ambient gas, large scale structure (and the associated gas) is also expected to determine the lobe dynamics in giant radio sources. For example, if the parent optical galaxy resides at a boundary between an over-density and a void, then there may be a gradient in the gas density surrounding the radio source. If this gradient is aligned with the radio jet axis, then we would expect an asymmetry to arise in the lobe lengths, due to differing ram-pressure limitations in the ambient gas on the two sides. If, on the other hand, the gradient in the gas density is transverse to the jet axis, then we would expect the light synchrotron lobes to evolve transverse to the jet axis and in the direction of decreasing ambient density and pressure due to buoyancy. \\par The diffuse gas in the environs of giant radio galaxies, outside of rich clusters, most likely pertains to the warm-hot phase of the IGM, whose existence at low-redshifts has been predicted by large-scale cosmological hydrodynamical simulations of galaxy formation \\citep{cen99, cen06, dave01}. In the simulations, the warm-hot gas follows the filamentary galaxy distribution on large scales and represents unvirialized over-densities in the range 10-30. Little is known about this warm-hot gas since  its thermal emission is not detectable by present day X-ray telescopes. However, the shape taken by the lobes of a giant radio galaxy is a visible manifestation of their interaction with the surrounding gas. These galaxies, therefore, make ideal probes of the unseen ambient medium, which may be the warm-hot intergalactic medium (WHIM). \\par Recently, the 1.8~Mpc lobes of the giant radio galaxy MSH 05$-$2{\\it 2} were used to constrain the properties of its ambient thermal gas \\citep{subrahmanyan08}. The lobes in this galaxy appear to be relicts and there is an observed asymmetry, which appears to be related to an anisotropy in the local galaxy distribution. A comparison of the properties of the radio lobes and those of the ambient gas, which were derived from the surrounding galaxy distribution, indicated that the lobes were highly overpressured despite their relict appearance. Alternatively, the density-temperature product for the ambient IGM might be an order of magnitude larger than that predicted by structure formation models indicating significant feedback in the IGM. \\par In this paper, we present a similar study of another giant radio galaxy \\rg (also referred to previously as MSH 03$-$4{\\it 3} and PMN J0321$-$4510) and its environment. The 2.5 Mpc radio source is located within a galaxy filament of the Horologium-Reticulum super-cluster \\citep{fleenor05} at a redshift of $z\\approx0.6$. Its 26\\arcmin size on the sky, make the measurement of the surrounding galaxy distribution possible with multi-object fibre instruments such as AAOmega on the Anglo-Australian Telescope. \\rgns, therefore, provides a further opportunity to study the morphology of a giant radio galaxy and its relationship to the surrounding galaxy distribution. Such studies are a first step in using giant radio galaxies as probes of the unseen intergalactic gas associated with large-scale galaxy structures.  \\par The giant radio source \\rg was previously observed at 843-MHz with the Molonglo Observatory Synthesis Telescope (MOST), and with the Australia Telescope Compact Array (ATCA) at 20, 13 and 6 cm \\citep{jones89, saripalli94}. In the latter work, Saripalli et al. presented a detailed study of the morphology of the giant radio galaxy pointing out a number of features that make the radio source unusual: a unique configuration of five compact hot spots in one of the lobes, a prominent jet and counter-jet detected out to exceptional distances, and lobes that are asymmetrically shaped and positioned about the core. Considering the likely expansion velocity of the source, it was argued that light travel-time effects were not responsible for the asymmetry in the lobe separations from the core. Based on an examination of the projected  galaxy distribution around the host, Saripalli et al. attributed this feature to corresponding asymmetries in the ambient intergalactic medium.  \\par Given the Mpc size and pronounced asymmetries in the radio morphology, \\rg is an excellent example of a radio galaxy where the morphology is shaped by the IGM. In this paper, we examine the interaction between the lobes and the ambient IGM using our new radio observations made with the specific purpose of imaging the radio lobes fully.  Our new improved radio images of \\rg have been made using an observing mode designed to have high surface brightness sensitivity and in full polarization. We use these in conjunction with redshift measurements of the surrounding galaxy distribution, in an attempt to model the interaction between the giant lobes with their IGM environment. Our radio observations include polarization measurements for the first time. \\par The paper is organized as follows: In Sect.~\\ref{s:radio} we present our new radio continuum images of \\rgns. In Section 3, we describe the galaxy redshift data; in the sections that follow an attempt is made towards understanding the lobe-IGM interaction assuming that the galaxy distribution traces the gas distribution on large scales. Throughout we adopt a flat cosmology with parameters $\\Omega_{\\rm 0} = 0.3$, $\\Omega_ {\\Lambda }= 0.7$ and a Hubble constant $H_{0}=71$ km s$^{-1}$ Mpc$^{-1}$. The host elliptical galaxy,  ESO~248-G10 (also cataloged as AM~0319$-$452),  has an R-band absolute magnitude $M_R = -23.7$  \\citep{saripalli94}. Previously, the redshift of the host was estimated to be $z=0.0633$ \\citep{jones89}; our new data, presented herein, give an estimate of $z = 0.0622$.  At this redshift, 1\\arcsec = 1.2 kpc. ",
        "conclusions": "\\label{s:con}  We have presented new high sensitivity radio images of the powerful radio galaxy \\rg together with observations of the surrounding large-scale structure. Using these we have sought to understand the influence of the ambient gas and gravity on the evolution of  the jets and post hot spot plasma in this remarkable source. Our observations are qualitatively consistent with a model wherein the galaxies trace inhomogeneities in the intergalactic gas, and these in turn determine the evolution of the giant radio source. In this respect, our study has reached similar conclusions to those of \\citet{subrahmanyan08}, who also showed that asymmetries in the morphology of the giant relict source, MSH J0505$-$2835, may be governed by inhomogeneities in the ambient intergalactic gas distribution, whose distribution follows the large scale structure in galaxies. In particular, both studies have found a foreshortening of the lobes in the direction of greater galaxy density, and off-axis distortions in the direction of decreasing galaxy density.  \\par In \\rgns, over-densities in the galaxy distribution are found north of the radio core. These are related to the loose galaxy group of which the host is a member. The advance of the NE jet and expansion of the corresponding lobe are ram-pressure limited by the intra-group medium of this galaxy concentration. Furthermore, there is less gas density to the south of the radio core, explaining the rapid advance of the SW jet and increased expansion losses in the post hot spot flow. A density gradient, perpendicular to the jet axis in the SW lobe, may also influence the asymmetric lobe expansion to the NW. \\par We further hypothesize that the SW backflow may be moving away from the source axis and NW due to buoyancy. We find that buoyant forces (due to a gravitational field transverse to the radio source axis) can deflect the backflow of the SW jet plasma, provided there has been minimal entrainment of ambient gas into the lobe. This leads to a model in which the medium external to the SW radio lobe has the same properties as those inferred for the WHIM component of the IGM. \\par Future deep imaging of the gas associated with the large-scale structure would help to further our interpretation. This work demonstrates the usefulness of giant radio galaxies as probes of the IGM through the relationship between source morphology and the ambient large-scale structure distribution."
    },
    "0812/0812.2264_arXiv.txt": {
        "abstract": "We observed at very high spectral resolution the prototype Z-source \\cygx2 twice along its entire X-ray spectral variation pattern. In this preliminary analysis we find an extended accretion disk corona exhibiting Lyman $\\alpha$ emissions from various H-like ions, as well as emissions from He-like ions of Fe and Al, and Be-like ions of Fe. The brightest lines show a range of line broadening: H-like lines are very broad with Doppler velocities between 1100 and 2700 km s$^{-1}$, while some others are narrower with widths of a few hundred km s$^{-1}$. Line diagnostics allow us for the first time to determine coronal parameters. The line properties are consistent with a stationary, extended up to $10^{10}$ cm, dense ($1\\times10^{15}$ cm$^{-3}$), and hot (log$ \\xi \\approxgt 3; T \\approxgt 10^6$ K) accretion disk corona. We find ongoing heating of the corona along the Z-track and determine that heating luminosities change from about 0.4$\\times L_{Edd}$  on the horizontal to about 1.4$\\times L_{Edd}$ on the flaring branch. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.1392_arXiv.txt": {
        "abstract": "In support of the new Sloan III survey, which will measure the baryon oscillation scale using the luminous red galaxies (LRGs), we have run the largest N-body simulation to date using $4120^3 = 69.9$ billion particles, and covering a volume of $(6.592~ h^{-1} {\\rm Gpc})^3$. This is over 2000 times the volume of the Millennium Run, and corner-to-corner stretches all the way to the horizon of the visible universe. LRG galaxies are selected by finding the most massive gravitationally bound, cold dark matter subhalos, not subject to tidal disruption, a technique that correctly reproduces the 3D topology of the LRG galaxies in the Sloan Survey. We have measured the covariance function, power spectrum, and the 3D topology of the LRG galaxy distribution in our simulation and made 32 mock surveys along the past light cone to simulate the Sloan III survey. Our large N-body simulation is used to accurately measure the non-linear systematic effects such as gravitational evolution, redshift space distortion, past light cone space gradient, and galaxy biasing, and to calibrate the baryon oscillation scale and the genus topology. For example, we predict from our mock surveys that the baryon acoustic oscillation peak scale can be measured with the cosmic variance-dominated uncertainty of about 5\\% when the SDSS-III sample is divided into three equal volume shells, or about 2.6\\% when a thicker shell with $0.4<z<0.6$ is used.  We find that one needs to correct the scale for the systematic effects amounting up to 5.2\\% to use it to constrain the linear theories. And the uncertainty in the amplitude of the genus curve is expected to be about 1\\% at 15 $h^{-1}$Mpc scale. We are making the simulation and mock surveys publicly available. ",
        "introduction": "A leading method for characterizing dark energy is measuring the baryon oscillation scale using Luminous Red Galaxies (LRGs).  Baryon oscillations are the method ``least affected by systematic uncertainties'' according to the Dark Energy task Force report (Albrecht et al. 2006).  The baryon oscillation scale provides a ``standard ruler'' (tightly constrained by the WMAP observations), which has been detected as a bump in the covariance function of the LRG galaxies in the Sloan Survey (and also in the 2dFGRS survey) (see Percival et al. 2007 for a review).  This observation has prompted a team to propose the Baryon Oscillation Spectroscopic Survey (BOSS) for an extension of the Sloan Survey (Sloan III), which will obtain the redshifts of 1,300,000 LRG galaxies over 10,000 square degrees selected from the Sloan digital image survey.  This will make it possible to measure the baryon bump in the covariance function of the LRG galaxies to a high accuracy as a function of redshift out to a redshift of z = 0.6. It is claimed that this allows, in principal, measurement of the Hubble parameter to an accuracy of 1.7\\% at $z=0.6$, and allows us to place narrower constraints on $w(z)$ (the ratio of dark energy pressure to energy density as a function of redshift). Systematic uncertainties are claimed to be of the order of 0.5\\% (Eisenstein et al 2007$a$; Crocce \\& Scoccimarro 2007).  These systematics are only as well controlled as knowledge of all non-linear effects associated with LRG galaxies and thus one needs large N-body simulations.  Small N-body simulations--$(256)^3$ particles, ($512 h^{-1}$Mpc)$^3$ (Seo \\& Eisenstein 2005), show that non-linear effects do change the covariance function of the LRG galaxies relative to that in the initial conditions.  The baryon bump in the covariance function is somewhat washed out, but can be recovered by reconstruction techniques moving the galaxies backward to their initial conditions using the Zeldovich approximation (Eisenstein et al. 2007$b$).  However, these simulations are inadequate because of their small box size which is only a factor of 4.7 larger than the baryon oscillation scale which is of order of $108 h^{-1}$Mpc.   It is very important to  model the power spectrum accurately at large scales and to have a large box size so that the statistical errors in the power spectrum are small.  After all, we need to measure the acoustic peak scale to an accuracy better than 1\\%.  Concerning the Baryon Acoustic Oscillation method applied to dark energy, Crotts et al. (2005) note that ``systematic effects will inevitably be present in real data; they can only be addressed by means of studying mock catalogues constructed from realistic cosmological volume simulations'' just as we are setting about to do here. According to Crotts et al. (2005) such modeling would typically improve the accuracy of the estimated dark energy parameters by 30-50\\%. In support of this goal we have previously completed a $2048^3$ particle ($4915 h^{-1}$Mpc)$^3$ Cold Dark Matter Simulation (Park et al. 2005$a$). We have developed a new technique to identify LRG galaxies by finding the most massive bound subhalos not subject to tidal disruption (Kim \\& Park 2006). This $2048^3$ simulation successfully models the 3D topology of the LRG galaxies in the Sloan Survey, correctly modeling within approximately $1\\sigma$ level the genus curve found in the observations (Gott et al. 2009). By contrast the semi-analytic model of galaxy assignment scheme applied to the Millennium run by Springel et al. (2005) differed from the observational data on topology by $2.5\\sigma$ indicating either a need for an improvement in its initial conditions (it used a bias factor that is too low and an $n_s = 1$ power spectrum that have relatively low power at large scales according to WMAP-3 (Spergel et al 2007)) or its galaxy formation algorithm (cf. Gott et al. 2008).  The fact that our $2048^3$ N-body simulation modeled well the observational data on LRG topology suggests that the formation of LRG galaxies is a relatively clean problem that can be well modeled by large N-body simulations using WMAP 5-year initial conditions.  Building on this success in modeling the formation and clustering of LRG galaxies, in this paper we will show first results from our very large Horizon Run N-body simulation in support of the BOSS in the Sloan III survey. ",
        "conclusions": "Baryon oscillations are believed to be the method of characterizing dark energy ``least affected by systematic uncertainties'' according to the Dark Energy Task Force Report.   But the baryon oscillation is affected by all kinds of non-linear effects, so comparison with large N-body simulations is essential to control the systematics. Gott et al. (2009)'s study of the topology of the LRG galaxy clustering in the current SDSS shows that large N-body simulations are accurately modeling these non-linear gravitational and biased galaxy formation effects.  To aid in calibrating non-linear gravitational and biasing effects in the future Sloan-III baryon oscillation scale survey (BOSS), we have made here a large N-body simulation: $4120^3 = 69.9$ billion particles in a $(6592 h^{-1}$Mpc$)^3$ box, which models the formation of LRGs in this region. Using improved values over those adopted by the Millennium Run, which was based on WMAP-1 initial conditions (Spergel et al. 2003), our simulation has more power at large scales ($n_s = 0.953$ versus $n_s = 1$, which should make for a higher frequency of large scale structures), and more accurate normalization at $8 h^{-1}$Mpc scale ($\\sigma_8 = 0.796$ while Millennium Run used $\\sigma_8=0.9$).  We have made 32 mock surveys of the Sloan III survey, which should allow us to correct for small systematic effects in the measurement of the physical scale of the baryon oscillation peak as a function of redshift. We predict from our mock surveys that the BAO peak scale can be measured with the cosmic variance-dominated uncertainty of about 5\\% when the SDSS-III sample is divided into three equal volume shells, or about 2.6\\% when a thicker shell with $0.4<z<0.6$ is used.  We find that one needs to correct the scale for the systematic effects amounting up to 5.2\\% to use it to constrain the linear theories.  The amplitude of the genus curve measured in redshift space in the mock SDSS III survey is, in the mean, expected to be equal to the linear regime prediction to an accuracy of about 2\\% at the smoothing scale of 20 $h^{-1}$Mpc. So the systematic evolution correction in the amplitude of the genus curve will be tiny. Also, the RMS uncertainty in the amplitude in the genus curve due to the cosmic variace as measured in one SDSS-III survey is only 1.5\\% at 20 $h^{-1}$Mpc scale. Since the uncertainty scales with inverse square root of the number of resolution elements (or equivalently with $R_G^1.5$), this implies the uncertainty of about 1.0\\% at 15 $h^{-1}$Mpc scale, which will be the mean separation of the SDSS-III LRGs. In subsequent papers we shall show how measuring the amplitude of the genus curve at different smoothing length makes possible an independent measure of scale complimenting the baryonic oscillation bump method for measuring $w(z)$.  As dark energy makes up three quarters of the universe today, making such improvements in characterizing dark energy are particularly exciting. Much ancillary science can be done with this large N-body simulation; detailed modeling of the integrated Sachs-Wolf effect and comparison with counts of LRG galaxies in the sky, for example. We can search for Great Walls and voids to estimate their frequency and size (Gott et al. 2005). Here, our large volume size and WMAP-5 based initial conditions (Komatsu et al. 2008) should be of great benefit. We will make the simulation data available to the community (See http://astro.kias.re.kr/Horizon-Run/ for information). Please cite this paper when the simulation data is used."
    },
    "0812/0812.2028_arXiv.txt": {
        "abstract": "We present for the first time the complete matter power spectrum for $R^n$ gravity which has been derived from the {\\it fourth order} scalar perturbation equations. This leads to the discovery of a characteristic signature of fourth order gravity in the matter power spectrum, the details of which have not seen before in other studies in this area and therefore provides a crucial test for fourth order gravity on cosmological scales. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.2358_arXiv.txt": {
        "abstract": "We present the discovery of two new late-T dwarfs identified in the UKIRT Infrared Deep Sky Survey (UKIDSS) Galactic Clusters Survey (GCS) Data Release 2 (DR2). These T dwarfs are nearby old T dwarfs along the line of sight to star--forming regions and open clusters targeted by the UKIDSS GCS\\@. They are found towards the $\\alpha$\\,Per cluster and Orion complex,respectively, from a search in 54 square degrees surveyed in five filters. Photometric candidates were picked up in two--colour diagrams, in a very similar manner to candidates extracted from the UKIDSS Large Area Survey (LAS) but taking advantage of the $Z$ filter employed by the GCS\\@. Both candidates exhibit near--infrared $J$-band spectra with strong methane and water absorption bands characteristic of late--T dwarfs. We derive spectral types of T6.5$\\pm$0.5 and T7$\\pm$1 and estimate photometric distances less than 50 pc for UGCS J030013.86$+$490142.5 and UGCS J053022.52$-$052447.4, respectively. The space density of T dwarfs found in the GCS seems consistent with discoveries in the larger areal coverage of the UKIDSS Large Area Survey, indicating one T dwarf in 6--11 square degrees. The final area surveyed by the GCS, 1000 square degrees in five passbands, will allow expansion of the LAS search area  by 25\\%, increase the probability of finding ultracool brown dwarfs, and provide optimal estimates of contamination by old field brown dwarfs in deep surveys to identify such objects in open clusters and star--forming regions. ",
        "introduction": "\\label{dT:intro}  The advent of large--scale sky surveys has revolutionised our knowledge of ultracool dwarfs (refered to as dwarfs with spectral types of M7 and later). The first spectroscopic brown dwarfs were confirmed in 1995: Gl229B, a T dwarf orbiting an M dwarf \\citep{nakajima95} and Teide~1 in the Pleiades open cluster \\citep{rebolo95}. More than twenty years later, about 530 L dwarfs with effective temperatures (T$_{\\rm eff}$) between $\\sim$2200 and $\\sim$1400\\,K \\citep{basri00,leggett00} have now been identified (as of May 2008), along with around 140 T dwarfs with lower temperatures \\citep[T$_{\\rm eff}\\simeq$ 1400--700 K;][]{golimowski04a,vrba04}. The full catalogue of L and T dwarfs is available on the DwarfArchives.org webpage\\footnote{http://dwarfarchives.org, a webpage dedicated to M, L, and T dwarfs maintained by C.\\ Gelino, D.\\ Kirkpatrick, and A.\\ Burgasser.}. There are currently 63 T5 or later dwarfs at the time of writing, refered to here as late-T dwarfs. The spectral classification of T dwarfs follows the unified scheme by \\citet{burgasser06a} and is based on the strength of methane and water absorption bands present in the near-infrared. This sample of ultracool dwarfs is now large enough to characterise the binary properties of brown dwarfs \\citep{close02b,burgasser03a,bouy03,burgasser06d,liu06,burgasser07d} and investigate the influence of gravity and metallicity on their spectral energy distributions \\citep{kirkpatrick05,jameson08b}.  The UKIRT Infrared Deep Sky Survey (UKIDSS) is defined in \\citet{lawrence07}. UKIDSS uses the Wide Field Camera \\citep[WFCAM;][]{casali07} installed on the UK InfraRed Telescope (UKIRT) and a photometric system described in \\citet{hewett06}. The pipeline processing is made at the Cambridge Astronomical Unit (CASU; Irwin et al., in prep)\\footnote{http://casu.ast.cam.ac.uk/surveys-projects/wfcam/technical} and the data are available from the WFCAM Science Archive \\citep{hambly08}. The survey is now well underway with several ESO-wide releases: the Early Data Release (EDR) in February 2006 \\citep{dye06} and the Data Release 1 \\citep[DR1;][]{warren07a} in 2006, the Data Releases 2 \\& 3 in March and December 2007 \\citep{warren07b}\\footnote{DR2 is now a public worldwide release since September 2008}, and DR4 in July 2008, the latest to date.  The Galactic Cluster Survey (GCS), one of the five components of UKIDSS, will cover 1000 deg$^2$ in $ZYJHK$ in ten clusters and star--forming regions down to a 5$\\sigma$ sensitivity limit of $K\\simeq$ 18.7 mag after two epochs at $K$. The main science goal of the GCS is to derive the very low--mass stellar/substellar mass function in several open clusters and star--forming regions over large areas to investigate the (possible) role of environment and time on the mass function. We have already described the selection procedure of low--mass stars, brown dwarfs, and planetary-mass members of the Upper Sco association \\citep{lodieu06,lodieu07a} and confirmed many photometric candidates as young L dwarf members \\citep{lodieu08a}. Additionally, we have derived the mass function in the Pleiades down to 30 Jupiter masses over 12 square degrees and estimated the photometric brown dwarf binary fraction \\citep{lodieu07c}.  Although the GCS is not sensitive to young T dwarfs at the distance of the targeted clusters, it is possible to look for nearer, old field T--type dwarfs in a similar manner to the LAS \\citep{kendall07,warren07c,lodieu07b,pinfield08,burningham08} which is beneficial both in the context of field and cluster brown dwarf studies. It is a headline science goal of UKIDSS to find the so-called Y dwarfs: current searches have focused on the LAS but ignored the GCS as a ground to uncover field T dwarfs. Ultimately, the GCS will provide an additional 25\\% coverage in five filters to the 4000 deg${^2}$ planned by the LAS\\@.  In this paper we describe the first search using the GCS DR2 in this context and report the discovery of a T6.5$\\pm$0.5 and a T7$\\pm$1 towards $\\alpha$ Per and Orion, respectively. In Section \\ref{dT_GCS:selection} we describe the photometric selection of late--T dwarfs from two--colour diagrams. In Section \\ref{dT_GCS:properties} we present the near--infrared spectra for both targets, assign spectral types based on the unified classification scheme by \\citet{burgasser06a}, and derive their main properties. In Section \\ref{dT_GCS:contamination} we examine the issue of field dwarf contamination towards young clusters and star--forming regions targeted by the GCS\\@. In Section \\ref{dT_GCS:prospects} we discuss future prospects offered by the GCS to extend this preliminary search for field T dwarfs and further constrain levels of field dwarf contamination.  \\begin{figure} \\includegraphics[width=\\linewidth, angle=0]{fig1.ps} \\caption{Histograms for the number of objects in the LAS (black) and the GCS (red) per magnitude bin as a function of magnitudes for the $Y$ (top left), $J$ (top right), $H$ (bottom left), and $K$ (bottom right) filters. The 100\\% completeness limits are quoted in Table \\ref{tab_dT_GCS:depth_compare}. } \\label{fig_dT:histograms} \\end{figure} ",
        "conclusions": "\\label{dT:conclusions}  We have presented the spectroscopic confirmation of two new nearby old late--T dwarfs identified in 54 square degrees surveyed as part of the Second Data Release of the UKIDSS Galactic Cluster Survey. Those sources lie in front the Alpha Per cluster and Orion complex at distances less than 50 pc. These sources have been identified in regions previously ignored by earlier searches for cool brown dwarfs.  This work opens new prospects to (i) extend the search area of the UKIDSS LAS and increase the number of late--T dwarfs by 25\\%, and (ii) estimate the contamination in deep optical and near--infrared surveys of young open clusters and star--forming regions by nearby old L and T dwarfs and thereof correct the faint end of the photometric mass functions."
    },
    "0812/0812.2450_arXiv.txt": {
        "abstract": "The Global Magneto-Ionic Medium Survey (GMIMS) is a project to map the diffuse polarized emission over the entire sky, Northern and Southern hemispheres, from 300 MHz to 1.8 GHz. With an angular resolution of 30 - 60 arcmin and a frequency resolution of  1 MHz or better, GMIMS will provide the first spectro-polarimetric data set of the large-scale polarized emission over the entire sky, observed with single-dish telescopes. GMIMS will provide an invaluable resource for studies of the magneto-ionic medium of the Galaxy in the local disk, halo, and its transition. ",
        "introduction": "Polarization of the diffuse synchrotron emission from the Galaxy provides a sensitive probe of the magnetic fields and warm and hot plasmas making up the Galactic magneto-ionic medium (MIM). Polarization maps at decimetre wavelengths thus reveal intriguing rotation measure structures that are otherwise undetectable. These structures contain information about the Galactic magnetic field and the distribution of the ionized gas. In short, the physics of the MIM is imprinted on the diffuse polarized emission. So far, single-frequency (or narrow-band) polarization surveys have only allowed rather qualitative studies due to ambiguities in the data. Wide-band polarization data, as they are now underway through various survey projects, can resolve these ambiguities and will allow quantitative analysis of the diffuse polarized emission of the Galaxy.  Figure 1 summarizes our current knowledge of the northern polarized sky at 408 MHz (left) and 1.4 GHz (right). The polarization vectors at 408 MHz were measured more than 30 years ago. Similar data at 1.4 GHz were only recently superseded by the new DRAO Low-Resolution Polarization Survey. Data with higher angular resolution have been observed by, for example, the EMLS \\citep{2004mim..proc...45R} and CGPS (Landecker et al., in prep.) in the northern sky, as well as SGPS \\citep{2006ApJS..167..230H} and S-PASS \\citep{2008arXiv0806.0572C} on the southern sky. GMIMS aims at superseding and complementing these data through its frequency coverage (or, more importantly for rotation measure studies, through its $\\lambda^2$-coverage).  \\begin{figure}[t] \\includegraphics[scale=0.35]{wolleben_fig01.eps} \\includegraphics[scale=0.35]{wolleben_fig02.eps} \\caption{Polarization maps of the northern sky in Galactic coordinates at 408 MHz \\citep[left, from][]{1976A&AS...26..129B} and 1.4 GHz \\citep[right, from][]{2006A&A...448..411W}. See Reich \\& Reich (this issue) for an all-sky version of the 1.4 GHz survey.} \\end{figure} ",
        "conclusions": ""
    },
    "0812/0812.2720_arXiv.txt": {
        "abstract": "\\emph{Chandra} observations of large samples of galaxy clusters detected in X-rays by \\emph{ROSAT} provide a new, robust determination of the cluster mass functions at low and high redshifts. Statistical and systematic errors are now sufficiently small, and the redshift leverage sufficiently large for the mass function evolution to be used as a useful growth of structure based dark energy probe. In this paper, we present cosmological parameter constraints obtained from \\emph{Chandra} observations of 37~clusters with $\\langle z\\rangle=0.55$ derived from 400~deg$^2$ \\emph{ROSAT} serendipitous survey and 49 brightest $z\\approx 0.05$ clusters detected in the All-Sky Survey.  Evolution of the mass function between these redshifts requires $\\OmegaL>0$ with a $\\sim 5\\sigma$ significance, and constrains the dark energy equation of state parameter to $\\wo=-1.14\\pm0.21$, assuming constant $w$ and flat universe.  Cluster information also significantly improves constraints when combined with other methods. Fitting our cluster data jointly with the latest supernovae, WMAP, and baryonic acoustic oscillations measurements, we obtain $w_0=-0.991\\pm0.045$ (stat) $\\pm0.039$ (sys), a factor of 1.5 reduction in statistical uncertainties, and nearly a factor of 2 improvement in systematics compared to constraints that can be obtained without clusters.  The joint analysis of these four datasets puts a conservative upper limit on the masses of light neutrinos, $\\sum m_\\nu<0.33$~eV at 95\\% CL.  We also present updated measurements of $\\OmegaMh$ and $\\sigma_8$ from the low-redshift cluster mass function. ",
        "introduction": "\\label{sec:intro}  Recent accelerated expansion of the Universe detected in the Hubble diagram for distant type~Ia supernovae is one of the most significant discoveries of the past 10 years \\citep{1999ApJ...517..565P,1998AJ....116.1009R}. The acceleration can be attributed to the presence of a significant energy density component with negative pressure, hence the phenomenon is commonly referred to as Dark Energy. For a recent review of the dark energy discovery and related theoretical and observational issues, see \\cite*{2008arXiv0803.0982F} and references therein. Perhaps the simplest phenomenological model for dark energy is non-zero Einstein's cosmological constant. The supernovae data indicated (and other cosmological datasets now generally agree) that a cosmological constant term currently dominates energy density in the Universe.  The next big question is whether Dark Energy really is the cosmological constant. The properties of dark energy are commonly characterized by its equation of state parameter, $w$, defined as $p=w\\rho$, where $\\rho$ is the dark energy density and $p$ is its pressure. A cosmological constant in the context of General Relativity corresponds to a non-evolving $w=-1$. It is proposed that departures from the cosmological constant model should be sought in the form of observed $w$ being either $\\ne -1$, or evolving with redshift.  Combination of supernovae, cosmic microwave background, and baryonic acoustic oscillations data currently constrain $|1+w|<0.15$ at 95\\% CL \\citep{2008arXiv0803.0547K}.  Observational signatures of such deviations of $w$ from $-1$ are very small, and hence the measurements are prone to systematic errors. For example, variations of $w$ between $-1$ and $-0.9$ change fluxes of $z=0.75$ supernovae in a flat universe with $\\OmegaM=0.25$ by only $0.03$~magnitudes.  Therefore, it is crucially important that the dark energy constraints at this level of accuracy are obtained from combination of several independent techniques.  This not only reduces systematics but also improves statistical accuracy by breaking degeneracies in the cosmological parameter constraints.  One of the methods that has been little used so far is evolution in the number density of massive galaxy clusters. Evolution of the cluster mass function traces (with exponential magnification) growth of linear density perturbations. Growth of structure and distance-redshift relation are similarly sensitive to properties of dark energy, and also are mutually highly complementary methods \\citep[e.g.,][]{2003MNRAS.346..573L}.  Mapping between the linear power spectrum and cluster mass function relies on the model for nonlinear gravitational collapse.  This model is now calibrated extensively by $N$-body simulations (see \\S\\,\\ref{sec:theory}). The cluster mass function models also use additional assumptions (e.g., that the mass density is dominated by cold dark matter in the recent past, and that the fluctuations have Gaussian distribution). However, corrections due to reasonable departures from these assumptions are negligible compared to statistical uncertainties in the current samples (we discuss these issues further in \\S\\,\\ref{sec:theory}).  It is important also that the theory of nonlinear collapse is insensitive to the background cosmology. For example, the same model accurately describes the relation between the linear power spectrum and cluster mass function in the $\\OmegaM=1$, $\\OmegaL=0$, low-density $\\OmegaM=0.3$, $\\OmegaL=0$, and ``concordant'' $\\OmegaM=0.3$, $\\OmegaL=0.7$ cosmologies \\citep{2001MNRAS.321..372J}.  Fitting cosmological models to the real cluster mass function measurements uses not only growth of structure but also the distance-redshift information because observed properties for objects of the same mass generally depend on the distance. Therefore, constraints on $w$ derived from the cluster mass function internally make a combination of growth of structure and distance based cosmological tests, and thus potentially can be very accurate and competitive with any other technique \\citep[e.g.,][]{2006astro.ph..9591A}.   Previous attempts to use evolution of the cluster mass function as a cosmological probe were limited by small sample sizes and either poor proxies for the cluster mass (e.g., the total X-ray flux) or inaccurate measurements (e.g., temperatures with large uncertainties). Despite these limitations, reasonable constraints could still be derived on $\\OmegaM$ \\citep[e.g.,][]{2001ApJ...561...13B,2004ApJ...609..603H}. However, constraints on the dark energy equation of state from such studies are weak.  For example, \\citet{2004ApJ...609..603H} derived the best-fit $w=-0.42$, only marginally inconsistent with $w=-1$, using the temperature function of the \\emph{Einstein} Medium Sensitivity Survey clusters; \\citet{2007arXiv0709.4294M} determine $w=-1.4\\pm0.55$ with a larger sample of distant clusters \\citep[MACS survey, see][]{2001ApJ...553..668E} but using the X-ray luminosity as a mass proxy.  The situation with the cluster mass function data has been dramatically improved in the past two years. A large sample of sufficiently massive clusters extending to $z\\sim 0.9$ has been derived from \\emph{ROSAT} PSPC pointed data covering 400~deg$^2$ \\citep[][Paper~I~hereafter]{2007ApJS..172..561B}.  Distant clusters from the \\400d{} sample were then observed with \\emph{Chandra}, providing high-quality X-ray data and much more accurate total mass indicators. \\emph{Chandra} coverage has also become available for a complete sample of low-$z$ clusters originally derived from the \\emph{ROSAT} All-Sky Survey. Results from deep \\emph{Chandra} pointings to a number of low-$z$ clusters have significantly improved our knowledge of the outer cluster regions and provided a much more reliable calibration of the $\\Mtot$ vs.\\ proxy relations than what was possible before. On the theoretical side, improved numerical simulations resulted in better understanding of measurement biases in the X-ray data analysis \\citep{2007ApJ...655...98N,2006MNRAS.369.2013R,2007arXiv0708.1518J}. Even more importantly, results from these simulations have been used to suggest new, more reliable X-ray proxies for the total mass \\citep{2006ApJ...650..128K}. We discuss all this issues in the previous paper \\citep[Paper~II hereafter]{vikhlinin08}. The cluster mass functions derived in this paper are reproduced in Fig.\\,\\ref{fig:mfun}. Overall, these results are an important step forward in providing observational foundation for cosmological work with the cluster mass functions.  \\defcitealias{2007ApJS..172..561B}{Paper~I}  \\defcitealias{vikhlinin08}{Paper~II}  In this work, we present cosmological constraints from the data discussed in \\citetalias{vikhlinin08}. The cosmological information contained in the cluster mass function data and relevant to dark energy constraints can be approximately separated into 3 quasi-independent components:  (1)\\hspace*{0.5em}Changes in the comoving number density at a fixed mass threshold constrain a combination of the perturbations growth factor and relative distances between low and high-$z$ samples; this by itself is a dark energy constraint (\\S\\ref{sec:w0}).  (2)\\hspace*{0.5em}The overall normalization of the observed mass function constrains the amplitude of linear density perturbations at $z\\approx0$, usually expressed in terms of the $\\sigma_8$ parameter.  Statistical and systematic errors in the $\\sigma_8$ measurement are now sufficiently small, and the ratio of $\\sigma_8$ and the amplitude of the CMB fluctuations power spectrum gives the total growth of perturbations between $z\\approx1000$ and $z=0$ --- a second powerful dark energy constraint (\\S\\,\\ref{sec:other:data}).  (3)\\hspace*{0.5em}The slope of the mass function measures $\\OmegaM\\times h$; this by itself is not a dark energy probe but can be used to break degeneracies present in other methods.   Our dark energy constraints were derived for the following cases. Assuming constant $w$ and flat universe, we measure $\\wo=-1.14\\pm0.21$ using only cluster data (i.e., evolution of the mass function between our two redshift samples) and the HST prior on $h$ (\\S\\,\\ref{sec:wo:clusters}). Combining cluster and WMAP data, we obtain $\\wo=-1.08\\pm0.15$ but ($\\wo$ is constrained much more tightly for a fixed $\\OmegaM$ (\\S\\,\\ref{sec:wo:clusters+cmb}). Finally, adding cluster data to the joint supernovae + WMAP + BAO constraint, we obtain $\\wo=-0.991\\pm0.045$ (\\S\\,\\ref{sec:wo:clusters+others}), significantly reducing statistical and especially systematic (\\S\\,\\ref{sec:w:systematics}) uncertainties compared to the case without clusters. A large fraction of the extra constraining power comes from contrasting $\\sigma_8$ with normalization of the CMB power spectrum; this procedure is sensitive to non-zero mass of light neutrinos. Allowing for $m_\\nu>0$, we obtain a new conservative upper limit $\\sum m_\\nu<0.33$~eV (95\\% CL) while still improving the $w_0$ measurement relative to the SN+WMAP+BAO-only case ($\\wo=-1.02\\pm0.055$, \\S\\,\\ref{sec:mnu}). Adding clusters also improves equation of state constrains for evolving $w$ in flat universe (\\S\\,\\ref{sec:w0-wa}) and constant $w$ in non-flat universe (\\S\\,\\ref{sec:w0-Ok})  The paper is organized as follows. We start with a short summary of cluster data and systematic uncertainties (\\S\\,\\ref{sec:data}), discuss issues relevant for computing theoretical mass function models (\\S\\,\\ref{sec:theory}) and describe our fitting procedure (\\S\\,\\ref{sec:fitting}). We then discuss constraints that can be obtained from low-redshift mass function only ($\\OmegaMh$ in \\S\\,\\ref{sec:OmegaMh} and $\\sigma_8$ in \\S\\,\\ref{sec:s8}). We then consider as an example constraints from the cluster evolution in non-flat \\LCDM{} model (i.e., $w$ fixed at $-1$); $\\OmegaL>0$ is required with $\\sim 5\\sigma$ confidence (\\S\\,\\ref{sec:O-L}). Constraints on the dark energy equation of state are considered in \\S\\S\\,\\ref{sec:w0}--\\ref{sec:w:general}. Systematic errors are discussed in \\S\\,\\ref{sec:w:systematics}. ",
        "conclusions": "We presented constraints on the cosmological parameters from a new measurement on the galaxy cluster mass function in the redshift range $z=0-0.9$. All major sources of information contained in the cluster mass function --- its overall normalization and slope at $z=0$, and evolution at high redshifts --- are determined with our new data with a higher statistical accuracy and smaller systematic errors than before. This leads to much improved and more reliable constraints on the cosmological parameters.  From the normalization of the mass function estimated at low redshifts, we derive the $\\sigma_8$ parameter degenerate with $\\OmegaM$: $\\sigma_8(\\OmegaM/0.25)^{0.47} =0.813\\pm 0.013$ (stat) $\\pm0.024$ (sys). The slope of the low-$z$ mass function is a measure of $\\OmegaMh$: $\\OmegaMh=0.184\\pm0.037$; combined with the HST prior on $h$, this is an independent measurement of $\\OmegaM=0.255\\pm0.043$. The matter density can be independently measured with our cluster data using evolution of the temperature function, yielding consistent results, $\\OmegaM=0.30\\pm0.05$ in a flat \\LCDM{} model and $0.34\\pm0.08$ in a general cosmology.  Evolution of the mass functions between $z=0$ and $0.5$ (median redshift for our high-$z$ sample) constrains $\\OmegaL=0.83\\pm0.15$ in non-flat \\LCDM{} cosmology, or the dark energy equation of state parameter, $\\wo=-1.14\\pm0.21$, in a spatially flat Universe. Inclusion of the information provided by our cluster data also significantly improves the equation of state constraints obtained from combination of multiple cosmological datasets. For example, by combining the 5-year WMAP, most recent supernovae measurements, and detection of baryonic acoustic oscillations in the SDSS with our cluster data, we obtain $\\wo=-0.991\\pm0.045$ (stat) $\\pm0.040$ (sys); both the statistical and systematic errors in the combined constraint are a factor of $1.5-2$ smaller than those without clusters. Including cluster information also improves results for an evolving equation of state parameter and for constant $w$ in a non-flat universe. A spatially flat \\LCDM{} model is within the 68\\% CL interval from the best fit in all cases that we tested.  A good agreement between the geometric and growth of structure-based measurements of $w$ in principle can be used to place limits on modified gravity theories which attempt to explain cosmic acceleration without dark energy \\citep[e.g.,][]{2007PhRvD..76f3503W}. When self-consistent models of non-linear collapse in such theories become available, it sould be straightforward to use our cluster data in such tests also.  Comparison of the power spectrum normalization at $z=0$ obtained from clusters with the amplitude of the CMB fluctuations is a sensitive measure of the mass of light neutrinos. We constrain $\\sum m_\\nu < 0.33$~eV at 95\\% CL, at the expense of slightly weakening the measurement of $\\wo$ obtained assuming that the neutrino masses are negligibly small.  To facilitate the use of our cluster results in our cosmological studies, we provide at the project WWW site\\footnote{\\texttt{http://hea-www.harvard.edu/400d/CCCP}} machine readable tables of the likelihood function computed on several cosmological grids."
    },
    "0812/0812.2516_arXiv.txt": {
        "abstract": "We present the first results from two-dimensional simulations of radiatively-efficient accretion of metal-free gas onto intermediate-mass black holes.  We fix the shape of the spectral energy distribution of the radiation produced near the event horizon and study the structure of the irradiated low-angular-momentum accretion flow over three orders of magnitude in radius from the black hole, $10^{14}-10^{17}\\textrm{ cm}$ for a $100~M_\\odot$ black hole.  The luminosity of the central source is made to be proportional to  the rate at which gas accretes across the inner boundary, which we set just inside the sonic radius. We find that photoionization heating and radiation pressure modify the structure of the flow.  When the ambient gas density is $10^7\\textrm{ cm}^{-3}$, accretion is intermittent and on average reduced to $32\\%$ of the Eddington-limited rate, two orders of magnitude below the ``Bondi'' rate evaluated ignoring radiation, in agreement with simplified theoretical models.  Even if the vicinity of the black hole is supplied with high density gas, accretion is rendered inefficient through heating and radiation pressure. ",
        "introduction": "\\label{sec:introduction}  Known stellar-mass and intermediate-mass black hole candidates accrete through mass transfer from a stellar companion without exception. Rapidly-accreting massive black hole candidates, such as in Seyfert galaxies and quasars, appear to be situated close to dusty, molecular, and highly inhomogeneous gas reservoirs, where cold, dense clumps are in free fall or are embedded in a hot, ionized, and tenuous confining medium.  In neither case does the accretion flow resemble the laminar flow of the Bondi-Hoyle-Littleton solution. There is a distinct possibility, however, that massive black hole progenitor ``seeds,'' with masses $M_{\\rm bh}\\gtrsim 100~M_\\odot$ \\citep{Madau:01}, did pass through an early phase of quasiradial accretion in protogalaxies, prior to significant dust enrichment and the emergence of a pervasive cold phase in the interstellar medium \\citep[e.g.,][]{Volonteri:05,Alvarez:06,Alvarez:09,Johnson:07,Pelupessy:07,DiMatteo:08,Greif:08}. This was the period immediately following the onset of atomic cooling, the epoch when protogalaxies experienced rapid growth due to cosmic infall and yet had potential wells too shallow to retain all of their newly synthesized metals.  Since the accretion rate in quasiradial, Bondi-like accretion is proportional to gas density, Bondi-like accretion presents an appealing formation process for massive black holes: given the presence of sufficiently dense gas in a protogalactic nucleus, Eddington-limited accretion and mass exponentiation on a Salpeter time scale would proceed; stellar-size seed black holes would reach masses $\\gtrsim 10^9 M_\\odot$ inferred in high-redshift quasars in the cosmic time available to them \\citep[e.g.,][]{Haiman:01,Volonteri:03,Tanaka:08}.  A potential flaw in this reasoning was noticed early, namely, that in the event of radiatively-efficient accretion, radiative heating and radiation pressure may reduce the accretion rate \\citep{Shvartsman:71,Buff:74,Hatchett:76,Ostriker:76,Cowie:78,Stellingwerf:82,Begelman:85,Ricotti:08,Alvarez:09}.  The possibility of a transition to radiative-feedback-driven temporal intermittency, which has been evident in one-dimensional simulations of Compton-heated accretion onto massive black holes in galaxy clusters \\citep{Ciotti:01,Ciotti:07,Sazonov:05}, has further complicated the assessment of quasiradial accretion's viability as a massive black hole formation process.  If the accretion is indeed intermittent and episodic, how long is the duty cycle, how large is the average accretion rate, and does it permit evolution of seed black holes into quasars? Does episodic accretion proceed in bursts, where episodes of Eddington-limited rapid accretion alternate with near-complete radiative quenching and outflow, as one-dimensional models appear to suggest, or does the multidimensional nature of realistic irradiated accretion flows imply a damping of the fluctuations and convergence to a quasi-steady state characterized by a reduced accretion rate?  Here, we present the first results from our program to investigate radiatively-efficient accretion onto black holes embedded in a realistic galactic gaseous medium with the help of high-resolution hydrodynamic simulations. We do not attempt to simulate fluid flow and radiation emission mechanisms near the event horizon of the black hole; instead, we simply fix the shape of the spectral energy distribution of the radiation produced near the event horizon and simulate the irradiated fluid flow on radial scales spanning the sonic radius, where the photoheated fluid falls supersonically toward the black hole, and a much larger radius in the gas cloud that is entirely unaffected by the black hole's radiative output, where boundary conditions for the accretion flow are set.  This work is organized as follows. In \\S~\\ref{sec:numerical} we concisely describe our numerical algorithm. A more detailed description and a suite of tests will be provided in a longer follow-up paper (Couch et al., in preparation). In \\S~\\ref{sec:results} we describe our results with particular focus on the intermittency of the accretion flow.  In \\S~\\ref{sec:conclusions} we compare our results to analytical models \\citep[][and references therein]{Milosavljevic:08} and discuss the outlook for massive black hole formation through quasiradial accretion. ",
        "conclusions": "\\label{sec:conclusions}  We have simulated accretion from a uniform, high-density ($10^7\\textrm{ cm}^{-3}$)\\footnote{Atomic densities have been found to reach comparable high levels in simulations of protogalactic collapse \\citep[e.g.,][]{Bromm:03,Wise:07,Wise:08}.} metal-free protogalactic cloud onto an intermediate-mass black hole. At radii from the black hole smaller than the size of the excised region, $10^{14}\\textrm{ cm}$, the accretion was assumed to be radiatively efficient. We find that a compact \\ion{H}{2} region forms around the black hole.  The accretion is intermittent due to alternating photoionization radiation pressure-driven expulsion and external pressure-driven fallback. The average accretion rate is $\\lesssim 1\\%$ of the Bondi accretion rate calculated ignoring the radiation's influence, and $\\sim \\frac{1}{3}$ of the Eddington-limited rate.  The numerical results are consistent with the predictions of our toy model for episodic quasiradial accretion \\citep{Milosavljevic:08}, where we suggested that photoionization heating and photoionization radiation pressure within the compact  \\ion{H}{2} region are the dominant accretion-suppression mechanisms.  We also suggested that Ly$\\alpha$ line resonance radiation pressure, which we do not account for in the simulations presented here, should further reduce the accretion rate. The intermittency is qualitatively similar to that seen in one-dimensional simulations of the accretion of hot irradiated intergalactic medium onto black holes in central cluster galaxies by  \\citet{Sazonov:05} and \\citet{Ciotti:07}, with the notable caveat that photoionization radiation pressure was unimportant in that context.  The simulation in which the accreting fluid was endowed with small angular momentum corresponding to maximum circularization radius equal to $r_{\\rm hole}$  exhibited similar intermittency and only a slightly reduced average accretion rate, consistent with \\citet[and references therein]{Krumholz:05}.  In contrast with the nonrotating simulations that did not contain axial outflows, narrow polar outflows form that resemble some of those found by \\citet{Proga:07} and \\citet{Proga:08} in simulations of accretion flows irradiated by a quasar.  The accretion intermittency and inflow-outflow behavior that we observe is qualitatively similar to that seen in Proga et al., but additional tests are needed to ascertain whether out axial outflows are physical.  The overall goal is to combine the highly-resolved simulations presented here with large-scale cosmological simulations of how the first galaxies formed at the end of the cosmic dark ages \\citep[e.g.,][]{Greif:08}.  Those simulations can resolve the central gas flows to within $\\sim 0.01\\textrm{ pc}$, where we can implement sub-grid prescriptions based on the calculations done here.  An important challenge will be to derive prescriptions for the feedback from black hole accretion that are applicable to a wide range of protogalactic conditions.  We will report on these efforts elsewhere (Couch et al., in preparation)."
    },
    "0812/0812.3726_arXiv.txt": {
        "abstract": "Antarctica offers unique conditions for ground-based observations, such as low sky background in the infrared, improved seeing, and low turbulence and scintillation noise. These properties are particularly beneficial for imaging, precision photometry and infrared observations. It may be less clear if Antarctica offers equally compelling advantages for spectroscopy, in particular in the optical domain. However, scientific programmes that make use of imaging (or 3D) spectroscopy for selected follow-up studies of IR surveys, long-term monitoring of extended targets and resolved stellar population studies in crowded fields, also benefit from the site conditions at Dome C. ",
        "introduction": "The major motivations to conduct astronomical observations in Antarctica seem to be (wide-field) infrared surveys, diffraction-limited imaging and long-term accurate photometric monitoring. While for these programmes the conditions in space are superior, a ground-based site in Antarctica offers lower costs and the option to deploy larger telescopes and instruments. As spectroscopy requires larger collecting areas than imaging, it needs to be evaluated for which cases a median-sized telescope at Antarctica can outperform a larger one at temperate sites.  However, if spectroscopy were combined with imaging, as in 3D- or Integral-Field Spectroscopy (IFS), which provides multiple spectra for each point of a 2-dimensional field, the case for Antarctica becomes more compelling.  For extended targets, such as clusters, HII regions, nebulae, jets, or galactic environments, the use of integral-field techniques becomes necessary to overcome the source confusion in these complex fields. Integral field spectrophotometry is a useful technique for the study of resolved stellar populations in nearby galaxies (\\cite{Roth2004}). The spectroscopy of individual extragalactic stars is limited by the accuracy to subtract the bright, spatially and wavelength-dependent non-uniform background of the underlying galaxy. Imaging spectroscopy is able to improve the limited accuracy of background subtraction that one would normally obtain with slit spectroscopy. Resolved stellar populations are foreseen targets for the 2nd generation instrumentation at the VLT such as MUSE (\\cite{Bacon2006}), and provide a major science case for the planned ELTs (\\cite{Hook2007}).   A further case was made by Andersen (2007) for high-resolution imaging in the optical. Assuming a seeing of less than a 1/4 arcsecond for a significant fraction of the time, a 2.4m PILOT-like telescope (\\cite{Storey2007}) with tip-tilt correction is diffraction limited. Such a facility could pickup some of the science cases of HST after its de-commissioning. This statement equally holds for imaging spectroscopy, as long as the spatial sampling of the integral-field-unit (IFU) does not degrade the input conditions. Most likely such an IFU will be based on an image slicer type, as this technique preserves the spatial information best. Image slicers at cryogenic temperatures are being used in SPIFFI; advanced image slicers are planed for MUSE, which will operate in the wavelength regime of 465-930nm. After the de-commissioning of HST, an integral-field spectrograph even at a median-sized telescope in Antarctica would be a unique facility for diffraction-limited spatial spectroscopy. Given its southern location, high resolution spectroscopic studies in the SMC/LMC are possible objectives.   Many targets discovered in future IR-surveys will need spectroscopic follow-up and only a small fraction can be done from space. Given the wide-fields of IR-imaging surveys, integral field  spectroscopy over a large area would be desirable. A proposal was made by Maillard (2007) for a near-IR Imaging Fourier-Transform Spectrometer, that combines a wide-field with a high spectral resolution and can be used for Spitzer survey follow-ups. Less complicated, conventional IFUs can be used to study the details of complex star formation regions, for example in the Carina nebula or towards the Galactic Center (\\cite{Zinnecker2007}).   At Dome C, time-series data can be taken over months with duty cycles that are otherwise only possible using a network of telescopes. While the nights at the South Pole are longer, the seeing and the related scintillation noise are likely to be better at Dome C. Therefore, it shall be possible to obtain spectrophotometric data with greater precision than from any other ground-based site.   The spatial position of gamma ray bursts or X-ray sources is only known with a limited accuracy of a few arcseconds. Here, imaging spectroscopy provides an increased error budget to ensure that the target is not missed altogether. The measured 2-dimensional PSF allows one to disentangle the spectra of the central source and the surrounding host galaxy.   To detect and to measure the dynamics of low surface brightness objects, sufficient flux collection, rather than spatial resolution, is an issue. IFUs with high instrumental grasp, i.e. light collecting power (e.g. PPak-IFU, \\cite{Kelz2006}) and spatial binning of spectra can be applied to improve the signal-to-noise ratio.   For distant emission line sources (e.g. Ly-$\\alpha$ galaxies) with a priori unknown redshifts, 3DS is the only technique that can reliable detect these. The 3D-data cube corresponds to a volume in space, which otherwise can only be recorded with time-consuming scanning techniques.  Given the above scenarios, two strawman models for imaging spectrographs may be considered:  an IFU with high (diffraction-limited) spatial resolution but a small field of view and one with high light collecting power over a larger field.  \\begin{table}[h!] \\begin{tabular}{l|l|l} \\hline characteristics \t& high resolution IFU \t& high-grasp IFU \\\\ \\hline field of view\t\t& several arcseconds\t& few arcminutes  \\\\ spatial sampling\t& $\\approx 0.2$ arcsec\t& coarse ($>1''$)\\\\ & diffraction limited \t& high flux collection \\\\ IFU type\t\t& image slicer  \t& fiber bundle \\\\ & \t\t\t& (+ OH-surpression) \\\\ spectrograph\t\t& fixed setting \t& fiber-coupled \\\\ & winterised at telescope & remote, bench-mounted \\\\ & \t\t\t& multi-unit \\\\ \\hline application \t\t& crowded field spectroscopy & detection \\& dynamics of \\\\ & spectrophotometry & low surface brightness objects \\\\ & resolved populations & star formation regions \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "In summary, scientific motivations for imaging spectroscopy in Antarctica are found i) in regions where spectral windows are present that are not accessible anywhere else from the ground (\\cite{Burton2005}); ii) for either time-consuming or frequent spectroscopic measurements, for which sufficient time from space will not be available; iii) for high-precission spectrophotometry that needs uninterrupted time series or 2D-monitoring (e.g. planets); iv) for imaging spectroscopy in the optical with high spatial resolution (e.g. resolved population studies, AGN/SN/GRB host galaxies).  For the above applications, innovative imaging spectroscopy offers technical advantages, as compared to classical spectroscopy (\\cite{Kelz2004}). It relaxes operational requirements (e.g. acquisition accuracy, dispersion compensation), which simplifies the instrumental design and reduces the number of movable parts as potential sources of failure. Fiber-coupled instruments can be placed remotely from the telescope in warm conditions. Finally, fiber-spectrographs, similar to the ones used in astronomical spectroscopy, may be coupled to advanced fiberoptical sensors and can provide data for a broad range of other scientific research areas, such as geophysics, environmental monitoring, atmospheric and climate research. Therefore, techniques and infrastructure that are being developed in Astrophysics could be used for other, interdisciplinary projects at Dome-C as well. \\\\  \\noindent AK gratefully acknowledges support from ARENA to attend this conference."
    },
    "0812/0812.3984_arXiv.txt": {
        "abstract": "For very slow white dwarf accretors in CV's Townsley and Bildsten (2004) found a relation between the accretion rate $\\dot{M}$ and the central temperature of the white dwarf T$_c$. According to this relation for $\\dot{M}$ less than $10^{-10} M_{\\odot} yr^{-1}$ T$_c$ is much lower than $10^{7}$ K. Motivated by this study we follow the thermonuclear runaway on massive white dwarfs ($M_{WD}=1.25 - 1.40$ \\m) with T$_c$ lower than $10^{7}$ K, accreting matter of solar composition. We demonstrate that in that range of the relevant parameter space ($Tc,M_{WD}$ and $\\dot{M}$) the slope of the relation between the peak temperatures achieved during the runaway and T$_c$ becomes much steeper than its value for T$_c$ above $10^{7}$ K. The peak temperatures we derive can lead to nuclear breakout from the conventional \"hot carbon-nitrogen-oxygen\" cycle. When breakout conditions are achieved the heavy element abundances can show a much wider variety than what is possible with the common enrichment mechanisms. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.2046_arXiv.txt": {
        "abstract": "We present the modeling of partial frequency redistribution (PRD) effects for the fluorescent emission lines of molecular hydrogen, the general computational approximations, and the applications to planetary atmospheres, as well as interstellar medium. Our model is applied to $FUSE$ observations of Jupiter, Saturn, and reflection nebulae, allowing an independent confirmation of the H$_{2}$ abundance and the structure of planetary atmospheres. ",
        "introduction": "\\paragraph{Frequency Redistribution} $FUSE$ observations of planetary atmospheres and reflection nebulae have revealed the need to include PRD in H$_{2}$ fluorescence radiative transfer models. The angle-averaged laboratory frame redistribution function describes the conditional probability that a photon absorbed at x$_{i}$ Doppler widths from the center of line $i$, is emitted at x$_{f}$ Doppler widths from the center of line $f$ \\citep{Mihalas:1978}. Complete redistribution (CRD) refers to the photons being re-emitted in the line core according to the Voigt profile, while PRD takes into account the changes in the line profile due to coherent scattering in the line wings. PRD becomes important for fast transitions and integrated optical depths larger than 10$^{3}$. Observable effects include shifts in the peak wavelength in emission due to variations of the exciting spectrum over the absorbing line profile, a decrease in the line-to-continuum contrast, and a decrease in the number of photons scattered in subordinate lines (cross redistribution, XRD).\\\\[-0.25in]  \\paragraph{Radiative Transfer} The PRD radiative transfer problem is more complex than the CRD case, due to the heavy couplings in frequency in addition to spatial correlations and detailed balance \\citep{HubenyLites:1995,Uiten:2001}. Even restricting the geometry to the plane-parallel case, and assuming fixed level populations, the problem remains computationally expensive due to intrinsically large optical depths and fine frequency grids. We investigated a set of numerical methods (see Table~\\ref{table}) that would allow the full treatment of the H$_{2}$ molecule (10$^{4}$ transitions) to become feasible. We find that the multilayer approximation is the best approach, with a computing time independent of optical depth, and no matrix inversions required. This method is a layer-by-layer extension of the optically thin single layer solution of \\citet{LiuDalgarno:1996}. The results are consistently within 10\\% of the Feautrier lambda iteration \\citep{Mihalas:1978} for grids of at least 10 points per dex in optical depth.  \\begin{table}[h] \\begin{tabular}{lrrrr} \\hline & \\tablehead{1}{r}{n}{Lambda Iteration \\\\(Feautrier solution)} & \\tablehead{1}{r}{n}{Multilayer} & \\tablehead{1}{r}{n}{Single\\\\layer}   \\\\ \\hline Resources & High & Moderate & Low\\\\ Consistency & Can be unstable & Constrained & Constrained\\\\ Convergence & Scales with optical depth & 2-pass layer-by-layer & Single step\\\\ Spatial variations  & Yes & Yes & No \\\\ \\hline \\end{tabular} \\caption{Characteristics of the radiateve transfer methods used.} \\label{table} \\end{table} ",
        "conclusions": ""
    },
    "0812/0812.0333_arXiv.txt": {
        "abstract": "{In the study of the process of cosmic structure  formation numerical simulations are     crucial tools to interface observational data to theoretical models and to investigate issues otherwise unexplored.    Enormous advances have been achieved in the last years thanks to the availability of sophisticated  codes. The ever improving performances of large supercomputing facilities coupled with this efficiency of   codes are now allowing to tackle the problem of cosmic structure formation  and subsequent evolution by covering larger and larger dynamical ranges and to  provide  a progressively more realistic account of related complex astrophysical and cosmological processes. Moreover, computational cosmology is the ideal interpretative framework for the overwhelming amount of new data from   extragalactic surveys and from  large sample of  individual objects. The Workshop Novicosmo 2008 \"The Impact of Simulations in Cosmology and Galaxy Formation' held  in SISSA was  aimed at providing the state-of-the-art on the latest numerical simulations in Cosmology and in Galaxy Formation. Particular emphasis  was given to the implementation of new physical processes in simulation codes,  to the comparison between different codes and numerical schemes and  how to use  best  supercomputing facilities of the next generation. Finally, the impact on our knowledge on  the Physics of the Universe brought by this new channel of investigation has also  been focused. The Workshop was divided  in  three sections corresponding (roughly) to  three main areas of study: Reionization and Intergalactic medium;  Dark and Luminous matter in galaxies;  Clusters of galaxies and Large scale Structures.  This paper   will   provide  i) a short resume' of the  scientific results of the  Workshop ii) the complete  list of the talks and  the instructions on how  to  retrieve the .pdf of the related   (powerpoint)  presentations iii) a  brief presentation  of the associated Exhibition  \"Space Art\" }  \\vspace{0.2cm} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.3040_arXiv.txt": {
        "abstract": "We present results from integral field spectroscopy with the Potsdam Multi-Aperture Spectrograph of the head of the Herbig-Haro object \\hh\\ with a spatial sampling of 1$\\arcsec\\times$1$\\arcsec$. We have obtained maps of different emission lines, physical conditions --such as electron temperature and density-- and ionic abundances from recombination and collisionally excited lines. We present the first map of the Balmer temperature and of the temperature fluctuation parameter, \\tf. We have calculated the \\tf\\ in the plane of the sky, which is substantially smaller than that determined along the line of sight. We have mapped the abundance discrepancy factor of \\ioni{O}{2+}, \\adfo, finding its maximum value at the \\hh-S position. We have explored the relations between the \\adfo\\ and the electron density, the Balmer and [\\ion{O}{iii}] temperatures, the ionization degree as well as the \\tf\\ parameter. We do not find clear correlations between these properties and the results seem to support that the ADF and \\tf\\ are independent phenomena. We have found a weak negative correlation between the \\ioni{O}{2+} abundance determined from recombination lines and the temperature, which is the expected behaviour in an ionized nebula, hence it seems that there is not evidence for the presence of super-metal rich droplets in \\hii\\ regions. ",
        "introduction": "\\label{intro} The Orion Nebula is the nearest, most observed and studied Galactic \\hii\\ region. It is an active star-formation region where phenomena associated with the early stages of stellar evolution such as protoplanetary disks (proplyds) and Herbig-Haro (HH) objects can be observed in detail.\\\\ HH objects are bright nebulosities associated with high-velocity gas flows. The most prominent high-velocity feature in the nebula is the Becklin-Neugebauer/Kleinmann-Low (BN/KL) complex, which contains several HH objects (HH~201, 205, 206, 207 and 208). But also, there are other important high-velocity flows that do not belong to the BN/KL complex, as is the case of \\hh, 203 and 204. The origin of these flows has been associated with infrared sources embedded within the Orion Molecular Cloud 1 South \\citep[see][and references therein]{odellhenney08}.\\\\ \\hh\\ was one of the first HH objects identified in the Orion Nebula by \\cite{cantoetal80} in their [\\ion{S}{ii}] and [\\ion{N}{ii}] images as an emission line object showing two bright knots --the so-called \\hh-N and \\hh-S (see Figure \\ref{f1})-- embedded in a more extended nebulosity with a long concave form. After its discovery, the spectroscopic work by \\cite{meaburn86} showed that the arc-shaped nebulosity emits in [\\ion{O}{iii}] lines.\\\\ Further high-spectral resolution spectroscopy in several emission lines ([\\ion{O}{i}], [\\ion{S}{ii}], [\\ion{N}{ii}], [\\ion{O}{ii}], [\\ion{S}{iii}] and H$\\alpha$) has been performed in \\hh\\ by \\cite{odelletal91} who studied \\hh-S and found two velocity components. More recently, we can find in the literature extensive works on the gas kinematics in the Orion Nebula as those by \\cite{doietal02} and \\cite{odelldoi03}, where tangential velocities have been measured, or \\cite{doietal04}, who provide with accurate radial velocity measurements, detecting up to three kinematic components in \\hh-S. On the other hand, the most detailed imaging study of the Orion Nebula in the [\\ion{S}{ii}], [\\ion{O}{iii}] and H$\\alpha$ lines has been carried out by \\cite{odelletal97} with the Hubble Space Telescope ($HST$). These authors also detect an intense [\\ion{O}{iii}] emission in the extended nebulosity of \\hh\\ and a strong [\\ion{S}{ii}] emission at its knots. The extended [\\ion{O}{iii}] emission found in all the works indicates that the main excitation mechanism of \\hh\\ is photoionization, most probably from the brightest and hottest star of the Trapezium cluster, $\\theta^1$ Ori C. This particular property --not very common in other HH objects, which are usually ionized by shocks-- permits to derive the physical conditions and chemical abundances of the high velocity gas associated with \\hh\\ making use of the standard techniques for the study of ionized nebulae.\\\\ Our group has studied in detail the chemical composition of the Orion Nebula in some previous papers \\citep{estebanetal98,estebanetal04,mesadelgadoetal08} and is especially interested in studying the so-called abundance discrepancy (AD) problem in \\hii\\ regions, which is the disagreement between the abundances of the same ion derived from collisionally excited lines (CELs) and recombination lines (RLs). From intermediate and high resolution spectroscopy of Galactic and extragalactic \\hii\\ regions, our group has found that the \\ioni{O}{2+}/\\ioni{H}{+} ratio calculated from RLs is always between 0.1 and 0.3 dex higher than the value obtained from CELs for the same region \\citep[see][]{garciarojasesteban07}. The results obtained for \\hii\\ regions are quite different to those obtained for planetary nebulae (PNe), where the AD shows a much wider range of values and could be substantially larger in some objects \\citep[$e.g.$][]% {liuetal00,liuetal06,tsamisetal03b,tsamisetal04}. As argued in \\cite{garciarojasesteban07}, the AD problem in \\hii\\ regions seems to be consistent with the predictions of the temperature fluctuaction paradigm proposed by \\cite{peimbert67} and parametrized by the mean square of the spatial distributions of temperature, the so-called \\tf\\ parameter. In this scenario, the AD would be produced by the very different temperature dependence of the emissivities of both kinds of lines. In any case, the existence and origin of such temperature fluctuations are still controversial. \\cite{tsamispequignot05} and \\cite{stasinskaetal07} have proposed a different hypothesis in order to explain the AD in \\hii\\ regions. Based on the model for heavy-element mixing of \\cite{tenoriotagle96}, the presence of cold metal-rich droplets of supernova ejecta still not mixed with the ambient gas of the \\hii\\ regions would produce most of the RL emission while CEL emission would be produced by the ambient gas, wich would have the typical electron temperature of an \\hii\\ region and the expected composition of the local interstellar medium. In such case, the abundances from CELs would be the truly representative ones of the nebula.\\\\ Several authors have studied the spatial distribution of the physical conditions in the Orion Nebula. \\cite{odelletal03} have obtained a high-spatial resolution map of the electron temperature in a field centered at the southwest of the Trapezium cluster from narrowband images taken with the $HST$. They found small spatial scale temperature variations, which seem to be in quantitative agreement with the \\tf\\ parameter determined using the AD factor found by \\cite{estebanetal04}. On the other hand, \\cite{rubinetal03} have taken long-slit spectroscopy making use of $STIS$ at the $HST$ in several positions and do not find substantial spatial variations or gradients of the physical conditions along their slits positions. More recently, \\cite{sanchezetal07} have obtained a mosaic of the Orion Nebula from integral field spectroscopy with a spatial resolution of 2\\farcs7, but lacking a confident absolute flux calibration due to bad transparency during the observations. Their electron temperature map shows clear spatial variations and the electron density map is very rich in substructures, some of which could be related to HH objects. A more detailed view --but based on long-slit spectroscopy-- has been obtained by \\cite{mesadelgadoetal08}, who obtained spatial distributions --at spatial scales of 1\\farcs2-- of electron temperature, density, line intensities, ionic abundances and the AD factor (ADF, defined as the difference between the abundances derived from RLs and CELs) for five slit positions crossing different morphological zones of the Orion Nebula, finding spikes in the distribution of the electron temperature and density which are related to the position of proplyds and HH objects. In particular, these authors found that the ADF shows larger values at the HH objects.\\\\ The use of the integral field spectroscopy is still in its infancy in the study of ionized nebulae and the number of works available is still rather small. In the case of PNe, \\cite{tsamisetal08} have carried out the first deep study using this technique with FLAMES at the VLT taking spectra of three PNe of the Galactic disc and covering the spectral range from 3964 to 5078 \\AA. For the Orion Nebula, \\cite{tsamisetal08b} have presented preliminary results of integral field spectrocopy with FLAMES for several proplyds with the aim of studying the behaviour of the physical conditions, chemical abundances and the ADF in and outside these kinds of objects. The first results show that the temperature measured in the proplyds can be affected by collisional deexcitation as has been also suggested in previous works \\citep{rubinetal03,% mesadelgadoetal08}.\\\\ The main goal of the present paper is to use integral field spectroscopy at spatial scales of 1$\"$ in order to explore the effect of HH objects on the derived ADF considering the results of \\cite{mesadelgadoetal08}. \\\\ In \\S\\ref{obsred} we describe the observations obtained with PMAS and the reduction procedure. In \\S\\ref{medpmas} we describe the emission line measurements and the reddening correction as well as some representative bidimensional maps of those quantities. In \\S\\ref{resul} we describe the determination of the physical conditions, chemical abundances --from both kinds of lines, CELs and RLs-- and the ADF of \\ioni{O}{2+} as well as their corresponding maps. In \\S\\ref{discu} we calculate the \\tf\\ parameter from different methods and present the first bidimensional map of this quantity. Also, we show the possible correlations among the ADF and the rest of properties determined. Finally, in \\S\\ref{conclu} we summarize our main conclusions. \\begin{figure*} \\centering \\includegraphics[scale=0.7]{f1.ps} \\caption{$HST$ image of the central part of the Orion Nebula, which combines WFPC2 images taken in different filters \\citep{odellwong96}. The white square corresponds to the field of view of PMAS IFU used, covering the head of \\hh. The separate close-up image on the right shows the H$\\alpha$ map obtained with PMAS. The original map is 16$\\times$16 pixels of 1$\\arcsec\\times$1$\\arcsec$ size and has been rebinned to 160$\\times$160 pixels. Note the remarkable similarity between the $HST$ image and our rebinned PMAS H$\\alpha$ map.} \\label{f1} \\end{figure*} ",
        "conclusions": "\\label{conclu} In this paper, we present results from integral field spectroscopy of an area of 15$\\arcsec\\times$15$\\arcsec$ covering the head of \\hh\\ in the Orion Nebula. The FOV comprises the bright regions known as \\hh-S and \\hh-N. We have obtained maps of relevant emission line ratios, physical conditions and ionic abundances, including \\ioni{O}{2+} and \\ioni{C}{2+} abundances determined from recombination lines (RLs). Additionally, we have obtained maps of other interesting nebular properties, as the temperature fluctuation parameter, \\tf, and the abundance discrepancy factor of \\ioni{O}{2+}, \\adfo, which is defined as the difference between the \\ioni{O}{2+}/\\ioni{H}{+} ratios determined from RLs and collisionally excited lines (CELs).\\\\ We find that the flux distributions of the [\\ion{O}{iii}] and \\ion{O}{ii} lines are rather similar and that the HH object is comparatively much brighter in the lines of low-ionization potential ionic species and especially in [\\ion{Fe}{iii}]. The \\nel\\ is about 4000 \\cmc\\ in most of the FOV and higher --about 10000 \\cmc-- at \\hh-S and \\hh-N. The \\te([\\ion{O}{iii}]) map shows a narrow range of variation, but the values are higher at \\hh-S. On the other hand, \\te([\\ion{N}{ii}]) shows larger variations and a very different spatial distribution, being higher at the northern and eastern edges of \\hh-N and \\hh-S, respectively, a likely consequence of the ionization stratification or the presence of some shock excitation in the knots of the HH object. We have obtained --for the first time in an ionized nebula-- the \\te(Bac) map, which follows closely that of \\te([\\ion{O}{iii}]).\\\\ The \\ioni{O}{+}/\\ioni{H}{+} ratio map reaches the highest values just on the arc that delineates the north of \\hh\\ from the east and west edges of the FOV, whereas the map of \\ioni{O}{2+} abundance obtained from CELs shows an inverse behaviour, probably produced by the higher densities in the shock gas that increase the recombination rate of \\ioni{O}{2+}. The spatial distributions of the \\ioni{O}{2+} abundance obtained from CELs and RLs agree in showing the lower values at the positions of \\hh-S and \\hh-N and the arc connecting both features. However, for a given spaxel, the \\ioni{O}{2+}/\\ioni{H}{+} ratios obtained from RLs are always about 0.10 dex higher than those obtained from CELs. The map of the \\adfo\\ shows the highest values at and around \\hh-S.\\\\ We have determined --for the first time in an ionized nebula-- the \\tf\\ map of the FOV from the comparison of \\te(Bac) and the other \\te\\ values determined from collisionally excited line ratios, finding that it does not match the \\adfo\\ map. This result supports that the AD and temperature fluctuations are independent phenomena. \\\\ Finally, we have found weak correlations between the \\adfo\\ and \\te\\ values of the different spaxels of the FOV and between the \\ioni{O}{2+}/\\ioni{H}{+} ratios obtained from RLs and CELs and \\te([\\ion{O}{iii}]). The negative correlation between the \\ioni{O}{2+} abundances obtained from RLs and \\te([\\ion{O}{iii}]) does not support the predictions of the chemical inhomogeneity models of \\cite{tsamispequignot05} and \\cite{stasinskaetal07}."
    },
    "0812/0812.3795_arXiv.txt": {
        "abstract": "We repeat and extend the analysis of Eriksen et al 2004 and Hansen et al 2004 testing the isotropy of the Cosmic Microwave Background (CMB) fluctuations. We find that the hemispherical power asymmetry previously reported for the largest scales $\\ell=2-40$ extend to much smaller scales. In fact, for the full multipole range $\\ell=2-600$, significantly more power is found in the hemisphere centered at $(\\theta=107^\\circ\\pm10^\\circ,\\phi=226^\\circ\\pm10^\\circ)$ in galactic co-latitude and longitude than in the opposite hemisphere consistent with the previously detected direction of asymmetry for $\\ell=2-40$. We adopt a model selection test where the direction and amplitude of asymmetry as well as the multipole range are free parameters. A model with an asymmetric distribution of power for $\\ell=2-600$ is found to be preferred over the isotropic model at the $0.4\\%$ significance level taking into account the additional parameters required to describe it. A similar direction of asymmetry is found independently in all six subranges of 100 multipoles between $\\ell=2-600$ and none of our 9800 isotropic simulated maps show a similarly consistent direction of asymmetry over such a large multipole range. No known systematic effects or foregrounds are found to be able to explain the asymmetry. ",
        "introduction": "\\label{sect:intro}  In the first public release of the WMAP data, we reported a significant asymmetry of the distribution of large scale power on the sky\\citep{asymm1,asymm2}. This finding as well as other statistical anomalies were confirmed by several other authors using other methods \\citep{alignment,ilc2,park,vielva,minkowski,axisofevil,bispectrum,spots1,directionalw,curvature,parameters}. Some of these have been confirmed again in the 3 and 5 year data \\citep{wmap32,wmap34,cruz,directionalw2,eriksenmod,pietrobon}. In \\citep{asymm2} we found that the independent multipole ranges $\\ell=2-19$ and $\\ell=20-40$ were both particularly asymmetric but with two different axes of asymmetry, the former being a galactic north-south asymmetry and the latter being an east-west asymmetry. Using the full range $\\ell=2-40$ we found the highest  significance of the asymmetry with the axis pointing in the direction of $(100^\\circ,237^\\circ)$ in galactic co-latitude and longitude (this is the convention we will use for all sky positions in this paper).  The position of the non-Gaussian cold spot found in \\citep{vielva} is found to be positioned close to the center of the hemisphere with large fluctuation power. Furthermore, the CMB signal also demonstrated non-Gaussian statistical properties in the direction where the power spectrum amplitude was low \\citep{park,minkowski,curvature}. Our own attempts to provide an explanation to the asymmetry invoked a class of homogeneous models that include anisotropic expansion (shear) and global rotation (vorticity) - the Bianchi type VIIh models\\citep{tess1,tess2}. Unfortunately, this fails as a physical explanation\\citep{tess3,bridges} since the best-fit parameter space is then inconsistent with a wide range of other evidence.  Recent theoretical developments\\citep{ackerman,erickcek,gordon} have proposed new mechanisms for generating the imprint of a preferred direction in the CMB. Some of these models would lead to hemispherical asymmetry also for multipoles much larger than $\\ell=40$ (angular scales smaller than 5 degrees). A recent analysis\\citep{nicolaas} of the Ackerman et al. model indicates a good fit to the WMAP data up to a multipole moment $\\ell=400$.  However, this model is not able to address the issue of the hemispherical power asymmetry.  In this paper we reinvestigate the nature of the power asymmetry using the recent 5 year data release from WMAP \\citep{wmap5}, and specifically assess whether the signal manifests itself on smaller angular scales $\\ell>>40$. Instead of randomly searching for an asymmetry in arbitrary multipole ranges, we adopt a model selection procedure that searches for the best fit asymmetry direction and multipole interval. An asymmetric model of the CMB needs more free parameters (i.e. direction of the axis of asymmetry) than a simpler isotropic model. The model selection procedure tests whether an asymmetric model of the CMB is actually preferred by the data taking into account the introduction of additional free parameters. We introduce a set of general asymmetric models with three to five additional parameters. These models are not based on any physical theory but are rather general parametrization of hemispherical asymmetry. We use these models in order to investigate how the asymmetry is distributed in harmonic space as well as on the sphere.   In section \\ref{sect:data} we describe the data and masks used in the analysis. The methods used to assess the asymmetry are outlined in section \\ref{sect:method}. In section \\ref{sect:res} we show the results of the isotropy tests applied to the 5 year WMAP data and in section \\ref{sect:conclusions} we conclude. ",
        "conclusions": "\\label{sect:conclusions}  We have reassessed the asymmetry in the distribution of CMB fluctuation  power on the sky reported in \\cite{asymm1} and \\cite{asymm2}. In order to test whether an anisotropic model of the CMB fluctuations is actually preferred over an isotropic model taking into account the additional parameters required, we implement a new model selection procedure. We model the asymmetric distribution of power on the sky as a dipole in the power distribution. Note that this is not to say that the CMB fluctuation field has a dipole, rather that the power for a certain scale (multipole) has a dipole distribution on the sky. We use a model where there is a common dipole component in the power distribution for a set of multipoles in a range $[\\ell_\\mathrm{min},\\ell_\\mathrm{max}]$ where $\\ell_\\mathrm{min}$ and $\\ell_\\mathrm{max}$ are free parameters as well as the direction ($\\theta,\\phi$) and the amplitude $A_0$ of the dipole. We use a $\\chi^2$ approach to find the best fit model parameters among these 5 parameters.  We first investigated a model where we assume the asymmetry to start at $\\ell_\\mathrm{min}=2$, reducing the number of free parameters in the model to 4. Using power spectra estimated on hemispheres in the combined V+W band with the KQ85 galactic cut,  we find a strong asymmetry in the multipole range $\\ell=2-221$ with an axis pointing in the direction $(\\theta=107^\\circ\\pm10^\\circ,\\phi=226^\\circ\\pm10^\\circ)$ (which is the direction where the power is largest). Only $0.3\\%$ of the simulated isotropic maps show a similarly strong asymmetry. Performing the same test on the Q, V and W bands individually as well as with different galactic cuts, a similar asymmetry is found. The significance is reduced for larger cuts but for an extended KQ75 cut (excluding $37\\%$ of the sky), only $1.5\\%$ of the simulated isotropic map show a simiarly strong common dipole component in the range $\\ell=2-300$. Using more localized power spectra estimated on smaller disks, we perform the same tests on the V+W band using the KQ85 cut. Smaller disks allow faster spherical harmonic transforms and allows the analysis to include multipoles up to $\\ell=800$. We find that the range $\\ell=2-600$ is asymmetric with dipole direction  $(\\theta=107^\\circ\\pm11^\\circ,\\phi=216^\\circ\\pm10^\\circ)$ using the smallest $22.5^\\circ$ diameter disks. Only $0.4\\%$ of the simulated maps show similar asymmetry. In figure \\ref{fig:cl} we show the spectra in the two opposite parts of the sky. The spectra are clearly different for the largest scales as well as around the two first peaks.   Including $\\ell_\\mathrm{min}$ as a free parameter, the same anisotropic model is favored and the best fit value for the first multipole with a common dipole is $\\ell_\\mathrm{min}=2$. We therefore concluded that the 4 parameter model was sufficient to describe the asymmetry.  Testing models with a multipole dependent dipole amplitude, we found that a model where the asymmetry is maximum for small $\\ell$ and decreasing with increasing $\\ell$ vanishing at about $\\ell\\sim650$ gives a good fit to the data ($0.4\\%$).  To check whether the asymmetry is present in the full range $\\ell=2-600$ or only for some multipoles, we performed a second simpler test of asymmetry. We found the dipole direction for the power distribution in multipole ranges of 100 multipoles, $\\ell=2-101$, $\\ell=102-201$, etc. to $\\ell=502-601$. We thus obtained 6 dipole directions from 6 independent multipole ranges. The power distribution in these ranges are shown in figure \\ref{fig:directions}. We found that these 6 dipoles were much more aligned in the WMAP data than in isotropic simulations. In fact, using power spectra estimated on $22.5^\\circ$ diameter disks from the combined V+W bands with KQ85 galactic cut we found that none of our 9800 simulations show a similarly strong alignment between these 6 dipoles. The dipole directions for these 6 ranges are shown in figure \\ref{fig:directions3}. We find that the spatial distribution of CMB fluctuations is strongly correlated between small and large angular scales.  The fact that all the frequency channels,  all years of observations and also tests using cross power spectra show a similarly asymmetric distribution strongly disfavors an explanation in terms of systematic effects and residual galactic foregrounds. Further, in the 5 year WMAP data a different approach to foreground subtraction than for the first year data was applied, still the asymmetry remained significant at the same level for the large scale asymmetry $\\ell=2-40$. In \\citep{wifit} we also showed that a blind approach to foreground subtraction did not change this result for large scales.  Our results indicate that the reported common asymmetric axis extending over a large range in scales is highly unlikely to be a statistical fluke. Foregrounds and systematic effects do not seem to be probable explanations. The CMB does seem to have an uneven power distribution on the sky over a large range of angular scales. An important task for further research is to find a physical explanation for this asymmetry which can predict possible effects on CMB polarization to be tested in future experiments."
    },
    "0812/0812.2471_arXiv.txt": {
        "abstract": "We present fully three-dimensional local simulations of compressible MRI turbulence with the object of studying and elucidating the excitation of the non-axisymmetric spiral density waves that are observed to always be present in such simulations. They are potentially important for affecting protoplanetary migration through the action of associated stochastic gravitational forces and producing residual transport in MHD inactive regions through which they may propagate. The simulations we perform are with zero net flux and produce mean activity levels corresponding to the Shakura \\& Sunyaev \\mbox{$\\alpha{}\\sim{}5\\times{}10^{-3}$}, being at the lower end of the range usually considered in accretion disc modelling. We reveal the nature of the mechanism responsible for the excitation of these waves by determining the time dependent evolution of the Fourier transforms of the participating state variables. The dominant waves are found to have no vertical structure and to be excited during periodically repeating swings in which they change from leading to trailing. The initial phase of the evolution of such a swing is found to be in excellent agreement with that expected from the WKBJ theory developed in a preceding paper by Heinemann \\& Papaloizou. However, shortly after the attainment of the expected maximum wave amplitude, the waves begin to be damped on account of the formation of weak shocks. As expected from the theory the waves are seen to shorten in radial wavelength as they propagate. This feature enables nonlinear dissipation to continue in spite of amplitude decrease. As a consequence the waves are almost always seen to be in the non linear regime.  We demonstrate that the important source terms causing excitation of the waves are related to a quantity that reduces to the potential vorticity for small perturbations from the background state with no vertical dependence. We find that the root mean square density fluctuations associated with the waves are positively correlated with both this quantity and the general level of hydromagnetic turbulence. The mean angular momentum transport associated with spiral density waves generated in our simulations is estimated to be a significant fraction of that associated with the turbulent Reynolds stress. ",
        "introduction": "\\label{sec:introduction} This is the second of two papers which study the excitation of spiral density (SD) waves in accretion discs occurring through the action of stochastic forcing due to turbulence. In practice this has been taken to arise from the magneto-rotational instability (MRI) but excitation resulting from stochastic forcing produced by other mechanisms is expected to lead to similar results. In the preceding paper \\citep[][paper I]{2008arXiv0812.2068H} we developed a WKBJ theory of the excitation process. In this paper we study the wave excitation process directly as it is manifested in fully nonlinear three dimensional numerical simulations, with the object of elucidating it further and comparing the results with the predictions of the WKBJ theory developed in paper~I.  The outline of this paper is as follows. In Section~\\ref{sec:preliminaries} we give a brief review of the shearing box model, the basic equations solved in the simulation, as well as the equations governing the dynamics of SD waves. We also describe the procedure for performing the Fourier transforms of the state variables participating in the excited waves in shearing coordinates. This allows us to identify and then follow the evolution of wave amplitudes in order to make a detailed comparison with the WKBJ theory presented in paper~I.  In Section~\\ref{sec:simulations}, after briefly describing the simulation set up and the numerical method used, we present a fully three-dimensional simulation of compressible MRI turbulence in which we directly observe the excitation of large scale, non-axisymmetric SD waves. We discuss their main characteristics and the correlation of the root mean square density fluctuations associated with them with the level of turbulent activity. We go on to study the source terms driving the wave excitation and motivate the simplifying assumption that those proportional to a quantity that is related to the potential vorticity dominate. This assumption enabled us to derive the linear theory of wave excitation presented in paper~I.  A detailed comparison between this linear theory and simulation outcomes is presented in Section~\\ref{sec:comparison}. There we follow the time dependent evolution of the appropriate Fourier amplitudes through a swing cycle, during which the wave excitation occurs, in detail. We obtain very good agreement with predictions from the theory developed in paper~I. Finally, we discuss and summarize our results in Section~\\ref{sec:discussion}, indicating the magnitude and the scaling of the wave angular momentum flux with the Reynolds stress and also the wavelength in the azimuthal direction for which the wave excitation phenomenon is optimal. ",
        "conclusions": "\\label{sec:discussion} In this paper we have presented fully three-dimensional simulations of compressible MRI turbulence in which we observed the excitation of large scale, non-axisymmetric SD waves. The root mean square density fluctuations associated with these waves were shown to be correlated with the degree of turbulent activity measured by the \\citet{1973A&A....24..337S} $\\alpha$ parameter, \\begin{displaymath} \\alpha = \\frac{\\left<B_x B_y - \\rho u_x u_y\\right>_V}{\\rho_0 c^2}. \\end{displaymath}  We performed a detailed analysis of the excitation of these waves by following the time dependent evolution of the Fourier transforms of the participating state variables. Their evolution during the excitation phase of the associated SD wave was found to be in excellent agreement with the WKBJ theory developed in paper~I. However, shortly after the attainment of the maximum wave amplitude, the waves are rapidly damped. This damping is consistent with the sharp density profiles observed in the simulations which are indicative of the existence of shocks. Shocks are expected even under conditions of weak excitation because as shearing waves propagate their radial wavelength shortens, resulting in the onset of nonlinearity and shock formation \\citep[e.g.][]{2001ApJ...552..793G}. Thus the waves are always likely to be observed in the nonlinear regime. Our results indicate that for activity levels corresponding to \\mbox{$\\alpha\\sim{}5\\times{}10^{-3}$}, being at the lower range of magnetic activities normally considered, shock formation occurs when their radial wavelength is comparable to the scale height.  The excited waves are trailing and associated with an outward angular momentum flux. The calculations performed here are consistent with the WKBJ theory presented in paper~I in indicating that this outward flux scales in the same way as the Reynolds stress. This is supported by the fact that the density fluctuations associated with the waves are correlated with it. The level of angular momentum transport associated with SD waves is estimated to be a significant fraction of that due to the turbulent Reynolds stress $\\left<-\\rho u_x u_y\\right>_V$.  By consideration of the source terms for the wave excitation, we identified those associated with the pseudo potential vorticity (PPV) as the dominant ones as was assumed in paper~I. Wave like structure was not seen in the PPV field, see Fig.~\\ref{fig:zaverage}, indicating that this could act consistently as an independent source generated by background turbulent fluctuations.  For the simulation discussed in detail, SD wave excitation is found to be most effective for waves with the smallest azimuthal wave number available in the computational domain. This can be understood from the theory presented in paper~I. There it was found that the amplitude of an excited SD wave is directly proportional to the Fourier amplitude of the pseudo potential vorticity at the swing stage, $\\hat{\\zeta}_\\mathrm{s}$, while also having an exponential dependence on the azimuthal wave number $k_y$ through the parameter \\begin{equation*} \\epsilon = \\frac{q\\Omega k_y c}{k_y^2 c^2 + \\kappa^2}. \\end{equation*} If the latter dependence alone is important it implies that, for a Keplerian rotation law, SD wave production will be most effective for $k_y\\sim{}k_y^\\mathrm{opt}=\\Omega/c$. This is longer than the longest possible azimuthal wavelength in the shearing box considered here with $L_y=4H$, so that SD wave production should most effective at the longest azimuthal wavelength.  The other important dependence is that of the PPV spectrum on $k_y$ at the time of swing. This spectrum is difficult to predict theoretically, but easily determined from the simulation data. We do so by taking for each $k_y$ an average over the ensemble of shearing waves that we find to swing from leading to trailing during the course of the simulation. The result is shown in Fig.~\\ref{fig:zeta-spectrum}. \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{figures/zeta-spectrum.pdf} \\end{center} \\caption{The spectrum of swing pseudo potential vorticity averaged over all swings as a function of azimuthal wave number $k_y$. The spectrum peaks at the smallest $k_y$, remains rather shallow at moderately small $k_y$, but falls off for high $k_y$.} \\label{fig:zeta-spectrum} \\end{figure} We see that the spectrum is rather shallow on large scales and falls off at small scales. A nearly flat spectrum at small $k_y$ means that the dependence of the wave amplitude on $k_y,$ at the most effective values, is determined by the WKBJ exponential factor appearing in e.g.\\ (\\ref{eq:limiting-value-xi}), and should be maximal at either \\mbox{$k_y=1/H$}, or the smallest possible $k_y$ in the box should that be larger. Accordingly we expect the smallest $k_y$ to be dominant for $L_y=4H$ but this dominance should begin to be lost once $L_y$ starts to exceed $2\\pi{}H.$ We checked that when $L_y$ was extended to $8H,$ the contribution from the smallest $k_y$ no longer exceeded that from the next smallest confirming the discussion given above.  Differentially rotating discs undergoing MRI driven turbulence are ubiquitous in astrophysics and the results presented here and in paper~I indicate that the production of density waves is generic. The density perturbations associate with these may play an important role in driving protoplanetary migration \\citep[e.g.][]{2004MNRAS.350..849N,2005A&A...443.1067N} and may play a role in exciting quasi periodic oscillations in discs around close binary stars. In the case of protoplanetary discs it has been suggested that only the high altitude regions may be turbulent while the mid-plane region lies in a `dead zone' \\citep[e.g.][]{1996ApJ...462..725G}. We comment that we found that the SD waves excited by the turbulence have very weak vertical dependence. Although we considered an unstratified disc, SD waves with weak vertical dependence also exist in the isothermal stratified case \\citep[e.g.][]{2007A&A...468....1F}, so that although they may be excited in the upper layers propagation throughout the vertical structure is to be expected with the consequence that some small but non zero transport may result throughout the disc. In this context we note that, consistently with these ideas, recent simulations by \\citet{2007ApJ...670..805O} that incorporated dead zones were found to have wavelike motions excited in the mid-plane regions associated with a small residual angular momentum flux that was uniformly distributed in the dead zone with a typical $\\alpha$ value between $10^{-4}$ and $10^{-5}$. We remark that a simple application of the WKBJ theory developed in paper~I suggests, since the source of wave excitation is a vertical average of the pseudo potential vorticity, and the wave angular momentum flux is the square of this, that the value of $\\alpha$ in the dead zone should scale as ratio of the square of the surface density in the high altitude turbulent regions to the surface density of the whole disc.  However, this may be a lower bound because the dependence on surface density may actually be weaker on account of weaker dissipation in the linear regime allowing the build up of a relatively larger amount of wave angular momentum flux. This aspect remains to be investigated and decided by future much higher resolution simulations that we are unable to perform at the present time."
    },
    "0812/0812.4247.txt": {
        "abstract": "We calculate the primordial black hole (PBH) mass spectrum produced from a collapse of the primordial density fluctuations in the early Universe using, as an input, several theoretical models giving the curvature perturbation power spectra ${\\cal P}_{\\cal R}(k)$ with large ($\\sim 10^{-2} \\div 10^{-1}$) values at some scale of comoving wave numbers $k$. In the calculation we take into account the explicit dependence of gravitational (Bardeen) potential on time. Using the PBH mass spectra, we further calculate the neutrino and photon energy spectra in extragalactic space from evaporation of light PBHs, and the energy density fraction contained in PBHs today (for heavier PBHs). We obtain the constraints on the model parameters using available experimental data (including data on neutrino and photon cosmic backgrounds). We briefly discuss the possibility that the observed 511 keV line from the Galactic center is produced by annihilation of positrons evaporated by PBHs. ",
        "introduction": "It is well known that for a sufficient production of primordial black holes (PBHs) in the early Universe \\cite{Zeldovich1967, Hawking1971, Carr:2005zd, Khlopov:2008qy} the spectrum of the density perturbations set down by inflation must be ``blue'', i.e., it must have more power on small scales. This implies that the spectral index of the scalar perturbations must be larger than 1, in strong contradiction with the latest WMAP results \\cite{Spergel:2006hy, Dunkley:2008ie, Komatsu:2008hk}. In particular, standard inflationary models of hybrid type, in which the inflaton is trapped in a local minimum of the potential and which predict blue spectra seem to be excluded as a possible source of PBHs.  Such a conclusion is correct, however, only in rather special case: namely, it is based on the prediction of slow-roll single field inflationary scenario, according to which the power spectrum of curvature perturbations is nearly scale-invariant, i.e., the spectral index $n$ is close to unity and the variation in the spectral index $dn/d \\log k$ is small.  Although a prediction of the approximate scale invariance of the primordial power spectrum is a necessary requirement to any inflationary model, some deviations from pure scale invariance are consistent with the observational data. These deviations are described by adding localized features to the primordial spectrum (see, e.g., \\cite{Hoi:2007sf} and references therein) and/or by introducing spectral features modifying a single power law. Models with such peculiarities (sometimes called broken-scale-invariant (BSI) models) were proposed, in main aspects, in eighties \\cite{Starobinsky:1986fxa, Kofman:1985aw, Kofman:1986wm, Silk:1986vc, Kofman:1988xg, Salopek:1988qh}. Such models generally include, in addition to the usual inflaton field, other scalar fields driving successive stages of inflation and triggering phase transitions.  Evidently, the BSI models of inflation could predict, generically, the essential production of primordial black holes at small and medium scales. In particular, in \\cite{Starobinsky:1992ts} the inflationary model with a steplike power spectrum based on a nonanalyticity in the inflaton potential was proposed and in subsequent works the inflationary potentials of such kind were used for making predictions of PBH production \\cite{Ivanov:1994pa, Ivanov:1997ia, Bullock:1996at}.  A step feature in the potential is an effective field theory description of a phase transition, so, in a more realistic approach an inflationary model should certainly involve more than one scalar field, and such rapid phenomena as phase transitions \\cite{Starobinsky:1998mj}. Second order phase transitions during inflationary expansion had been first considered in \\cite{Kofman:1986wm, Kofman:1988xg} in models with two scalar fields. In scenarios of such type, during a short stage, corresponding to the beginning of a phase transition, the mass of the trigger field becomes negative, and adiabatic perturbations are exponentially amplified resulting in the formation of a narrow spike in the primordial spectrum, and, as a consequence, in a copious production of PBHs \\cite{GarciaBellido:1996qt, Randall:1995dj}. There are many multiple field scenarios predicting the existence of spike- or bumplike features in the primordial spectrum (e.g., supersymmetric double hybrid models \\cite{Lesgourgues:1999uc}, multiple inflation models based on supergravity \\cite{Adams:1997de}, etc). Some of these models are specially constructed to predict efficient PBH production \\cite{Yamaguchi, Kawasaki}.  Further, large curvature perturbations (especially on subhorizon scales) leading to features in the primordial spectrum and possible PBH production, can arise in multiple field scenarios at the end of inflation (during preheating era) or between two consecutive stages of inflation, as a result of parametric resonance or tachyonic instability. Estimates of PBH production possibilities had been obtained in  \\cite{Green:2000he, Bassett:2000ha, Finelli:2001db} for two-field models and in \\cite{Suyama:2006sr} for models with a self-interacting scalar field. The more complicated model (containing three fields), based on supergravity, is considered in \\cite{Kawasaki}.  An existence of the narrow spikes in the primordial spectrum is possible not only in multiple field inflationary scenarios. Such a feature can, in principle, exist even in single field models (see, e.g., \\cite{Saito:2008em, Bugaev:2008bi}). If, in particular, the inflationary potential has an unstable maximum at origin (e.g., the double-well potential) then, with some fine-tuning of parameters and initial conditions, the inflation process may have two stages, with a temporary stay at the maximum, that may lead to the corresponding peak in the primordial spectrum and, depending on the amplitude of the peak, to the PBH production.  %As an alternative one can consider the inflation models in which the power %on small scales is boosted by means of a bump in the power %spectrum of primordial fluctuations, as suggested, e.g., in %\\cite{Ivanov:1994pa, Randall:1995dj, GarciaBellido:1996qt, Bringmann:2001yp}. Such a % bump is predicted in many scenarios of two-step inflation %\\cite{Yamaguchi, Kawasaki} and can, in principle, exist even in %single-field inflationary models (see, e.g., \\cite{Saito:2008em, Bugaev:2008bi}).  Another example of an inflationary model predicting large amplitudes of the density perturbations at small scales is the inflationary model with the running mass potential. More generally, large amplitudes of the primordial curvature perturbation spectrum at small scales are possible in models of \"hilltop hybrid inflation\" \\cite{Kohri:2007gq}. Potentials of these models have concave-downward form at cosmological scales (corresponding to red primordial spectrum as required by observations), but can be much flatter near the end of inflation. Such a modification of a standard hybrid inflation scenario is discussed, e.g., in \\cite{NeferSenoguz:2008nn, Rehman:2009wv}. The running mass potential belongs to the class of hilltop potentials; in models with such potentials, due to spectral index scale-dependence, the amplitude of the perturbation spectrum at small scales depends on the value of $dn/ d\\log k$ at cosmological scales \\cite{Kohri:2007gq}. It is shown in our previous work \\cite{Bugaev:2008bi} that possibilities of noticeable production of PBHs in running mass model are still open in spite of the rather severe experimental constraints \\cite{Kohri:2007qn, Peiris:2008be} on the value of the spectral index running.  In this paper we limit ourselves to the situations where PBHs form from the density perturbations, induced by quantum vacuum fluctuations during inflationary expansion. The details of the PBH formation had been studied in \\cite{Carr:1975qj, Khlopov:1980mg}, the astrophysical and cosmological constraints on the PBH density had been obtained in many subsequent works (see, e.g., the recent review \\cite{Khlopov:2008qy}). The order of magnitude of the corresponding constraint on the value of the density perturbation amplitude (for the PBH mass region $10^{11} - 10^{18}$ g, which we are interested in) is well known \\cite{Carr:1994ar}, but, if the primordial spectrum contains the peaklike feature, the concrete value of the PBH constraint clearly depends on the parameters characterizing the form of this feature (in particular, on the width of the peak). Such an information may be rather useful for the model makers.  One must note that the formation of PBHs in models with the primordial spectrum having the features had been studied earlier, in works \\cite{Yokoyama:1998xd, Blais:2002gw}. In \\cite{Yokoyama:1998xd} the case of the spectrum that is sharply peaked on a single mass scale was considered. The $\\beta(M)$ function (the probability of a region of mass $M$ to form a PBH) was calculated and it was shown that the form of this function depends on the type of the gravitational collapse. In \\cite{Blais:2002gw} two cases were studied: a pure step in the primordial spectrum and the spectrum produced in an inflationary model \\cite{Starobinsky:1992ts} with a jump in a first derivative of the inflaton potential at some scale. It was shown that the corresponding $\\beta$-functions have the pronounced bumps and, in connection with these bumps, authors of \\cite{Blais:2002gw} discuss the possibility that PBHs can be a significant part of dark matter.   In the present work we reconsider the problem of constraining the power spectrum of the primordial fluctuations (with accent on the peaklike features)  calculating the process of the formation of PBHs having small masses ($\\sim 10^{11} - 10^{18}$g). Products of evaporation of these PBHs contribute, in particular, to extragalactic photon and neutrino diffuse backgrounds (which are measured experimentally), and the constraints are calculated using the standard procedure \\cite{Page:1976wx}. We do not calculate the $\\beta$-functions preferring to constrain the power spectrum directly. The concrete constraints are obtained for two special cases (the power spectrum with a peak and the spectrum of the running mass model), although the general formulas of Sec. \\ref{sec-form} and \\ref{sec-nu-gamma} can be used for any form of the spectrum.  At the end of the paper we use our approach for a checking, once more, the idea proposed in \\cite{Frampton:2005fk, Bambi:2008kx}, namely, we study the possibility that evaporating PBHs (having mass spectrum calculated with formulas of Sec. \\ref{sec-form}) cluster in the Galactic center and produce the observed $511$ keV photon line.  The plan of the paper is as follows. In Sec. \\ref{sec-form} we present the general formalism of the PBH mass spectrum calculation with taking into account the explicit time dependence of the gravitational potential. The results of such a calculation for the concrete example, in which the curvature perturbation spectrum has a sharp peak, are given. In Sec. \\ref{sec-nu-gamma} we give some details of the calculation of the  neutrino and the photon extragalactic diffuse spectra from the PBH evaporations. In Sec. \\ref{sec-constr-peak} and Sec. \\ref{RM-sec} the constraints on the curvature spectrum from PBHs are given, for the model of ${\\cal P}_{\\cal R}(k)$ with a peak and for the running mass model, respectively. In Sec. \\ref{sec-DM} we consider the possibility of the explaining the observed $511$ keV photon line from the Galactic center by PBHs clusterizing there.  Section \\ref{sec-concl} contains our conclusions. ",
        "conclusions": "\\label{sec-concl}  {\\bf 1.} We have developed a method of PBH mass spectrum calculation, which is based on the Press-Schechter formalism and takes into account the dependence of the gravitational potential on time. We found that it is more convenient and natural to use in the formula for the mass variance (Eq. (\\ref{SigR})) the gravitational potential $\\Psi_k(t)$ calculated as a function of time rather than the transfer function $T(k,t)$ (the latter is defined, e.g., in \\cite{Blais:2002gw}). To derive the PBH mass spectrum, the power spectrum of curvature perturbations calculated (using inflationary models) at the end of inflation is needed as an input.  {\\bf 2.} PBH mass spectra are calculated in two inflationary models predicting large amplitudes of the curvature perturbation spectrum at some wave number $k$, namely, in the model with the power spectrum with a peak and in the running mass model. Examples of such a calculation are given in Figs. \\ref{nBHfig}, \\ref{nBH-sig-fig} and \\ref{nBH-rm-fig}.  {\\bf 3.} Constraints on parameters of both these models are obtained using available data on the extragalactic photon diffuse background and the upper limit on the extragalactic neutrino background, obtained in the Super-Kamiokande neutrino experiment.  The results are shown in Figs. \\ref{const} and \\ref{constr-rm-fig}. It is seen that for the power spectrum with a peak, the height and the width of the peak can be constrained in rather wide region of wave numbers determining its position in the spectrum. For the running mass model we obtain the constraint on the spectral index running at cosmological scales, $n_0' \\lesssim (3 \\div 6) \\times 10^{-3}$ (the exact constraint value depends on $n_0$ and $T_{RH}$).  {\\bf 4.} It follows from Fig. \\ref{const} that constraints obtained from neutrino experiments are stronger than those obtained from extragalactic photon background, in the region of horizon masses from $\\sim 10^{11}$ to $\\sim 10^{13}$ g. So, constraints from neutrino experiments and from measurements of diffuse photon background are complementary.  {\\bf 5.} PBHs can be the main contributor to the dark matter in the Universe. It is shown that if, in this case, the PBH mass distribution is peaked near $\\sim 8 \\times 10^{16}$ g, as in Fig. \\ref{nBH-sig-fig} with $\\Sigma=3$, the experimentally observed galactic $511$ keV photon line can, in fact, be caused by annihilating positrons produced by evaporating PBHs clustered in the Galactic center. The total PBH mass needed for this to happen must be about $10^9 M_\\odot$, and clusterization factor needed is $\\zeta \\sim 3\\times 10^7$.  \\bigskip"
    },
    "0812/0812.0018_arXiv.txt": {
        "abstract": "A popular paradigm to explain the rapid temporal variability observed in gamma-ray burst (GRB) lightcurves is the internal shock model.  We propose an alternative model in which the radiating fluid in the GRB shell is relativistically turbulent with a typical eddy Lorentz factor $\\gamma_t$.  In this model, all pulses in the gamma-ray lightcurve are produced at roughly the same distance $R$ from the center of the explosion.  The burst duration is $\\sim R/c\\Gamma^2$, where $\\Gamma$ is the bulk Lorentz factor of the expanding shell, and the duration of individual pulses in the lightcurve is $\\sim R/c\\Gamma^2\\gamma_t^2$. The model naturally produces highly variable lightcurves with $\\sim\\gamma_t^2$ individual pulses.  Even though the model assumes highly inhomogeneous conditions, nevertheless the efficiency for converting jet energy to radiation is high. ",
        "introduction": "Our understanding of gamma-ray bursts (GRBs) has improved enormously in the last 15 years, thanks to observations by dedicated $\\gamma$-ray/X-ray satellites such as Compton-GRO, BeppoSAX, HETE-2 and Swift, and follow-up observations from the ground in optical and radio (for recent reviews see M\\'esz\\'aros, 2002; Piran 2005; Woosley \\& Bloom, 2006; Fox \\& M\\'esz\\'aros, 2006; Zhang, 2007).  As a result of this work, it is now established that at least some long-duration bursts are produced in the collapse of a massive star (as suggested by Woosley 1993; Paczynski 1998), accompanied by the ejection of a highly relativistic jet.  Afterglow observations as well as energy considerations indicate that the jet is well collimated (Rhoads 1999, Frail et al. 2001, Panaitescu \\& Kumar, 2001).  The presence of a relativistic jet has also been directly confirmed in the nearby burst GRB 030329 which exhibited ``superluminal'' motion in its radio afterglow (Taylor et al. 2004).  GRB lightcurves are known to be highly variable (Meegan et al. 1992), and this has led to the development of the {\\it internal shock model} (Piran, Shemi \\& Narayan 1993; Rees \\& Meszaros 1994; Katz 1994). According to this model, the Lorentz factor of the jet varies with time, and as a result, faster portions of the jet catch up with slower portions.  In the ensuing collision, a fraction of the kinetic energy of the jet is converted to thermal energy.  The observed radiation is then produced via synchrotron and inverse-Compton processes.  In a seminal paper, Sari and Piran (1997) provided a general argument why GRB lightcurve variability cannot be explained by simply appealing to a highly inhomogeneous source.  They showed that the efficiency for converting jet energy to the observed radiation is extremely poor when variability arises purely from inhomogeneity; the essence of their argument is summarized in \\S2.  The argument is both powerful and compelling, and it has led to wide acceptance of the internal shock model.  A basic feature of the model is that different pulses within the lightcurve of a GRB are produced in distinct internal collisions/shocks, generally at different distances from the central explosion.  There are, however, a number of observations and/or theoretical considerations that pose difficulties for the internal shock model. Internal shocks have only a modest efficiency $\\sim 1$--10\\% for converting jet energy to the radiation observed in the $20$\\,keV--1\\,MeV band (Kumar, 1999; Panaitescu, Spada \\& Meszaros, 1999; Lazzati et al. 1999)\\footnote{Beloborodov (2000) and Kobayashi \\& Sari (2001) reported a much higher efficiency $\\sim100$\\%. However, their estimates were based purely on the kinematics of colliding shells, where shells of Lorentz factor $\\gamma\\sim1$ collided with high $\\gamma$ shells.  They did not take into consideration the efficiency of the emergent radiation in the commonly observed 20\\,keV--1\\,MeV band.}.  Even a $\\sim$10\\% radiative efficiency is low compared to the burst efficiency implied by measurements of the jet kinetic energy through modeling GRB afterglow lightcurves (Panaitescu \\& Kumar, 2002).  Another difficulty with the internal shock model is the large distance from the central explosion that one estimates ($R\\ \\gta\\ 10^{16}$cm) for the $\\gamma$-ray-producing region in a number of GRBs (Kumar et al. 2007).  This distance is significantly larger than what one expects in the internal shock model.  Moreover, the estimated distance is within a factor of a few of the deceleration radius where the jet begins to interact with the external medium.  This coincidence between two unrelated radii is unexpected.  These difficulties, along with the problem of avoiding excessive baryon loading, motivate us to consider an alternative to the internal shock model.  We show in \\S3 that the argument of Sari \\& Piran (1997) against an inhomogeneous source can be successfully overcome if we consider source inhomogeneities that move with {\\it relativistic velocities}.  In \\S4, we calculate a model lightcurve using this relativistic model and demonstrate that it is consistent with observations of GRBs.  We conclude in \\S5 with a discussion. ",
        "conclusions": "The turbulent model described in this paper has several attractive features.  \\smallskip\\noindent (i) The model naturally produces a highly variable GRB lightcurve with large amplitude variations across individual pulses and a duty cycle of order unity (\\S4, Fig. 1).  \\smallskip\\noindent (ii) According to equation (\\ref{tburst}), the quantity $R/\\Gamma^2c$ is equal to $t_{\\rm burst}$ rather than $t_{\\rm var}$ (eq. \\ref{eqstd}).  Relative to the standard model, the turbulent model can thus accommodate much larger values of $R$ and smaller values of $\\Gamma$.  This eliminates a problematic constraint which leads to difficulties when attempting to fit GRB observations using equation (\\ref{eqstd}) from the standard model (e.g., Kumar \\& Narayan 2008).  \\smallskip\\noindent (iii) The larger value of $R$ obtained with the turbulent model is compatible with ideas described in Lyutikov \\& Blandford (2003), according to which the jet energy is primarily in the form of Poynting flux.  This energy is converted to radiation through plasma instabilities near the deceleration radius ($\\lta\\ 10^{17}$ cm).  The same instabilities may also produce the relativistic turbulence invoked in our model.  \\smallskip\\noindent (iv) The turbulent model neatly avoids the efficiency argument of Sari \\& Piran (1997).  Any particular observer receives radiation from only a fraction $\\sim1/\\gamma_t$ of the available eddies.  However, the additional relativistic boost of the received radiation because of turbulent motion makes up for the missing eddies.  Thus, each observer receives a fair share of the emission from the shell, and there is no radiative inefficiency in the model.  \\smallskip\\noindent (v) The model has a clear prediction for the variability parameter ${\\cal V}$: \\begin{equation} {\\cal V} \\sim \\gamma_t^2, \\qquad \\gamma_t \\sim {\\cal V}^{1/2}. \\end{equation} Since a typical long GRB has ${\\cal V}\\sim100$, the turbulent eddy Lorentz factor $\\gamma_t$ needs to be $\\sim 10$.  For simplicity, we have assumed in this paper that the eddy motions are isotropic in the frame of the GRB shell.  This is, however, not essential.  We could, for instance, have random relativistic motions which are concentrated primarily in a plane perpendicular to the radius vector of the shell, e.g., parallel to the local tangential magnetic field.  Such a model would give qualitatively similar results, though some of the scalings may be a little different.  Also, we have simplified matters by assuming that all eddies have the same Lorentz factor, which is unlikely in a real turbulent medium.  In fact, a likely scenario is that a part of the fluid moves relativistically with a range of Lorentz factors, and a part resides in a (more-or-less) stationary inter-eddy medium which is produced when eddies collide with one another in shocks.  This would slightly modify the radiative properties of the medium (Kumar \\& Narayan 2008), but it will not change the key features of the model as described in the present paper.  Finally, the relevant energy-bearing eddies may not be as large as their comoving causality size but may be smaller by a numerical factor, and the solid angle swept by an eddy in the shell frame may be different from $\\sim1/\\gamma_t$ (for instance, eddies might be hardly accelerated at all once they are formed).  These effects will modify equations (\\ref{tburst}), (\\ref{tvar}) and (\\ref{npulse}).  However, the results will remain qualitatively the same."
    },
    "0812/0812.3641_arXiv.txt": {
        "abstract": "{This paper presents estimations on the possibility of detection of ellipsoidal variations by means of measuring brightness of the star distorted by a close massive planet using Wilson-Devinney method. The problem was already discussed by Phafl et al. (2008) and earlier by Loeb \\& Gaudi (2003). The effect is well known in the case of binary stars where it can produce light curves with amplitutudes of ellipsoidal variations of about 0.1 mag for distorted stars. For planets the effect is very small and usually less than 0.0001 mag. The detection of an exoplanet, by searching for small amplitude ellipsoidal variations, will be very difficult and affected by other photometric effects; however, it maybe possible for some extreme cases. Observations of ellipsoidal variations can provide additional constraints on the model of the system. Light curves for few star/planet systems have been calculated using PHOEBE eclipsing binary software based on Wilson-Devinney method. As an example of ellipsoidal variations the synthetic light curve for $\\tau$Bootis is presented. The amplitude of ellipsoidal variation is 0.01 mmag. The companion is massive (7.3 M$_{\\mathrm{Jup}}$) and short-period hot Jupiter.} ",
        "introduction": "\\label{s1}  During the last decade we have observed a very fast development of extrasolar planet research. The new methods of detection have been successfully used. The observational techniques give us the possibility to detect more and more shallow effects like transits of a planet in front of the hosting star, a direct imaging of planets and measuring the chemical abundance in the planet atmosphere. In this paper I will discuss the question whether we can observe light curve variations of the hosting star distorted by the gravity field of a massive planet for some of the already known systems. The detailed analytical formulation of ellipsoidicity effect caused by substellar companion is given by Pfahl et al. (2008). A close massive companion of the star can produce tides analogue to the ones present in close binary stars. The effect is of course much weaker and difficult for detection. Actually, the presence of ellipsoidal variations among the transit candidates is treated as an evidence of a non planetary behaviour of the eclipses. Patzold et. al (2004) describe orbital period changes in planetary systems caused by gravity tides induced by exoplanets. Some short-period exoplanets, for example $\\tau$Bootis, are probably in tidal lock that demonstrates the significance of the tidal interaction. We can model ellipsoidal variations using a code based on Wilson-Devinney (WD) method developed for eclipsing binary stars. This method has already been used for modeling eclipses and radial velocity curves of transiting planets - one of them being OGLE-TR-56 (Voccaro and Van Hamme 2005). ",
        "conclusions": "\\label{s4} The future mass photometry missions can be used for detection of exoplanet candidates by searching for effects weaker than transits like induced spots, ellipsoidal variations or scattered light. The WD code is already used for modeling transits of exoplanets. This method can be usful for calculating weaker effects like ellipsoidal variations or spots on the host star induced by the planetary companion, and maybe applied for the scattered light after some modification of the code. The method can also be useful for modeling the Rossiter-McLaughlin effect in radial velocity curves."
    },
    "0812/0812.3688_arXiv.txt": {
        "abstract": "The chemistry in the central regions of galaxies is expected to be strongly influenced by their nuclear activity. To find the best tracers of nuclear activity is of key importance to understand the processes taking place in the most obscured regions of galactic nuclei. In this work we present multi-line observations of CS, C$^{34}$S, HNCO and C$^{18}$O in a sample of 11 bright galaxies prototypical for different types of activity. The $^{32}$S/$^{34}$S isotopic ratio is $\\sim10$, supporting the idea of an isotopical $^{34}$S enrichment due to massive star formation in the nuclear regions of galaxies. Although C$^{32}$S and C$^{34}$S do not seem to be significantly affected by the activity type, the HNCO abundance appears highly contrasted among starburst. We observed HNCO abundance variations of nearly two orders of magnitude. The HNCO molecule is shown to be a good tracer of the amount of molecular material fueling the starburst and therefore can be used as a diagnostics of the evolutionary state of a nuclear starburst. ",
        "introduction": "The nuclear regions of active galaxies can harbor different energetic phenomena such as starbursts (SBs) and Active Galactic Nuclei (AGN) which are responsible for producing the bright emission stemming from their central regions. The heating of the interstellar medium (ISM) in nuclei with different type of activity is expected to be dominated by a variety of mechanisms, such as UV, X-rays, and/or shocks, which will ultimately drive their observed chemical richness \\citep[][and references therein]{Martin06b}. Disentangling which of this mechanisms is the main power source of galactic nuclei is complicated by the high obscuration affecting these objects. This fact becomes critical in the case of Ultra Luminous Galaxies (ULIRGs) and high-z sources. Therefore the observation of the dense molecular material has become an essential tool to get an insight in the evolution and classification of the nuclear activity in galaxies.  Finding appropriate molecular species to accurately trace each of these heating mechanisms is crucial to stablish the nature of the central engine. Molecular species like HCN and HCO$^+$, and in particular the ratios HCN/CO and HCN/HCO$^+$, have been claimed to be appropriate to differentiate between the SB and the AGN contribution in galactic nuclei \\citep{Kohno99,Kohno05,Krips08}. Similarly, species such as HNC and CN are thought to show enhanced abundances in extremely irradiated environments \\citep{Aalto02,Aalto07}. A dichotomy still remains concerning the different evolutionary state of nuclear starburst. While at an early stage, the heating of the ISM is thought to be dominated by shocks affecting the molecular clouds fueling the starburst, the late stages of starburst are vastly dominated by the UV radiation originated in the newly formed massive stars. This scenario has been inferred from the extensive observation of the prototype sources NGC\\,253 and M\\,82 in a number of key molecular tracers such as SO$_2$, H$_2$S, OCS \\citep{Martin03,Martin05}, HOC$^+$, and CO$^+$ \\citep{Fuente05,Fuente06}. Unfortunately, these tracers are generally too faint to be detected, but for the brightest prototypical nearby galaxies.  In a recent observational study carried out towards a sample of molecular clouds dominated by different heating mechanisms within the Galactic center region \\citep{Martin08}, we have found that the CS/HNCO abundance ratio is highly contrasted between molecular clouds illuminated by the UV radiation from massive star clusters and the giant molecular cloud complexes shielded from photodissociation and mostly heated by shocks. The origin of the large changes in the CS/HNCO ratio is due to the enhancement of CS in photon-dominated regions (PDRs) through reactions involving S$^+$ \\citep{Stern95} as opposed to the highly photodissociable HNCO, which is destroyed by the UV photons. Both observations of photodissociation regions within our Galaxy and photochemical models show the suitability of CS as a PDR tracer \\citep{Goico06}. In addition, HNCO is enhanced in the presence of shocks due to its injection in the gas phase from the grain mantles \\citep{Zinchen00}. This paper presents a follow up study of the CS/HNCO ratio towards a sample of prototype galaxies with different nuclear activity. Although CS has been extensively observed in external galaxies \\citep{Baan08}, HNCO had only been detected in four galaxies prior to this work \\citep{Nguyen91,Wang04}. ",
        "conclusions": "From our study of the CS and HNCO abundances in a selected sample of extragalactic sources, we have found differences of nearly two orders of magnitude in the HNCO/CS abundance ratios among the starburst galaxies. These differences are in agreement with the results from the observations of Galactic center clouds dominated by various heating mechanisms suggesting that the observed HNCO abundances could be related to the evolutionary stage of their nuclear starburst. Shocks affecting the massive dense molecular clouds in the early stages and strong UV fields in the evolved stages dominate the HNCO abundance in galaxies. We do not find significant differences in the molecular abundances of HNCO in galaxies with different type of nuclear activity (i.e. AGN or SB) which implies this species are not affected differently by these processes at the scales observed by single dish beam telescopes.  As derived from the HNCO abundances we can point out prototypes of these nuclear starburst evolutionary stages. M\\,83 and the nuclear spiral arms in IC\\,342 would be representative of the early stage, while M\\,82 and NGC\\,5194 nuclei would be in the later stages of evolution. NGC\\,253, claimed to be a prototype of early starburst, shows abundances below the average suggesting an important influence of photodissociation in the heating of its ISM, and therefore chemically closer to M\\,82 than to M\\,83.  The $^{32}$S/$^{34}$S isotopic ratio, derived from the C$^{32}$S/C$^{34}$S abundance ratio, is found to be $\\sim10$, and thus similar for all the observed sources. This ratio, found to be well below the value of $\\sim24$ in the Galactic disk, supports the idea of the $^{34}$S isotope being enriched due to massive star formation events in the nuclear region of galaxies.  The CS emission is suggested to be a good tracer of the molecular gas affected by photodissociation in the Galactic center \\citep{Martin08}. We do not observe a significant contrast in our sample. Though its abundance is enhanced in UV radiated regions, its emission is also favored in high density molecular clouds. Thus the overall emission over the central few hundred parsecs of galactic nuclei is averaged out, and no significant trends in the CS abundances is found in the observed galaxies."
    },
    "0812/0812.1472_arXiv.txt": {
        "abstract": " ",
        "introduction": "Light polarisation from astronomical objects (apart from the Sun) has become more and more popular since the first polarimetric observations of comets (Wright\\cite{wright81}), of a ``Nebulosity about Nova Persei'' (Perrine\\cite{perrine02}), of the moon (Barabascheff \\cite{barabascheff27}) or the discoveries of the polarised poles of Jupiter (Lyot\\cite{lyot29}). Nowadays, there exist many observatories providing the  possibility for polarimetry, e.g. at the VLT: NACO, FORS1, ISAAC; at the NTT: SOFI and EFOSC2; at Keck 1: LRIS etc. This short list aims not to be complete but rather to show that polarimetry is offered at the world's largest observatories.  However, polarimetric measurements are often affected by telescope and/or instrumental polarisation and a main effort during data reduction is to get rid of these ``disturbing'' effects. Especially, observations with Nasmyth instruments are particularly affected by strong and variable telescope polarisation. The Nasmyth mirror M3 deflects the light by $90\\degr$ to the Nasmyth focus and introduces polarisation. Metallic mirrors introduce polarisation dependent on the surface coatings, the incidence angle ($45\\degr$ for M3) and the wavelength. For example, an aluminium coated flat mirror with an incidence angle of $45\\degr$ introduces up to $5\\,\\%$ of polarisation at 800\\,nm and about $1\\,\\%$ at 1600\\,nm. Although the induced polarisation is intrinsically not dependent on the pointing direction of the telescope (always 45$\\degr$ for M3) the measured polarisation directions in a Nasmyth instrument depend strongly on the elevation of the target because the induced polarisation due to M3 is decoupled from the field. Thus, the polarisation breakdown into Stokes components is variable (dependent on pointing direction).   Due to the mentioned reasons but also with the future more complex polarimetric instruments (e.g. polarimetric mode of VLT/SPHERE, see Schmid et al. \\cite{schmid06a}) one aims to get the polarimetric data reduction integrated in a pipeline. The pipeline data reduction should be able to correctly handle the telescope/instrument polarisation without the need of time consuming calibration measurements during the night. This can be achieved with a polarimetric instrument model based on Mueller matrices describing all polarimetric effects from the sky to the polarimetric analysing system. This model needs only few calibration measurements during the night to adjust/monitor the model parameters.  The next section gives a short description on how to get polarimetric data and on an optimised way of performing polarimetric observations. A good observing strategy is always the first step in achieving good polarimetric data. With the optimal observing technique several instrumental effects can be corrected without applying specific  calibration measurements (self-calibration). The following section describes the standard way of calibrating polarisation measurements. In section \\ref{muellersection} the proposed calibration based on a polarisation model described by Mueller matrices is explained in detail and verified with an example measurement.  The whole paper treats the measurement of linear polarisation only. Circular polarisation is not subject of this publication. ESO instruments/telescopes are used here as examples since the authors have performed several own polarimetric observations at these instruments (e.g. Joos \\& Schmid \\cite{joos07}, Schmid et al. \\cite{schmid06}) Nevertheless, the examples are also valid for similar telescope/instrument configurations. ",
        "conclusions": "Polarimetry is a powerful technique taking into account also the orientation of the electric field and not only the intensity of light. Nevertheless, polarimetry is rather a niche science. A reason for this is among others the fact that polarimetric measurements are often ``contaminated'' by telescope/instrumental effects which are demanding to calibrate.  To get good polarimetric observations we strongly recommend the described technique with a combination of temporal and spatial differential measurements (see section \\ref{optimum}). To benefit at most from the advantages of this technique a rotatable half-wave plate in front of the polarisation analysis is very important (not offered in NTT/SOFI). If it is not provided, the whole instrument must be rotated to measure at different polarisation angles. This has strong impacts on the polarimetric precision, especially for extended objects, since different detector zones are illuminated differently due to rotating. Many rotations of the whole instrument by large angles in short time may also have an impact onto sensitive parts of the instrument itself, like e.g. the cryogenic system. In addition dithering becomes difficult since the orientation relatively to the occulting masks has to be taken into account. Nevertheless, in several instrument hand books (like for the NTT/SOFI) a different and less satisfactory observing technique is proposed which leads to less precise polarimetric measurements.  With our method of calibrating polarimetric measurements by means of Mueller calculus we have demonstrated that polarimetric calibration and data reduction can be implemented into a pipeline and hence, make polarimetric observations more attractive. We have implemented a very simple model for describing the polarimetric properties at NTT/SOFI. This model could also be upgraded by taking into account all optical devices instead only the folding mirror M3. On the other hand we reached with the simple model already an accuracy which is well inside the error of the measurements.  We would therefore strongly encourage to implement the polarimetric data reduction and calibration in a pipeline and benefit not only from the pipeline service itself but also reducing much overhead time up to now ``wasted'' for calibration measurements."
    },
    "0812/0812.4582_arXiv.txt": {
        "abstract": "The blazar 3C~454.3 was revealed by the Fermi Gamma-ray Space Telescope to be in an exceptionally high flux state in July 2008. Accordingly, we performed a multi-wavelength monitoring campaign on this blazar using IR and optical observations from the SMARTS telescopes, optical, UV and X-ray data from the Swift satellite, and public-release gamma-ray data from Fermi. We find an excellent correlation between the IR, optical, UV and gamma-ray light curves, with a time lag of less than one day. The amplitude of the infrared variability is comparable to that in gamma-rays, and larger than at optical or UV wavelengths. The X-ray flux is not strongly correlated with either the gamma-rays or longer wavelength data.  These variability characteristics find a natural explanation in the external Compton model, in which electrons with Lorentz factor  $\\gamma \\sim 10^{3-4}$ radiate synchrotron emission in the infrared-optical and also scatter accretion disk or emission line photons  to gamma-ray energies, while much cooler electrons ($\\gamma \\sim 10^{1-2}$) produce X-rays by scattering synchrotron or other ambient photons. ",
        "introduction": "\\label{sec:intro} Blazars are understood to be active galactic nuclei (AGN) with aligned relativistic jets \\citep{urry95}, so they offer a unique laboratory for studying the physics of astrophysical jets. The spectral energy distributions (SEDs) of blazars have a characteristic double-humped shape with a low-energy component peaking anywhere from radio to X-rays, and a high-energy component peaking at MeV to GeV energies \\citep{fossati98}. Flat Spectrum Radio Quasars (FSRQs) like 3C~454.3 have SED peaks at radio-IR wavelengths and $\\sim 1$~GeV \\citep{urry95, sambruna96}. The low-energy component is well modeled as synchrotron emission from relativistic electrons in the jet \\citep{konigl81, urry82}, while the origin of the second SED peak at high energies is not fully understood. Current explanations for the gamma-ray emission fall into two categories, leptonic and hadronic. Leptonic models produce high-energy flux by inverse-Compton scattering of low-energy seed photons, either the synchrotron photons themselves \\citep[Synchrotron Self-Compton, ][]{jones74} or photons from an external source, such as thermal accretion disk emission or broad-line emission \\citep{sikora94, dermer93, ghisellini96, celotti08}. In hadronic models, protons that are accelerated to very high energies in the jet produce gamma-rays from neutral pion decay, proton synchrotron emission, and synchrotron emission from pair production \\citep{mucke01, mucke03, boettcher07}. Both leptonic and hadronic models can adequately fit single-epoch blazar SEDs, but variability offers a test of either model.  \\thesource\\ was among the more intense and variable FSRQs detected with CGRO EGRET \\citep{hartman99}, varying over several years by factors of up to five, with a flare-state flux of $F_{\\rm >100~MeV} \\sim$0.5$\\times$10$^{-6}$~photons/s/cm$^2$ \\citep{hartman93,hartman99,aller97}. Long-term optical variability has also been reported, with up to $\\sim$3 mag changes over several years \\citet{djorgovski08}. During a 2005 optical flare to R=12 \\citep{villata06}, \\thesource\\ was detected with INTEGRAL at a flux of $F_{\\rm 3-200~keV} \\sim$3$\\times$10$^{-2}$~photons/s/cm$^2$ \\citep{pian06}; a radio flare followed about a year later \\citep{villata07}. \\thesource\\ has been detected with the AGILE gamma-ray satellite \\citep{tavani08}, flaring in July 2007 and again in July 2008 \\citep{vercellone08,gasparrini08} with associated flaring at optical and longer wavelengths \\citep{ghisellini07,villata08}. On 24 July 2008, \\citet{tosti08} confirmed the high gamma-ray flux state of the source with a detection by the Fermi Large Area Telescope (LAT) while still in its post-launch commissioning phase. In the Fermi/LAT first light image released on 26 August 2008, \\thesource\\ was among the brightest sources in the gamma-ray sky, at the high end of its recorded gamma-ray intensity, $F_{\\rm 0.1-300~GeV} \\sim$4.4$\\times$10$^{-6}$ photons/s/cm$^2$.  Here we present data from our multi-wavelength optical and infrared monitoring program of \\thesource\\ from June to December 2008 with the Small and Moderate Aperture Research Telescope System (SMARTS). We correlate these data with Target of Opportunity observations carried out with the Swift X-ray Telescope (XRT) and Ultraviolet and Optical Telescope (UVOT), as well as with 0.1--300 GeV fluxes made public by the Fermi Science Support Center. The observations are described in Section~\\ref{sec:obs}. The light curves, correlation functions, and SED are discussed in Section~\\ref{sec:results}. ",
        "conclusions": "\\label{sec:results}   The correlated variability across all observed wavebands except the X-ray is readily apparent in Figure~\\ref{fig:lc}. All observed bands save the X-ray show two prominent peaks around JD~2454715 and a short flare near JD~2454740. The amplitude is largest in the gamma-rays and J band. Figure~\\ref{fig:dcf} shows the discrete correlation function \\citep[DCF,][]{edelson88,white94} calculated for the gamma-ray (0.1--300 GeV) flux versus light curves in the optical B band\\footnote{For the B band, we include optical fluxes from both the Swift and SMARTS telescopes in order to have complete coverage over gaps in the individual light curves.}, which has the best temporal coverage, and infrared J-band, which shows the strongest variations. The DCF shows a peak correlation amplitude $\\sim$0.7 at $\\tau = 0$, indicating no detectable lag between IR/optical and gamma-ray fluxes. Given the sampling, this means any lag is less than or about 1~day. Similar results were reported by \\citet{vercellone09} for the earlier flare observed with AGILE, though with much lower significance. The optical versus IR DCF shows even stronger correlation (amplitude $\\sim$0.8), also with $0 \\pm 1$~day lag.    Table~1 shows the fractional root mean square (rms) variability amplitude \\citep{vaughan03} for each band. The IR, optical, and UV variability amplitudes decrease toward shorter wavelengths, suggesting the possible presence of steady thermal emission (UV accretion disk emission plus Balmer continuum, Fe~{\\sc ii}, and \\mgii\\ in the V and B bands) added to the steeper-spectrum jet. Evidence for `big' and `little' blue bumps was found previously in the SED of \\thesource\\ during periods of low emission \\citep{raiteri07}. The colors of \\thesource\\ are redder at brighter levels, historically \\citep{villata06} and in the present data, also supporting the presence of thermal emission beneath the much brighter non-thermal jet.  The closely correlated IR/gamma-ray variability of \\thesource\\ supports a model in which relativistic electrons in the jet radiate IR/optical synchrotron photons and inverse Compton scatter thermal photons to X- and gamma-ray energies. The observed gamma-ray flares must be caused by changes in the injection luminosity of the higher energy electrons, rather than variability of the ambient thermal photons, since in that case there would be higher amplitude variations in the UV than the infrared. The implication of the short lag time (Fig.~\\ref{fig:dcf}) is that electrons of similar energy produce IR and gamma-ray emission.   Figure~\\ref{fig:sed} shows the SED in the high state (JD~2454680--2454750) and at the lower final intensity (JD~2454750--2454820). The SED of the high flux state prior to JD~2454750 shows an optical/IR flux level similar to that of the May 2007 flare \\citep{raiteri08}, intermediate between the high and low states reported by \\citet{raiteri07} (and references therein), so not surprisingly, the basic model parameters are similar. The optical/UV emission is due to the highest energy electrons (Lorentz factors $\\sim10^{3-4}$) radiating via synchrotron in a field of $\\sim10$~Gauss, while the gamma-rays come from inverse Compton scattering on the broad-line photons. The bulk Lorentz factor is $\\Gamma \\sim \\delta$ $\\sim 10-15$  (where $\\delta = [\\gamma(1-\\beta\\cos\\theta)]^{-1}$  is the Doppler beaming factor).   The lack of correlation seen in the DCF for 2--10~keV X-rays with respect to the other wavebands finds a natural explanation in the external Compton scenario, with the X-rays coming from low-energy electrons ($\\gamma\\sim$10--100) inverse-Compton scattering external UV photons, rather than higher energy electrons ($\\sim 10^{3-4}$) scattering synchrotron photons. An SSC component in X-rays would introduce correlation between X-rays and gamma-rays, which is not seen. The highest energy electrons (producing the IR/optical and gamma-ray emission) vary more rapidly (the radiative timescales are shorter) while the low energy electrons act as a reservoir and vary more slowly.  More precise SED modeling is needed to determine detailed model parameters, such as the energy density and location of the thermal photons, the location and size of the scattering region, the electron distribution, the bulk Lorentz factor and jet orientation, etc. This detailed analysis will be deferred to a later paper. Still, some additional conclusions can be made. The overall stability of source parameters and the correlation imply that the emission region is stable on time scales of $\\sim 1$~month. If the electrons are localized in a fast moving knot (which might become visible in VLBI maps in a few months), it moves a distance $\\gamma^2 c \\Delta t$, roughly 1-10~pc, i.e., the jet parameters cannot change dramatically on this scale. However, the \\citet{sikora94} model for \\thesource\\ can be ruled out as the source of the rapid variations discussed here, since their assumed source size of $10^19$~cm implies $\\Delta t \\gtrsim 1$~year. Instead, their model might explain a slowly changing, much larger region of the jet.   In conclusion, \\thesource\\ shows very strong, correlated variability between the peak of the synchrotron component (at infrared, optical and UV wavelengths) and the peak of the gamma-ray component. No such correlation is seen between X-rays and any other band. These results suggest that the variability arises from changes in the electron luminosity at a compact location in the jet. The highly variable infrared through UV emission, particularly in the brightest state, is dominated by synchrotron emission from a compact region of high-energy electrons in the jet, with a smaller contribution from a relatively steady accretion disk. The slowly varying low-energy part of the electron spectrum gives rise to relatively stable X-ray emission via scattering. The gamma-rays vary in a correlated way because they result from  the same high-energy electrons up-scattering ambient UV photons."
    },
    "0812/0812.2856_arXiv.txt": {
        "abstract": "We present the results from our search for \\civ\\ in the intergalactic medium at redshifts $z = 5.3 - 6.0$.  We have observed four $z \\sim 6$ QSOs with Keck/NIRSPEC in echelle mode.  The data are the most sensitive yet taken to search for \\civ\\ at these redshifts, being 50\\% complete at column densities $\\log{N_{\\mciv}} \\approx 13.4$.  We find no \\civ\\ systems in any of the four sightlines.  Taking into account our completeness, this translates into a decline in the number density of \\civ\\ absorbers in the range $13.0 < \\log{N_{\\mciv}} < 15.0$ of at least a factor 4.4 (95\\% confidence) from $z \\sim 2-4.5$, where the number density is relatively constant.  We use our lack of detections to set limits on the slope and normalization of the column density distribution at $z = 5.3-6.0$.  The rapid evolution of \\civ\\ at these redshifts suggests that the decrease in the number density may largely be due to ionization effects, in which case many of the metals in the $z \\sim 4.5$ IGM could already be in place at $z \\sim 5.3$, but in a lower ionization state.  The lack of weak systems in our data, combined with the presence of strong \\civ\\ absorbers along at least one other sightline, further suggests that there may be large-scale variations in the enrichment and/or ionization state of the $z \\sim 6$ IGM, or that \\civ\\ absorbers at these redshifts are associated with rare, UV-bright star-forming galaxies. ",
        "introduction": "\\label{sec:civ_intro}  Metal absorption lines are one of our most versatile tools for studying the high-redshift Universe.  Just as the ionization history of the IGM reflects the output of ionizing photons from star-forming galaxies and quasars, the metal content of the IGM can be used to infer the integrated and instantaneous rates of global star formation. At very high redshifts, where the direct detection of galaxies becomes increasingly challenging, metal absorption lines can serve as sign posts to the ``typical'' galaxies that may be responsible for reionizing the Universe at $z > 6$.  Perhaps most intriguingly, metal lines can also serve as a probe of the reionization process itself.  A considerable amount of work on the high-redshift Universe has focused on cosmic reionization, yet the timing and mechanisms by which the IGM becomes permanently ionized remain poorly understood.  The three main observational constraints on reionization are (1) the electron optical depth measured from the polarization of the cosmic microwave background (CMB), (2) the evolution of the high-redshift \\lya\\ forest, and (3) the luminosity functions of high-redshift \\lya-emitting galaxies (LAEs).  Five-year {\\it WMAP} measurements of the CMB are consistent with an instantaneous reionization at $z_r \\sim 11$, but are also compatible with an extended reionization process \\citep{dunkley08}.  The rapid evolution of the mean transmitted \\lya\\ flux in the spectra of $z \\gtrsim 5.7$ quasars has been used to infer that reionization may have ended as recently as $z \\sim 6$ \\citep{becker01,white03,fan02,fan06,gnedin06}.  The quasar results remain controversial, however \\citep[e.g.,][]{songaila04,ohfur05,lidz06,becker07}.  \\lya\\ saturates for even a tiny neutral fraction \\citep[$f^{V}_{\\mhi} \\sim 10^{-3}$;][]{fan02}, making it extremely challenging to measure the ionization state of the IGM at these redshifts directly from the \\lya\\ forest.  Finally, the lack of a sharp decline in the luminosity function of LAEs from $z \\approx 5.7$ to $z \\approx 6.5$ suggests that most of the volume of the IGM is already highly ionized by $z \\sim 6.5$.  \\citep[e.g.,][but see Kashikawa et al.~2006]{mr04,mr06,hu06,dijkstra07}.  Interpreting the LAE results is somewhat complicated, however, as \\lya\\ photons may escape through locally ionized bubbles \\citep[e.g.,][]{haiman05,wyithe05,furlanetto06} or by redshifting via galactic winds before passing through the IGM \\citep{santos04}.  Quasar metal absorption lines provide a complimentary tool for studying the ionization state of the IGM.  In a reionization scenario where galaxies provide most of the ionizing flux, overdense regions of the IGM should be the first to become enriched, yet may remain significantly neutral until the end of reionization due to the short recombination times at high densities.  ``Forests'' of low-ionization absorption lines such as \\oi~$\\lambda 1302$, \\siii~$\\lambda 1260$, and \\cii~$\\lambda 1334$ may therefore be present prior to the completion of reionization \\citep{oh02}.  Our recent search for \\oi\\ using HIRES spectra of $z > 5$ QSOs yielded mixed results \\citep{becker06}.  While an overdensity of low-ionization absorbers at $z \\sim 6$ was found along at least one sightline, SDSS~J1148$+$5251 ($z_{\\rm Q} = 6.42$), no enhancement was seen towards SDSS~J1030$+$0524 ($z_{\\rm Q} = 6.30$), which has a complete Gunn-Peterson trough spanning $\\sim 80$~comoving Mpc \\citep{white03}.  These differences suggest there may be large-scale variations in the enrichment and/or ionization state of the IGM at $z \\sim 6$, and perhaps a strong redshift evolution in the ionization state of metal absorption systems.  The metal-enriched, low-ionization regions towards SDSS~J1148$+$5251 may represent the final stages of reionization.  Alternatively, they may arise by chance from multiple pockets of enriched gas near galaxies aligned along a filament.  In order to effectively use metal lines as probes of reionization, the roles of enrichment and ionization must be separately understood.  A quasar sightline that passes through a largely neutral IGM may not show an \\oi\\ forest if the IGM has not yet been enriched.  At lower redshifts, highly-ionized species such as \\civ\\ and \\siiv\\ have been the primary tracers of IGM metallicity \\citep[e.g.,][]{meyer87, ssb88, petitjean94, cowie95, ellison99, schaye03, pettini03, aguirre04, simcoe04, scannapieco06}.  Both the comoving mass density of and the column density distribution of \\civ\\ remain roughly constant over at least $1.5 \\lesssim z \\lesssim 4.5$ \\citep{songaila01}.  Since the fraction of carbon in the form of \\civ\\ varies depending on the ionization state of the gas, however, a constant \\civ\\ density does not necessarily mean that the total IGM metallicity remains unchanged \\citep[e.g.,][]{oppenheimer06}.  Nevertheless, some fraction of IGM metals are clearly in place by $z \\sim 4.5$.  The current constraints on the abundance of metals in the IGM at $z > 5$ are mixed.  \\citet{songaila05} found tentative evidence for a decrease in the comoving mass density of \\civ\\ from $z \\sim 4$ to 5. To measure the metal content of the IGM at even redshifts the \\civ~$\\lambda1548,1551$ doublet must be pursued into the near-infrared.  Preliminary searches by \\citet{simcoe06} and \\citet{r-w06} found evidence for a comoving \\civ\\ mass density at $z \\sim 6$ that is consistent with measurements at $z \\sim 2-4.5$.  In addition, \\citet{simcoe06} found tentative evidence that the \\civ\\ column density distribution may also remain constant out to $z \\sim 6$.  These works examined only two sightlines, however, and were generally sensitive only to the rare, high-column density ($\\log{N_{\\mciv}} \\gtrsim 13.8$) systems that dominate the comoving mass density.  A complete census of \\civ\\ at $z \\sim 6$ requires a broader, more sensitive search.  In this series of papers, we present the results from a two-part search for intergalactic metals at $z \\sim 6$.  In this paper, we describe our survey for \\civ\\ performed with Keck/NIRSPEC in high-resolution mode.  This is the most sensitive search for \\civ\\ at $z > 5.3$ yet performed, partly due to the high resolution ($R \\approx 13,000$), which is better suited to search for narrow \\citep[$b \\sim 10$~\\kms,][]{rauch96} lines than previous studies.  By covering multiple sightlines with high sensitivity, we are able to place strong constraints on the abundance of the kinds of highly-ionized absorbers that are common at lower redshifts.  In the following paper, we will expand our search for low-ionization systems initially presented in \\citet{becker06}.  The combined results will provide unique constrains not only on the ionization state of the IGM, but on the star formation history of the Universe at $z > 5$.  The rest of the paper is organized as follows.  The data are presented in \\S2.  In \\S3 we give the results of our search for \\civ, along with our completeness estimates.  We show that the number density of \\civ\\ absorbers at $z > 5.3$ has declined sharply from $z \\sim 2-4.5$ in \\S4, and give constraints on the column density distribution and the contribution of \\civ\\ to the closure density.  In \\S5 we discuss possible reasons for the decrease in \\civ\\ with redshift and the significance of our lack of detections in light of the strong systems previously reported along other sightlines.  Finally, our conclusions are presented in \\S6.  Throughout this paper, we assume $\\mOm = 0.274$, $\\mOl = 0.726$, and $H_0 = 70.5$~\\kms\\ \\citep{komatsu08}.  \\begin{deluxetable*}{lccccc} \\tablecolumns{6} \\tablecaption{Summary of NIRSPEC Echelle Observations} \\tablehead{\\colhead{QSO\\tablenotemark{a}} & \\colhead{redshift} & \\colhead{$t_{\\rm exp}$~(hrs)} & \\colhead{$S/N$\\tablenotemark{b}} & \\colhead{$\\Delta z_{\\mciv}$\\tablenotemark{c}} & \\colhead{$\\Delta X_{\\mciv}$\\tablenotemark{d}} } \\startdata SDSS~J0002$+$2550 & 5.82 & 12.0 & $11.4-16.7$ & 0.458 & 2.33 \\\\ SDSS~J0818$+$1722 & 6.00 & 11.0 & $10.3-14.2$ & 0.638 & 3.29 \\\\ SDSS~J0836$+$0054 & 5.80 & 11.5 & $14.5-21.5$ & 0.438 & 2.23 \\\\ SDSS~J1148$+$5251 & 6.42 & 12.0 & $10.9-15.6$ & 0.679 & 3.44 \\enddata \\tablenotetext{a}{QSO names given as SDSS Jhhmm$+$ddmm} \\tablenotetext{b}{Middle 50\\% range of signal-to-noise per resolution element values for pixels within the wavelength interval used to search for \\civ} \\tablenotetext{c}{Total redshift interval used to search for \\civ} \\tablenotetext{d}{Total absorption path length interval used to search for \\civ \\vspace{0.05in}} \\label{tab:ne_obs} \\end{deluxetable*} ",
        "conclusions": "We have presented a survey for \\civ\\ at $z = 5.3-6.0$ towards four $z \\sim 6$ QSOs using Keck/NIRSPEC in echelle mode.  This is the most sensitive search for \\civ\\ at these redshifts to date, yet we find no \\civ\\ systems in any of our spectra.  Based on our completeness estimates and the column density distribution of \\citet{songaila01}, which remains largely invariant over $2 < z < 4.5$, we would have expected to detect roughly 13 systems in the range $13.0 < \\log{N_{\\mciv}} < 14.5$ in the case of no evolution out to $z \\sim 6$. Our lack of detections implies that the number density of \\civ\\ absorbers to which our data are sensitive declines by at least a factor of $\\approx 4.4$ (95\\% confidence) from $z \\sim 4.5$.  Our null result places strong constraints on the shape and normalization of the \\civ\\ column density distributions at $ z = 5.3 - 6.0$.  For a power-law distribution, $f(N_{\\mciv}) \\propto N_{\\mciv}^{-\\alpha}$, with a slope equal to that measured at $z \\sim 3$ ($\\alpha \\approx 1.8$), the normalization must decline by at least a factor of 4.4 (95\\%).  The distribution may also flatten or steepen, but only a narrow range of normalizations are allowed for each $\\alpha$ by our data and the fact that at least two strong ($\\log{N_{\\mciv}} > 13.8$) systems are seen in the two previously studied sightlines \\citep[][and confirmed herein]{simcoe06,r-w06}. Our limits on \\Ociv\\ are less strong due to the fact that the comoving density can be domiated by rare, strong systems.  Within the range $13.0 < \\log{N_{\\mciv}} < 15.0$, however, the constraints on \\fN\\ imply that \\Ociv\\ must decline by at least a factor of two from $z \\sim 2-4.5$ to $z = 5.3-6.0$ (see also Ryan-Weber et al., in prep). The fact that \\fN\\ and \\Ociv\\ are relatively constant at $z \\lesssim 4.5$ further means that most of the buildup of \\civ\\ in the IGM must occur over an interval $\\Delta z \\lesssim 1$, or $\\Delta t \\lesssim 300$~Myr.  The fact that \\civ\\ in the IGM rapidly increases between $z \\sim 5.3$ and $z \\sim 4.5$ suggests that ionization effects may be playing an important role.  That is, a significant fraction of intergalactic metals we see as \\civ\\ at $z \\sim 4.5$ may already be in place at $z > 5.3$, but in a lower ionization state.  This may especially be true if the metals are confined to relatively high overdensities \\citep[e.g.,][]{oppenheimer06}, but other factors may also be at work. An increase in the overall IGM metallicity, a wider dispersal of metals by galactic winds, and an increase in the \\civ/C ratio due to an increase in the intensity or a hardening of the UV background would all help to explain the rapid increase in \\civ\\ at $z < 5.3$.  To evaluate the relative importance of these effects, it is critical to measure the abundance of metals in lower ionization states.  We will do this in a subsequent work.  The contrast between our lack of detections along four sightlines and the presence of two strong \\civ\\ absorbers towards SDSS~J1030$+$0524 suggests that there may be large-scale variations in IGM enrichment and/or in the ionization state of the metal-enriched gas.  We are currently planning follow-up observations of SDSS~J1030$+$0524 to look for weak \\civ\\ systems.  The presence of additional \\civ\\ would more strongly indicate that this sightline is either more enriched or more highly ionized than those in our current sample.  Alternatively, the strong \\civ\\ systems may be associated with rare, UV-bright star-forming galaxies, in which case the enhancement of \\civ\\ would be a local effect.  This is supported by the overabundance of $i$-dropouts in the field of SDSS~J1030$+$0524 \\citep{stiavelli05,kim08}, although spectroscopic follow-up is required to be certain that these galaxies are associate with the \\civ\\ systems.  Finally, we note that the decline of \\civ\\ at $z > 5.3$ poses an interesting challenge to the use of low-ionization metal lines as tracers of reionization.  If the number density of \\civ\\ systems had remained constant out to $z \\sim 6$, then we would be more certain that the IGM is sufficiently enriched at early times for metal lines to serve as meaningful probes of the ionization state of the gas.  The disappearance of \\civ\\ at high redshifts, however, means that if an ``\\oi\\ forest'' is not found at $z \\gtrsim 6$, the result could be attributed to a deficit of metals rather than to a highly-ionized IGM. On the other hand, the detection of a significant number of neutral metal absorption systems at $z \\sim 6$ would have a double significance: that the IGM is at least partially metal-enriched by that redshift, and that pockets of the IGM were not permanently ionized until well after the early reionization epoch favored by CMB studies.  This will be the focus of the following paper.   \\vspace*{5.8mm}  The authors would like to thank Emma Ryan-Weber and Max Pettini for many stimulating conversations over the course of this work; as well as Martin Haehnelt, Bob Carswell, and Rob Simcoe for their comments on the first draft of this paper.  G.~B. was supported by the Kavli Institute for Cosmology at the Institute of Astronomy in Cambridge. M.~R. was supported by the National Science Foundation through grant AST 05-06845."
    },
    "0812/0812.0124_arXiv.txt": {
        "abstract": "High-energy emission from gamma-ray bursts (GRBs) is widely expected but had been sparsely observed until recently when the \\textit{Fermi} satellite was launched. If $>$~TeV gamma rays are produced in GRBs and can escape from the emission region, they are attenuated by the cosmic infrared background photons, leading to regeneration of $\\sim$~GeV-TeV secondary photons via inverse-Compton scattering. This secondary emission can last for a longer time than the duration of GRBs, and it is called a pair echo. We investigate how this pair echo emission affects spectra and light curves of high energy afterglows, considering not only prompt emission but also afterglow as the primary emission. Detection of pair echoes is possible as long as the intergalactic magnetic field (IGMF) in voids is weak. We find (1) that the pair echo from the primary afterglow emission can affect the observed high-energy emission in the afterglow phase after the jet break, and (2) that the pair echo from the primary prompt emission can also be relevant, but only when significant energy is emitted in the TeV range, typically ${\\mathcal{E}}_{\\gamma, >0.1~{\\rm TeV}} > Y {(1+Y)}^{-1} \\epsilon_{e} {\\mathcal{E}}_{k}$. Even non-detections of the pair echoes could place interesting constraints on the strength of IGMF. The more favorable targets to detect pair echoes may be the ``naked\" GRBs without conventional afterglow emission, although energetic naked GRBs would be rare. If the IGMF is weak enough, it is predicted that the GeV emission extends to $> 30-300$ s. ",
        "introduction": "High-energy emission from gamma-ray bursts (GRBs) has been expected and various theoretical possibilities have been discussed by numerous authors \\citep[see e.g.,][and references there in]{FP08}. In fact, EGRET detected several GRBs with GeV emission \\citep[e.g.,][]{Hur+94}. Recently, the \\textit{Fermi} satellite was launched and the onboard Large Area Telescope (LAT) is widely expected to detect high-energy ($>$ GeV) emission from a fraction of GRBs. In addition, other space- and ground-based gamma-ray observatories such as AGILE, MAGIC, VERITAS and HESS also regard GRBs as one of the main scientific targets. Theoretically there are the two main classes as high-energy emission mechanisms, i.e., leptonic and hadronic mechanisms. The leptonic mechanisms include synchrotron self-Compton (SSC) emission and external inverse-Compton emission, which are the most discussed scenarios for both the prompt and the afterglow emission components. High-energy SSC emission is produced by relativistic electrons that radiate seed synchrotron photons \\citep[e.g.,][]{SE01,ZM01,GG03}. In addition, there are various possibilities for external inverse-Compton emission. For example, prompt gamma-ray photons or the X-ray flare photons may act as seed photons for the relativistic electrons accelerated during the afterglow phase in the external shocks \\citep[e.g.,][]{Bel05,Wan+06}. The hadronic mechanisms include synchrotron radiation of high-energy baryons, synchrotron radiation of the secondary leptons generated in photohadronic interactions, as well as the photons directly produced from $\\pi^0$ decays. In order to see the baryon synchrotron radiation, sufficiently strong magnetic fields are typically required \\citep[e.g.,][]{GZ07,Mur+08a}. Otherwise, photohadronic components would dominate over the baryon synchrotron component as long as the photon density is high enough. Hadronic gamma rays can be observed only when the nonthermal baryon loading is large enough \\citep[e.g.,][]{MN06,AI07}. So far, both emission mechanisms have been widely considered in the standard scenario \\citep[see reviews, e.g.,][]{Mes06,Zha07}, i.e., the internal shock model for the prompt emission and the external shock model for the afterglow emission, respectively.  Both mechanisms can in principle produce $>1$ TeV photons, although high-energy photons may not escape from the source due to two-photon pair production, especially during the prompt emission phase \\cite{LS01,GZ08,MI08,Gra+08}. Even if such super-TeV photons can escape from the source, they still suffer from pair creation due to the interaction with the cosmic infrared background (CIB) or the cosmic microwave background (CMB). In particular, the direct detection of TeV photons would be difficult for GRBs with redshift $z > 1$. On the other hand, the electron-positron pairs resulting from the pair creation are still energetic, so that they up-scatter numerous CMB photons via the inverse-Compton process. Such secondary photons are able to reach the observer in a longer duration than the duration of primary emission, and a significant fraction of them may be observed with a time delay due to several effects such as magnetic deflection and angular spreading. Therefore, this emission is called \"pair echo\" emission, with a typical energy in the range of $\\sim (1-100)$~GeV. This pair echo emission is not only indirect evidence of the intrinsic TeV emission but also a clue to probe the weak intergalactic magnetic field (IGMF) of $B_{\\rm IG} < {10}^{-16}$ G \\cite{Pla95}.  The Plaga's method is hitherto the only one to probe very weak magnetic fields of $B_{\\rm IG} < {10}^{-16}$ G. Other methods utilizing Faraday rotation or cosmic microwave background are sensitive to magnetic fields of order $B_{\\rm IG} \\sim 1 ~ {\\rm nG}$ \\cite{Kro94}. The presence of very weak IGMFs has been predicted by several mechanisms, such as inflation \\citep[e.g.,][]{TW88}, reionization \\citep[e.g.,][]{GFZ00} and density fluctuations \\citep[e.g.,][]{Tak+05,Ich+06}. Observations of IGMFs in voids would give important information on the origin of the galactic magnetic fields \\citep{Wid02}, although they may be contaminated by astrophysical sources such as galactic winds or quasar outflows \\cite{FL01}.  In this paper, we reinvestigate the observational effects of the possible pair echo emission of GRB high-energy emission in the afterglow phase. Three criteria should be satisfied to detect pair echo emission: (1) the object must emit $\\sim$ TeV gamma rays leading to pair echoes; (2) the pair echo flux must be higher than the detector's flux sensitivity; and (3) the pair echo emission component must not be masked by other emission components. Concerning the point (1), TeV photons from GRBs can be emitted during both the prompt and the afterglow phases. Here we consider both as the primary emission components for the echoes, by acknowledging that during the prompt phase strong TeV gamma rays are expected only for a small fraction of GRBs due to the large $\\gamma \\gamma$ optical depth, as has been studied by various authors \\cite{DL02,MAN07,RMZ04,Tak+08}. Concerning the point (2), we need to evaluate the pair echo flux quantitatively. This flux depends on the amount of the CIB photons, the IGMF strength, and the source distance. As for the CIB, we use the acceptable CIB models given by Kneiske et al. (2002, 2004). In order to take into account of the effects of the IGMF properly, we adopt the formulation developed by Ichiki et al. (2008), which enables us to calculate the time-dependent spectra better than the previous works \\cite{Dai+02,DL02,RMZ04,Wan+04,MAN07}. In addition, we have also taken into account up-scatterings of the CIB photons as well as the CMB photons. This effect was neglected in the previous work for simplicity \\cite{Tak+08}, but it can be also important \\cite{MAN07}. In this work, we focus on the detectability of the \\textit{Fermi} LAT, which is the most suitable one for our purpose, but also touch upon the capabilities of other ground based TeV detectors such as MAGIC and VERITAS. Concerning the point (3), we pay special attention to the high-energy afterglow emission, which is the main competitor of the pair echoes, and compare the its strengths with respect to the echo components. Such a comparison was not done for previous researchers who studied the pair echo. At present, a detailed comparison between the pair echoes and high-energy afterglows is highly uncertain, as both have never been clearly detected. Since various predictions of high-energy emission rely on many model assumptions, they should be tested by observations of \\textit{Fermi}, MAGIC, VERITAS and other detectors. Despite of these uncertainties, we think it would be interesting and important to study effects of pair echoes that can affect high-energy emission, especially in the late phase \\cite{DL02,RMZ04,MAN07,Tak+08}. ",
        "conclusions": "In this paper, we have calculated the time-dependent spectra of the secondary pair echoes from the GRB prompt and afterglow TeV emission components that are attenuated by the CIB, applying a recently developed formalism to properly describe the temporal evolution of the pair echoes. We have compared the flux of the pair echoes to that of the afterglow, taking into account up-scattering of the CIB photons. In particular, we have demonstrated (1) that afterglow-induced pair echoes can be important after the jet break for long GRBs with a canonical afterglow; and (2) that prompt-induced pair echoes may also outshine the afterglow emission, if the prompt TeV emission is intense, typically with $\\mathcal{E}_{\\gamma, >0.1~{\\rm TeV}} > Y {(1+Y)}^{-1} \\epsilon_e \\mathcal{E}_k$.  Weak but non-zero IGMFs can be crucial for detectability, since they make the duration of the pair echo emission much longer than the time scale of primary emission (see Figs. 2, 4, 6, and 8). Although the detectability itself also depends on both of the spectral evolution of the primary emission and detector sensitivities, such non-zero IGMFs can make it easier to detect secondary photons at late times when the pair echo emission remains shallow compared to the afterglow emission. Concerning with the detection of pair echo signals, ``naked\" (short) GRBs without a significant afterglow emission could be more promising. The pair echo should be observed as extended emission with the time scale of $t > 30-300$ s. The observational prospects of such pair echoes are quite interesting for the recently launched \\textit{Fermi}. Successful detections may be possible for nearby, bright events, and would open a new window to study the poorly unknown IGMF. Even in the case of non-detections, lower limits on the IGMF of $B_{\\rm IG} \\cdot {\\rm min[\\lambda_{\\rm coh}^{1/2}, \\lambda_{\\rm IC}^{1/2}]} \\sim {10}^{-20}-{10}^{-21}~{\\rm G}~{\\rm Mpc}^{1/2}$ may be obtained.  The main caveat in hunting afterglow-induced pair echoes and pair echoes from short GRBs is that nearby bright GRBs do not seem frequent. Although there is large uncertainty on the nearby burst rate, the rate of bursts occurring within $z \\sim 0.3$ is estimated as $\\sim$ a few events per year \\citep[e.g.,][]{GPW05,GP06,LZVD07}. The actual detection rate also depends on several factors such as the detector sensitivity and field of view (e.g., $\\sim 2.4$ sr for the \\textit{Fermi}/LAT detector), so that only a fraction of them would be detected. If all the bursts are ideal TeV emitters, we can expect pair echoes for these bursts in the near future. However, it is unlikely that all the bursts are bright TeV emitters (and it seems more plausible for prompt emission due to significant attenuation by the pair creation). Although it is currently impossible to predict how many bursts can be bright TeV emitters in both the prompt and afterglow phases, the expected detection rate for $z < 0.3$ bursts would be ``at most\" $\\sim 1-2$ events per year. There may be further complications about nearby GRBs. Some of the nearby long bursts detected so far seem somewhat dimmer than classical GRBs occurring at $z > 1$, but their local rate may be higher than the estimated local rate of classical GRBs \\citep[e.g.,][]{GD07,LZVD07}. Hence, we may have more nearby bursts that can be detected in the keV-MeV band by detectors with better sensitivities (e.g., \\textit{EXIST}). But, since the typical luminosity of such low luminosity bursts seems small, it is not so easy to see pair echoes from them. In addition, energetic short GRBs assumed in Figs. 9 and 10 would also be rare, whose radiation energy is larger than the typical one ($\\mathcal{E}_{\\gamma}^{\\rm iso} \\sim {10}^{50-51}$ ergs). Nevertheless, possible detections of pair echoes would bring us a big impact in understanding GRB physics and IGMF, even though the bright TeV GRBs that can lead to such detections are rare. The current on-orbit \\textit{Fermi} satellite is suitable for such a purpose. MAGIC and VERITAS can also provide valuable data via follow-up observations, since the pair echo emission can last for a long duration of time. In the near future, some constraints on the models may be achieved even for non-detections.  We must also beware of the uncertainties in the intrinsic primary spectra since the pair echo flux depends on the amount of TeV photons. As for afterglow emission, we only consider the conventional forward shock model with energy injection. Although other parameter sets or other models such as the varying $\\epsilon_e$ model can be considered, we expect that the qualitative features of the pair echoes themselves will not be changed significantly, as long as the light curve of high-energy emission is similar to that of X-rays and the amount of TeV photons is not too different from that invoked in our case. As for the prompt emission, possible uncertainties may come from the intrinsic emission properties such as $E_{\\gamma}^{\\rm cut}$, as discussed in Murase et al. (2007).  The contamination by other high-energy emission components might complicate the picture further. There are many possibilities of high-energy gamma ray emission during the afterglow phase \\citep[see, e.g.,][and references therein]{Zha07,FP08}. For example, high-energy emissions associated with X-ray flares are expected at $\\sim$ GeV energies. GeV photons can be produced by both of the leptonic mechanisms \\citep[e.g.,][]{Wei+06,Wan+06,YD08} and the hadronic mechanisms \\cite{MN06}. In addition, the reverse shock electrons can also provide high-energy photons during the early afterglow phase. Nonetheless, it is in principle possible to distinguish the pair echo emission from other possibilities, given an ideal broad-band (optical, X-ray, MeV and GeV) observational campaign."
    },
    "0812/0812.2910_arXiv.txt": {
        "abstract": "{Stars in globular clusters (GCs) exhibit a peculiar chemical pattern with strong abundance variations in light elements along with a constant abundance in heavy elements. These abundance anomalies can be explained by a primordial pollution due to a first generation of fast rotating massive stars which released slow winds into the ISM from which a second generation of chemically anomalous stars can be formed. In particular the observed ratio of anomalous and standard stars in clusters can be used to constrain the dynamical evolution of GCs as around 95\\% of the standard stars need to be lost by the clusters. We show that both residual gas expulsion during the cluster formation and long term evolution are needed to achieve this ratio.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.3916_arXiv.txt": {
        "abstract": "Large-scale radial transport of solids appears to be a fundamental consequence of protoplanetary disk evolution based on the presence of high temperature minerals in comets and the outer regions of protoplanetary disks around other stars.  Further, inward transport of solids from the outer regions of the solar nebula has been postulated to be the manner in which short-lived radionuclides were introduced to the terrestrial planet region and the cause of the variations in oxygen isotope ratios seen in primitive materials.  Here, both outward and inward transport of solids are investigated in the context of a two-dimensional, viscously evolving protoplanetary disk.  The dynamics of solids are investigated to determine how they depend on particle size and the particular stage of protoplanetary disk evolution, corresponding to different rates of mass transport.  It is found that the outward flows that arise around the disk midplane of a protoplanetary disk aid in the outward transport of solids up to the size of CAIs and can increase the crystallinity fraction of silicate dust at 10 AU around a solar mass star to as much as $\\sim$40\\% in the case of rapidly evolving disks, decreasing as the accretion rate onto the star slows. High velocity, inward flows along the disk surface aid in the rapid transport of solids from the outer disk to the inner disk, particularly for small dust. Despite the diffusion that occurs throughout the disk, the large-scale, meridonal flows associated with mass transport prevent complete homogenization of the disk, allowing compositional gradients to develop that vary in intensity for a timescale of one million years.  The variations in the rates and the preferred direction of radial transport with height above the disk midplane thus have important implications for the dynamics and chemical evolution of primitive materials. ",
        "introduction": "The results of the Stardust mission, which collected roughly 0.001 g of material ejected from Comet Wild 2 \\citep{brownlee06}, suggests that large-scale transport in the solar nebula was a fundamental process in shaping the properties of the planetesimals that formed within it, and by extension, the planets that accreted from these building blocks.  Among the materials returned to Earth were crystalline Mg-rich silicate grains which require temperatures in excess of 1000 K to form from the amorphous precursors expected to have been inherited from the natal molecular cloud \\citep{hallenbeck00}, or even higher temperatures to condense directly from a gas phase \\citep[][and references therein]{wooden07}.  Further, a refractory grain called Inti, whose mineralogy and oxygen isotopic compositions is reminiscent of the Calcium, Aluminum-rich Inclusions (CAIs) found within chondritic meteorites was also found within the Stardust aerogel \\citep{brownlee06,zolensky06,mckeegan06}.  These findings have led many to conclude that large amounts of materials that formed or were processed in the hot, inner regions of the solar nebula migrated outwards tens of AU to be accreted by comets. The exact distance that materials were transported is still the matter of debate as comets likely originate from beyond 20 AU \\citep{charnoz07}, but the cometesimals from which these objects formed likely originated anywhere from $\\sim$5 AU \\citep{supulverlin00} to $\\sim$100 AU \\citep{weidenschilling97,weidenschilling04}.  That dust was mobile in the solar nebula is not surprising, as we have long recognized that transport processes must have operated in our solar nebula.  For example, the chondritic meteorites, which are relatively unaltered products of our solar nebula, come in a variety of types, each characterized by their bulk chemistry, oxygen isotopic ratios, as well as chondrule sizes and abundances \\citep[see review by][]{weisberg06}.  That each of these bodies record such different formation environments despite forming in close proximity to one another (a scale of 1-3 AU in the solar nebula, in all likelihood) and over a time period of 1-3 million years \\citep[e.g.][]{kita05} suggests that the solar nebula was either spatially heterogeneous or that its chemical and isotopic properties changed with time due to the transport of materials within \\citep{stevenson88,cyr98,cuzzi05,cieslacuz06}.  Further, the components of chondritic meteorites--the refractory inclusions, chondrules, and matrix which contains pre-solar grains--all record very different chemical and isotopic environments as well as different thermal histories prior to their incorporation into common meteorite parent bodies.  It is difficult to imagine how such diverse materials could be brought together without invoking solar nebular transport of some kind.  Fortunately, we can gain insight into the details and driving mechanisms for transport in our solar nebula by examining other protoplanetary disks.  Crystalline silicates are abundant in a number of disks around other stars at levels greater than those seen in molecular clouds, particularly in the hot inner regions, indicating that their formation occurs at high temperatures, while their presence in the cool outer regions of disks is taken as evidence that large-scale radial transport of these materials occurred \\citep{vanboekel04,apai05,watson07,wooden07,wooden08}.  That evidence for this processing and transport is seen in disks  around stars of all masses suggests that this is a fundamental consequence of disk evolution.  Thus by understanding the processes that control or operate during protoplanetary disk evolution we may identify what processes controlled how solids were transported in our own solar nebula.  Protoplanetary disks are known to be dynamic objects that transport mass inward to be accreted by their central stars as part of their stars' last stages of pre-main sequence evolution \\citep[see review by][and references therein]{dullemond07}.  Small dust particles that are entrained in the gas would also be subjected to these motions, thus providing a means of large-scale movement.  The exact driving force of this evolution remains the subject of ongoing work, but likely involves instabilities of some kind, such as the magnetorotational instability (MRI) of \\citet{bh91} or gravitational torques arising from massive clumps in marginally gravitationally unstable disks \\citep[e.g.][]{boss02,boley06}.  \\citet{hartmann98} demonstrated that the observed accretion rates of T-Tauri stars can be matched by adopting the classical $\\alpha$-viscosity model in which the evolution of the disk is treated as arising due to viscous stresses in the differentially rotating gas \\citep{ss73,lbp74}.  These stresses serve to drive mass inward with time and allow the disk to spread in the radial direction to conserve angular momentum.  The $\\alpha$-viscosity model is a turbulent viscosity model as it assumes that turbulence within the disk acts as the source of the viscosity, because the molecular viscosity of the gas is orders of magnitude too small to explain the observed mass accretion rates.  A similar model, the $\\beta$-viscosity model, has also been proposed in which the turbulence specifically arises due to hydrodynamic instabilities in the disk \\citep{richard99}.  This latter model assigns a viscosity to the disk whose functional dependence on disk properties differs slightly from the  $\\alpha$-viscosity model.  Because of the turbulent nature of the gas, the random motions associated with its evolution lead to a diffusivity of material within it, allowing dust particles to be redistributed in a manner to smooth out concentration gradients.  This diffusion has previously been investigated  to evaluate whether particles that formed or were processed in the hot, inner regions of the solar nebula could have been carried outwards against the inward flows associated with mass transport to be delivered to where comets formed \\citep{gail01,gail04,bockelee02}.  In the case of \\citet{bockelee02}, they found that such a scenario was possible, however, in order to get a large fraction of the silicate grains in the outer nebula to be crystalline, in agreement with the observations of comets Halley and Hale-Bopp, the nebula must have started out ``warm'' and compact.  That is, the entire nebula mass was initially distributed within 10 AU of the Sun, resulting in mass accretion rates in excess of 10$^{-5}$ M$_{\\odot}$/year.  Such a mass accretion rate is greater than observed for disks around roughly solar mass stars \\citep[e.g.][]{calvet05}, though the data is lacking for ages less than 10$^{5}$ years old.  \\citet{gail01,gail04} also found that high mass accretion rates favored the outward transport of high temperature materials, though the fractions of crystalline forsterite in these models were small ($\\sim$1\\%).   Previous studies of dust transport in protoplanetary disks, such as those described above, have largely relied on one-dimensional models where the changes in the surface (or column) density of the dust are calculated as a function of time and location as the disks evolved.  Thus dust at a given location of a disk is expected to follow the same dynamical behavior wherever it is located in that column. However, as discussed in detail in the next section, the dynamics of dust particles depend on the local gas density, pressure, and their respective gradients, all of which vary with height in a protoplanetary disk.  The importance of such considerations for the chemical evolution of protoplanetary disks has been recognized in recent work by \\citet{tschar07}, \\citet{cieslasci07}, and \\citet{wehrstedt08} in which the rates as well as the preferred direction of radial transport were found to vary with height above the disk midplane.   Here, the model of \\citet{cieslasci07} is extended to examine the roles that these variations played in the transport of primitive materials in our solar nebula as well as other protoplanetary disks.  The model of \\citet{cieslasci07} is used as it accounts for the motions that develop as a result of gas drag and vertical settling, which are not included in \\citet{tschar07} and \\citet{wehrstedt08}, and therefore can be applied to cases beyond vapor species and sub-micron dust particles.  In the next section, the two-dimensional (radial and vertical) model of particle transport in protoplanetary disks is described, with the key differences with one-dimensional (radial) models highlighted.  Section 3 describes the protoplanetary disks to which this two-dimensional model is applied, with the specific numeric approaches described in Section 4.  Section 5 presents the results for transport calculations for disks of the various structures described.   The implications that these results have for the evolution of primitive materials in our solar nebula and other protoplanetary disks are then discussed. ",
        "conclusions": "The results of the different models presented here demonstrate that the vertical variations in the rates and preferred directions of radial transport of solids in a protoplanetary disk are important to consider when talking about the early evolution of primitive materials.  These variations arise due to the fact that the gas densities, pressure, and their respective gradients, all of which are responsible for driving motions in a protoplanetary disk, are strong functions of location.  In terms of the viscous evolution of the disk, the vertical variations in gas density are responsible for driving a flow structure in which the gas moves toward the star rapidly at high altitudes in the disk, but moves outward, away from the star around the disk midplane.  The inward motions of the gas are responsible for the mass accretion onto the central star, representing its final stages of pre-main sequence growth, whereas the outward motions are responsible for the outward transport of angular momentum to compensate.  The relative fraction of material transported outward at a given location can range from $\\sim$5 to 30\\% of the mass that is transported inward at a given time \\citep{kg04}.  Further, the gas drag induced velocities of solids in the solar nebula will also vary with height as the radial pressure gradient, which is responsible for the gas rotating at a non-Keplerian rate, also varies with height.  Unlike the large-scale flows caused by viscous evolution, however, the gas drag induced velocities will reach their largest, negative (towards the central star) velocities near the disk midplane.  At higher altitudes, the inward velocities diminish, and in fact, will become positive as the radial pressure gradient switches signs, leading the gas to orbit at super-Keplerian rates at higher altitudes.  These variations in the radial transport models are overlooked in traditional one-dimensional models of material transport in protoplanetary disks \\citep[e.g.][]{stepval97,gail01,bockelee02,bitw03,cieslacuz06}.  Instead,  the motions of the particles due to the viscous evolution of the disk in these models are found by assuming that those particles are uniformly suspended throughout the height of the disk (thus ignoring vertical settling) and their net velocity is found by calculating the vertically averaged motions of the gas at that location ($V_{r}$=-3$\\nu$/2$r$ under steady-state condtions).  For those 1D models that account for the radial motions due to gas drag \\citep[e.g.][]{stepval97,bitw03,cieslacuz06}, these velocities are generally found by calculating the inward velocity at the disk midplane and applying that velocity to all particles, regardless of where they are located in the vertical column above the disk midplane.  While this may be appropriate for particles which are tightly clustered around the disk midplane due to vertical settling ($\\sim$meter-sized bodies which do not experience significant vertical diffusion, for example), this would overestimate the inward gas drag velocities of particles lofted to higher altitudes.  Thus while 1D models are instructive for investigating the general dynamical evolution of materials in protoplanetary disks, consideration of the 2D effects can lead to important differences in the dynamical evolution of primitive materials.   Here, as in \\citet{cieslasci07}, it was shown that the outward flows that are generated around the midplane in a viscously evolving disk aid in the outward transport of materials by providing an environment in which materials can diffuse outwards without being hindered by the accretional flows associated with protoplanetary disk evolution.  This provides a path by which high temperature materials (refractory inclusions and crystalline grains) collected from comet Wild 2 or seen in comet Halley and Tempel 1 could be delivered to the outer solar nebula.  Predicting the particular fraction of silicate grains in the outer disk  that would be crystalline in each of the models presented here is not straightforward, unfortunately.  This fraction depends on the amount of material that is delivered to the outer disk as well as the amount of unprocessed material that is present.  The abundance of this material will be determined by its dynamical interaction with the gas in the outer disk, the extent of the disk (the initial inventory of unprocessed materials) and the rate at which these unprocessed materials coagulate into larger bodies.  A first order approximation, however, would be that the total processed and unprocessed silicate grain abundance would add up to the solar ratio ($C_{crystal}+C_{amorphous} \\sim$ 0.005).  This would suggest that the crystalline grain fraction at 10 AU at the end of each simulation would be $\\sim$40\\% for each of the $\\dot M=10^{-6} M_{\\odot}$/yr cases, $\\sim$10\\% for the $\\dot M=10^{-7} M_{\\odot}$/yr cases, and $<1\\%$ for the $\\dot M=10^{-8} M_{\\odot}$/yr cases.  These numbers should only be taken as rough estimates, however, and future work should follow both the unprocessed and processed materials simultaneously.  These numbers are generally comparable to those found by \\citet{wehrstedt08} in their two-dimensional models, though direct comparisons are not straightforward due to different model treatments.   For example, \\citet{wehrstedt08} assumed that grains could become crystalline at temperatures as low as 800 K, allowing for a larger volume of grains to be processed at any given time.  More importantly, however, as discussed above, these authors focused on the dynamics of dust particles in the region around the midplane, rather than considering the dynamics of solid particles at very high altitudes.  These higher altitudes still contribute significantly to the inward mass flux in the disk.  Thus these models likely overestimate the outward flux of high temperature materials and underestimate the inward flux of unprocessed materials.  It should be noted that the differences due to the particular choice of viscosity model ($\\alpha$-viscosity in this work, $\\beta$-viscosity in \\citet{wehrstedt08}) needs to also be quantified in future studies.   Not only does the rate of mass accretion in the disk determine the amount of outward transport that occurs in a disk, but the timing of this transport as well. For disks with mass accretion rates in excess of $\\dot M >$10$^{-7} M_{\\odot}$/yr,  high-temperature materials can be transported beyond 20 AU on timescales of $\\sim$10$^{5}$ years.  Such rapid outward transport is likely necessary given that  the formation timescale of the giant planets via core accretion is approximately 1-2 Myr \\citep{hubickyj05}.  Delivery of high temperature materials on longer timescales, such as a few million years, would be problematic as the outward transport of the grains would likely be blocked by the young giant planets, either because the cores would accrete the grains themselves or because the planets would begin to open a gap in the disk through which particle transfer would likely be difficult.  At lower mass accretion rates, with  $\\dot M \\sim$10$^{-8} M_{\\odot}$/yr, typical of 10$^{6}$ year old T-Tauri stars \\citep{hartmann98}, delivery of crystalline silicates or CAIs to the outer disk is difficult as high temperatures are generated only at the very inner edge of the disk.  The lower outward flow velocities along the disk midplane and the smaller volume of the disk in which high temperature materials can be formed or processed means that delivery of crystalline grains and refractory inclusions to beyond 20 AU will not occur as quickly as in the higher mass accretion rate models.  Thus it is likely that the high temperature material seen in comets were processed and delivered to the outer solar nebula during the early stages of disk evolution.  However, even at lower mass accretion rates, such as $\\dot M \\sim$10$^{-8} M_{\\odot}$/yr, the flows at the disk midplane are likely still important in terms of understanding the dynamical evolution of sub-cm sized objects, as the outward flows velocities in the inner disk are $>$10 cm/s in the model presented here.  For the $a$=0.5 mm particles (which is roughly the size of chondrules) considered here, this means that the outward flows at the disk midplane exceed the inward drift velocities at the disk midplane that arise due to gas drag (Figure 3) out to about 20 AU.  In fact, given that inward drift velocity for small particles is roughly proportional to its stopping time \\citep{cuzzweid06} the outward flows at the disk midplane would exceed the inward drift velocities of objects with radii of 5 mm (comparable to the larger CAIs found in CV chondrites) out to a few AU, offering a way of preserving these objects in the nebula for millions of years before the formation of chondrules and accretion of chondrites \\citep{amelin02,kita05}.    The variations in flow structure also are important to consider when discussing the inward transport of materials from the outer parts of the protoplanetary disk.  Materials that would be injected into the solar nebula, rain down after being lofted upward in a protostellar jet, or are created due to photochemistry on the disk surface would initially be subjected to rapid inward flows in the very upper regions of the disk.  This would remain true until the particles are transported down towards the midplane where the inward flows diminish in magnitude or the flows reverse direction.  This removal from the upper part of the disk occurs more rapidly for larger particles due to their greater settling velocities, meaning that smaller particles can be transported to the inner disk on shorter timescales than larger ones.  In fact, because smaller particles can remain lofted at high altitudes for long periods of time, they are more likely to be accreted onto the central star in accretionary flows as compared to the larger bodies that settle around the disk midplane as demonstrated in the simulations presented above.  This remains true as long as the velocities due to viscous evolution of the disk exceed those of that arise due to gas drag. This result is not limited to materials that begin at high altitudes.  Vertical diffusion would loft small particles to higher altitudes than larger ones, thus introducing them to the regions of rapid inward flow more readily.   This runs counter to the previous ideas that millimeter to centimeter sized objects would be more difficult than fine dust to retain in the solar nebula due to their inward drift velocities caused by gas drag.  Instead, it appears that aggretation into bodies of this size may actually have allowed them to be retained in the solar nebula more effectively.   It should be noted that millimeter-sized grains themselves do not necessarily serve as good analogs for the materials that would be injected into the solar nebula by a supernova \\citep{ouellette07} or for the heavy ($^{16}$O-poor) water ice that would be created by CO self-shielding in the outer nebula \\citep{ly05}, as these materials are likely contained within or adsorbed by grains much smaller and closer to the $a$=5 $\\mu$m cases discussed above. However, it should be kept in mind that these smaller grains would likely have experienced collisions with other grains that led to sticking and growth.  \\citet{weidenschilling97} showed that micron-sized grains would coagulate into millimeter-sized grains at 30 AU on timescales of less than 10,000 years.  Those coagulation calculations assumed a non-turbulent and non-evolving disk; growth at these small sizes may actually be accelerated by turbulence as the random motions of the dust particles would lead to more frequent collisions \\citep{weidenschilling84,dd05,ccoag07,ormcuzzi07}.  Growth to very large sizes may be inhibited due to turbulence, as larger particles will develop relative velocities that become destructive in collisions, rather than accretional, but millimeter-sized grains appear to be below this threshold \\citep[see discussion in][]{cuzzweid06}.  Thus, the results of the $a$=0.5 mm models are likely more reflective of how solids from the upper layers of the outer disk get transported to the inner disk.   The issue of retention is particularly important for the case of injection of short-lived radionuclides directly into the solar nebula \\citep{ouellette07}.   If this material was injected, it would likely have been done in one or more Rayleigh-Taylor fingers \\citep{boss07}, with materials being peppered along the surface of the nebula as assumed here.  \\citet{ouellette07} demonstrated the ratios of $^{26}$Al/$^{27}$Al and $^{60}$Fe/$^{56}$Fe would be close to the inferred meteoritic values if the solar nebula were located 0.15 pc from a supernova.  These calculations, however, assumed that these materials became well mixed in the nebula.  However, as shown here, injected material that was present as small grains ($<$10 $\\mu$m, as expected for the supernova grains) at the surface of the nebula would be lost to the Sun faster than larger solids, which would have settled around the midplane.  Thus to ensure that these injected materials become well mixed in the solids, they would quickly have to be accreted into the larger bodies that had already formed.  If they were not, the injected material would be lost at a faster rate than the native solids, thus leading to lower ratios of the abundance of short-lived radionuclides to stable isotopes.  This would be particularly true if solids were injected closer to the Sun than the $\\sim$20 AU distance considered here as accretion onto the sun would happen on shorter timescales.  This would require the solar nebula to be closer to the supernova to achieve the meteoritic ratios.  Such considerations should be made in future studies.  An important point to note in each of these simulations is that the concentration at the end of each of the ``Inward Transport'' simulations shown here is not uniform throughout the disk, but rather, gradients exist.  This is particularly seen in panel C of Figures 12, 15, and 18, which show that the concentrations of the species at the disk midplanes vary with distance from the star.  Radial mixing timescales for cases such as these are typically defined as the diffusive timescale, $t_{mix}$=$r^{2}$/$\\nu$,  which for the inward transport calculations presented here would be 6$\\times$10$^{5}$, 1.5$\\times$10$^{6}$, and 2.5$\\times$10$^{6}$ years at 20 AU for each of the disks considered.  These numbers represent estimates for the amount of time it would take for the ``injected'' material tracked here to mix into the very inner part of the protoplanetary disk.  While the model simulations are not run to these times for the cases considered here, it can still be seen that delivery of significant amounts of material occurred on shorter timescales.  However, complete homogenization (little to no concentration gradients) of the injected material had not been achieved at the end of the simulations, nor is it expected to be achieved at any reasonable time afterward.  The evidence for this is that gradients exist at every point in the disk, regardless of the distance from where the materials originated.  Because of the variation in the direction and magnitudes of the large-scale flows in the disk, conditions are constantly changing, not allowing diffusion to work to smooth out the concentration gradients completely.  The agreement between Pb-Pb ages and Al-Mg ages of CAIs and chondrules \\citep{amelin02} suggests that $^{26}$Al was uniformly distributed relative to the stable isotope $^{27}$Al throughout the inner solar nebula where these materials formed, implying that if injected, $^{26}$Al was homogeneously distributed throughout the inner solar system where these objects formed. Such apparent homogenization in the injection model could be achieved if mixing had occurred in such a way that these objects formed near the concentration maximum, where the radial gradients in concentration are the lowest (but non-zero). Alternatively, near homogenization of injected materials could have been achieved if the materials were contained in small dust grains and they were injected at a late stage of nebular evolution.  Radial concentration gradients are smallest in the cases presented for the 5 $\\mu$m grains in the disk accreting at 10$^{-8} M_{\\odot}$/year, as the large-scale flows slowed to the point that diffusion became more effective at smoothing out gradients.  However, this low mass accretion rate also would not produce very high temperatures in the disk, meaning that CAIs are unlikely to form or their formation and redistribution was limited to only regions very close to the Sun.  Further, the small dust grains would likely have coagulated into larger objects, such as the 0.5 mm grains considered here, which show comparable or stronger radial gradients in all cases considered.  Therefore, if $^{26}$Al was injected directly into the nebula, it would likely not have been perfectly homogenized, but rather gradients, or variations in its concentration, would have existed. These variations would lead to apparent age differences between chondritic components, the magnitude of which depended on the particular values of $^{26}$Al/$^{27}$Al that were recorded by those components.  For example, a 10\\% variation in the $^{26}$Al/$^{27}$Al  ratio could easily develop over a scale of a few AU based on the results of the simulations presented here, and also is in line with the variations found by \\citet{boss08} for materials injected into a marginally gravitationally unstable disk.  If two solids formed contemporaneously, one with $^{26}$Al/$^{27}$Al =5$\\times$10$^{-5}$, the canonical value within CAIs \\citep[see][for the most recent discussion on this issue]{jacobsen08}, and one with $^{26}$Al/$^{27}$Al =4.5$\\times$10$^{-5}$, this would be interpreted as an age difference of $\\sim$100,000 years, which is less than the uncertainty on the Pb-Pb ages.   If this were the case, and $^{26}$Al was injected directly into the solar nebula, the inherent uncertainty in Al-Mg ages would be at least $\\sim$100,000 years.   Recently, measurements of $^{60}$Fe, whose presence in solar system materials \\citep[e.g.][]{tachibana03} has been among the motivating factors for looking at the direct injection of short-lived nuclides into the young solar system from a supernova, have helped to shed light on how well homogenized potential supernova derived materials were in the solar nebula. \\citet{bizzarro07} originally  reported measurements of $^{60}$Ni deficits in meteorites and terrestrial samples  that indicated $^{60}$Fe was not present in planetesimals that accreted within the first 10$^{6}$ years of the formation of the solar system.  This finding supported the model of direct injection into the solar nebula by a supernova \\citep{ouellette07}.  In this scenario, the 10$^{6}$ year interval represented the time it took for $^{60}$Fe to be injected and mixed in to the inner solar nebula.  More recently however, \\citet{dauphas08} analyzed similar samples as \\citet{bizzarro07}, but found that  $^{60}$Fe was homogenized in the terrestrial planet region of the solar nebula such that its concentration varied by less than 10\\%.  Again, these results could be achieved in the injection model if the maximum in the concentration of $^{60}$Fe had migrated to the inner solar system at the time that planetesimals there began to form.  If this were not the case, or if the variation in $^{60}$Fe is much less than 10\\% in planetary materials, the origin of short-lived radionuclides may be best understood as resulting from the injection into the natal molecular cloud prior to (or as the cause of) collapse \\citep{boss08sn}.   The results presented here differ somewhat from those of \\citet{boss08} who showed that the concentration of ``dye'' that was injected to a marginally gravitationally unstable disk would be mixed to the point that there would be less than 10\\% variation in its concentration throughout the \\emph{entire} disk after $\\sim$1000 years.  This may represent a fundamental difference in how transport occurs in gravitationally unstable disks and viscous disks, or it may represent differences in the ways that calculations were performed.  For example, while the study described here tracked the evolution of solids, the dye in \\citet{boss08} was more akin to vapor in the disk, meaning that the dynamics of vertical settling and gas drag were not considered.  While this may approximate the dynamical evolution of small dust particles which are well coupled to the gas, as described by \\citet{hagboss03}, these processes are likely important as solids grow beyond a few microns in size.  Future work aimed at determining the differences in how solids evolve in marginally gravitationally unstable and viscous disks will help to identify which model better describes the early evolution of our solar nebula.   \\emph"
    },
    "0812/0812.1855_arXiv.txt": {
        "abstract": "We present near- and mid-infrared observations on the shock-cloud interaction region in the northern part of the supernova remnant HB21, performed with the InfraRed Camera (IRC) aboard \\akari{} satellite and the Wide InfraRed Camera (WIRC) at the Palomar 5 m telescope. The IRC 7 \\um{} (S7), 11 \\um{} (S11), and 15 \\um{} (L15) band images and the WIRC \\HH{} 2.12 \\um{} image show similar shock-cloud interaction features. We chose three representative regions, and analyzed their IRC emissions through comparison with \\hh{} line emissions of several shock models. The IRC colors are well explained by the thermal admixture model of \\hh{} gas---whose infinitesimal \\hh{} column density has a power-law relation with the temperature $T$, d$N\\sim T^{-b}dT$---with \\nhh{} $\\sim10^3$ \\ncm, $b\\sim3$, and \\NhhIRC{} $\\sim3\\times10^{20}$ \\Ncm. The derived $b$ value may be understood by a bow shock picture, whose shape is cycloidal (cuspy) rather than paraboloidal. However, this picture raises another issue that the bow shocks must reside within $\\sim0.01$ pc size-scale, smaller than the theoretically expected. Instead, we conjectured a shocked clumpy interstellar medium picture, which may avoid the size-scale issue while explaining the similar model parameters. The observed \\HH{} intensities are a factor of $\\sim17-33$ greater than the prediction from the power-law admixture model. This excess may be attributed to either an extra component of hot \\hh{} gas or to the effects of collisions with hydrogen atoms, omitted in our power-law admixture model, both of which would increase the population in the $\\upsilon=1$ level of \\hh. ",
        "introduction": "\\label{intro} Shock-cloud interactions in supernova remnants (SNRs) have been studied for their wide astrophysical importance, such as the followings. Firstly, we can study the physics and chemistry of molecular shocks \\citep{Draine(1983)ApJ_264_485,Hollenbach(1989)ApJ_342_306}, to be used as a basis for interpreting astrophysical phenomena. Secondly, the evolution of shocked molecular clouds is closely related to the star formation occurring in them \\citep{Elmegreen(1978)ApJ_220_1051,Vrba(1987)ApJ_317_207,Reynoso(2001)AJ_121_347}. Lastly, the evolution of the SNRs interacting with molecular clouds is presumed to be significantly different from that of SNRs that do not interact with molecular clouds \\citep[e.g.][]{Rho(1998)ApJ_503_L167,Shinn(2007)ApJ_670_1132}. In addition, SNR shocks may have a simpler configuration in their interaction region than do shocks in outflows, since the latter can be generated by a diverse range of bullets, clumps, and/or winds \\citep[cf.~Fig.~11 in][]{Hollenbach(1989)inproc}. Hence, SNRs provide geometrically simpler conditions to study the shock-cloud interaction.  The shock-cloud interaction manifests itself through radiative means \\citep{Draine(1993)ARA&A_31_373,Wardle(2002)Science_296_2350}: high-velocity wings on low-$J$ CO rotational lines \\cite[e.g.][]{Koo(1997)ApJ_485_263,Koo(2001)ApJ_552_175}; high-$J$ CO rotational lines and OH rotational lines in the far-infrared \\cite[e.g.][]{Watson(1985)ApJ_298_316,Melnick(1987)ApJ_321_530}; radio maser lines of OH \\cite[e.g.][]{Frail(1994)ApJ_424_L111}; and infrared emission lines of \\hh{} \\cite[e.g.][]{Neufeld(2007)ApJ_664_890}. Among these, the infrared emission lines of \\hh{} are useful in studying the excitation condition of shocks, since \\hh{} is the most abundant molecule in interstellar clouds \\citep{Snow(2006)ARA&A_44_367} and the lines, originating from quadrupole transition, are optically thin and closely spaced through $\\sim1-30$ \\um. In addition, high resolution imaging is easier for \\hh{} infrared emission lines, than for radio emission lines.  Infrared \\hh{} emission lines have been frequently observed in SNRs interacting with nearby clouds. One noticeable common result is that the \\hh{} level-population diagram shows an \\emph{ankle-like} curve (cf.~Fig.~\\ref{fig-pop}) in the range of energy levels, $0-25,000$ K (\\citealt{Oliva(1990)A&A_240_453,Richter(1995)ApJ_454_277,Richter(1995)ApJ_449_L83,Rho(2001)ApJ_547_885,Neufeld(2008)ApJ_678_974}; see Fig.~7a in \\citealt{Rho(2001)ApJ_547_885} for the \\emph{ankle-like} curve), which has been usually interpreted as a population of two-temperature ($\\sim10^2$ and $\\sim10^3$ K) \\hh{} gas in Local Thermodynamic Equilibrium (LTE). Such a population cannot be explained by a single planar shock model \\citep[e.g.~Table~2 of][]{Wilgenbus(2000)A&A_356_1010}. Hence, several mechanisms have been proposed in order to explain this level population distribution, such as a partially dissociative J-shock \\citep{Brand(1988)ApJ_334_L103,Burton(1989)inproca,Moorhouse(1991)MNRAS_253_662}, a bow shock \\citep{Smith(1991)MNRAS_248_451}, and a non-stationary shock \\citep{Chieze(1998)MNRAS_295_672,Cesarsky(1999)A&A_348_945}.  However, the interpretation of the \\hh{} population diagram is still under debate; for instance, even for the same target SNR, IC443, all the above three models were preferred by different authors \\citep{Richter(1995)ApJ_454_277,Cesarsky(1999)A&A_348_945,Neufeld(2008)ApJ_678_974}. One restriction that hinders a clear understanding of the \\hh{} population diagrams, and subsequently of the shock-cloud interaction, is that the observations hitherto performed were not able to resolve the interaction features sufficiently to distinguish their structures, such as planar shocks, bow shocks, shocked clumps etc. Our target SNR, G89.0+4.7 (conventionally known as HB21), provides an excellent example to study the shock-cloud interaction, avoiding this restriction, because of its large angular size and proximity.  HB21 is a large ($\\sim120\\arcmin\\times90\\arcmin$), middle-aged ($\\sim5000-7000$ yr, \\citealt{Lazendic(2006)ApJ_647_350,Byun(2006)ApJ_637_283}) SNR, located at the distance of $\\sim0.8-1.7$ kpc \\citep{Leahy(1987)MNRAS_228_907,Tatematsu(1990)A&A_237_189,Byun(2006)ApJ_637_283}. It has been suspected of interacting with a molecular cloud, because of its deformed shell-like shape in the radio and the existence of nearby giant molecular clouds \\citep{Erkes(1969)AJ_74_840,Huang(1986)ApJ_309_804,Tatematsu(1990)A&A_237_189}. Such an interaction was confirmed with the detection of broad CO emission lines near the edge and the center of the remnant \\citep{Koo(2001)ApJ_552_175,Byun(2006)ApJ_637_283}. This detection also supported the suggestion that the cloud evaporation might be responsible for the enhanced thermal X-rays in the central region \\citep{Leahy(1996)A&A_315_260}.  We have obtained infrared images of two specific regions of the remnant, where the shock-cloud interaction is under way. The near- and mid-infrared band images ($\\sim2-27$ \\um), obtained from two infrared cameras respectively onboard the \\akari{} satellite and the Palomar 5 m Hale telescope, show several shock-cloud interaction features. We derive their brightness ratios and analyze them on the basis of the \\hh{} emission model from shocked gas. We find that their infrared properties are well described by shocked \\hh{} gas with a power-law distribution of temperatures. ",
        "conclusions": "We observed the shock-cloud interaction region in the SNR HB21 at near- and mid-infrared wavelengths, with the WIRC at the Palomar telescope and the IRC aboard the \\akari{} satellite. The IRC S7, S11, and L15 band images and the WIRC \\HH{} image reveal similar diffuse features, such as a blast wave, a bow shock, and a shocked clump. We chose three representative regions---N1wake, N2clump, and N2front---and analyzed their infrared characteristics with several different shock models, on the basis that \\hh{} emission lines are the radiation source of the shock-cloud interaction features.  We found that the IRC colors are well explained by an admixture model of \\hh{} gas temperatures, whose infinitesimal column density varies as d$N\\sim T^{-b}dT$. Three physical parameters---\\nhh, $b$, and \\NhhIRC---were derived from this thermal admixture model (see Table~\\ref{tbl-par}). The derived $b$ value ($\\sim3$) can be understood by a bow shock picture, whose shape is cycloidal rather than paraboloidal. However, the bow shock interpretation has a size-scale problem for the N2front, because the observations require that the bow shock be hidden \\emph{within} $\\sim0.01$ pc $\\sim3\\times10^{16}$ cm, to be unresolved in the WIRC \\HH{} image ($\\Gamma\\sim1.2$\\arcsec); this scale is smaller than the expected size of bow shocks. Instead of the bow shock picture, we propose a shocked clumpy ISM picture to simultaneously explain the obtained model parameters and the mid-infrared appearances. To confirm this picture, more robust theoretical studies and observations are required. We also compared the observed \\HH{} intensity to the predicted intensity from the power-law admixture model. The observed \\HH{} intensities are a factor of $\\sim17-33$ greater than the predicted ones. This excess might be caused by either an additional component of hot, dense H2 has (which has low total column density), or through the omission of collisions with hydrogen atoms in the power-law admixture model (which results in an under-prediction of the near-IR line intensity)."
    },
    "0812/0812.1593_arXiv.txt": {
        "abstract": "{ Disk galaxies are the most frequent objects among the galaxy population in the local universe. However, the formation of their disks and substructures --  in particular their bars -- is still a matter of debate. } { We present a physical model of the formation of J033239.72-275154.7, a galaxy observed at $z=0.41$ and characterized by a big young bar of size $6$ kpc. The study of this system is particularly interesting for understanding the connection between mergers and bars as well as the properties and fate of this system as it relates to disk galaxy formation. } { We compare the morphological and kinematic properties of J033239.72-275154.7, the latter obtained by the GIRAFFE spectrograph, to those derived from the merger of two spiral galaxies described by idealized N-body simulations including a star formation prescription. } { We found that the general morphological shape and most of the dynamical properties of the object can be well reproduced by a model in which the satellite is initially put in a retrograde orbit and the mass ratio of the system is 1:3. In such a scenario, a bar forms in the host galaxy after the first passage of the satellite where an important fraction of available gas is consumed in an induced burst. In its later evolution, however, we find that J033239.72-275154.7, whose major progenitor was an Sab galaxy, will probably become a S0 galaxy. This is mainly due to the violent relaxation and the angular momentum loss experienced by the host galaxy during the merger process, which is caused by the adopted orbital parameters. This result suggests that the building of the Hubble sequence is significantly influenced by  the last major collision.  In the present case, the merger leads to a severe damage of the disk of the progenitor, leading to an evolution towards a more bulge dominated galaxy. } {% } ",
        "introduction": "The formation of disk galaxies remains an outstanding puzzle in contemporary astrophysics (see Mayer et al. 2008 for a review). According to hierarchical models of structure formation, mergers and interaction of galaxies are an essential ingredient of galaxy formation and evolution. Earlier works and numerical simulations show that the remnants of mergings of purely stellar progenitors are more likely to be elliptical galaxies (Toomre 1977; Barnes 1988; Barnes \\& Hernquist 1992; Hernquist 1992; Lima-Neto \\& Combes 1995; Balcells \\& Gonz{\\'a}lez 1998; Naab et al. 1999 etc.) and recent studies extended this result to gas-rich progenitors (Springel \\& Hernquist 2005; Robertson et al. 2006; Hopkins et al. 2008). Whether this is compatible, or incompatible, with the fraction of disk galaxies present in the local universe depends on the typical number of expected mergers by galaxy (see for instance Kazantzidis et al. 2007). Such predictions seems to be incompatible with observations which suggest that disk galaxies represent the majority (70\\%) of the galaxy population observed in the local universe (see Hammer et al. 2005; Nakamura et al. 2004 and references therein). To help resolve such issues, Hammer et al. (2005) suggested that disks can be rebuilt during the encounters of gas rich spirals. Indeed such proposition was guided by the remarkable coincidence of the redshift increase, up to z=1, of the merger rate, of the fraction of actively star forming galaxies (including luminous IR galaxies, LIRGs), of the fraction of galaxies with peculiar galaxies (including those with compact morphologies) etc... This has been followed by simulations of gas rich mergers (Springel \\& Hernquist 2005; Robertson et al. 2006; Governato et al. 2007; Hopkins et al. 2008), evidencing that under certain conditions, a disk may be re-built after a merger. Indeed such a conclusion had been already reached by Barnes (2002), but these simulations are less convincing since they do not include a prescription for star formation. Lotz et al. (2006) have also analysed a large suite of simulated equal mass rich mergers and find that most merger remnants appear disc-like and dusty. Moreover, this scenario is also consistent with results of other simulations and semi-analytical models which claim that, without merger processes, most of galaxies and their host dark matter halos cannot acquire the required angular momentum to form disks (see Peirani et al. 2004; Puech et al. 2007 and references therein). The formation of bars is also a fundamental issue in the evolution of disk galaxies, particularly since it has now been shown that only about one third of them were in place at $z$ = 0.8 (Sheth et al. 2008).   Our aim here is to reproduce with numerical modelling the general morphology (presence of a bar, of substructures, etc...), the dynamical properties of the gas component (velocity fields), the photometric properties of stars (colors, star formation rate, etc...) of J033239.72-275154.7, a galaxy located at $z=0.41$ for which we have high quality imaging and kinematics. This work will also provide useful input to disk formation models, since it gives information on potential progenitors of the present-day galaxy disks, as well as constraints on their formation (initial orbital configuration, mass ratio of the system, etc...) and possible fate. This work is an offspring of a VLT large program entitled IMAGES (``Intermediate-MAss Galaxy Evolution Sequence'', Yang et al. 2008) which gathers high quality kinematics for a representative sample of $\\sim 100$ massive galaxies at $z=0.4-0.75$ and with $M_J(AB)\\leq-20.3$. Using the GIRAFFE spectrograph at the VLT, the kinematic properties of 65 of these galaxies, J033239.72-275154.7 for instance, have been derived. This galaxy lies at $z=0.41$, is classified as a merger from analysis lead by Neichel et al. (2008) and presents a big, young bar. This bar has a size of 6 kpc,  i.e. is quite big since 72\\% of barred galaxies from a  sample of 2000 galaxies from the SDSS have a size lower than this (Barazza et al. 2008). This bar has also an extremely blue color, consistent with a starburst, i.e. with ages well below few 100s of Myr. This contrasts with many bars in the local universe which are known to include relatively old and red stars (Sheth et al. 2003). This paper is organized as follows: in Section 2 we present the general morphological and kinetic properties of J033239.72-275154.7 and in Section 3 a short description of the numerical modelling and of the results obtained. We conclude in Section 4. ",
        "conclusions": "In this paper, we presented photometric and kinematic data of J033239.72-275154.7 and made a detailed comparison with an $N$-body simulation. A first interesting and important result is that among all the simulations that we have realized  with specific orbital parameters and progenitor mass ratios, the combination of parameters such as the star formation history, the VF and $\\sigma$ map put strong constraints on models and can remove the degeneracy of good candidates based on morphological criteria only.   We found that our fiducial model can roughly reproduce the general morphological shape and the total stellar mass of the object as well as the observed SF rate, VF and intensity of the $\\sigma$ map. This specifically includes the reproduction of the giant young bar, its location and shape, the relative projected location of the companion, the overall morphological shape of the galaxy, its low rotation and its off-center dynamical axis. The model predicts a slow-down of the velocity of the host galaxy due to exchanges of angular momentum with the companion  (see Fig. \\ref{simulated_galaxy}). However, using such idealized simulations, we were unable to reproduce many features, such as the blue arm or tidal tail on the upper right part, the morphological patterns of the two bright knots of the satellite remnant and the location of the $\\sigma$ peaks. Such discrepancies with observations can be easily understood as due to the huge parameter space to be covered by the simulations. These include the properties of the progenitors (halo:disk:bulge:gas mass ratios and relative extents, shape of their density profiles and kinematics), the geometry of the encounter and of the viewing, the gas properties, the SF, cooling and feedback modeling. On the other hand, the discrepancies with observations may also have astrophysical origin and/or be due to a more complex formation scenario. Several such possibilities -- such as multiple encounters or progenitors with specific properties -- come to mind, but would increase substantially the already too large free parameter space.  In  spite of these difficulties, it's encouraging to note that our simple numerical modelling is able to build a consistent picture of the formation of J033239.72-275154.7. This system is consistent with having been formed from a merger of two objects with a mass ratio 1:3. The simulation indicates that a bar is forming in the host galaxy after the first passage of the satellite where an important fraction of available gas is consumed in the induced burst giving a plausible explanation of  the observed blue colors of the bar and of the satellite remnant. Moreover, both VF and $\\sigma$ map derived from the simulation match with the observational values and thus support this scenario. In its later evolution, we found that J033239.72-275154.7 may become a S0 galaxy, as suggested by the results of our simulations. This is mainly explained  by the fact that the host galaxy experiences a violent relaxation and looses some angular momentum during the merger process due to the retrograde orbit of the satellite. Moreover, by loosing some angular momentum, some of the gas can sink toward the center of the galaxy where it can be converted into new stars and then accelerate the growth of the bulge.   To finish, it is interesting to point out that recent numerical studies have demonstrated that gas-rich mergers can produce remnant disks provided strong feedback processes and also if both stellar and gas component do not experience a significant angular momentum loss (Springel \\& Hernquist 2005; Robertson et al. 2006; Governato et al. 2007, Hopkins et al. 2008). We propose in this work an extreme example where i) the initial orbital configuration of the merger event do not permit to satisfy the latter criteria  and ii) the progenitors of the system (in our numerical model) may not be sufficiently gas rich to reform a significant disk according those past studies. However, the present results are  consistent with  the previous ones in the emphasis on the fundamental role played by the  last major event in building the Hubble sequence (see e.g. Hammer et al., 2005 and 2007). It would thus be interesting to determine the frequency of each orbital configuration during a merger event and to compare with the galaxy population in the local universe to confirm this galaxy formation scenario."
    },
    "0812/0812.3769_arXiv.txt": {
        "abstract": "The eclipsing and double-lined spectroscopic binary system V453\\,Cygni consists of two early B-type stars, one of which is nearing the terminal age main sequence and one which is roughly halfway through its main sequence lifetime. Accurate measurements of the masses and radii of the two stars are available, which makes a detailed abundance analysis both more interesting and more precise than for isolated stars. We have reconstructed the spectra of the individual components of V453\\,Cyg from the observed composite spectra using the technique of spectral disentangling. From these disentangled spectra we have obtained improved effective temperature measurements of $27\\,900 \\pm 400$\\,K and $26\\,200 \\pm 500$\\,K, for the primary and secondary stars respectively, by fitting non-LTE theoretical line profiles to the hydrogen Balmer lines. Armed with these high-precision effective temperatures and the accurately known surface gravities of the stars we have obtained the abundances of helium and metallic elements. A detailed abundance analysis of the primary star shows a normal (solar) helium abundance if the microturbulence velocity derived from metallic lines is used. The elemental abundances show no indication that CNO-processed material is present in the photosphere of this high-mass terminal age main sequence star. The elemental abundances of the secondary star were derived by differential study against a template spectrum of a star with similar characteristics. Both the primary and secondary components display elemental abundances which are in the ranges observed in the Galactic OB stars. ",
        "introduction": "High-mass stars are major producers of ultraviolet light and byproducts of thermonuclear fusion, so are important objects in the study of star formation, and galactic chemical and kinetic evolution (e.g.\\ Samland 1998). The large luminosities of high-mass stars means that they are important distance indicators in both a galactic and cosmological context (e.g.\\ Guinan et al.\\ 1998; Ribas et al.\\ 2005; Southworth et al.\\ 2005, 2007).  In the last decade new theoretical models have been constructed which incorporate a number of physical processes important to the structure and evolution of high-mass stars, including convective core overshooting, semiconvection, rotational mixing and mass loss. Overshooting and rotation can both have a large effect on the predicted lifetimes and luminosities of high-mass stars (c.f.\\ review by Maeder \\& Meynet 2000). However, empirical constraints on these processes remain hard to come by, despite a steady improvement in observational techniques and capabilities (see Hilditch 2004).  Detached eclipsing binaries (dEBs) are vital objects for obtaining observational constraints on the structure and evolution of high-mass stars, since they are the primary source of directly measured stellar properties (Andersen 1991). Unfortunately, accurate (2\\% or better) physical properties are available for only eight high-mass dEBs\\footnote{An up-to-date compilation of the properties of well-studied dEBs is available at {\\tt http://www.astro.keele.ac.uk/$\\sim$jkt}}, and most of these objects have no observational constraints on their chemical composition. Chemical abundances are difficult to determine from the spectra of high-mass dEBs for several reasons. Firstly, they tend to display only a small number of spectral lines. Secondly, the often high rotational velocities of these objects means their spectral lines are wide and shallow so high signal-to-noise ratio (S/N) spectra are needed to obtain useful results. Thirdly, in dEBs the spectral lines from one component interfere with the line profiles of the other component (`line blending').  In a seminal work, Simon \\& Sturm (1994) introduced the technique of {\\em spectral disentangling} ({\\sc spd}), by which {\\em individual} spectra of the component stars of a double-lined spectroscopic binary system can be deduced from a set of composite spectra observed over a range of orbital phases. {\\sc spd} can be used to measure spectroscopic orbits which are not affected by the line blending which afflicts other methods, such as cross-correlation (see Southworth \\& Clausen 2007). The resulting disentangled spectra also have a much higher S/N than the original observations, making them very useful for chemical abundance analyses. As a bonus, the strong degeneracy between effective temperature (\\Teff) and surface gravity (\\logg) is not a problem for dEBs because surface gravities can be measured to high accuracy (0.01\\,dex or better) for both components. A detailed study of these possibilities is given by Pavlovski \\& Hensberge (2005; hereafter PH05).   \\subsection{The eclipsing binary system V453\\,Cyg}   \\label{sec:intro:v453}  V453\\,Cyg belongs to rather sparse group of high-mass close binary systems for which accurate physical properties have been measured for both components. In the most recent study of V453\\,Cyg, Southworth, Maxted \\& Smalley (2004b; hereafter SMS04) analysed new spectroscopic observations and previously published light curves, obtaining masses of $M_{\\rm A} = 14.36 \\pm 0.20$\\Msun\\ and $M_{\\rm B} = 11.11 \\pm 0.13$\\Msun, and radii of $R_{\\rm A} = 8.55 \\pm 0.06$\\Rsun\\ and $R_{\\rm B} = 5.49 \\pm 0.06$\\Rsun. The high precision of the radius measurements was helped by the geometry of the V453\\,Cyg system, which exhibits deep and total eclipses. SMS04 derived effective temperatures of ${\\Teff}_{\\rm A} = 26\\,500 \\pm 800$\\,K and ${\\Teff}_{\\rm B} = 25\\,300 \\pm 600$\\,K from the equivalent widths of several \\ion{He}{I} and \\ion{He}{II} lines, using an analysis based on local thermodynamic equilibrium (LTE) synthetic spectra but with a correction for non-LTE effects.  V453\\,Cyg is known to display the phenomenon of apsidal motion, which can be used as a probe of the internal structure of the components of an eccentric binary system (e.g.\\ Claret \\& Gim\\'enez 1993). SMS04 found an apsidal period of $U = 66.4 \\pm 1.8$\\,yr, which gives the structural parameter\\footnote{Note that the $\\log k_2$ value in the abstract of SMS04 is incorrect: the correct value is given in SMS04 (their Section\\,7.2) and here.} $\\log k_2 = -2.254 \\pm 0.024$. The masses, radii, \\Teff s and $\\log k_2$ of V453\\,Cyg match theoretical predictions for an age of $10.0 \\pm 0.2$\\,Myr and an approximately solar metal abundance. The primary component (the hotter and more massive star) is reaching the terminal age of its main sequence lifetime (TAMS) whilst the secondary star is roughly halfway through its MS life. The significant difference between the properties of the two stars means that V453\\,Cyg is a very useful system for testing evolutionary models.  In this work we analyse the spectroscopy of V453\\,Cyg obtained by SMS04 and by Simon \\& Sturm (1994) in order to obtain new \\Teff s and accurate chemical abundances for the stars. V453\\,Cyg was targeted by SMS04 as part of a program to study dEBs which are members of open clusters (see Southworth, Maxted \\& Smalley 2004a; Southworth et al.\\ 2004c). It is a probable member of NGC\\,6871, so our results can be compared to published analyses of single stars which are also cluster members. ",
        "conclusions": "\\label{sec:discussion}  Due to the serious problems of line blending and high rotational velocities, abundances have been measured for only a few high-mass stars in double-lined binary systems. Prior to the invention of {\\sc spd} (Simon \\& Sturm 1994) some helium abundances in high-mass binaries were determined from equivalent widths measured using Gaussian fitting (Lyubimkov 1998; Pavlovski 2004 and references therein). The use of {\\sc spd} has allowed helium abundances to be measured for four close binaries: DH\\,Cep (Sturm \\& Simon 1994), Y\\,Cyg (Simon et al.\\ 1994), V578\\,Mon (Hensberge et al.\\ 2000), and DW\\,Car (Southworth \\& Clausen 2007). Preliminary results have also been announced for V380\\,Cyg (Pavlovski et al.\\ 2005), and V478\\,Cyg and CW\\,Cep (Pavlovski \\& Tamajo 2007). Finally, Freyhammer et al.\\ (2001) have derived helium abundances for the components of CPD\\,$-$59$^\\circ$2628 by fitting theoretical spectra.  Lyubimkov et al.\\ (2004) found an enrichment of helium during the MS evolution of their sample of stars, although many theoretical evolutionary models (e.g.\\ Schaller et al.\\ 1992) predict no surface enrichment of helium or other CNO-cycle elements. However, more recent theoretical studies (Maeder \\& Meynet 2000; Heger \\& Langer 2000) have found that rotationally induced mixing can affect the chemical composition of the surface layer of high-mass stars even whilst they are on the MS. Due to CNO processing in the stellar core, He and N should be enriched and C and O depleted. The strength of these effects depends on initial rotational velocity, so observational constraints can be used to both calibrate theoretical models and check their predictions (e.g.\\ Venn et al.\\ 2000).  The observations of 100 Galactic B stars by Lyubimkov et al.\\ (2000, 2004) have confirmed that helium overabundance is correlated with stellar age expressed as a fraction of its MS lifetime. Large spectroscopic surveys of OB stars in our Galaxy and the Large and Small Magellanic Clouds (Hunter et al.\\ 2008) have further confirmed this effect at a range of metallicities. Unfortunately, investigations using single stars suffer from the low accuracy with which their masses, radii and ages can be measured. Further work must therefore concentrate on binary stars, where all these properties can be measured to within a few percent (e.g.\\ SMS04).  The subsample of high-mass stars in the sample of OB stars studied by Lyubimkov et al.\\ (2004) show helium enrichment during their MS lifetimes, which has been found previously for close binaries (see Lyubimkov 1998 and references therein). In two recent studies in which large samples of B stars were analysed, progressively increasing helium enrichment from the ZAMS to the TAMS was found (Lyubimkov et al.\\ 2004, Huang \\& Gies 2006). Both studies saw this effect over their whole mass ranges, with large enrichments for high-mass stars. Moreover, Huang \\& Gies noted that the effect is large among the faster rotators (see their fig.\\ 12). Lyubimkov et al.\\ derived \\micro\\ both from helium lines and from \\ion{O}{II}, while Huang \\& Gies (working with only a limited spectral range around H$\\gamma$) accepted $\\micro\\ = 2$\\kms\\ for all their sample after considering the discussion of \\micro\\ in Lyubimkov et al.\\ (2004). Therefore, the agreement achieved by these studies may depend on their use of a common \\micro.  So far, in the sample of OB binaries studied by {\\sc spd} and cited above, no MS star has been found to have a helium enrichment. To this list we add now both components of V453\\,Cyg. The ages of the stars on the list lie between 0 and 0.8 of their MS lifetimes, with V453\\,Cyg\\,A contributing the upper limit. Possible explanations for the non-detection of any changes in elemental abundances are moderate rotational velocity and tidal effects between the components. In this context it is interesting that detailed abundance analyses of about a dozen $\\beta$ Cephei stars (Morel et al.\\ 2006) has given an average helium abundance of $\\epsilon({\\rm He}) = 0.087 \\pm 0.012$, with a range from $0.075$ to $0.100$. These objects are slowly-rotating stars in advanced phases of their MS lifetimes and do not show enhanced helium abundances. The \\Teff\\ and \\logg\\ of V453\\,Cyg\\,A fall within the range of their sample, and this star also does not show enhanced helium. These two results are in excellent agreement and corroborate previous finding of Herrero et al.\\ (1992), who found a very clear dependence of helium enrichment on rotational velocity, but not on fractional lifetime on the MS, for O-type stars.  In Section\\,\\ref{sec:he} we determined the He abundance in V453\\,Cyg\\,A by two approaches: (i) simultaneous line-profile fitting allowing \\micro\\ to be a free parameter; (ii) fixing \\micro\\ to the value obtained from \\ion{O}{II} lines. When using a fixed \\micro\\ we found a normal helium abundance ($\\epsilon({\\rm He)} = 0.086 \\pm 0.012$), but when \\micro\\ was a free parameter we found a helium overabundance ($\\epsilon({\\rm He}) = 0.123 \\pm 0.017$) and also erratic metal abundances. We therefore reject the helium enrichment scenario and caution that it is crucial to carefully consider the value of \\micro\\ in similar studies."
    },
    "0812/0812.0306_arXiv.txt": {
        "abstract": "We present a hadronic model of activity for Galactic \\gr-loud binaries, in which the multi-TeV neutrino flux from the source can be much higher and/or harder than the detected TeV \\gr\\ flux. This is related to the fact that  most neutrinos are produced in pp interactions close to the bright massive star, in a region optically thick for the TeV \\gr s. Considering the specific example of \\lsi, we derive upper bounds for neutrino fluxes from various proton injection spectra compatible with the observed multi-wavelength spectrum. At this upper level of neutrino emission, we demonstrate that ICECUBE will not only detect this source at $5\\sigma$ C.L. after one year of operation, but, after 3 years of exposure, will also collect a sample marginally sufficient to constrain the spectral characteristics of the neutrino signal, directly related to the underlying source acceleration mechanisms. ",
        "introduction": "The recently discovered \"\\gr -loud\" binaries (GRLB) form a new sub-class of Galactic binary star systems which emit GeV-TeV \\gr s  \\cite{ls5039,lsi,psrb, HESSJ0632,cygx1}. These are high mass X-ray binaries (HMXRB) composed of a compact object (a black hole or a neutron star) orbiting a massive star. The detection of \\gr s with energies up to 10~TeV from these systems shows that certain HMXRBs host powerful particle accelerators producing electrons and/or protons with energies above 10~TeV.  At the moment it is not clear which physical process leads to the very high energy (VHE) particle acceleration and whether the VHE particle acceleration is a generic feature of the HMXRBs or the result of specific physical conditions in a restricted HMXRB sub-class. It is possible that  particle acceleration is taking place in a large number of X-ray binaries, but the \\gr s can be detected only from the objects with preferred orientations w.r.t. the line of sight (e.g. with jets pointing toward the observer) \\cite{mirabel}. Alternatively, it may be that the conventional accretion-powered binaries are not capable of accelerating particles and that the \\gr\\ activity of a binary is powered by a different mechanism (see e.g. \\cite{maraschi,tavani97}).  The theoretical models of \\gr\\ activity of HMXRBs \\cite{maraschi,tavani97,bosch-ramon,dubus,bednarek,khangulyan,romero,chernyakova06a,chernyakova06b} all assume \\gr\\ emission  from the interaction of a relativistic outflow from the compact object (jet from a black hole, or wide angle wind from a pulsar) with the wind and radiation emitted by the companion massive star. The basic properties of the outflows, such as the composition (e$^+$e$^-$ pairs or electron-nuclei plasma), anisotropy (a collimated jet or a wide angle outflow) etc. are as yet poorly constrained by the data.  In one of the three persistent (periodic) \\gr-loud binaries, PSR B1259-63, the relativistic outflow is known to be produced by a young pulsar.  In principle, a similar mechanism could power the activity of other sources (except for Cyg X-1), although direct proof of the presence of the young pulsar in these systems is not possible: the radio emission from the compact object is free-free absorbed. Recent puzzling detection of a short soft \\gr\\ flare from one of the \\gr-loud binaries, LSI +61 303 \\cite{lsi_burst}, may indicate the presence of a neutron star in the system.  The outflows from LSI +61 303 and LS 5039, extending to the distances $\\sim 10^{14}$~cm (far beyond the binary orbit) are revealed by the radio observations \\cite{dhawan,ribo08}. The nature of the outflows (a collimated jet or a wide-angle outflow) is still debated.  With the exception of Cyg X-1, all the known GRLBs have similar spectral energy distributions (SED), peaking in the MeV-GeV energy band. The physical mechanism producing the MeV-GeV bump in the spectra is, however, not clear. The \\gr\\ emission from GRLBs is supposed to come from internal and/or external shocks formed in the relativistic outflow either as a result of the development of intrinsic instabilities or through interactions with the stellar wind of the massive star. The MeV-GeV emission can be the synchrotron emission from electrons with energies much above TeV \\cite{tavani97,bosch-ramon}, produced locally at the shock, or else, be produced via inverse Compton scattering of the UV thermal emission from the massive star by electrons of the energies $E\\sim 10$~MeV \\cite{maraschi,chernyakova06a,chernyakova06b}.  The available multi-wavelength data do not constrain the composition of the relativistic wind from the compact object. On the one hand, the multi-TeV or 10~MeV electrons, responsible for the production of the MeV-GeV bump in the SED, could be injected into the shock region from the e$^+$e$^-$ pairs loaded wind. On the other, these electrons could be secondary particles produced in e.g. proton-proton collisions, if the relativistic wind is proton-loaded. The only direct way to test if relativistic protons are present in the \\gr\\ emission region would be the detection of multi-TeV neutrinos.  The possibility of detection of neutrinos from GRLBs by a km$^3$-class neutrino telescopes was considered in several references \\cite{aharonian07n,aharonian06n,orellana,halzen,kappes,christiansen}. If the sources are assumed to be transparent to the TeV \\gr s, the estimation of the neutrino flux is straightworward, given a known source \\gr\\ flux and spectrum. The assumption of source transparency was adopted e.g. in the estimates of the number of detected neutrinos as in Ref. \\cite{kappes}. However, the TeV \\gr\\ flux from the GRLBs can be significantly attenuated by the pair production on the UV photon field of the massive star \\cite{dubus,bednarek}. In addition, if the \\gr s and neutrinos are produced close to the compact object, the \\gr\\ flux can be further suppressed by pair production on the soft photons emitted by the accretion flow \\cite{aharonian06n}. The derivation of the estimate of the neutrino flux and spectral characteristics based on the observed TeV \\gr\\ emission is inconclusive due to the uncertain attenuation of the \\gr\\ flux in the TeV band (see \\cite{halzen} for s specific discussion of \\lsi).  In the absence of a direct relation between the characteristics of the observed TeV \\gr\\ and the neutrino emission from a GRLB, the only way to constrain possible neutrino signals from the source is via the detailed modelling of the broad-band spectrum of the source within the hadronic model of activity. The idea is that the pp interactions, which result in the production of neutrinos, also result in the production of e$^+$e$^-$ pairs and the subsequent release of their energy via synchrotron, inverse Compton and Bremsstrahlung emission. The electromagnetic emission from the secondary e$^+$e$^-$ pairs is readily detectable. The total power released in the pp interactions determines the overall luminosity of the emission from the secondary pairs. The known electromagnetic luminosity and broad-band spectral characteristics of the source can be used to constrain the power released in pp interactions as well as the spectrum of the primary high energy protons.  In the following we develop the hadronic model of activity of GRLBs and derive the constraints on the spectrum and overall luminosity of neutrino emission from the analysis of the broad-band spectral characteristics of GRLBs.  Although the following discussion is generically applicable for the GRLBs as a class, we concentrate on the particular example of the \\lsi\\ system, because it is the only known persistent GRLB in the Northern hemisphere, available for observations with the ICECUBE neutrino telescope \\cite{icecube}. The existing observational data are found to be consistent with a possible very strong neutrino emission from \\lsi\\ (with a flux at the level of $10^{-10}$~erg/cm$^2$s), close to the best reported AMANDA upper limit, $\\Phi(1.6\\mbox{ TeV}<E_\\nu<2.5\\mbox{ PeV})\\le 1.26\\times 10^{-10}\\mbox{ cm}^{-1}\\mbox{ s}^{-1}\\mbox{ TeV}^{-1}$) \\cite{amanda} (a more recent reference \\cite{a2-7yrs} combining more data reported a slightly degraded upper limit for \\lsi\\ over approximately the same range of energy, while the sensitivity had been improved by roughly a factor two; this however translates a non significant but large background fluctuation at the \\lsi\\ source location). Moreover, the assumption of an almost arbitrarily hard neutrino spectrum (e.g. a powerlaw $dN_\\nu/dE\\sim E^{-\\Gamma_\\nu}$ with index $\\Gamma_\\nu\\sim 0$) is not ruled out by the available multi-wavelength observational data. This means that, in principle, the soure could be readily detectable  with ICECUBE.  The plan of the paper is as follows: In section \\ref{sec:model} we describe the general features of the hadronic model of activity of GRLBs. In particular, we stress that when the massive companion star is a Be type star, the presence of a dense equatorial decretion disk around the massive star can boost the pp interaction rate. At the same time, the \\gr\\ emission from the pp interactions in the disk would be strongly suppressed, because of the large density of the soft photon background in the direct vicinity of the star. In section \\ref{sec:numeric} we perform the detailed numerical modeling of the broad-band spectra of GRLBs in the hadronic model and show, on the particular example of \\lsi, that the model provides a suitable explanation of the typical shape of the GRLB SEDs. Satisfactory fits to the observed SEDs can be achieved under quite arbitrary assumptions about the shape of the initial high energy proton spectrum. In particular, the initial spectra as different as an $E^{-2}$ power law with a high energy cut-off and a monochromatic proton spectrum peaked at $\\sim$PeV energy could both explain the observed X-ray-to-\\gr\\ spectrum of the source. Finally, in section \\ref{sec:icecube} we work out the predictions for the detection of neutrinos from \\lsi\\ with ICECUBE. We show that in an optimistic scenario, when the anisotropy of the neutrino emission of the source does not result in a suppression of the neutrino flux in the direction of the Earth, \\lsi\\ should be detectable within a year of exposure. Nevertheless, the spectral characteristics of the emission can only be marginally delineated after three years of exposure, given the wash out of the original neutrino spectrum from the measurement of the muon spectrum. ",
        "conclusions": "We have estimated the neutrino flux from GRLBs expected within the hadronic model of activity of these sources. Within such a model, the measured spectral characteristics of \\gr\\ emission from the source in the TeV energy band are not directly related to the spectral characteristics of the neutrino emission, because of absorption of the TeV \\gr s on the thermal photon background produced by the massive star in the system. The uncertainty of the calculation of the attenuation of the TeV \\gr\\ flux introduces a large uncertainty to the estimate of the neutrino flux based on the measured TeV \\gr\\ flux.  Taking this uncertainty into account, we have adopted a different approach for the estimate of the neutrino flux from a GRLB. Namely, we have noted that the energy output of proton-proton interactions, and hence the neutrino flux, can be constrained by the broad-band spectrum of the source. The idea is that the ${\\rm e}^+{\\rm e}^-$ pairs, produced in the decays of charged pions, release their energy in the form of the synchrotron, IC and Bremsstrahlung emission, which contributes to the observed broad-band SED of the source. The requirement that the power of electromagnetic emission from the secondary ${\\rm e}^+{\\rm e}^-$ pairs in different energy bands is not higher than the observed source luminosity imposes a constraint on the neutrino luminosity of the GRLB.  We have worked out the numerical model of the broad-band emission from a GRLB, assuming that pp interactions take place in the vicinity of the bright massive star, in the innermost and densest part of the stellar wind. We have shown that the interactions of high energy protons (with energies reaching $\\sim$PeV) in the dense stellar wind can explain the observed broad-band emission, with the MeV-GeV \"bump\" of the SED being due to the synchrotron emission from the secondary ${\\rm e}^+{\\rm e}^-$ pairs produced either in the decay of charged pions or resulting from the absorption of 10~GeV-TeV \\gr s on the thermal photons from the massive star (see Figs. 1-4).  Although the observed bolometric luminosity of the source (i.e. the height of the MeV-GeV bump of the SED) constrains the overall neutrino luminosity, the shape of the neutrino spectrum and the overall normalization of the neutrino flux are only mildly constrained by the multi-wavelength data. The problem is that the properties of the broad-band emission from the secondary ${\\rm e}^+{\\rm e}^-$ pairs are mostly determined by the effects of the radiative cooling and the escape of the pairs from the source, rather than by the initial injection spectrum of the pairs. This means that quite different injection spectra of the primary high energy protons (and hence of the neutrinos) can result in approximately the same broad-band SEDs, as is clear from Figs. 1-4. In addition, the anisotropy pattern of neutrino emission is expected to differ from that of the synchrotron emission from the ${\\rm e}^+{\\rm e}^-$ pairs. This results in an uncertainty  of the anisotropy-related suppression of the neutrino flux compared to the MeV-GeV synchrotron flux.  Taking into account these uncertainties of the neutrino emission spectrum, we have estimated the expected number of neutrinos which will be detected by ICECUBE, assuming that the neutrino flux saturates the upper bound imposed by the observed \\gr\\ flux in the MeV-GeV energy band. Considering the particular example of \\lsi, we have found that if the spectrum of high energy protons in the source extends to the PeV energies, the source would be readily detectable within roughly one year of exposure with ICECUBE.  We have also explored the potential of the full ICECUBE detector for the measurement of the spectral characteristics of the neutrino signal from \\lsi. We find that in the case when the neutrino flux is at the level of the upper bound imposed by the observed MeV-GeV \\gr\\ flux, an exposure time longer than 3 years will be required to reliably constrain the spectral index of the primary high energy proton spectrum  via observations of neutrino signal in ICECUBE. The mere fact of a possible spectral characterization of neutrino sources within the lifetime of ICECUBE represents however a most exciting prospect for the future of mutimessenger astronomy with the possible emergence of a new field, the astrophysics of neutrinos."
    },
    "0812/0812.0901_arXiv.txt": {
        "abstract": "In supernova remnants, the nonlinear amplification of magnetic fields upstream of collisionless shocks is essential for the acceleration of cosmic rays to the energy of the \\lq\\lq knee\\rq\\rq\\ at $10^{15.5}\\,$eV.  A nonresonant instability driven by the cosmic ray current is thought to be responsible for this effect. We perform two-dimensional, particle-in-cell simulations of this instability.  We observe an initial growth of circularly polarized non-propagating magnetic waves as predicted in linear theory. It is demonstrated that in some cases the magnetic energy density in the growing waves, can grow to at least 10 times its initial value. We find no evidence of competing modes, nor of significant modification by thermal effects. At late times we observe saturation of the instability in the simulation, but the mechanism responsible is an artefact of the periodic boundary conditions and has no counterpart in the supernova-shock scenario. ",
        "introduction": "Diffusive shock acceleration (DSA) at supernova remnant shocks is widely considered to be the primary source of galactic cosmic rays \\citep[for a recent review see][]{Hillas05}. A crucial aspect of DSA is the self-excitation of hydromagnetic waves due to currents produced by the streaming energetic ions. These provide the pitch-angle scattering necessary to allow spatial diffusion of the relativistic particles. The maximum acceleration rate that can be obtained corresponds to the smallest possible spatial diffusion coefficient. In this \\lq Bohm limit\\rq\\ the maximum particle energy is determined either by the age of the remnant \\citep{lag83}, or by its geometry \\citep{Berezhko}. Even adopting optimistic values for the upstream parameters of a supernova remnant shock, particles are limited to energies below the knee of the cosmic-ray spectrum at $\\sim 10^{15.5}$~eV. However, nonlinear amplification of the upstream magnetic field due to cosmic-ray currents may provide a possible solution to this problem \\citep{Quest, luc00}.  There is increasing observational evidence that the magnetic fields immediately downstream of several young supernova remnant shocks are much stronger than would be obtained by compressing the ambient interstellar field at an MHD shock front \\citep{Vink,ber03,bam05,uch07}. These observations give additional motivation to the development of theoretical models of nonlinear field generation by both plasma instabilities in the precursor~\\citep{luc00,bel04} and fluid-type instabilities in the downstream plasma~\\citep{gia07}. These two mechanisms can, in principle, operate simultaneously \\citep{zir08}. However, to reach the highest energies, particles must be speedily returned to the shock from excursions into both the upstream and downstream plasmas. The generation of an amplified magnetic field in the precursor that is subsequently advected into the downstream region is sufficient to ensure this, whereas amplification of only the downstream field is not.  Using a hybrid kinetic-MHD analysis, it was demonstrated by \\cite{bel04} that efficient cosmic-ray acceleration can, in the linear regime, drive a short wavelength, almost aperiodic, instability with wave-vector parallel to the ambient field. The growth of this mode can be considerably more rapid than that of the resonantly driven modes previously considered \\citep{mcken,acht83}. The instability is driven by the reaction of the thermal plasma, as it attempts to compensate the current produced by the streaming cosmic rays. Numerical investigations of this instability have been performed using both MHD \\citep{bel04, bel05, rev08, zpv08} and Particle-In-Cell (PIC) simulations \\citep{nie08, riq08}. While the results of MHD simulations show that the perturbed magnetic field $\\delta B$ becomes much larger than the initial uniform magnetic field $B_0$, the PIC simulations of \\cite{nie08} did not observe the parallel mode predicted in the linear analysis, and only moderate amplification of the magnetic field was observed. On the other hand, the recent results reported by \\citet{riq08} do identify the parallel mode and propose an explanation of its saturation in the non-linear stage.  In this paper we report new PIC simulations of this instability. We extend the range of parameters beyond that investigated by \\cite{nie08}, and select values that more closely represent the environment in an SNR precursor. Although both 2D and 3D simulations in a similar parameter range have been performed \\citep{riq08}, we restrict our work to 2D. Our results --- obtained independently and with some technical differences in the numerical treatment --- can be regarded as complementary, and our physical interpretation, at least of the early stages of the non-linear development, are congruent. In section~2, we recall the linear kinetic analysis of the nonresonant instability, in order to elucidate the conditions that must be satisfied in the PIC simulations. The details of our simulation parameters are described in section~3. We present the results of the simulations in section~4 followed by a discussion of the saturation and its implications in section~5. We conclude with a summary of the results. ",
        "conclusions": "\\label{discussion}  We find that the nonresonant cosmic-ray driven instability develops in all cases where the background plasma is magnetized. From Figure \\ref{fig2}(a) and \\ref{fig2}(b), it can be seen that there are considerable differences between the results of run~A and run~B although the observed wavelength is similar to $\\lambda_{\\rm NR}$. The most unstable mode in run~B is the Weibel type instability, occurring between the counterstreaming electrons and ions. Neglecting thermal effects and the magnetic field, the growth rate $\\gamma_{\\rm WI}^{\\rm max}$ and the wavelength $\\lambda_{\\rm WI}$ of the Weibel instability are \\begin{equation} \\gamma_{\\rm WI}^{\\rm max}=\\frac{n_{\\rm cr}}{n_{\\rm e}}\\frac{v_{\\rm s}}{c}\\omega_{\\rm pi}, \\ \\ \\lambda_{\\rm WI}=2\\pi \\frac{c}{\\omega_{\\rm pe}}, \\end{equation} where the direction of the wave vector is perpendicular to the drift direction. Comparing the two theoretical growth rates and the wavelengths we see that \\begin{equation} \\frac{\\gamma_{\\rm WI}^{\\rm max}}{\\gamma_{\\rm NR}^{\\rm max}} = \\frac{n_{\\rm e}}{n_{\\rm i}} \\simeq 1, \\ \\ \\frac{\\lambda_{\\rm WI}}{\\lambda_{NR}}=\\frac{1}{2}\\frac{n_{\\rm cr}}{n_{\\rm i}}\\frac{V_{\\rm d,cr}}{c}\\frac{\\omega_{\\rm pe}}{\\Omega_{ce}} \\simeq 1. \\end{equation} For the choice of parameters used in run~B,  the wavelengths of the different instabilities are quite similar, even though the mechanism is quite different. This is consistent with the results of \\citet{nie08}, and we also find the fastest growing mode to be slightly oblique. Hence we may be observing a mixed-mode between the Weibel and the nonresonant current driven instability or some other type of mixed-mode instability similar to those discussed in \\cite{Bret}.  In the late stages of the simulations we see the emergence of a non-thermal power law tail in the ion distribution. The use of a constant external cosmic-ray current prevents the development of the Buneman instability, a source of ion heating in the linear development. This dramatically reduces the heating in the initial stages, that was previously observed in \\cite{nie08}. Figure 6 shows the final electron and ion energy distribution for run~A, C and D. The distributions are essentially isotropic, and since the bulk drift velocity is small in comparision with the thermal velocities, we can safely calculate the spectrum in the box frame. The electron energy distribution can be well fitted with a Maxwellian distribution with temperature $10$ keV for run~A and $40$ keV for run~C and D. A non-thermal tail appears in the distribution of the ions with a very soft power-law index of $\\sim 7.3$ for Run-A, but a relatively hard $\\sim2$ for run~C and D. The mechanism responsible for the formation of this power-law distribution is not clear, and although our simulations have a relatively large number of particles per cell, we cannot rule out a numerical artefact. As regions of oppositely polarized magnetic field collide each other, an anti-parallel configuration of the magnetic fieldlines appears.  It is, therefore, possible that collisionless magnetic dissipation may also play an important role in plasma heating mechanism and the saturation of the field growth as well as the reduction in the drift velocity between the electrons and ions. This is an interesting process in its own right, and further investigation is warranted. However, in the current context, it is not clear what role these mechanisms play, since the late-stage nonlinear behavior has no counterpart in the supernova-shock scenario.   In agreement with \\cite{nie08} and \\cite{riq08}, we observe in our simulations that the magnetic field growth ceases when the plasma bulk velocity becomes comparable to the cosmic-ray drift velocity. Although we have fixed the cosmic ray current in our simulations, the observed saturation is clearly due to the reduction in the relative speeds of plasma and cosmic rays. Indeed, if the cosmic rays are isotropic in the plasma frame, no streaming instability can operate. The simulations are performed in a box with periodic boundary conditions in space. The plasma speed is initially zero, and the cosmic ray streaming speed is maintained constant in time and space. It represents the speed with which the shock front approaches the simulation box. Within this picture, the instability can saturate only when the entire plasma moves with the drift speed of the cosmic rays.  In the precursor of a shock front, conditions are different. There, the interaction between plasma and cosmic rays defines the diffusion length scale of the cosmic rays. This is the scale on which the cosmic ray current, or, equivalently, pressure, decays with distance ahead of the shock front. In the case of efficient acceleration, it is also the length scale on which the precursor plasma is compressed. Simulations can model this situation provided the box size is small compared to the diffusion length scale, and the time scales are short compared to the time on which the cosmic ray current and plasma density change because of the approaching shock front. However, this restriction means that the simulations no longer model a precursor when they are followed until the plasma speed approaches the shock speed (which equals the cosmic-ray drift speed).  This can be seen explicitly from the equation of conservation of mass in a stationary precursor, $\\rho(x) u(x) = \\mbox{constant}$, which dictates that any change in the flow velocity must be associated with a change in the density of the fluid. In the present context, the simulation box is in the upstream plasma frame, approaching the shock with velocity $v_{\\rm s}$. Defining the density and bulk velocity inside the box to be $\\rho_b$ and $u_{\\rm b}$ respectively, we see that \\begin{equation} \\rho_{\\rm b}= \\frac{v_{\\rm s}\\rho_0}{v_{\\rm s}-u_{\\rm b}} \\end{equation} where $\\rho_0$ is the initial density of the thermal plasma. Thus, as the flow speed inside the box increases, the density should also increase. Since the total density in the box is fixed throughout the simulation, the observed growth rate is only accurate provided the bulk velocity inside the box is small compared to the shock velocity, or, equivalently, the streaming speed of the cosmic rays, i.e., $u_{\\rm b} \\ll v_{\\rm s}$. This corresponds to $\\tau_{\\rm grow}< 20$ for run A and $\\tau_{\\rm grow}< 23$ for runs C and D. The amplification of the energy density in the total field in run~A is not yet substantial, but in runs~B and C roughly 10 times the initial energy density is reached. The final saturation levels are much higher, but our simulations do not necessarily imply that these can be reached in a precursor scenario, and a PIC simulation of this case is currently out of reach."
    },
    "0812/0812.0130_arXiv.txt": {
        "abstract": "Since its discovery in 1990, UW CrB (also known as MS1603+2600) has remained a peculiar source without firm classification. Our current understanding is that it is an Accretion Disc Corona (ADC) low mass X-ray binary. In this paper we present results from our photometric campaign dedicated to studying the changing morphology of the optical light curves. We find that the optical light curves show remarkable evidence for strongly evolving light curve shapes. In addition we find that these changes show a modulation at a period of $\\sim$ 5 days. We interpret these changes as either due to strong periodic accretion disc warping or other geometrical changes due to disc precession at a period of 5 days. Finally, we have detected 11 new optical bursts, the phase distribution of which supports the idea of a vertically extended asymmetric accretion disc. ",
        "introduction": "UW CrB (also known as MS1603+2600) was discovered in X-rays during the \\textit{Einstein} medium sensitivity survey (Morris et al. 1990). Morris et al. (1990) also identified the optical counterpart of the X-ray source. A star with blue, accretion disc-like spectrum with Balmer emission lines was found close to the X-ray position. Subsequent optical photometry revealed a period of 111 minutes. Morris et al. (1990) also noted that  the shape of the optical light curve seemed to change from one night to the other. The system is located at a high galactic latitude of 47$^o$. This, together with its low X-ray flux of 1.14$\\times$10$^{-12}$ erg/s (Morris et al. 1990), suggests that either the source is underluminous in X-rays for a low mass X-ray binary (LMXB) or that it is actually located in the galactic halo.  Earlier models for UW CrB  have involved both LMXB and magnetic CV scenarios (Morris et al. 1990, Ergma \\& Vilhu 1993). Based on evolutionary calculations Ergma \\& Vilhu prefer the LMXB option. Hakala et al. (1998) compared the ROSAT spectra of UW CrB with soft X-ray spectra of various classes of interacting binaries. Their conclusion was that the only class where the spectral fits seemed to match those of the source was that of the soft X-ray transients (SXT) in quiescence. This could also explain the relatively low \\textit{F$_{X}$/F$_{opt}$} ratio the system has. The main complication for this explanation arises from the fact that, according to our knowledge, the optical light curves of other SXT's do not vary like they do in UW CrB. The same is true for the magnetic CV hypothesis (AM Herculis type). In addition, Hakala et al. (1998) report a negative result from their search for circular polarisation.  Hakala et al. (1998) also published the results of further optical photometry together with the ROSAT observations. They confirmed the extreme variability seen in the optical light curve pulse shapes. Even if there are dramatic changes  in the optical light curve shape, the overall optical flux level does NOT seem to vary much. The analysis of ROSAT data proved that the soft X-ray spectrum of UW CrB can be modelled with a single blackbody component with a temperature of  0.24 keV. The ROSAT data also showed that there is no significant interstellar absorption in the direction of the source.  Mukai et al. (2001) presented ASCA observations of UW CrB. They detect a single type I X-ray burst, which rules out a black hole as the primary component in the system (which was suggested by Hakala et al. 1998). These bursts have also been seen in the optical (Muhli et al. 2004, Hynes et al. 2004.). Mukai et al. (2001) propose that UW CrB is a short period X-ray dipper, like 4U1916-05. They also note that the system would probably in this case have to reside in the outer galactic halo, and could have formed in the globular cluster Palomar 14, which is located 11$^o$ from UW CrB at a distance of 73.8 kpc. The {\\it Chandra} observations of UW CrB by Jonker et al.(2003) modelled the X-ray spectrum with a simple powerlaw of spectral index $\\alpha$= 2.0. They conclude that the source is an accretion disc corona (ADC) system at a distance of 11--24 kpc. Finally, two {\\it XMM-Newton} observations, separated by 2.5 days, have shown that in addition to the changes in the shape of the optical light curve, also the X-ray and UV light curves change significantly on a time scale of days (Hakala et al. 2005). They also detect several X-ray bursts confirming the neutron star nature of the primary. The overall {\\it XMM-Newton} EPIC X-ray spectrum cannot be fitted with a simple model. In addition to the blackbody component and the powerlaw used to model the {\\it ROSAT} and {\\it Chandra} spectra Hakala et al. (2005) find that an additional thermal plasma (MEKAL) component is required in order to fit the spectrum.  The phase resolved X-ray spectra suggest that the most likely cause for the X-ray modulation is changing partial covering of the X-ray emission components over the orbital cycle. This strongly suggests that UW CrB is an ADC source with a thick or warped asymmetric accretion disc.   In this paper we will present further (high time resolution) photometry of UW CrB. We will discuss the observations, followed by a statistical analysis of our data. We will also analyse the 11 optical bursts detected in our data set, together with the bursts detected earlier in various data sets (Hynes et al. 2004, Hakala et al. 2005). Finally we discuss the possible origin of the modulation detected and its implication on our understanding of UW CrB, and accretion discs in LMXB's in general. ",
        "conclusions": "It is rather obvious from our photometry, that UW CrB is most likely a high inclination system, where at least a partial eclipse of an accretion disc is always visible, and that the bulk of the optical orbital modulation is caused by either varying self-shadowing and projected area of a non-axisymmetric accretion disc and/or a partial eclipse of a such disc by the secondary star. Hakala et al. (2005) noted that the X-ray light curves can be best modelled by varying partial covering of the X-ray emitting region over the orbital period. This suggests that the accretion disc either has substantial asymmetric vertical structure or it is clearly warped out of the orbital plane. Our current data cannot distinguish between these models. Our interpretation of the optical light curves does not make any assumptions on the nature of the compact object itself. It is still worth noting that precession of accretion discs is a well known phenomenon in high mass ratio systems (Whitehurst, 1988). Given the magnitude of effects seen in UW CrB, it is possible that this system is an extreme example of the superhump phenomenon seen in nonmagnetic cataclysmic variables with similar orbital periods, i.e. SU UMa systems (Patterson, 2001)."
    },
    "0812/0812.2904_arXiv.txt": {
        "abstract": "We present a comprehensive treatment of the spectrum of electric dipole emission from spinning dust grains, updating the commonly used model of Draine \\& Lazarian.  Grain angular velocity distributions are computed using the Fokker-Planck equation; we revisit the drift and diffusion coefficients for the major torques on the grain, including collisions, grain-plasma interactions, and infrared emission. We use updated grain optical properties and size distributions. The theoretical formalism is implemented in the companion code, {\\sc SpDust}, which is publicly available. The effect of some environmental and grain parameters on the emissivity is shown and analysed. ",
        "introduction": "Observational cosmology has entered an area of high precision, exemplified by the most recent temperature results from sensitive cosmic microwave background (CMB) experiments \\citep{Dickinson04,Readhead04,Kuo07,Hinshaw08}. However, foreground separation and removal remains a major challenge for any CMB measurement (e.g. \\citealt{Eriksen08,Leach08}). In addition to the standard Galactic foregrounds, free-free, synchrotron and thermal dust emission, an unknown ``anomalous'' dust-correlated emission has been observed over the last decade, in the microwave region of the spectrum. The anomalous emissions was first interpreted as free-free emission from shock-heated gas by \\citet{Leitch}, but \\citet{DL98a} showed that this would require an extremely high plasma temperature and a corresponding unrealistic energy injection rate. They proposed instead two possible mechanisms to explain the anomalous microwave emission. One of them is the magnetic dipole emission from thermal fluctuations in the magnetisation of interstellar dust grains \\citep{DL99}. The other possible mechanism, on which the present work focuses, is electric dipole radiation from the smallest carbonaceous grains, described in \\citet{DL98b}, hereafter DL98b. The physical principle is quite straightforward: dust grains are presumably asymmetric, and thus will have a nonzero electric dipole moment.  These grains will spin due to interaction with the ambient interstellar medium (ISM) and radiation field, and thus radiate electromagnetic waves due to the rotation of their electric dipole moment. To get the electric dipole radiation spectrum, one thus needs three ingredients: the quantity of small grains, then their dipole moment, and finally their rotation rates.  Although the observational interest in electric dipole radiation from spinning dust grains has only grown in the last decade, there is a long standing history of theoretical work on the subject. \\citet{Erickson} was the first to consider the possibility that rotating dust grains could be the source of non-thermal radio-noise. \\citet{hoyle} showed that this process was dominated by grains with radius $a \\lesssim 10^{-6} \\ \\cm$ and could lead to radio emission around $10$ GHz. \\citet{Ferrara} estimated the spinning dust emissivity for thermally rotating grains. The first to provide a detailed treatment of rotational excitation of small grains were \\citet{Rouan}. They considered the effect of collisions with gas atoms and absorption and emission of radiation. \\citet{AW93} evaluated the effect of collisions with ions and ``plasma drag'' (torques due to the electric field of passing ions).   DL98b provided the first comprehensive study of the rotational dynamics of small grains, including all the previous effects. They evaluated, as a function of grain radius and environmental conditions, rotational damping and excitation rates through collisions, ``plasma drag'', infrared emission, emission of electric dipole radiation, photoelectric emission and formation of H$_2$ molecules.  The spectra they provided are now widely used in interpreting ISM microwave emission (e.g. \\citealt{Finkbeiner04,Watson05,Casassus06,Casassus07,Casassus08,Dickinson07,Dickinson08,Dobler08}) and for CMB foreground analyses (e.g. \\citealt{Banday03,Davies06,Bonaldi07,Hildebrandt07,Gold}).  Given that the DL98b models are now a decade old, and the recent surge in interest in anomalous emission, it is timely to revisit the theory of spinning dust emission, including the approximations made in DL98b.  This is the purpose of this paper.  As in DL98b, we concentrate on the rotation rate of the grains; the size distribution has been reconsidered by other authors, and the grain dipole moment distribution should be regarded as a model parameter since one cannot compute it from first principles. We first review and generalize DL98b rotational excitation and damping rates. We modify the rotational excitation and damping rates by collisions with neutral species, such that it respects detailed balance in the case where the evaporation temperature is equal to the gas temperature. We include the electric dipole potential when evaluating the effect of collisions with ions. Full hyperbolic trajectories and rotating grains are used when computing the effect of plasma drag. We correct the infrared emission damping rate which was underestimated for a given infrared spectrum. Finally, we use these excitation and damping rates to calculate the grain rotational distribution function by solving the Fokker Planck equation. Updated grain optical properties and size distribution are used throughout this analysis. An Interactive Data Language (IDL) code implementing the formulas in this paper, {\\sc SpDust}, is available on the web\\footnote{http://www.tapir.caltech.edu/$\\sim$yacine/spdust/spdust.html}, and will hopefully allow for a more thorough exploration of the parameter space, as well as model fitting to observations.  The paper is organized as follows. In Section \\ref{section : electric dipole} we remind the reader of the electric dipole radiation formula and give the resulting expected emissivity. In Section \\ref{grain prop} we discuss the size distribution and dipole moments, along with other grain properties. We then turn to the main thrust of this study, which is the computation of the angular velocity distribution function. The theoretical formalism is exposed in Section \\ref{Fokker}, which presents the Fokker-Planck equation. Sections \\ref{collisions}--\\ref{H2} discuss the various rotational damping and excitation processes : collisions with ions and neutral species, plasma drag, infrared emission, photoelectric emission, and random H$_2$ formation. The reader interested primarily in the predicted emission may wish to proceed directly to Section \\ref{emissivity}, where we present the resulting emissivity and the effect of various parameters and environment conditions. Our conclusions are given in Section \\ref{conclusion}. Appendix A exposes the techniques used to numerically evaluate integrals of rapidly oscillating functions involved in the plasma drag calculation. Appendix B presents an alternate, quantum mechanical derivation of the rotational damping rate through infrared emission. ",
        "conclusions": "\\label{conclusion}  We have presented a detailed analysis of the rotational excitation and damping of small carbonaceuous grains. We have refined DL98b results in the case of collisions, accounting properly for the centrifugal potential which increases the net damping rate. In the case of collisions with ions, we accounted for the effect of the electric dipole potential on the collision cross section. We found that this is a small effect in the case of charged grains, but that it may significantly increase the damping and excitation rates in the case of neutral grains. We evaluated the contribution of ``plasma drag'' by considering hyperbolic trajectories and rotating grains in the case of charged grains, and straightlines in the case of neutral grains. We corrected DL98b results for the damping through infrared emission. Finally, we calculated the rotational distribution function by solving the Fokker-Planck equation.  We believe our model provides a much more accurate description of the spinning dust spectrum than previous work. However, we would like to remind the reader of its uncertainties and limitations. First, our model only computes the total intensity of the emitted radiation and not the polarization, which would require an additional study of the alignment mechanisms for PAHs. Secondly, the dust grains properties are poorly known: \\begin{itemize} \\item The size distribution and abundance of the smallest grains is uncertain, and in particular the nature of the cutoff at small grain sizes $a\\sim a_{\\rm min}$ can have a large effect on the spectrum. \\item The permanent electric dipole moments of dust grains are not directly constrained by other dust observables.  Given that it cannot be computed from first principles, one may regard it as a free parameter (or parameters) of the spinning dust model. \\end{itemize} Thirdly, we made some simplified calculations in some cases, as an accurate calculation would have been intractable numerically or substantially complicated the code: \\begin{itemize} \\item We used the Fokker-Planck approximation, which starts to break down for our smallest grains because a single collision suffices to change the rotational state.  We expect that the main consequence of a full treatment would be a tail in the emission spectrum extending to high frequencies, because impulsive collisions would be able to increase the rotation velocities of the grains to $>2\\nu_{\\rm peak}$ before dissipative forces had time to act (an effect missed by the Fokker-Planck treatment).  Therefore one should not place too much confidence in the many order-of-magnitude falloff at $\\sim 100\\ $GHz seen in most of our models.  (In many cases this will be unimportant observationally since at high frequencies the vibrational dust contribution is dominant.) \\item In the plasma drag calculation, we neglected the electric dipole potential when evaluating the trajectory of ions, taking the straight-line (neutral grain) or hyperbolic (charged grain) approximation. Relying on the study of collisions with ions, we may expect the dipole moment to have a small effect in the case of a charged grain. On the other hand, its effect in the case of a neutral grain may be more important, as in that case the electric dipole potential provides the dominant interaction. \\item We assumed the evaporation temperature for the smallest grains was the ``temperature'' of the grain just after it has absorbed a UV photon.  This is a physically motivated assumption but its validity is not established. The evaporation temperature can have a significant effect on the spectrum, as can be seen from Fig. \\ref{CNM Tev effect} and one should be aware of the uncertainty in this parameter. Also, we assumed that collisions transition from being sticking to elastic, as the density exceeds a given threshold. Our model is therefore inaccurate in the transition region. \\item When calculating the infrared emission spectrum of the grains, we used DL01 ``thermal continuous'' approximation, which is not very accurate to describe the low energy part of the spectrum. Whereas these uncertainties are not important if one only wants the spectrum $F_{\\nu}$ in the mid-infrared, they may lead to significant errors when calculating the corresponding damping and excitation rates, which are proportional to $\\int \\nu^{-2}F_{\\nu} \\rmd \\nu$ and $\\int \\nu^{-1}F_{\\nu} \\rmd \\nu$ respectively. \\item We ignored systematic torques, although this may not be a major omission for the smallest dust grains. \\end{itemize}  Despite these uncertainties, we believe that this model is the most complete thus far, and will be a useful tool for comparison to observations and testing the spinning dust hypothesis for anomalous microwave emission in various ISM phases."
    },
    "0812/0812.2245_arXiv.txt": {
        "abstract": "Black hole-neutron star (BHNS) binary mergers are candidate engines for generating both short-hard gamma-ray bursts (SGRBs) and detectable gravitational waves.  Using our most recent conformal thin-sandwich BHNS initial data and our fully general relativistic hydrodynamics code, which is now AMR-capable, we are able to efficiently and accurately simulate these binaries from large separations through inspiral, merger, and ringdown. We evolve the metric using the BSSN formulation with the standard moving puncture gauge conditions and handle the hydrodynamics with a high-resolution shock-capturing scheme.  We explore the effects of BH spin (aligned and anti-aligned with the orbital angular momentum) by evolving three sets of initial data with BH:NS mass ratio $q=3$: the data sets are nearly identical, except the BH spin is varied between $a/M_{\\rm BH}=-0.5$ (anti-aligned), 0.0, and 0.75. The number of orbits before merger increases with $a/M_{\\rm BH}$, as expected. We also study the nonspinning BH case in more detail, varying $q$ between 1, 3, and 5.  We calculate gravitational waveforms for the cases we simulate and compare them to binary black-hole waveforms.  Only a small disk ($<0.01M_\\odot$) forms for the anti-aligned spin case ($a/M_{\\rm BH}=-0.5$) and for the most extreme mass ratio case ($q=5$).  By contrast, a massive ($M_{\\rm disk}\\approx 0.2M_\\odot$), hot disk forms in the rapidly spinning ($a/M_{\\rm BH}=0.75$) aligned BH case. Such a disk could drive a SGRB, possibly by, e.g., producing a copious flux of neutrino-antineutino pairs. ",
        "introduction": "Mergers of black hole-neutron star (BHNS) binaries are expected to be among the most promising sources of gravitational waves detectable by ground-based laser interferometers like LIGO~\\cite{LIGO1,LIGO2}, VIRGO~\\cite{VIRGO1,VIRGO2}, GEO~\\cite{GEO}, TAMA~\\cite{TAMA1,TAMA2}, and AIGO~\\cite{aigo}, as well as by the proposed space-based interferometers LISA~\\cite{LISA} and DECIGO~\\cite{DECIGO}. Theoretical models indicate that a neutron star-neutron star (NSNS)~\\cite{HMNS1,HMNS2,ShiUIUC,STU1,STU2,ST,SST} or black hole-neutron star (BHNS)~\\cite{FBSTR,FBST,SU1,SU2,SST,ST07} binary merger may result in a hot, massive disk around a black hole, whose temperatures and densities could be sufficient to trigger a short-hard gamma-ray burst (SGRB).  Indeed, SGRBs have been repeatedly associated with galaxies with extremely low star formation rates (see~\\cite{SGRB_local} and references therein for a review), indicating that the source is likely to involve an evolved population, rather than main sequence stars. The number of detectable BHNS mergers in the observable universe is still an open question.  The uncertainties arise from some aspects of population synthesis calculations that are only partially understood. The estimated event rate of BHNS mergers observable by an Advanced LIGO detector typically falls in the range ${\\cal R}\\sim 1-100~{\\rm yr}^{-1}$~\\cite{KBKOW}.  Motivated by the significance of BHNS binaries for both GRB and gravitational-wave physics, many simulations of BHNS systems have been performed over the past years. In most dynamical simulations of BHNS binaries to date, the self-gravity of the neutron star (NS) and/or the tidal gravity of the black hole (BH) are treated in a Newtonian or post-Newtonian framework (see, e.g., \\cite{Lee00,RSW,Rosswog,Koba,RKLRasio}; see also~\\cite{LRA}, who provide fully relativistic head-on collision simulations).  Our first BHNS merger simulations were performed in an approximate relativistic framework~\\cite{FBSTR,FBST}, which assumed conformal flatness of the spatial metric throughout the evolution (see~\\cite{Isen,WMM}).  In particular, we evolved extreme mass ratio ($q=M_{\\rm BH}/M_{\\rm NS}$$\\gg$1) initial data developed earlier by our group~\\cite{BSS,TBFS05}.  Though this simulation technique only allows for crude estimates, we found that mergers of irrotational BHNS binaries may lead to disks with masses up to 0.3 $M_{\\odot}$, with sufficient heating to emit the neutrino fluxes required to launch an SGRB.  To date the only self-consistent, general relativistic, dynamical simulations of BHNS inspiral and coalescence have been those of Shibata {\\it et al.}~\\cite{SU1,SU2,ST07,yst08}, Etienne {\\it et al.}~\\cite{eflstb08} (hereafter Paper~I), and Duez {\\it et al.}~\\cite{dfkpst08}. In many earlier calculations, especially those with initial mass ratios $q\\lesssim 3$, significant disks are formed after the NS is disrupted, and, for very stiff nuclear equations of state (EOSs), the core of the NS may survive the initial mass transfer episode and remain bound.  These findings contrast with some semi-analytic relativistic arguments, which suggest that disks with appreciable masses are difficult to form via the merger of BHNS binaries~\\cite{MCMiller}.  In their first set of fully relativistic BHNS simulations, Shibata and Ury\\=u found disk masses in the range of 0.1 -- 0.3 $M_{\\odot}$ for corotating neutron stars~\\cite{SU1,SU2}. Subsequently, Shibata and Taniguchi (\\cite{ST07}, hereafter ST) found smaller disk masses of 0.04 -- 0.16 $M_{\\odot}$ for irrotational neutron stars, which are more realistic. However, our own fully general relativistic simulations of irrotational BHNSs in Paper~I suggest that the disk mass is no more than 0.04 $M_{\\odot}$.  Such a small disk mass, if true for generic cases, might disfavor BHNS mergers as possible central engines for SGRBs.  Recently, Yamamoto, Shibata, and Taniguchi presented a new code which implements an adaptive mesh refinement (AMR) algorithm~\\cite{yst08}. In one of their test simulations of BHNS binaries, they use the same initial data as in ST, but they find that the disk mass is less than 0.001 $M_{\\odot}$. This result disagrees with that in ST, but is consistent with our result in Paper~I. They attribute the discrepancy to different grid structures and computational methods.  Most recently, Duez {\\it et al.} developed a new fully general relativistic hydrodynamics code~\\cite{dfkpst08} that evolves the generalized harmonic formulation of Einstein's equations using a pseudospectral method, and the equations of hydrodynamics via shock-capturing, finite-difference techniques. They use this code to perform a simulation of a single equal-mass BHNS and find that the disk mass is less than 0.03$M_{\\odot}$, which agrees with our result in Paper~I.  As we report in this paper, our code has been modified to use the moving-box Carpet~\\cite{Carpet} AMR infrastructure. With AMR, our code is now capable of performing BHNS simulations with resolutions equivalent to Paper~I at only about $\\sim1\\%$ the total computational cost. This enables us to simulate strong-field regions at higher resolution while simultaneously extending our grid's outer boundaries much farther into the wavezone for more accurate wave extraction. We also evolve initial data at larger binary separations than in Paper I.  Some of the BHNS initial data evolved in this paper are also new, as the BHs are now spinning. All of our initial data are constructed in the conformal thin-sandwich (CTS) formalism. The BH spin is incorporated by imposing boundary conditions on the BH horizon as in~\\cite{CP04,Caudill}.  The initial NS is irrotational and is modeled by an $n=1$ polytropic EOS.  In this paper, we present a new set of BHNS simulations that probes how the initial BH spin and the binary mass ratio affect the dynamics and outcome of the merger.  In particular, we study mergers of nonspinning BHNS binaries with mass ratios $q=1$, $3$, and $5$. For the mass ratio $q=3$, we study three cases in which the BH spin parameter ($\\tilde{a}\\equiv J_{\\rm BH}/M_{\\rm BH}^2$) is $-0.5$ (counter-rotating), 0 and 0.75.  As in Paper~I, we focus on binaries in which the NS tidally disrupts before reaching the innermost stable circular orbit (ISCO) and plunging into the BH. These systems are the most likely to yield a significant disk after merger. We thus do not evolve extreme mass-ratio binaries in this paper.  Starting with approximately the same initial orbital angular velocity $\\Omega$, we find that when the BH is spinning, the merger time is delayed when the BH spin aligns with the orbital angular momentum, and is shortened when the BH spin is anti-aligned with the orbital angular momentum.  For example, evolutions of initial data with $M\\Omega \\approx 0.033$ [$M$ is the initial Arnowitt-Deser-Misner (ADM) mass of the system],  require about 4.25 orbits before merger for $\\tilde{a}=-0.5$, but 6.5 orbits before merger for $\\tilde{a} = 0.75$.  This finding is expected, since for large, aligned spins there is more angular momentum that must be radiated away to bring the binary to merger.  This result is consistent with similar findings in the case of BHBHs (see e.g.~\\cite{clz06}).  After the merger, we find a remnant disk with rest mass between $\\lesssim 0.01M_\\odot$ and $\\approx 0.2M_\\odot$ (assuming the NS has a rest mass of $M_0=1.4M_\\odot$). When the initial BH is nonspinning, the disk mass is $\\approx 0.06M_\\odot$ for $q=3$, $\\approx 0.03M_\\odot$ for $q=1$ and tiny ($\\lesssim 0.01M_\\odot$) for $q=5$. For a fixed mass ratio $q=3$, the disk mass is $\\approx 0.2 M_\\odot$ for $\\tilde{a}=0.75$ and tiny ($\\lesssim 0.01M_\\odot$) for $\\tilde{a}=-0.5$. The increase in disk mass with increasing BH spin agrees qualitatively with the semi-relativistic, smoothed particle hydrodynamics (SPH) simulations of large mass-ratio ($q\\approx 10$) BHNS mergers reported in~\\cite{RKLRasio}.  We find that the remnant disks are hot. A rough estimate suggests that the temperature in the disks is $T\\sim 5\\times 10^{10}$K. The results of BH disk simulations in~\\cite{SRJ} suggest that the disk could generate a gamma-ray energy $E\\sim 10^{47}$--$10^{50}$erg from neutrino-antineutrino annihilation, which is promising for BHNS mergers as plausible central engines for SGRBs.  We compute gravitational waveforms and the effective wave strain in the frequency domain for our mergers.  As in Paper~I, we find measurable differences between BHNS waveforms and those produced by BHBH mergers within the Advanced LIGO band.  These differences appear at frequencies at which the NS is tidally disrupted and accreted by the black hole. Extracting the masses and spins of the objects from the inspiral data, the merger waveforms could provide information about the radius of the NS, which in turn may be used to constrain the NS equation of state.  In contrast with Paper~I, we find substantial disks in two of our nonspinning BH cases.  The reason is that in Paper~I, we imposed limits on pressure to prevent spurious heating/cooling in the low-density artificial ``atmosphere''.  We find that imposing these artificial limits spuriously removes angular momentum, especially during the NS tidal disruption phase, thereby suppressing disk formation (see Section~\\ref{sec:Pceiling} for further details).  As we cautioned in Paper~I, our findings are fundamentally limited by the physics we do {\\it not} model in these preliminary simulations (e.g., neutrino transport, magnetic fields, etc), as well as by our choice of a simple $\\Gamma$-law EOS to model the (cold and hot) nuclear EOS. Disk masses may depend sensitively on the initial structure of the neutron star, especially the low-density outer regions, which is determined by the adopted EOS.  So all BHNS merger results should be viewed in light of these caveats. Our main focus has been to establish that reliable simulations in full general relativity can now be launched to address these important issues.  This paper is organized as follows.  In Secs.~\\ref{sec:basic_eqns} and \\ref{sec:numerical}, we briefly outline the basic equations and their specific implementation in our general relativistic, hydrodynamics scheme. Here we also provide an overview of our initial data, gauge conditions, matter evolution equations, and diagnostics. Sec.~\\ref{sec:code_tests} presents code tests designed to validate our new AMR-based scheme. In Sec.~\\ref{sec:results}, we review the results of our BHNS merger simulations.  In Sec.~\\ref{sec:discussion} we summarize our findings and comment on future directions. ",
        "conclusions": "\\label{sec:discussion}  In this paper we perform a new set of fully general relativistic BHNS simulations using our code upgraded with AMR capability.  We vary the initial binary separations and confirm that the simulation outcomes do not change when the initial BHNS separation is large enough. In most cases, a toroidal disk forms around the black hole remnant after the merger.  This result revises our finding of Paper~I, where we artificially imposed a uniform pressure ceiling on all the matter present in the simulation. This ceiling removed a substantial amount of angular momentum in the NS matter during the merger phase, causing the matter to fall into the BH prematurely, and suppressing disk formation.  Fixing the mass ratio $q=3$, we study the effects of BH spin by evolving initial data with BH's spin parameter $\\tilde{a}=-0.5$ (count-rotating), 0 and 0.75 with nearly the same initial orbital angular velocity $M\\Omega \\approx 0.033$.  Not surprisingly, we find that the BHNS inspiral phase lasts longer as $\\tilde{a}$ increases, and the final disk mass increases from $< 1\\%$ ($\\tilde{a}=-0.5$) to $\\approx 4\\%$ ($\\tilde{a}=0$) and $\\approx 15\\%$ ($\\tilde{a}=0.75$) of the NS's rest mass.  Next, we study the effect of varying the mass ratio $q$ by evolving BHNS initial data with $q=1$ and 5. For the $q=5$ case almost all the NS matter plunges into the BH, leaving $< 1\\%$ of the NS's rest mass to form a disk at the end of the simulation. For the $q=1$ case, a low-density, hot spiral region of disrupted NS matter winds around the BH and smashes into the nascent NS tidal tail, causing strong shock heating before it rapidly falls into the BH.  Most of the remaining matter in the tail eventually falls into the BH, leaving about 2\\% of the NS's rest mass behind to form a disk.  The disks formed after the merger of Cases~\\A\\ and \\B\\ have a characteristic density of $3\\times 10^{11} \\mbox{g cm}^{-3}$ and a temperature of $T \\sim 5\\times 10^{10}$K, assuming the NS's rest mass is $M_0=1.4M_\\odot$.  Applying the results of the simulations of BH disks in~\\cite{SRJ}, the disk may produce a gamma-ray energy of $E\\sim 10^{47}$--$10^{50}$erg due to neutrino-antineutrino annihilation. Hence the merger of a BHNS binary is a promising central engine for a SGRB. But self-consistent simulations, taking into account the detailed microphysics and including the correct EOS, are necessary to fully explore the viability of and the mechanism for generating SGRBs from BHNS mergers.  We compute the BHNS gravitational waveforms and power spectra, and find that the signals are attenuated at frequencies roughly equal to double the orbital frequency at which tidal disruption begins. The resulting deviation of BHNS and BHBH waveforms is visible in the Advanced LIGO high frequency band out to distances $\\gtrsim 100~{\\rm Mpc}$. Should the chirp mass determination, combined with higher order PN waveform phase effects, allow for an independent determination of the component masses and spins of the binary companions, the measurement of the BHNS tidal disruption frequency should give a good estimate of the NS radius and, hence, insight into the nuclear EOS.  Carpet-based AMR has provided us with a great improvement in efficiency and accuracy, and future work will continue in this direction.  To improve code performance (which is usually memory-bound), we also plan to employ new techniques (e.g., loop-tiling) for minimizing cache misses.  To improve accuracy, we will implement the piecewise parabolic method (PPM) or weighted essentially non-oscillatory (WENO) reconstruction schemes in hydrodynamics/MHD and experiment with tapered grids to replace second order temporal prolongation (see~\\cite{taperedgrids}).  We plan to study techniques for minimizing refinement boundary crossings, by adding more refinement boxes around the disrupted NS.  This might help reduce our angular momentum conservation violation to below 1\\% for all cases.  Rich with the challenges of modeling extreme matter in the presence of intense gravitational fields, BHNS binary mergers will likely remain a topic at the forefront of theoretical astrophysics and numerical relativity for years to come. The most exciting possibility remains the simultaneous detection of a gravitational wave signal and a GRB. Our theoretical work on BHNSs is partly in anticipation and preparation for that discovery."
    },
    "0812/0812.2842_arXiv.txt": {
        "abstract": "{ Two-armed, grand design spirals and inner and outer rings in barred galaxies can be due to orbits guided by the manifolds emanating from the vicinity of the $L_1$ and $L_2$ Lagrangian points, located at the ends of the bar. We first summarise the necessary theoretical background and in particular we describe the dynamics around the unstable equilibrium points in barred galaxy models, and the corresponding homoclinic and heteroclinic orbits. We then discuss two specific morphologies and the circulation of material within the corresponding manifolds. We also discuss the case where mass concentrations at the end of the bar can stabilise the $L_1$ and $L_2$ and the relevance of this work to the gas concentrations in spirals and rings. ",
        "introduction": "Barred galaxies often have interesting morphological characteristics. These include two-armed, grand design spirals and rings, both inner and outer. Over the last five years we developed a theory which accounts for the formation and the properties of these structures, starting with the dynamics of the regions around the ends of the bar. Our main results are summarised in four papers [Romero-G\\'omez et al. 2006 (Paper I), 2007 (Paper II), 2008; Athanassoula, Romero-G\\'omez \\& Masdemont 2008 (Paper III)], while a fifth one focuses more on comparisons with observations (Athanassoula, Romero-G\\'omez, Bosma \\& Masdemont 2009, Paper IV). In these papers, the reader will also find other relevant references, concerning other aspects of this problem, either theoretical, or observational.  As discussed in \\cite{BinneyTremaine08}, barred galaxies have five equilibrium points, often referred to as Lagrangian points, of which two are unstable ($L_1$ and $L_2$) and three are stable ($L_3$, $L_4$ and $L_5$). The former are located on the direction of the bar major axis, while the latter are located at the centre of the galaxy and on the direction of the bar minor axis. From the region around the $L_1$ and $L_2$ emanate the manifolds which guide the chaotic orbits escaping this vicinity. These manifolds can be simply thought of as tubes that guide and confine the chaotic orbits in question. It is this confined chaos that can account for the grand design spiral arms and the inner and outer rings.  The structures we wish to understand are the spirals and the inner and outer rings in barred galaxies. Inner rings have roughly the size of the bar and are elongated along it. They are noted as $r$ in the classification of de Vaucouleurs and his collaborators \\citep[][and references therein]{ButaCO07}. Outer rings are similar, but of larger size. They are either elongated perpendicular to the bar ($R_1$ type), or along it ($R_2$ type). In some cases the two types of outer rings co-exist, and the structure is then noted $R_1R_2$. Thus, a barred galaxy with both an inner ring and an outer ring of the first type will be noted $rR_1$. A description of this classification, as well as more information on these structures has been given by \\cite{Buta95}. Spirals are also very common amongst barred galaxies \\citep{ElmegreenElmegreen89}. In most cases they start off from the ends of the bar, and they always wind outwards in a trailing sense.  This paper is organised as follows. In Sect.~\\ref{sec:models} we describe the bar models used in our calculations. In Sect.~\\ref{sec:theory} we present the theoretical basis of our theory. In Sect.~\\ref{sec:twoexamples} we discuss two specific morphologies, the spirals and the $rR_1$ rings, and the different patterns of matter circulation that they entail. We also discuss the types of potentials that can lead to these morphologies. In Sect.~\\ref{sec:stable} and \\ref{sec:gas}, we address two physical problems, namely the presence of two extra mass concentrations centred on the two ends of the bar and the effect of the gas, respectively. Finally, we conclude in Sect.~\\ref{sec:final}.  Response calculations in barred galaxy potentials were discussed in this meeting by Patsis in his talk (these proceeedings; see also \\cite{Patsis06} and references in both), while Efthymiopoulos et al. put up a poster discussing driving by spirals (these proceedings; see also \\cite{TsoutsisEV08} and references in both). ",
        "conclusions": ""
    },
    "0812/0812.3149_arXiv.txt": {
        "abstract": "We present a novel, fast and precise method for including the effect of light neutrinos in cosmological $N$-body simulations. The effect of the neutrino component is included by using the linear theory neutrino perturbations in the calculation of the gravitational potential in the $N$-body simulation. By comparing this new method with the full non-linear evolution first presented in \\cite{Brandbyge1}, where the neutrino component was treated as particles, we find that the new method calculates the matter power spectrum with an accuracy better than 1\\% for $\\sum m_\\nu \\lesssim 0.5 \\, {\\rm eV}$ at $z = 0$. This error scales approximately as $(\\sum m_\\nu)^2$, making the new linear neutrino method extremely accurate for a total neutrino mass in the range $0.05 - 0.3 \\, {\\rm eV}$. At $z = 1$ the error is below 0.3\\% for $\\sum m_\\nu \\lesssim 0.5 \\, {\\rm eV}$ and becomes negligible at higher redshifts. This new method is computationally much more efficient than representing the neutrino component by $N$-body particles. ",
        "introduction": "Next to photons neutrinos are the most abundant particles in our Universe. The fact that at least two of the three neutrino mass eigenstates have masses much larger than the current temperature, $m_i \\gg T_0$, means that neutrinos contribute to the matter density and are important for cosmological structure formation. The mass differences established by oscillation experiments, $\\Delta{m}^2_{12}\\simeq 7.6 \\times 10^{-5} \\, {\\rm eV}^2$ and $|\\Delta{m}^2_{23}| \\simeq 2.4 \\times 10^{-3} \\, {\\rm eV}^2$ (see e.g.\\ \\cite{Schwetz:2008er} for a recent data analysis), mean that the heaviest mass eigenstate must have a mass of at least $m \\sim 0.05$ eV. This in turn will have the effect of suppressing the power spectrum of matter fluctuations by $\\sim 5$\\%, an effect which is significantly larger than the precision with which the matter power spectrum can be measured in upcoming surveys.  This has two interesting consequences, first it means that neutrino mass {\\it must} be included in the analysis of future data in order to avoid seriously biasing the measurements of parameters such as the dark energy equation of state, $w$ \\cite{Hannestad:2005gj}. Second, it means that future large-scale structure surveys potentially have the power to probe neutrino masses as low as the current {\\it lower} bound from oscillation experiments \\cite{futuremass}.  Consequently this has led to a much increased interest in gaining a detailed understanding of how neutrinos affect structure formation. In linear theory the effect is extremely well understood and can be calculated to a precision much better than 1\\%. However, these results apply only on very large scales, $k \\ll 0.1 \\, h {\\rm Mpc}^{-1}$, whereas most of the cosmologically relevant information from large-scale structure surveys is in the range $k \\sim 0.1-0.5 \\, h {\\rm Mpc}^{-1}$. On these intermediate scales it is mandatory to correct for non-linear effects, even if surveys are carried out at intermediate redshifts where non-linearity is weaker.  Non-linear corrections can be found in a number of ways: The most accurate, but also most time-consuming method is to use full $N$-body simulations with neutrinos included in a self-consistent way. In a previous publication we carried out a large suite of simulations with neutrinos included \\cite{Brandbyge1}, and found a significant non-linear correction caused by neutrinos. The effect is large enough to be important even for masses as low as 0.05 - 0.1 eV. Another way is to use higher order perturbation theory which, however, is applicable only up to $k \\sim 0.1-0.2 \\, h {\\rm Mpc}^{-1}$ today, if the precision needs to be better than 1-2\\%. Finally it is possible to use semi-analytic methods such as the halo model which has been calibrated using $\\Lambda$CDM simulations. Neutrinos can be included fairly easily in this scheme, but the accuracy can only be tested against full $N$-body simulations.  In conclusion, a very large number of high-resolution $N$-body simulations with neutrinos included will be necessary for future data analysis. The method used in \\cite{Brandbyge1} is well suited for large neutrino masses with $\\sum m_\\nu \\gtrsim 0.5$ eV at low redshift, but for all masses the simulations inevitably become noise dominated at high redshift, unless extremely many neutrino $N$-body particles are used.  In the present paper we describe a novel, fast and precise method for simulating low mass neutrinos in $N$-body simulations which uses the fact that low mass neutrinos have very little non-linear clustering even at low redshift. Instead of simulating neutrinos as particles with individual velocities we describe the local neutrino density on a grid. This density is evolved forward in time by using linear theory. The CDM/baryon component is followed simultaneously using the full TreePM method, with the neutrino component contributing to the long range force calculated using the PM method. As will be shown below the accuracy of this method is better than 1\\% for the matter power spectrum today on {\\it all} scales for all neutrino masses below $\\sum m_\\nu \\sim 0.5$ eV.  The method is very powerful, both because it speeds up simulations with neutrinos very significantly (by factors of order $\\sim 10$), {\\it and} because it can be applied to any other component of energy density which is almost linear. This could for example be dark energy with linear perturbations. The only requirement is that the linear component can be included in a Boltzmann solver like CAMB \\cite{CAMB}.  We note that our new method is conceptually similar to the perturbative approaches presented in \\cite{saito1, Wong08}, except for the fact that we give the CDM component a full non-linear treatment.  In Section \\ref{IC} we describe our cosmological model and the set-up of initial conditions for our neutrino {\\it N}-body simulations. In Section \\ref{RESULTS} we present our results and in Section \\ref{CONVERGE} we show that our results have converged. Finally, Section \\ref{DISCUSSION} contains a discussion and our conclusions. ",
        "conclusions": "\\label{DISCUSSION}  We have presented a new method for implementing neutrinos in $N$-body simulations which works extremely well for a total neutrino mass below $\\sim 0.5$ eV. For such masses the difference in the matter power spectrum compared with simulations where neutrinos are treated as particles is always below 1\\% on all scales. The precision is better for smaller masses, with the difference scaling roughly as $m_\\nu^2$, and the gain in computational speed compared to representing the neutrinos as $N$-body particles is very large (scaling roughly as $m_\\nu^{-1}$). For all masses this is the only computationally feasible way to include neutrinos in simulations at the required level of precision, especially for high $N$-body starting redshifts.  We also note that the method presented here will work for any type of energy density which has almost linear perturbations.  Finally we caution that although the method is very accurate for calculating all observables related to matter fluctuations, i.e.\\ the power spectrum, halo mass functions etc, it is not accurate at the 1\\% level in describing the neutrino fluctuations alone. For predicting quantities such as the local density of relic neutrinos additional steps should be taken. One possibility is to use the 1-particle Boltzmann technique \\cite{Singh:2002de,Ringwald:2004np}. In its simplest form, however, this method represents the neutrino component only in an approximate way. Another path is to solve with the method presented here until $z$ becomes sufficiently low (in practise $z \\lesssim 4$) that the noise from the neutrino thermal velocity can be kept under control. At this point the grid simulation can be used as the initial condition for a simulation with neutrinos treated as particles."
    },
    "0812/0812.1466_arXiv.txt": {
        "abstract": "We revisit the possible turbulent sources of the solar dynamo. Studying axisymmetric mean-field dynamo models, we find that the large-scale poloidal magnetic field could be generated not only by the  famous $\\alpha$ effect, but also by the $\\mathbf{\\Omega}\\times \\mathbf{J}$ and shear-current effects. The inclusion of these additional turbulent sources alleviates several of the known problems of solar mean-field dynamo models. ",
        "introduction": "Most  solar dynamo models use the scenario proposed by Parker \\cite{par55}. In this scenario, the solar magnetic field is produced by an interplay between differential rotation and the collective action of  cyclonic turbulent convection flows. The latter is widely known as the $\\alpha$ effect \\cite{krarad80}, which is believed to be responsible for the generation of the poloidal component of the large-scale magnetic field (LSMF) of the Sun. There are, however, a number of problems with the $\\alpha$ effect \\cite{brasub05}, and recent developments of mean-field magnetohydrodynamics have stimulated a quest for additional or alternative turbulent dynamo mechanisms. The influence of turbulence on the evolution of the LSMF is expressed by the mean electromotive force (MEMF) $\\boldsymbol{\\mathcal{E}}=\\left\\langle \\mathbf{u}\\times\\mathbf{b}\\right\\rangle$, where $\\mathbf{u}$ and $\\mathbf{b}$ are the fluctuating parts of the velocity and magnetic field (angular brackets denote ensemble averages). It can be written in compact form as \\begin{equation} \\boldsymbol{\\mathcal{E}}=\\left(\\hat{\\alpha}+\\hat{\\gamma}\\right)\\circ\\mathbf{B}-\\left(\\hat{\\eta}+\\hat{\\Omega}+\\hat{W}+ \\dots \\right)\\circ\\left(\\nabla\\times\\mathbf{B}\\right). \\label{eq:emf}\\end{equation} Here $\\mathbf{B}$ is the mean, large-scale magnetic field. The tensors $\\hat{\\alpha}$, $\\hat{\\gamma}$ and $\\hat{\\eta}$ describe the $\\alpha$ effect, turbulent pumping and turbulent diffusion, $\\hat{\\Omega}$ the $\\mathbf{\\Omega}\\times\\mathbf{J}$ effect (or $\\delta^{(\\Omega)}$ effect) and $\\hat{W}$ the shear-current effect (or $\\delta^{(W)}$ effect). Discussions of the different contributions to $\\boldsymbol{\\cal{E}}$ may, for instance, be found in \\cite{radste06,rogkle04}. In this paper, we use  calculations of them as described in \\cite{pip08}, and construct a set of axisymmetric kinematic dynamo models in a convective spherical shell. These models are obtained as solutions of the mean-field dynamo equation \\begin{equation} \\frac{\\partial\\mathbf{B}}{\\partial t}=\\mathbf{\\nabla}\\times\\left(\\mathbf{V}\\times\\mathbf{B} +\\boldsymbol{\\mathcal{E}}\\right),\\label{eq:dyn} \\end{equation} where $\\mathbf{V}=r\\sin\\theta\\,\\Omega\\left(r,\\theta\\right)$, with $\\Omega(r,\\theta)$ modeling the differential angular rotation rate of the Sun ($r$ and $\\theta$ denote radius and colatitude); see Fig.~\\ref{fig:fig0}, left-hand side in left panel. \\begin{figure} \\begin{centering} \\includegraphics[height=6.6cm,angle=-90]{fig1_10} \\includegraphics[height=6.6cm,angle=-90]{fig1_2} \\caption{ \\small\\label{fig:fig0}Basic model quantities. Left panel: Contours of the rotation rate in the solar convection zone (left) and geometry of the pumping velocity of the toroidal LSMF (right); the maximum amplitude of the pumping velocity is about $1\\, \\mathrm{m}/\\mathrm{s}$. Right panel: Radial profiles of the Coriolis number $\\Omega^{*}$ (solid line) and  $\\alpha_{\\phi\\phi}\\sec\\theta$ (dashed line). $\\alpha_{\\phi\\phi}$ changes sign near the bottom of convection zone. Similar radial dependences are found for the other components of the $\\hat{\\alpha}$ tensor.} \\par \\end{centering} \\end{figure} The LSMF is modeled in the form $\\mathbf{B}=\\mathbf{e_{\\phi}}B\\left(r,\\theta\\right)+\\mathbf{\\nabla}\\times\\left({\\displaystyle \\frac{A\\left(r,\\theta\\right)\\mathbf{e_{\\phi}}}{r\\sin\\theta}}\\right)$ ($\\mathbf{e}_{\\phi}$ denotes the azimuthal unit vector), and the MEMF is calculated according to Eq.~(\\ref{eq:emf}). Details of the procedure are given in \\cite{pipsee08}. The radial profiles of characteristic quantities of the turbulence, such as the rms convective velocity $u_{c}$, the correlation length and time $\\ell_{c}$ and $\\tau_{c}$, and the density stratification parameter $G=\\nabla\\log\\rho$ % ($\\rho$ is the mass density), are computed on the basis of a standard model of the solar interior \\cite{sti02}. Some of the generation effects depend on the rms intensity of a fluctuating background magnetic field $\\mathbf{b}^{(0)}$, generated by a small-scale dynamo. We assume energy equipartition between $\\mathbf{b}^{(0)}$ and the turbulent velocity field, i.e. $\\sqrt{\\left\\langle{\\mathbf{b}^{(0)}}^{2}\\right\\rangle}/\\left(u_{c}\\sqrt{4\\pi\\rho}\\right)=1$. The relative strengths of the turbulence effects are controlled via parameters $C_{\\alpha}$ ($\\alpha$ effect), $C_{\\omega}$ ($\\mathbf{\\Omega}\\times \\mathbf{J}$ effect) and $C_{W}$ (shear-current effect); in addition, turbulent pumping is always present, depending merely on the turbulence level (with which all other contributions are also scaled). The integration domain is radially bounded by $r=0.71R_{\\odot}$ and $r=0.96R_{\\odot}$, where the boundary conditions are ${\\displaystyle \\frac{\\partial rB}{\\partial r}=0,A=0}$ at the bottom boundary, and vacuum conditions at the top boundary. Four basic model quantities, namely, the differential rotation, the turbulent pumping velocity of the toroidal component of the LSMF and the radial profiles of the Coriolis number $\\Omega^{*}=2\\Omega_{0}\\tau_{c}$ ($\\Omega_{0}=2.86\\cdot 10^{-6}s^{-1}$ is the surface rotation rate) and  of $\\alpha_{\\phi\\phi}\\sec\\theta$ ($\\alpha_{\\phi\\phi}$ measures the strength of the azimuthal $\\alpha$ effect) are shown in Fig.~\\ref{fig:fig0}. ",
        "conclusions": "The inclusion of the $\\mathbf{\\Omega}\\times \\mathbf{J}$ and shear-current effects in addition to the $\\alpha$ effect and differential rotation in axisymmetric kinematic dynamo models alleviates some of the known problems of current mean-field solar dynamo models. For instance, the simulated period of the dynamo is increased and comes, thus, closer to the observed activity period. Furthermore, the large-scale toroidal field is concentrated towards the equator, thus bringing the models in better agreement with the observation of two activity belts relatively close to the equator. Improved diagnostic tools based on both observation and theory are needed to clarify the roles of the different turbulence effects, as well as that of meridional circulation, which was not included in the present study."
    },
    "0812/0812.1699_arXiv.txt": {
        "abstract": "{ We report the discovery of JKCS\\,041, a massive near-infrared selected cluster of galaxies at $z_{phot} \\sim1.9$. The cluster was originally discovered using a modified red-sequence method and was also detected in follow-up Chandra data as an extended X-ray source. Optical and near-infrared imaging data alone allow us to show that the detection of JKCS\\,041 is secure, even in absence of the X-ray data. We investigate the possibility that JKCS\\,041 is not a galaxy cluster at $z\\sim1.9$, and find other explanations unlikely. The X-ray detection and statistical arguments rule out the hypothesis that JKCS\\,041 is actually a blend of groups along the line of sight, and we find that the X-ray emitting gas is too hot and dense to be a filament projected along the line of sight. The absence of a central radio source and the extent and morphology of the X-ray emission argue against the possibility that the X-ray emission comes from inverse Compton scattering of CMB photons by a radio plasma. The cluster has an X-ray core radius of $36.6^{+8.3}_{-7.6}$ arcsec (about 300 kpc), an X-ray temperature of $7.4^{+5.3}_{-3.3}$ keV, a bolometric X-ray luminosity within $R_{500}$ of $(7.6\\pm0.5)\\times 10^{44}$ erg s$^{-1}$, and an estimated mass of $M_{500}=2.9^{+3.8}_{-2.4}\\times 10^{14}M_\\odot$, the last derived under the usual (and strong) assumptions. The cluster is composed of $16.4\\pm6.3$ galaxies within 1.5 arcmin (750 kpc) brighter than $K\\sim 20.7$ mag. The high redshift of JKCS\\,041 is determined from the detection colour, from the detection of the cluster in a galaxy sample formed by $z_{phot}>1.6$ galaxies and from a photometric redshift based on 11-band spectral energy distribution fitting. By means of the latter we find the cluster redshift to be $1.84<z<2.12$ at 68 \\% confidence. Therefore, JKCS\\,041 is a cluster of galaxies at $z_{phot} \\sim 1.9$ with a deep potential well, making it the most distant cluster with extended X-ray emission known. } ",
        "introduction": "Clusters of galaxies are known to harbour red galaxies out to the highest redshifts explored thus far, $z=1.45$ (Stanford et al. 2006). They owe their colour mostly to their old stellar populations: their luminosity function is passively evolving (de Propris et al. 1999, Andreon 2006a, Andreon et al. 2008) and their colour-magnitude relation evolves in slope, intercept and scatter as expected for a passive evolving population (e.g. Stanford et al. 1999; Kodama et al. 1998). The presence of red galaxies characterises clusters irrespectively of the way they are selected:  X-ray selected clusters show a red sequence (see, e.g. Andreon et al. 2004a, 2005, for XMM-LSS cluster samples at $0.3<z<1.2$, Ebeling et al. 2007 for MACS, whose red sequence is reported in Stott et al. 2007 and Andreon 2008, see also Stanford et al. 2005, Lidman et al. 2008, Mullis et al. 2005 for some other individual clusters). The first Sunyaev-Zeldovich selected clusters (Staniszewski et al. 2008) have been confirmed thanks to their red sequence, and their photometric redshifts were derived from its colour. The same is often true for shear-selected clusters (e.g. Wittman et al. 2006), starting with the first such example (Wittman et al. 2001).  It is now established that bright red galaxies exist at the appropriate frequency in high redshift clusters and the discussion has now shifted to the faint end of the red population. Current studies aim at establishing whether their abundance evolves with lookback time, with proponents divide into opposite camps (e.g Andreon 2008; Lidman et al. 2008; Crawford et al. 2008, Tanaka et al. 2008 vs Gilbank \\& Balogh 2008, Stott et al. 2007). However, these galaxies are too faint for the purpose of discovering clusters, and thus such discussion is not relevant here.  On the contrary, bright red galaxies are a minority in the field population. Consequently, selections based on galaxy colour enhance the contrast of the large red galaxy population in clusters relative to bluer field galaxies.  The Balmer break is a conspicuous characteristic feature of galaxy spectra. Red colours are observed when a galaxy has a prominent Balmer break and is at the appropriate redshift for the filter pair used. In some circumstances a galaxy may look red, when another spectral break (e.g. Lyman break) falls between the filter pair, but these galaxies are numerically few in actual samples and are not clustered as strongly as the galaxies in clusters. Therefore, their clustering cannot be mistaken as a cluster detection. Indeed, if they were found to be so strongly clustered it would be an interesting discovery in itself. The colour index thus acts as a (digital) filter and removes galaxies with a spectral energy distribution inconsistent with the expected one. One can then select against objects at either a different redshift or having unwanted colours, making the colour selection a very effective way of detecting clusters. Red-sequence-like algorithms, pioneered by Gladders \\& Yee (2000) for detecting clusters rely precisely on identifying a spatially localised overdensity of galaxies with a pronounced Balmer break.   Several implementations of red-sequence-like algorithms have been used to detect clusters (e.g. Gladders \\& Yee 2000; Goto et al. 2002, Koester et al. 2007), each one characterised by different assumptions on the cluster model (i.e. on the expected properties of the ``true\" cluster). For example, some models require clusters to follow the Navarro, Frenk \\& White (1997) radial profile with a predicted radius scale. Alternatively the scale (or form of the profile) can be left free (or only weakly bounded). To enlarge the redshift baseline in such cluster surveys, one is simply required to change filters in order to ``follow\" the Balmer break to higher and higher  redshifts. One version of these algorithms has been applied in a series of papers by Andreon et al.' to cover the largest redshift range sampled in one study. In Andreon (2003) the technique was applied to the nearby ($z<0.3$) universe using SDSS $g$ and $r$ filters.  With the $R$ and $z'$ filters  the medium-distant universe ($0.3 \\la z < 1.1$) was probed (Andreon et al. 2004a,b, 2005, 2008, note that the same bands are used by the red sequence survey, Gladders \\& Yee 2005). To sample the $z\\approx 1.2$ universe, infrared bands must be used for the ``red\" filter, because the Balmer break moves into the $z'$ band at $z \\sim 1.2$. This is shown by Andreon et al. (2005), who were able to detect XLSSC 046 (called ``cluster h\" in that paper) at $z=1.22$ using $z'$ and $J$, but we missed it in $R$ and $z'$.  XLSSC  046 was later spectroscopically confirmed in Bremer et al. (2006).  Cluster detection by red sequence-like method is observationally cheap: it is possible to image large ($\\ga 1$ deg$^2$ per exposure) sky regions at sufficient depth to detect clusters up to $z=1$, with short ($\\la 2$ ks) exposure times on 4m ground-based telescopes.  All-sky X-ray surveys such as the RASS also provide efficient cluster detections, although they do require more expensive space-based observations. For distant clusters, currently operational X-ray observatories are less efficient than red-sequence methods due to the longer exposure times required (e.g. $> 10$ ks) and smaller fields of view (1/9 deg$^2$ at most). As a result, red-sequence-like studies are able to follow the cluster mass function down to lower mass limits than is possible with X-ray observations (e.g. Andreon et al. 2005). On the other hand, X-ray detected clusters are less affected by confusion issues, i.e. by the possibility that the detected structure is a line-of-sight blend of smaller structures. Cluster detection by gravitational shear, meanwhile, requires deeper optical data than red-sequence-like algorithms. This is because shear studies require the measuement of subtler quantities (small distortions in shape) of fainter galaxies. Furthermore, detection by shear is affected by confusion (we will return to this issue in sec. 3.6), and thus may be of limited use for detecting samples of clusters. However, shear measurements offer a more direct probe of cluster mass, provided the shear detection has a sufficient signal to noise to perform the measurement in survey data (although this is currently not common, see e.g. Schirmer et al. 2007).   In this paper, we present the detection of the very distant, $z_{phot}=1.9$ cluster JKCS\\,041\\footnote{Cluster names are acronyms indicating the filter pair used for the detection (gr,Rz,JK) followed by the string CS, for colour selected, followed by the order number in the catalogue. These names are IAU-compliant as the acronyms are registered.} obtained by adopting redder ($J$ and $K$) filters to follow the Balmer break to even higher redshifts.  Since the Balmer break enters the $J$ filter at $z\\sim1.9$, $J-K>2.3$ mag (in the Vega system) is a very effective criterion (Saracco et al. 2001, Franx et al. 2003) for selecting $z\\ga 2$ galaxies. The spectroscopic study by Reddy et al. (2005), by Papovich et al. (2006) and by Kriek et al. (2008), all confirme that the $J-K>2.3$ mag criterion selects mainly galaxies at $z\\ga2$.  The paper is organised as follow: Section 2 presents the original cluster discovery. In section 3 we reinforce the cluster detection using different methods and we determined the cluster photometric redshift. In that section, we also show that JKCS\\,041 is not a blend of two (or more) groups along the line of sight. JKCS\\,041 is also X-ray detected in follow-up Chandra observations. Section 4 describes these data, confirming the cluster detection, and presents our measurements of basic cluster properties (core radius, luminosity, temperature). We also show here that other possible interpretations of the X-ray emission (unlikely a priori) do not match our data. After a short discussion (section 5), section 6 summarises the results.  We adopt $\\Omega_\\Lambda=0.7$, $\\Omega_m=0.3$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$.  The scale, at $z=1.9$, is 8.4 kpc arcsec$^{-1}$. Magnitudes are quoted in the photometric system in which they were published (Vega for near-infrared photometry, AB for optical photometry), unless stated otherwise.  \\begin{figure*} \\centerline{\\psfig{figure=12299f1a.ps,width=5truecm,clip=} \\quad % \\psfig{figure=12299f1b.ps,width=5truecm,clip=} \\quad% \\psfig{figure=12299f1c.ps,width=5truecm,clip=}} \\caption[h]{Number density image of a region of 146 Mpc$^2$ (0.16 deg$^2$) area centred on JKCS\\,041. In the left panel, only $2.1<J-K<2.5$ mag galaxies are considered, and all the ancillary photometry ignored. This shows the original cluster detection. In the central panel, we have discarded foreground galaxies, as identified by their spectral energy distribution (SED), using spectroscopy, optical and near-infrared photometry. In the right panel, we only keep galaxies with SEDs similar to the Grasil 1.5 Gyr old ellipticals at $z=1.9$ and we also use Spitzer photometry. JKCS\\,041, at the centre of each panel, is clearly detected in all images, independently of the filtering applied. The images have been smoothed with a Gaussian with $\\sigma=54$ arcsec for displaying purpose. A 5 arcmin ruler is also marked. The large amount of white/light-gray space in each panel qualitatively shows that the detection is unlikely to result from chance background fluctuations. North is up, East is to the left. } \\end{figure*} ",
        "conclusions": "\\subsection{The redshift of the X-ray emission}  The redshift of the X-ray emitting gas, unless it is directly derived from the X-ray spectrum, is commonly assumed to be the same of the galaxy overdensity spatially coincident with it. This is the tacit assumption made in all but a couple of studies related to cluster gas intrinsic properties, like X-ray luminosity or temperature. We make the same assumption and we further check that no other cluster is in the JKCS\\,041 line of sight.  JKCS\\,041 lies in a region rich of photometric data, which have been exploited by several groups independently.  Using  CFHTLS data, Olsen et al. (2007), Grove et al. (2009) and Mazure et al. (2007) explored the range up to $z \\sim 1.1$ using several filters and detection methods. None of their candidate clusters spatially matches JKCS\\,041, the nearest being 6.5 arcmin away. We also found no matching detection by using $R$ and $z'$ bands and the very same algorithm used with $J$ and $K$. Therefore, we found no evidence of another cluster on the JKCS\\,041 line of sight.  \\subsection{JKCS\\,041 and cosmological parameters}  The detection of JCSK\\,041 is in line with the expectation of a $\\Lambda$CDM model, with $\\sigma_8=0.9$ and power spectral shape parameter $\\Gamma=0.6$ and a Jenkins et al. (2001) mass function. In the surveyed area (about $53 \\times 53$ arcmin$^2$) and within $1.8<z<2.0$ about 2.4 clusters with $M>10^{13.75} \\ h^{-1}$ are expected.   In the past, cosmological parameters have been constrained by the observation of a single high redshift cluster.  The argument used is quite simple: cosmological parameters that predict in the studied volume fewer than one cluster more massive than the observed cluster are disregarded in favour of those that make predictions close to the observed number of clusters (i.e. one). Because the first discovered high redshift object is usually extreme (being extreme makes its discovery easier), it is likely to fall in the tail of the distribution. Trying to invert the argument, and constraining cosmological parameter values after having observed a likely extreme object is quite dangerous. This assumes a perfect knowledge of the tail of the distribution from which the object is drawn, which is seldom true on general grounds, and it is certainly not true in this specific case. The halo mass function differs somewhat at the high mass end between different theoretical determinations, e.g. Jenkins et al. (2001), Press \\& Schechter (1974), Sheth \\& Tormen (1999) and numerical simulations (see Jenkins et al. 2001).  For this reason, we do not attempt to constrain cosmological parameters from the JKCS\\,041 discovery, and we simply note that its detection is in line with (model dependent) predictions."
    },
    "0812/0812.2715_arXiv.txt": {
        "abstract": "We report initial results of the first flight of the Antarctic Impulsive Transient Antenna (ANITA-1) 2006-2007 Long Duration Balloon flight, which searched for evidence of a diffuse flux of cosmic neutrinos above energies of $E_{\\nu} \\simeq 3 \\times 10^{18}$~eV. ANITA-1 flew for 35 days looking for radio impulses due to the Askaryan effect in neutrino-induced electromagnetic showers within the Antarctic ice sheets. We report here on our initial analysis, which was performed as a blind search of the data. No neutrino candidates are seen, with no detected physics background. We set model-independent limits based on this result. Upper limits derived from our analysis rule out the highest cosmogenic neutrino models. In a background horizontal-polarization channel, we also detect six events consistent with radio impulses from ultra-high energy extensive air showers. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.2186_arXiv.txt": {
        "abstract": "{ We have studied the supernova remnant \\G\\ with the Chandra X-ray observatory. The remnant's X-ray morphology correlates well with the double-arc structure seen at radio wavelength. The X-ray spectra of the northern and southern rim of \\G\\ exhibit emission line features of highly ionized metals, which suggests that most of the observed X-rays originate in a thermal plasma. We find magnesium , silicon, and sulphur overabundant relative to the solar values. Gaussian emission lines at $\\sim4$ keV and $\\sim7$ keV are detected. The $\\sim4$ keV line is consistent with K-emission lines from $^{44}$Ca and/or $^{44}$Sc whereas the $\\sim7$ keV line feature may arise from unresolved Fe-K lines. Chandra's sub-arcsecond angular resolution allowed us to detect four faint point sources located within $\\sim1.5$ arc-minutes of the geometrical remnant center. Among these objects, \\sten\\ and \\hard\\ do not have optical counterparts, leaving them as candidates for a possible compact stellar remnant. ",
        "introduction": "Thanks to the progress made in high-energy astrophysics in recent years other manifestations of neutron stars besides accretion- and rotation-powered pulsars have been found. Among the various different categories known today, central compact objects (CCOs) are the most enigmatic ones. These sources show up as relatively faint X-ray point sources located close to the expansion center of their host supernova remnants (SNRs), strongly suggesting that they are indeed the compact remnants formed in the supernovae. Despite their apparently young age, the emission properties of CCOs are found to be very different from those observed in young rotation-powered pulsars. In the latter the X-ray emission is generally dominated by magnetospheric emission whereas the X-ray spectra of CCOs can be described well by composite blackbody models of temperatures $(T_1,T_2)\\sim (3-7)\\times 10^{6} K$ and radii $(R_1,R_2)\\sim 0.3-3$~km of the projected blackbody emitting areas that are much smaller than the canonical neutron star radius (see Becker 2008 for a review). In contrast to the young Crab- and Vela-like pulsars, no plerionic emission has been found around CCOs (see Hui \\& Becker 2006 and references therein).  The temporal emission behavior of the known CCOs is another puzzle. So far, X-ray pulses have been found from three of the seven known sources. The detected periodic signals span a wide range from $P\\sim 0.1\\,\\mbox{s}$ for the compact object in SNR  Kes 79 (Halpern et al.~2007) to $P\\sim 6.27$ hrs in the case of the central point source in RCW 103 (De Luca et al.~2006). Different solutions have been proposed to explain their rotational dynamics. There is, however, no consensus so far on whether these objects e.g.~were born with a slow spin and a weak magnetic field (Halpern et al.~2007) or at the other extreme that resembles the magnetars with a very short initial spin period and a very strong magnetic field (Li 2007).  Currently, only seven SNRs have been confirmed as hosting a CCO (cf.~Tables 6.3 and 6.4 in Becker 2008). For a better understanding of their nature, the sample of CCOs needs to be increased. We therefore inspected known supernova remnants for centrally peaked X-ray emission by reanalyzing the ROSAT all-sky survey (RASS) data. \\G\\ showed up in RASS data to have centrally bright X-ray emission that could potentially be associated with a compact source. A faint radio source seen in NRAO/VLA Sky Survey (NVSS) data (Condon et al. 1998) close to the geometrical center of the remnant correlates well with the central X-ray emission (cf.~Fig.1). The latter finding raised the question of whether this emission may eventually come from a young powerful pulsar and/or its synchrotron nebula.  However, searching for a radio pulsar in \\G\\ did not yield any detection. This non-detection placed an upper limit of 800 $\\mu$Jy on pulsed radio emission at 600 MHz within \\G\\ (Lorimer et al.~1998). Although this limit is not too constraining, it may suggest that the central emission could probably be associated with a CCO rather than with a young powerful rotation-powered pulsar.  \\begin{figure*} \\centerline{\\psfig{figure=10789f1.eps,width=18cm,clip=}} \\caption[]{ROSAT PSPC RASS image of G67.7+1.8 ($1.0-2.0$ keV) overlaid with contours from the NVSS survey data. The remnant appears center-filled in this data and has a bilateral radio arc-like structure. The peak of the central X-ray emission is found to be roughly matching the geometric center of the remnant. Faint radio sources are seen in the NVSS data near to the remnant center.} \\end{figure*}  \\G\\ was first discovered in a Galactic plane survey at 327 MHz by Taylor et al.~(1992), using the Westerbork Synthesis Radio Telescope. The radio image has shown that \\G\\ has a double arc structure. The angular diameter is $\\sim$ $9^{'}$. Its distance is quite uncertain. The only distance estimate available comes from the $\\Sigma-$D relation, which puts the remnant at $\\sim17$ kpc (Case \\& Bhattacharya 1998). However, optical filaments associated with the northern radio arc of \\G\\ were found by Mavromatakis et al.~(2001) and indicated that the remnants distance might actually be smaller. In the same work Mavromatakis et al.~(2001) concludes that the energy released during the supernova explosion was significantly less than the canonical value of $10^{51}$ ergs.  X-ray emission in the direction towards \\G\\ was first observed in the ROSAT all-sky survey (Mavromatakis et al.~2001). Fitting the data with a thermal bremsstrahlung model indicated a plasma temperature in the range of $\\sim0.2-0.3$ keV.  In this paper we present a detailed X-ray study of \\G\\ and investigate the nature of its centrally peaked emission using Chandra ACIS-I observations. Its superior angular resolution enables us to perform a spectro-imaging study of the supernova remnant and to search for compact sources and/or a pulsar-wind nebula in its central part. ",
        "conclusions": "We have performed a detailed spectro-imaging X-ray study of the supernova remnant \\G\\ with Chandra. Various properties of \\G, including distance, age and explosion energy, are still poorly constrained. The type of the supernova and the nature of the progenitor are also unknown. With the Chandra observation it became possible for the first time to shed some more detailed light on the X-ray emission nature of \\G\\ and its central point sources.  The spectra of the remnants northern and southern rims can be described with a single temperature CIE model plus additional Gaussian components. Our analysis suggests a plasma temperature of $\\sim7\\times10^{6}~K$. Assuming that the ions and electrons are fully equilibrated in the shock, the temperature implies that the blast wave has a velocity of $\\sim800$ km/s which is at the same order of magnitude as inferred from the RASS data (Mavromatakis et al. 2001). Although the single temperature plasma model describes our observed spectra very well, most supernova remnants require multi-temperature models to depict their X-ray spectra (e.g.~Gaensler et al.~2008). A composite model can describe the shocked ejecta and the swept-up material from the surrounding with the components of high temperature and low temperature, respectively. However, for a more detailed modeling of the remnant spectrum, data with improved photon statistics are badly needed.  Utilizing the $\\Sigma-D$ relation, Case \\& Bhattacharya (1998) estimated the distance to \\G\\ to be $\\sim17$ kpc. However, the large uncertainties on the proportionality factor and the exponent of the $\\Sigma-D$ relation make the distance poorly constrained and results in a range of $7-27$ kpc. On the other hand, the observation of the optical filaments suggests a distance less than 17 kpc, otherwise the detection of the optical radiation would be very difficult to explain in view of the measured interstellar extinction (Mavromatakis et al. 2001). In the following calculations we assume the distance to be in the range of $7-17$ kpc.  The emission measure inferred from the spectral fits of the remnant emission allows us to estimate the hydrogen density, $N_{\\rm H}$, in the shocked regions. Assuming the post-shock densities of hydrogen and electrons are uniform in the extraction region, the normalization of the CIE model can be approximated by $N_{\\rm cie}\\simeq 10^{-14}N_{\\rm H}N_{e}V/4\\pi D^{2}$, where $D$ is the distance to \\G\\ in $cm$ and $V$ is the volume of interest in units of $cm^{3}$. Assuming a geometry of oblated spheroid, the volume of interest for the northern rim and the southern rim are $3.9\\times10^{55} d^{3}_{\\rm kpc}$ cm$^{3}$ and $2.6\\times10^{55}d^{3}_{\\rm kpc}$ cm$^{3}$, respectively. Here $d_{\\rm kpc}$ is the remnant distance in units of kpc. Assuming a fully ionized plasma, $N_{\\rm H}N_{e}$ can be approximated by $0.7N_{\\rm H}^{2}$. The best-fit normalizations imply a post-shock hydrogen density of $N_{\\rm H} \\sim0.06-0.1$ cm$^{-3}$ for the assumed range of distances.  Assuming \\G\\ is in a Sedov phase, the shock temperature can be estimated by: \\begin{equation} T_{s}\\simeq 8.1\\times 10^{6}\\left(\\frac{E_{51}}{N_{\\rm H_{-1}}}\\right)^{\\frac{2}{5}}t_{4}^{-\\frac{6}{5}}~K \\end{equation}  \\noindent where $t_{4}$, $E_{51}$ and $N_{\\rm H_{-1}}$ are the time after the explosion in units of $10^{4}~years$, the released kinetic energy in units of $10^{51}~ergs$ and the post-shock hydrogen density of $0.1~cm^{-3}$ respectively (McCray 1987). Mavromatakis et al.~(2001) has argued the explosion energy should be significantly less than the canonical value of $10^{51}$ ergs. Adopting an explosion energy of $E=10^{50}$ ergs and the plasma temperature inferred from our spectral analysis as an estimate of $T_{s}$ (i.e.~$\\sim7\\times10^{6}~K$), the range of $N_{\\rm H}\\sim0.06-0.1$ cm$^{-3}$ yields the age of \\G\\ to be $\\sim5000-6000$ yrs. On the other hand, an age bracket of $\\sim11000-13000$ yrs is suggested from $E=10^{51}$ ergs.  Taking the estimates of the age, explosion energy and the density of the interstellar medium, we calculate the theoretical size of the remnant by (Culhane 1977):  \\begin{equation} R_{s}\\simeq 21.5\\left(\\frac{E_{51}}{N_{\\rm H_{-1}}}\\right)^{\\frac{1}{5}}t_{4}^{\\frac{2}{5}}~{\\rm pc}. \\end{equation}  From the X-ray image, we estimated the angular radius of the remnant to be $\\theta_{s}\\simeq5$ arcmin. Comparing $R_{s}$ and $\\theta_{s}$, we can estimate the distance by $d=R_{s}/\\theta_{s}$.  For an explosion energy of $E=10^{50}$ ergs, the distance is estimated to be in a range of $\\sim7-12$ kpc for $N_{\\rm H}=0.06-0.1$ cm$^{-3}$ and $t=5000-13000$ yrs. On the other hand, if $E=10^{51}$ ergs is adopted, it implies a distance bracket of $\\sim12-20$ kpc for the same ranges of $N_{\\rm H}$ and $t$ in the previous calculation.  As aforementioned, the optical extinction suggests that the distance is at the lower side. On the basis of our estimation, for a preferable distance much smaller than $17$ kpc, we favor a scenario that the explosion is less energetic than a canonical supernova (i.e.~$E<10^{51}$ ergs), which is consistent with the conclusion drawn by Mavromatakis et al.~(2001) from their optical study.  Assuming that the remnant \\G\\ is the result of a low energy supernova explosion has interesting implications for the properties of its progenitor. In the context of the current understanding in stellar evolution, a supernova with kinetic energy of $\\sim10^{50}$ ergs can be a result of two different evolutionary tracks.  In the first scenario, a Oxygen-Neon-Magnesium (ONeMg) core can be formed before Neon and subsequent nuclear burning take place. If the neutrino cooling is efficient enough, the core temperature will be reduced and prevents further nuclear fusion. When the ONeMg core reaches the Chandrasakar mass, the electron degenerate pressure is no longer able to support the core. Furthermore, electron capture by $^{24}$Mg and $^{20}$Ne will further reduce the electron degenerate pressure in the core and trigger the core to collapse (Miyaji et al.~1980; Guti\\'errez, Canal \\& Garca-Berro 2005). This is known as the electron-capture supernova. Most simulations have shown that this type of explosion has a rather low energy (see Eldridge, Mattila, \\& Smartt 2007). The progenitor's mass of an electron-capture supernova is limited in a narrow range of $\\sim6-8$ M$_{\\odot}$ (Eldridge, Mattila, \\& Smartt 2007). For more massive stars, they will go through all stages to silicon burning. In less massive stars, ONeMg cores cannot be formed.  There is another possibility to produce a low energy supernova. The exact evolution of a collapsing star after neutrino re-heating depends on the rate of early infall of stellar material on the collapsed core and on the binding energy of the envelope. If both are large, which is the case in high-mass stars, the energy required to accelerate and heat up the ejecta is not available, preventing a successful explosion and resulting in a supernova which appears to be under-energetic (see Heger et al.~2003; Eldridge \\& Tout 2004). The progenitor of this evolutionary track is likely to be $\\gtrsim20M_{\\odot}$.  We have also discovered emission line features in the remnant spectrum. The feature at $E=4.0\\pm 0.2$ keV is detected only in the northern rim. The width of this feature ($\\sigma=0.3$~keV) is a little larger than the energy resolution of ACIS-I at 4~keV (i.e. $\\sigma\\sim0.1$~keV). This might indicate that the feature can be a blend of several lines. Its energy centroid is close to a number of K emission lines from $^{44}$Ca and/or $^{44}$Sc (cf. Table 5 in Iyudin et al.~2005 and references therein). Therefore, we speculate that the feature at 4~keV can possibly comprises these lines.  If the identification is correct, it suggests a possible presence of $^{44}$Ti because both $^{44}$Ca and $^{44}$Sc are produced in the decay chain $^{44}$Ti$\\rightarrow$$^{44}$Sc$\\rightarrow$$^{44}$Ca. The half-life of $^{44}$Ti is $\\sim60$ yrs. If the line feature at 4 keV is indeed from the decay of $^{44}$Ti, its short half-life implies that the remnant should be rather young. Also, the production of $^{44}$Ti in the supernova is very sensitive to the explosion mechanism and the ejecta dynamics (see discussion in Iyudin et al.~2005). Therefore, obtaining the yield and spatial distribution of $^{44}$Ti can help to further constrain the physical properties of the remnant as well as the nature of the progenitor.  On the other hand, the broadening of the feature can also result from the motion of the ejecta. However, this would imply the ejecta to have a velocity of $\\sim10000$~km/s. This is not consistent with \\G\\ being the result of a low energy supernova explosion as inferred from our X-ray data analysis as well from  optical observation. Furthermore, the limited energy resolution of imaging data precludes any accurate measurement of the line width. High resolution grating spectroscopy is required to disentangle the possible blend of lines as well as determine their widths.  Apart from the feature at 4 keV, we have also found an emission line feature at $\\sim7$ keV in both, the northern and southern rims. The large errors of the energy centroid and the width do not allow a firm identification. A deep observation is required to better constrain these parameters. We note that the approximate energy of the feature is rather close to the Fe-K emission. This leads us to speculate that the feature can be a result from unresolved Fe K$_\\alpha$ and K$_\\beta$ lines. This also requires data with high spectroscopic resolution for further investigation.  The Chandra observation of \\G\\ also leads to the detection of X-ray point sources at the proximity of the remnant's geometrical center. However, the limited photon statistics do not allow a constraining spectral analysis of these sources. The non-detection of optical counterparts for \\sten\\ and \\hard\\ leaves these two objects as possible candiates for a stellar compact remnant.  A dedicated deep optical observation is desirable to better constraining the nature of the point sources."
    },
    "0812/0812.4457.txt": {
        "abstract": "\\noindent {\\footnotesize We argue that both the positron fraction measured by PAMELA and the peculiar spectral features reported in the total electron-positron flux measured by ATIC have a very natural explanation in electron-positron pairs produced by nearby pulsars. While this possibility was pointed out a long time ago, the greatly improved quality of current data potentially allow to reverse-engineer the problem: given the regions of pulsar parameter space favored by PAMELA and by ATIC, are there known pulsars that explain the data with reasonable assumptions on the injected electron-positron pairs? In the context of simple benchmark models for estimating the electron-positron output, we consider all known pulsars, as listed in the most complete available catalogue. We find that it is unlikely that a single pulsar be responsible for both the PAMELA positron fraction anomaly and for the ATIC excess, although two single sources are in principle enough to explain both experimental results. The PAMELA excess positrons likely come from a set of mature pulsars (age $\\sim\\times 10^6$ yr), with a distance of 0.8-1 kpc, or from a single, younger and closer source like Geminga. The ATIC data require a larger (and less plausible) energy output, and favor an origin associated to powerful, more distant (1-2 kpc) and younger (age $\\sim5\\times 10^5$ yr) pulsars. We list several candidate pulsars that can individually or coherently contribute to explain the PAMELA and ATIC data. Although generally suppressed, we find that the contribution of pulsars more distant than 1-2 kpc could contribute for the ATIC excess. Finally, we stress the multi-faceted and decisive role that Fermi-LAT will play in the very near future by (1) providing us with an exquisite measurement of the electron-positron flux, (2) unveiling the existence of as yet undetected gamma-ray pulsars, and (3) searching for anisotropies in the arrival direction of high-energy electrons and positrons.} ",
        "introduction": "Recent measurements of the positron fraction (the fraction of positrons over the sum of positrons and electrons) reported by the PAMELA collaboration \\cite{Adriani:2008zr} and of the electron-plus-positron differential energy spectrum reported by ATIC \\cite{2008Natur.456..362C} and by PPB-BETS \\cite{Torii:2008xu} stirred up great interest both in the cosmic-ray community and among those who work in the field of indirect dark matter searches.  The Payload for Antimatter Matter Exploration and Light-nuclei Astrophysics (PAMELA) measured a statistically significant trend of increasing positron fraction with energy in the $10\\lesssim E_{e^\\pm}/{\\rm GeV}\\lesssim100$ GeV range, at odds with the standard expectation in the context of secondary positron production from interactions of cosmic-ray nuclei with the interstellar gas: if most positrons are of secondary origin, the positron fraction should fall as a function of increasing energy (see e.g. the recent comprehensive stuy of Ref.~\\cite{Delahaye:2008ua}). In the past, a similar trend had actually been repeatedly observed, although with much lower statistical significance,  by other experiments, starting  as early as 1969 \\cite{1969ApJ...158..771F}. More recent results pointing towards an increasing positron fraction above 10 GeV include the measurement reported in Ref.~\\cite{1975ApJ...199..669B} (1975), Ref.~\\cite{1987ApJ...312..183M}, Ref.~\\cite{1994ApJ...436..769G} (1994), the data reported by HEAT (1997, Ref.~\\cite{1997ApJ...482L.191B}) and those quoted by the Alpha Magnetic Spectrometer (AMS) collaboration (2000, Ref.~\\cite{Alcaraz:2000bf}).  The Advanced Thin Ionization Calorimeter (ATIC) recently reported a ``bump'' in the high energy flux of electrons and positrons \\cite{2008Natur.456..362C}, in excess of the standard expectation, and marginally in agreement with a statistically less significant excess also visible in the PPB-BETS data reported in Ref.~\\cite{Torii:2008xu}\\footnote{PPB-BETS measured 84 electron events above 100 GeV \\cite{Torii:2008xu}, while ATIC 1724 \\cite{2008Natur.456..362C}}. Outside the bump, the ATIC data are in substantial agreement with previous measurements, including those reported by AMS \\cite{Alcaraz:2000bf}, BETS \\cite{Torii:2001aw}, older PPB-BETS data \\cite{Torii:2003ed}, as well as emulsion chambers data \\cite{1997AdSpR..19..767N,1980ApJ...238..394N}. The anomalous spectral feature observed in the data from both ATIC flights, and which had never been reported before, is a sharp excess in the electron-positron flux, specifically in the energy range between 300 and 800 GeV \\cite{2008Natur.456..362C}.  In addition to the mentioned experimental results on cosmic-ray electrons and positrons, the atmospheric Cherenkov telescope (ACT) H.E.S.S. recently reported measurements of the flux of very high-energy ($E_{e^\\pm}>600$ GeV) electrons, with data extending all the way up to energies around 5 TeV \\cite{Collaboration:2008aa}. The data were taken during off-target observations of extragalactic objects. Since no conclusive distinction is possible with ACT's between $e^\\pm$ and gamma-rays, the data likely include a contamination from diffuse extra-galactic (and possibly, to some extent, diffuse galactic) gamma rays as well as from residual hadronic cosmic-ray events. The H.E.S.S. data should therefore properly be regarded as {\\em upper limits} to the $e^\\pm$ flux at large energies, and this is how they will be considered in the present study.  The nature of cosmic-ray positrons have long been believed to be dominantly of secondary origin: the observed cosmic-ray positrons, in that scenario, are the yield of inelastic interactions of (primary and secondary) cosmic-ray protons with the interstellar gas (see e.g. Ref.~\\cite{1964ocr..book.....G}, \\cite{1977JPhA...10..843G} and \\cite{1987MNRAS.228..843D}). Calculations of the cosmic-ray positron and electron flux as well as of the fraction of secondary positrons, assuming several propagation models and diffusion setups, have been carried out extensively over the last 30 years (see e.g. \\cite{1982ApJ...254..391P}, \\cite{1998ApJ...493..694M}, \\cite{Delahaye:2008ua}), with increasing degrees of sophistication, and relying on experimental data of better and better quality . The generic prediction of all these models is a slowly decreasing positron fraction with increasing energy, and of a steeply declining high energy differential primary electron flux (for instance, the conventional and optimized models of Strong et al. \\cite{Strong:2004de}, which give a diffuse galactic continuum gamma-ray flux compatible with EGRET data as well as cosmic-ray abundances compatible with a wide set of data, feature injection spectral indexes for high-energy electrons of 2.54 and 2.42, respectively).  As mentioned above, however, experimental data on cosmic-ray positrons have actually shown for quite a long time a discrepancy with a purely secondary origin for positron cosmic rays. For instance, Ref.~\\cite{1989ApJ...342..807B} pointed out almost 20 years ago that, in addition to diffuse sources such as secondary cosmic-ray production and to primary sources, such as electrons re-accelerated in supernova remnant (SNR) shocks, data on the positron fraction and on the $e^\\pm$ flux rather conclusively pointed towards the existence of a powerful, nearby source of energetic primary electrons and positrons above 20 GeV.  As early as 1995, and commenting on data from before 1990, Ref.~\\cite{1995PhRvD..52.3265A} stated that ``{\\em the measured content of positrons in the total electron flux, [...], at least at high energies [...], is a[n] [...] enigma awaiting an explanation. [...] it is obvious that some other source of positrons is needed}''.  Ref.~\\cite{1995PhRvD..52.3265A} also argued for the first time that nearby pulsars might seed a high-energy positron excess, as well as non-standard spectral features in the cosmic-ray electron-positron spectrum. The positron anomaly is therefore nothing {\\em qualitatively} new. We argue here, however, that the greatly improved quality of recent data allows us to {\\em quantitatively} push previous analyses to the level of associating electron-positron data to existing, observationally well-known galactic objects. This is all the more exciting given the apparently anomalous data reported in the total differential electron-positron flux.  A physical hallmark of energetic ($E_{e^\\pm}\\gtrsim10$ GeV) electrons and positrons is that they  loose energy very efficiently via two dominant processes: inverse Compton (IC) scattering of background radiation photons, and synchrotron radiative losses. The energy loss rate for both processes goes as the square of the $e^\\pm$ energy. Defining the lifetime $t$ of relativistic $e^\\pm$ in the interstellar medium as the time it would take for them to cool from an injection energy $E_{e^\\pm}$ to rest energy ($E_{e^\\pm}=m_e$) if they only lost energy via IC and synchrotron losses, one finds \\begin{equation} t\\sim\\ 5\\times 10^5\\left(\\frac{\\rm TeV}{E_{e^\\pm}}\\right)\\ {\\rm yr}. \\end{equation} This implies that TeV $e^\\pm$ must have been injected not much longer than $\\sim10^5$ yr ago. This also implies that, for conventional values of the diffusion coefficient $D$ for the propagation of electrons and positrons, they must have been produced somewhere within $\\sqrt{D\\times t}\\sim$100-500 pc. The efficient energy losses of energetic electrons therefore constrain the sources of high energy cosmic-ray electrons and positrons both in time and in space \\cite{1995A&A...294L..41A}. Any interpretation of the reported ``lepton anomalies'' must evidently be compatible with this fundamental physical property of cosmic-ray electron and positron propagation in the interstellar medium (ISM).  Several scenarios have been suggested to explain a deviation of the electron-positron data from the standard cosmic ray secondary positron model (for a recent model-independent analysis see e.g. Ref.~\\cite{Serpico:2008te}). A list of these includes: \\begin{enumerate} \\item A cutoff in the primary electron spectrum at large energies (see e.g.~ \\cite{1987ICRC....2...88T}). This possibility might help explain the positron fraction data, but is evidently not an explanation to the measured differential $e^\\pm$ flux, which would still be dominated, at high energy, by primary cosmic-ray electrons; \\item The production and re-acceleration of secondary electrons and positrons via hadronic processes in giant molecular clouds. For instance, for nominal model parameters, Ref.~\\cite{1990ICRC....4..109D} points out that those systems can source a significantly larger positron fraction in excess to the prediction of ordinary diffuse cosmic-ray models, starting at energies around 10 GeV. In this scenario, the results of Ref.~\\cite{1999APh....11..429C} indicate that it could be possible to reproduce the PAMELA data below 30 GeV, but no satisfactory explanation to ATIC or to the higher energy bins reported in the PAMELA data seems to be achievable in the context of giant molecular clouds models; \\item Ref.~\\cite{1991JPhG...17.1769A} pointed out that $e^\\pm$ might be pair-produced near compact gamma-ray sources, for instance as a result of processes involving, in the initial state, the source gamma rays and the optical and UV photons also produced locally by the gamma-ray emitter. Electron-positron pair production could also occur from compact gamma-ray sources in interactions of gamma rays with starlight photons in the interstellar medium, as observed in Ref.~\\cite{1991ICRC....2..145M}. This possibility appears, however, to be tightly constrained by current galactic gamma-ray data, and it will be further constrained and quite easily scrutinized by upcoming Fermi data. \\item Another possible source of primary positrons is the radioactive $\\beta^+$ decay of nuclei (such as ${}^{56}$Co) ejected during a supernova explosion \\cite{1990ICRC....4...68E, 1979ICRC....1..501L}. Supernova models, however, indicate that the spectrum of such nuclei (and therefore of the positrons produced in the decays) is cut off at energies around 100 GeV, ruling out this scenario as an explanation to the $e^\\pm$ spectrum measured by ATIC. Also, the energetics of the process indicate that this mechanism is likely insufficient to produce as large a flux of positrons as to explain the PAMELA positron fraction data too. \\item Pulsars are undisputed sources of electron-positron pairs, produced in the neutron star magnetosphere and, possibly, re-accelerated by the pulsar wind and/or in SNR shocks. Nearby pulsars have since long been considered as very likely dominant contributors to the local flux of high energy electrons and positrons (see e.g. Ref.~\\cite{1986ARA&A..24..285T,1988gera.book..480B}). As shown in Ref.~\\cite{Hooper:2008kg} and \\cite{Yuksel:2008rf}, the recent PAMELA data can rather well be interpreted as originating from one single nearby pulsar, e.g. Geminga, or by the coherent superposition of the $e^\\pm$ flux from all galactic pulsars \\cite{2001A&A...368.1063Z,Hooper:2008kg} (see also Ref.~\\cite{Malyshev:2009tw} for a recent re-assessment of pulsars as the origin of the cosmic-ray lepton anomalies appeared after the present manuscript was first posted);  \\item A further possibility that recently received enormous attention is that the excess positrons and electrons originate from the annihilation or decay of particle dark matter (for a possibly incomplete list of related studies appeared before the first version of the present manuscript was released see Ref.~\\cite{Hamaguchi:2008ta, Allahverdi:2008jm, Pohl:2008gm, Liu:2008ci,Nath:2008ch,MarchRussell:2008tu,Zhang:2008tb, Pospelov:2008rn,Lattanzi:2008qa,Chun:2008by,Hisano:2008ah,Ibe:2008ye,Taoso:2008qz,Ishiwata:2008qy,Zurek:2008qg,Cholis:2008wq,Morselli:2008uk,Hall:2008qu,Chen:2008qs,Kalinowski:2008iq,Baek:2008nz,Ibarra:2008jk,Pospelov:2008zw,Ponton:2008zv,Hur:2008sy,Hamaguchi:2008rv,Chen:2008md,Fox:2008kb,Bai:2008jt,Ishiwata:2008cv,Yin:2008bs,Feldman:2008xs,Harnik:2008uu,Nomura:2008ru,Cholis:2008qq,Bringmann:2008jn,Masip:2008mk,Cirelli:2008pk,Cirelli:2008jk,Nelson:2008hj,Pospelov:2008jd,Huh:2008vj,Cholis:2008hb,Barger:2008su,Bergstrom:2008gr,Bae:2008sp,Chen:2008xx}; further studies appeared between the first and the present version of this manuscript include \\cite{Grajek:2008pg, Huh:2008ez, Bi:2009md, Cui:2009xq, Gogoladze:2009kv, Cao:2009yy, Guo:2009aj, Takahashi:2009mb, Shepherd:2009sa, Nezri:2009jd, Chen:2009mf, Meade:2009rb, Mardon:2009rc, deBoer:2009rg, Brandenberger:2009ia, Banks:2009rb, Kribs:2009zy, Chang:2009dh, Kyae:2009jt, Hamaguchi:2009sz, Phalen:2009xw, Chen:2009dm, Chen:2009iu, Hooper:2009fj, Goh:2009wg, Ibe:2009dx, Cheung:2009qd, Allahverdi:2009ae, Bae:2009bz, Bringmann:2009ip, Cheung:2009si, Shirai:2009kh, Bi:2009uj, Ishiwata:2009vx, Frampton:2009yc, Finkbeiner:2009mi, Cassel:2009pu, Roszkowski:2009sm, Chen:2009ew, Rothstein:2009pm, Ishiwata:2009pt, Zant:2009sv, Brun:2009aj, McDonald:2009cc, Lingenfelter:2009kx, Barger:2009yt, Gogoladze:2009gi, Morrissey:2009ur, Arvanitaki:2009yb, Davoudiasl:2009dg, Kuhlen:2009is, Kawasaki:2009nr}). As opposed to all of the other possibilities mentioned above, a dark matter interpretation invokes an entity whose fundamental particle physics nature has yet to be unveiled, and whose existence in the form of a particle with a mass in the GeV-TeV range is yet to be demonstrated. In addition to this, there are at least three general weaknesses to the dark matter interpretation of the positron fraction and of the features detected in the high energy electron-positron flux: \\begin{enumerate} \\item the annihilation or decay products of dark matter generically include both leptons and hadrons. Unless the dark matter model is tuned in such a way as to make it ``lepto-philic'' \\cite{Fox:2008kb}, the hadronic products will typically overproduce both gamma rays and antiprotons, in contrast with experimental data (e.g. \\cite{Bergstrom:2006tk}); \\item the annihilation rate required to explain the cosmic-ray data under discussion is typically a couple orders of magnitude larger than what expected if the dark matter is cosmologically produced through thermal freeze-out in the right amount. This implies a further model-building intricacy. In addition, low-velocity enhancements, which would lead to a burst of annihilations in the first collapsed halos \\cite{Profumo:2006bv}, are generically very tightly constrained by observations of the cosmic diffuse background radiation \\cite{Kamionkowski:2008gj}; \\item The large amount of energetic positrons and electrons injected in our own Galaxy and in any dark matter structure in the Universe would yield secondary emission at all wave-lengths (for a recent comprehensive analysis see Ref.~\\cite{Bergstrom:2008ag}): this includes radio signals from synchrotron losses, X-rays and gamma-rays from inverse Compton scattering off the cosmic microwave background and other background radiation photons, as well as from bremsstrahlung, neutral pion decay and internal bremsstrahlung radiation \\cite{Beacom:2004pe}. Tight constraints exist at each one of these wavelengths, at all dark matter halo scales and distances, ranging from local dwarf galaxies to distant clusters (a very incomplete list of references on multi-wavelength constraints on dark matter annihilation includes e.g. \\cite{Totani:2004gy,Colafrancesco:2005ji, Profumo:2005xd,Profumo:2006hs,Colafrancesco:2006he,Profumo:2008fy,Jeltema:2008ax,Jeltema:2008hf,Jeltema:2008vu,PerezTorres:2008ug, Essig:2009jx, Natarajan:2007ja, Colafrancesco:2004sp, Colafrancesco:2009jq, Borriello:2008dt, Borriello:2008gy, Caceres:2008dr, Regis:2008ij, Hooper:2008zg, Taoso:2007qk, Hooper:2007gi}). \\end{enumerate} \\end{enumerate}  Recent studies, appeared after this manuscript was first released, discuss further possible explanations to the cosmic ray ``lepton anomalies''. These include: a gamma-ray burst \\cite{Ioka:2008cv}, inelastic collisions between high energy cosmic rays and background radiation photons in the astrophysical sources where cosmic rays are accelerated \\cite{Hu:2009bc, Diehl:2009kp}, inhomogeneities in the spatial distribution of sources of cosmic-ray leptons \\cite{Shaviv:2009bu}, hadronic processes producing secondary positrons in the very same sites where primary cosmic rays are accelerated \\cite{Blasi:2009hv, Blasi:2009bd} and s scenario where one (or more) recent suernova event(s) occurred in gas clouds close to the Solar System \\cite{Fujita:2009wk}.  Of all possibilities listed above, we focus in the present study on the one which we believe -- and intend to argue with the present analysis -- is the best motivated: nearby, known galactic pulsars. On the one hand, pulsars are predicted to produce $e^\\pm$ pairs with a spectrum that is suitable in principle to explain both the data on the positron fraction and, possibly, the overall electron-positron differential flux. On the other hand, pulsars are well known and well studied objects, and catalogues of galactic pulsars exist \\cite{2005AJ....129.1993M,2005yCat.7245....0M}, off of which we can compare theoretical speculations on the most plausible sources for the observed $e^\\pm$ fluxes with actual observational data.  Unlike in the dark matter hypothesis, the pulsar scenario does not require the introduction of an extra, exotic component to explain the data. This goes well with the celebrated quote of the 14th-century English logician and Franciscan friar, William of Ockham: ``{\\em Entia non sunt multiplicanda praeter necessitatem}'': there is no need to introduce extra entities beyond necessity.  The scope of the present analysis is to  ``dissect'' the PAMELA and ATIC data in the spirit of (or with) Occam's razor: we will pinpoint which regions of the pulsar parameter space (to be defined in the next section) are favored by experimental data; we will then turn to observations, and see if those parameter space regions correspond to existing astrophysical objects, and if those objects are predicted to have an $e^\\pm$ energy output in line with what is needed to explain the cosmic ray data. We shall argue that the recent positron and electron data are certainly very intriguing, but that they are {\\em not anomalous}, since a very natural interpretation exists in terms of known, existing astrophysical objects. We will also outline the major uncertainties that enter the process of reverse-engineering cosmic-ray lepton data to infer their possible origin. Among these, uncertainties in the pulsars' age and location at the point in time when the bulk of the high-energy lepton pairs where injected play an extremely relevant role.  In summary, highlights and novelties presented in this study include: \\begin{itemize} \\item accounting for all known, existing pulsars, as listed in the ATNF catalogue \\cite{2005AJ....129.1993M}: this approach allows us to (i) understand, qualitatively, which regions of the pulsar age-versus-distance parameter space contribute in given ranges of the cosmic ray electron-positron spectrum (see sec.\\ref{sec:all} and fig.~\\ref{fig:distant} and \\ref{fig:distageall}), (ii) analyze, through the determination of the relevant age-distance range, where the pulsar catalogue incompleteness can affect predictions for the high-energy cosmic-ray electron-positron spectrum, and thus assess where and how newly discovered gamma-ray pulsars can impact our understanding of the origin of high-energy cosmic-ray electrons (see e.g. sec.\\ref{sec:glast} and fig. \\ref{fig:CTA1}), (iii) address how the hereby proposed ``reverse-engineering'' approach (and in general any attempt at pinpointing a local high-energy electron-positron source as the main contributor to the leptonic cosmic-ray fluxes measured at Earth) is affected by uncertainties in the determination of the pulsars' age and distance, due to e.g. to trapping of the $e^\\pm$ in the pulsar wind nebula, and to pulsar kick velocities (see fig.~\\ref{fig:bestfit_distage}), (iv) estimate, by considering all ATNF pulsars, the theoretical uncertainties built in several current ``benchmark'' models for the $e^\\pm$ energy injection from pulsars, such as those described in the next section (see fig.~\\ref{fig:output} and \\ref{fig:modelsagedist}); \\item systematically analyzing and quantifying the role of uncertainties related to cosmic-ray energy losses and diffusion modeling (see e.g. fig.\\ref{fig:single_hard}-\\ref{fig:pamela}), and quantifying how these uncertainties impact the identification of pulsars that can source the positron fraction data and/or the $e^\\pm$ spectral features; \\item pointing out that there are regions of the pulsar age-distance parameter space that are more likely to contribute to the detected lepton anomalies (see e.g. fig.\\ref{fig:atic} and tab.~\\ref{tab:combination}). This point is relevant to the possibility of carrying out anisotropy studies, since it provides e.g. target directions for the possible detection of a dipolar asymmetry in the CRE arrival directions. In addition, as we point out with fig.\\ref{fig:glast}, Fermi data may soon tell us whether one pulsar or a combination of pulsars are at the origin of the high-energy end of the cosmic-ray electron-positron spectrum, and has the potential to actually detect the mentioned anisotropies; \\item theoretically studying the role of leaky-box boundary conditions  to the widely-used analytical expression \\cite{1995PhRvD..52.3265A} employed to calculate the spectrum, at Earth, of cosmic-ray electrons and positrons injected by a pulsar in the burst-like approximation; this physically means that we provide an estimate of the fraction of electrons and positrons injected by a given pulsar at a given point in time that freely escapes the galactic diffusive region, as a function of the pulsar and of the diffusion model parameters (see the last few paragraphs of sec.~\\ref{sec:models} and fig.~\\ref{fig:spherical}). \\end{itemize}  It goes without saying that the quest for the fundamental nature of particle dark matter, in this case indirect detection, needs to undergo the perhaps painful, or prosaic, process of deep scrutiny of any ``mundane'', standard interpretation of experimental data, including known astrophysical processes. The present study is a modest attempt in this direction. An important point we wish make at the end of this study is that the Fermi gamma-ray space telescope will play a decisive role, in the very near future, to elucidate the nature of the $e^\\pm$ sources, as well as to inform us on what particle dark matter can or cannot be \\cite{Baltz:2008wd}.  The outline of this study is as follows: in the next section we describe the processes that can seed $e^\\pm$ pair production in pulsars, estimate the $e^\\pm$ energy output and injection spectrum, and discuss the subsequent particle propagation, including a novel analytical formula to account for leaky-box boundary conditions in the propagation setup. Sec.~\\ref{sec:nearby} gives a few examples of nearby pulsars that could contribute to the $e^\\pm$ flux observed locally at the level of the measurements reported by ATIC and PAMELA. The following section addresses the question of whether those measurements can be explained by a {\\em single} nearby source. Sec.~\\ref{sec:pamela} and \\ref{sec:atic} are then devoted to outlining the pulsar parameter space favored by the PAMELA and ATIC data, respectively. Sec.~\\ref{sec:onemulti} contrasts a scenario where the measured $e^\\pm$ originate from a single source to one where they stem from multiple pulsars in well defined parameter space regions. We carry out a self-consistency check of our models in sec.~\\ref{sec:all}, where we compute the $e^\\pm$ summing over the full galactic pulsar catalogue. Finally, we outline in sec.~\\ref{sec:glast} the decisive role of the Fermi telescope in understanding the $e^\\pm$ data: (1) to provide an exquisitely detailed measurement of the $e^\\pm$ spectrum from 20 GeV to 1 TeV, and (2) to discover potentially relevant new gamma-ray pulsars, and , and (3) to search for anisotropies in the arrival direction of high-energy electrons and positrons. Sec.~\\ref{sec:concl} summarizes and concludes. ",
        "conclusions": "\\label{sec:concl}  We argued in this study that there is indeed no ``need to add entities beyond necessity'' to explain the PAMELA and the ATIC data on cosmic-ray electron and positron fluxes: existing and well-known pulsars easily and naturally (from an energy budget standpoint) account for both experimental results.  We summarize below the main results of the present theoretical study. \\begin{itemize} \\item Local pulsars are well-known sites of electron-positron pair production. Although a firm theoretical understanding of the pulsars' $e^\\pm$ energy output is lacking, we surveyed here a few possibilities, which are all approximately in line with one another and with estimates of the output energy in electrons from SNR's. Remarkably, by picking a few nearby known pulsars and employing the mentioned energy output models, we easily reproduce the spectral features and intensities reported by both PAMELA and ATIC, with nominal values for the $e^\\pm$ output efficiency; \\item The quality of the recent data on the positron fraction and on the differential $e^\\pm$ flux is good enough to allow the exploration of the regions of the pulsar parameter space favored by the two experimental observations; uncertainties in the determination of pulsars' ages and distances -- particularly, as far as the latter are concerned, stemming from the large observed pulsar velocities -- can however significantly compromise our ability to pinpoint single major contributors to the local cosmic-ray lepton flux \\item Although statistically and energetically plausible for existing objects, we find that the PAMELA and ATIC data are not well fitted by one single bright nearby source; \\item For a given $e^\\pm$ injection spectrum and diffusion setup, the PAMELA data point towards a corridor of correlated values of pulsar age and distance. Several known pulsars fall within this corridor, and could coherently add up to contribute to the positron excess over the standard secondary population observed by PAMELA; \\item The ATIC feature at around 600 GeV can also be explained by one or more existing pulsars, although this generically requires a rather high electron-positron energy output, and, unlike PAMELA, it does not seem to be a natural or necessary {\\em post}-diction of the pulsar contribution to energetic $e^\\pm$: while it is plausible to expect pulsars to produce a positron fraction like that measured by PAMELA, it is not as obvious, in the same context, to expect a spectral feature like that reported by ATIC. Nevertheless, relatively young (0.4-0.6 Myr) pulsars 1-2 kpc away can easily account for a spectral feature such as the one observed by ATIC. \\item The Fermi Space Telescope will play an extraordinarily important role in the understanding of the local flux of energetic electrons and positrons, by (1) providing us, in the very near future, with an exquisite measurement of the $e^\\pm$ flux, that could unveil several interesting features, and/or rule out recent experimental claims, and (2) by discovering gamma-ray pulsars that can add-up to give a self-consistent picture of the origin of the local energetic electrons and positrons, and (3) searching for anisotropies in the arrival direction of high-energy electrons and positrons. \\end{itemize} A convincing picture of the nature of energetic cosmic ray electrons and positrons evidently calls for a two-pronged effort: on the one hand, theorists working on cosmic rays and on compact objects should refine the theoretical understanding of $e^\\pm$ production in pulsars and possibly in other astrophysical galactic objects; on the other hand, further and deeper observations of candidate electron-positron sources at various wavelengths (including radio, X-ray, gamma-ray and very high energy gamma-ray frequencies) will be needed to complement and to drive the theoretical effort towards the construction of a tho\\-rough\\-ly satisfactory picture.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %\\newpage  %"
    },
    "0812/0812.3745_arXiv.txt": {
        "abstract": "{\\bf Abstract}. The Lorentz violating term in the photon sector of Standard Model Extension, $\\mathcal{L}_K = -\\mbox{$\\frac14$} (k_F)_{\\alpha \\beta \\mu \\nu} F^{\\alpha \\beta} F^{\\mu \\nu}$ (here referred to as the Kosteleck\\'{y} term), breaks conformal invariance of electromagnetism and enables a superadiabatic amplification of magnetic vacuum fluctuations during inflation. For a wide range of values of parameters defining Lorentz symmetry violation and inflation, the present-day magnetic field can have an intensity of order of nanogauss on megaparsec scales and then could explain the large-scale magnetization of the universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.2163_arXiv.txt": {
        "abstract": "We introduce a parameterized high-density equation of state (EOS) in order to systematize the study of constraints placed by astrophysical observations on the nature of neutron-star matter.  To obtain useful constraints, the number of parameters should be smaller than the number of neutron-star properties that have been measured or will have been measured in the next several years.  And the set must be large enough to accurately approximate the large set of candidate EOSs.  We find that a parameterized EOS based on piecewise polytropes with 3 free parameters matches to about 4\\% rms error an extensive set of candidate EOSs at densities below the central density of $1.4\\, M_\\odot$ stars. Adding observations of more massive stars constrains the higher density part of the EOS and requires an additional parameter. We obtain constraints on the allowed parameter space set by causality and by present and near-future astronomical observations. In particular, we emphasize potentially stringent constraints on the EOS parameter space associated with two measured properties of a single star;  and we find that a measurement of the moment of inertia of PSR~J0737-3039A can strongly constrain the maximum neutron-star mass. We also present in an appendix a more efficient algorithm than has previously been used for finding points of marginal stability and the maximum angular velocity of stable stars. ",
        "introduction": "Because the temperature of neutron stars is far below the Fermi energy of their constituent particles, neutron-star matter is accurately described by the one-parameter equation of state (EOS) that governs cold matter above nuclear density.  The uncertainty in that EOS, however, is notoriously large, with the pressure $p$ as a function of baryon mass density $\\rho$ uncertain by as much as an order of magnitude above nuclear density.  The phase of the matter in the core of a neutron star is similarly uncertain: Current candidates for the EOS include non-relativistic and relativistic mean-field models; models for which neutron-star cores are dominated by nucleons, by hyperons, by pion or kaon condensates, and by strange quark matter (free up, down, and strange quarks); and one cannot yet rule out the possibility that the ground state of cold matter at zero pressure might be strange quark matter and that the term ``neutron star'' is a misnomer for strange quark stars.  The correspondingly large number of fundamental parameters needed to accommodate the models' Lagrangians has meant that studies of astrophysical constraints (see, for example, \\cite{EngvikOsnesHjorthJensen1996constr, Lattimer:2000nx, lattimerprakash2006nuc, KlahnBlaschkeTypel2006constraints, PageReddy2006review} and references therein) present constraints by dividing the EOS candidates into an allowed and a ruled-out list.  A more systematic study, in which astrophysical constraints are described as constraints on the parameter space of a parameterized EOS, can be effective only if the number of parameters is smaller than the number of neutron-star properties that have been measured or will have been measured in the next several years.  At the same time, the number of parameters must be large enough to accurately approximate the EOS candidates.  A principal aim of this paper is to show that, if one uses phenomenological rather than fundamental parameters, one can obtain a parameterized EOS that meets these conditions.  We exhibit a parameterized EOS, based on specifying the stiffness of the star in three density intervals, measured by the adiabatic index $\\Gamma = d\\log P/d\\log\\rho$.  A fourth parameter translates the $p(\\rho)$ curve up or down, adding a constant pressure -- equivalently fixing the pressure at the endpoint of the first density interval.  Finally, the EOS is matched below nuclear density to the (presumed known) EOS. An EOS for which $\\Gamma$ is constant is a polytrope, and the parameterized EOS is then piecewise polytropic.  A similar piecewise-polytropic EOS was previously considered by Vuille and Ipser \\cite{VuilleIpser1999max}; and, with different motivation, several other authors \\cite{ZdunikBejgerHaensel2006phase, BejgerHaenselZdunik2005phase, HaenselPotekhin2004anal, shibata2005es} have used piecewise polytropes to approximate neutron-star EOS candidates. In contrast to this previous work, we use a small number of parameters chosen to fit a wide variety of fundamental EOSs, and we systematically explore a variety of astrophysical constraints. Like most of the previous work, we aim to model equations of state containing nuclear matter (possibly with various phase transitions) rather than pure quark stars, whose EOS is predicted to be substantially different.  As we have noted, enough uncertainty remains in the pressure at nuclear density, that one cannot simply match to a fiducial pressure at $\\rho_{\\mathrm{nuc}}$.  Instead of taking as one parameter the pressure at a fiducial density, however, one could match to the pressure of the known subnuclear EOS at, say, 0.1 $\\rho_{\\mathrm{nuc}}$ and then use as one parameter a value of $\\Gamma_0$ for the interval between $0.1\\rho_{\\mathrm{nuc}}$ and $\\rho_{\\mathrm{nuc}}$.  Neutron-star observables are insensitive to the EOS below $\\rho_{\\mathrm{nuc}}$, because the fraction of mass at low density is small.  But the new parameter $\\Gamma_0$ would indirectly affect observables by changing the value of the pressure at and above nuclear density, for fixed values of the remaining $\\Gamma_i$.  By choosing instead the pressure at a fixed density $\\rho_1 > \\rho_{\\mathrm{nuc}}$, we obtain a parameter more directly connected to physical observables.  In particular, as Lattimer and Prakash~\\cite{Lattimer:2000nx} have pointed out, neutron-star radii are closely tied to the pressure somewhat above nuclear density, and the choice $p_1 = p(\\rho_1)$ is recommended by that relation.  In general, to specify a piecewise polytropic EOS with three density intervals above nuclear density, one needs six parameters: two dividing densities, three adiabatic indices $\\Gamma_i$, and a value of the pressure at an endpoint of one of the intervals.  Remarkably, however, we find (in Sec.~\\ref{sec:fits}) that the error in fitting the collection of EOS candidates has a clear minimum for a particular choice of dividing densities.  With that choice, the parametrized EOS has three free parameters, $\\Gamma_1, \\Gamma_2$ and $p_1$, for densities below $10^{15}$~g/cm$^3$ (the density range most relevant for masses $\\sim 1.4 M_\\odot$), and four free parameters (an additional $\\Gamma_3$) for densities between $10^{15}$~g/cm$^3$ and the central density of the maximum mass star for each EOS.  With the parameterization in hand, we examine in Sect.~\\ref{sec:constraints} astrophysical constraints on the parameter space beyond the radius-$p_1$ relation found by Lattimer and Prakash~\\cite{Lattimer:2000nx}.  Our emphasis in this first study is on present and very near-future constraints:  those associated with the largest observed neutron-star mass and spin, with a possible observation (as yet unrepeated) of neutron-star redshift, with a possible simultaneous measurement of mass and radius, and with the expected future measurement of the moment of inertia of a neutron star with known mass. A companion paper~\\cite{nrda} will investigate constraints obtainable with gravitational-wave observations in a few years. The constraints associated with  the largest observed mass, spin, and redshift have a similar form, each restricting  the parameter space to one side of a surface: For example, if we take the largest  observed mass (at a 90\\% confidence level) to be 1.7 $M_\\odot$, then the allowed  parameters correspond to EOSs whose maximum mass is at least 1.7 $M_\\odot$.  We can regard $M_{\\rm max}$ as a function on the 4-dimensional EOS parameter space. The subspace of EOSs for which $M_{\\rm max} = 1.7 M_\\odot$ is then described by a  3-dimensional surface, and constraint is a restriction to the high-mass side of the surface. Similarly, the observation of a 716 Hz pulsar restricts the EOS parameter space to one side of a surface that describes EOSs for which the maximum spin is  716 Hz. Thus we can produce model-independent extended versions of the multidimensional constraints seen in~\\cite{Lackey:2005tk}.  The potential simultaneous observation of two properties of a single neutron star (for example, moment of inertia and mass) would yield a significantly stronger constraint: It would restrict the parameter space not to one side of a surface but to the surface itself. And a subsequent observation of two different parameters for a different neutron star would then restrict one to the intersection of two surfaces. We exhibit the  result of simultaneous observations of mass and moment of inertia (expected within  the next decade for one member of the binary pulsar J0737-3039~\\cite{Lattimer:2004nj, bejger-bulik-haensel}) and  of mass and radius.  Gravitational-wave observations of binary inspiral can again  measure two properties of a single star: both mass and an observable roughly described as the final frequency before plunge (the departure of the waveform from a point-particle inspiral); and related work in progress examines the accuracy with which one can extract EOS parameters from interferometric observations of gravitational waves in the inspiral and merger of a binary neutron star system \\cite{nrda}.  {\\sl Conventions}:  We use cgs units, denoting rest-mass density by $\\rho$, and (energy density)/$c^2$ by $\\epsilon$.  We define rest-mass density as $\\rho=m_{\\rm B}n$ where $m_{\\rm B}=1.66\\times 10^{-24}$~g and $n$ is the baryon number density. In Sec.~\\ref{sec:ppdef}, however, we set $c=1$ to simplify the equations and add a footnote on restoring $c$. ",
        "conclusions": "We have shown how one can use a parameterized piecewise polytropic EOS to systematize the study of  observational constraints on the EOS of cold, high-density matter.  We think that our choice of a 4-parameter EOS strikes an appropriate balance  between the accuracy of approximation that a larger number of parameters would  provide and the number of observational parameters that have been measured or are  likely to be measured in the next several years.   The simple choice of a piecewise  polytrope, with discontinuities in the polytropic index, leads to suitable accuracy  in approximating global features of a star.  But the discontinuity reduces the expected accuracy with which the parameterized EOS can approximate the local speed of  sound.  One can largely overcome the problem by using a minor modification of the parameterized EOS in which a fixed smoothing function near each dividing density is used to join the two polytropes.  We see that high-mass neutron stars are likely to provide the strongest constraints from a single measurement.  The work dramatizes the significantly more stringent constraints associated with measurements like this, if two (or more) physical features of the same star  can be measured, and an $n$-dimensional parameter space is reduced  by one (or more) dimension(s), to within the error of measurement.  In particular, a moment of inertia measuremement for PSR~J0737-3039 (whose mass is already precisely known) could strongly constrain the maximum neutron star mass.  The effect of EOS-dependent tidal deformation can modify the gravitational waves produced by inspiraling neutron stars. This modification is largely dependent on the radius of the neutron star.   Flanagan and Hinderer \\cite{FlanaganHinderer2007} investigate constraints on an EOS-dependent tidal parameter, the Love number, from observations of early inspiral. A companion to this paper~\\cite{nrda} uses the parametrized EOS in numerical simulations to examine the future constraint associated with expected gravitational-wave observations of late inspiral in binary neutron stars.  Finally, we note that the constraints from observations of different neutron star populations constrain different density regions of the EOS. For moderate mass stars such as those found in binary pulsar systems, the EOS above $\\rho_2= 10^{15.0}$~g/cm$^3$ is unimportant. For near-maximum mass stars, the EOS below $\\rho_1=10^{14.7}$~g/cm$^3$ has little effect on neutron star properties.  This general behavior is independent of the details of our parameterization."
    },
    "0812/0812.2480_arXiv.txt": {
        "abstract": "*{ In a series of papers we have recently studied the clustering of LRG galaxies in the latest spectroscopic SDSS data release,  which has 75000 LRG galaxies sampling 1.1 ${\\,\\rm Gpc^3/h^3}$ to z=0.47. Here we focus on detecting a local maxima shaped as a circular ring  in the bidimensional galaxy correlation function $\\xisp$, separated in perpendicular $\\sigma$ and line-of-sight $\\pi$ distances. We find a significant detection of such a peak at $r\\simeq 110$Mpc/h. The overall shape and location of the ring is consistent with it originating from the recombination-epoch baryon acoustic oscillations (BAO). This agreement provides support for the current understanding of how large scale structure forms in the universe.  We study the significance of such feature using large mock galaxy simulations  to provide  accurate errorbars.}  \\abstract{ In a series of papers we have recently studied the clustering of LRG galaxies in the latest spectroscopic SDSS data release,  which has 75000 LRG galaxies sampling 1.1 ${\\,\\rm Gpc^3/h^3}$ to z=0.47. Here we focus on detecting a local maxima shaped as a circular ring  in the bidimensional galaxy correlation function $\\xisp$, separated in perpendicular $\\sigma$ and line-of-sight $\\pi$ distances. We find a significant detection of such a peak at $r\\simeq 110$Mpc/h. The overall shape and location of the ring is consistent with it originating from the recombination-epoch baryon acoustic oscillations (BAO). This agreement provides support for the current understanding of how large scale structure forms in the universe.  We study the significance of such feature using large mock galaxy simulations  to provide  accurate errorbars.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.5096_arXiv.txt": {
        "abstract": "In this publication we investigate dynamics of a flat FRW cosmological model with a non-minimally coupled scalar field with the coupling term $\\xi R \\psi^{2}$ in the scalar field action. The quadratic potential function $V(\\psi)\\propto \\psi^{2}$ is assumed. All the evolutional paths are visualized and classified in the phase plane, at which the parameter of non-minimal coupling $\\xi$ plays the role of a control parameter. The fragility of global dynamics with respect to changes of the coupling constant is studied in details. We find that the future big rip singularity appearing in the phantom scalar field cosmological models can be avoided due to non-minimal coupling constant effects. We have shown the existence of a finite scale factor singular point (future or past) where the Hubble function as well as its first cosmological time derivative diverge. ",
        "introduction": "\\label{sec:1} Recently scalar fields have played a very important role in cosmology. They are used in many phenomenological models like quintessence \\cite{Wetterich:1987fm, Ratra:1987rm}. Scalar fields are also very important in description of dynamics in the loop quantum cosmology, which base on the background independent theory without the canonical notion of time. In this theory one scalar field is chosen as an internal clock for other fields \\cite{Ashtekar:2006rx}. The scalar fields with a potential function are also very important in modelling of inflation. For example a scalar field with the simplest quadratic potential function was assumed in Linde's conception of chaotic inflation \\cite{Linde:1983gd}. The 5 years of WMAP observations \\cite{Komatsu:2008hk} rejected many inflationary scenarios (potential functions $V(\\psi)$), while models with a simple quadratic potential are admitted on the $1\\sigma$ confidence level.  In the standard quintessence energy density of the minimally coupled to gravity scalar field mimics the effective cosmological constant. Of course the detailed evolution is dependent upon a specific form of the potential $V(\\psi)$ but the $\\psi^{2}$ contribution can be treated as a leading order term of expansion of the potential function.  We extend the quintessence scenario by incorporating the non-minimal coupling constant \\cite{Faraoni:2006ik, Faraoni:2000gx, Hrycyna:2007mq, Hrycyna:2007gd, Szydlowski:2008zza}. In this paper we present a phase space analysis of the evolution of a spatially flat Friedmann-Robertson-Walker (FRW) universe containing a non-minimally coupled to gravity scalar field, both canonical and phantom, with the simplest form of quadratic potential function. The similar analysis for the case of minimally coupled scalar field was performed in \\cite{UrenaLopez:2007vz}. We extend this analysis on models with a non-zero coupling constant $\\xi$ which plays the role of a control parameter for an autonomous dynamical system on the phase plane. Therefore the location of fixed points (physically representing asymptotic states of the system) as well as their character depends upon the value of $\\xi$. The values of parameter $\\xi$ for which the global dynamics changes dramatically are called bifurcation values. In our previous paper \\cite{Szydlowski:2006qn}, in case of conformal coupling and quadratic potential function we have shown that phantom cosmology can be treated as a simple model with a scattering of trajectories whose character depends crucially on the sign of the potential function. We also demonstrated that there is a possibility of chaotic behavior in the flat Universe with a conformally coupled phantom field in the system considered on the non-zero energy level.  The minimally coupled scalar field endowed with a quadratic potential function has a strong motivation in inflationary models and its generalizations with a simple non-minimal coupling term $\\xi R \\psi^{2}$ have been studied \\cite{Park:2008hz} in the context of the origin of the canonical inflaton field itself. The physical motivation to investigate a non-minimally coupled scalar field cosmological models could be possible application of this models to inflationary cosmology or to the present dark energy (see for example \\cite{Spokoiny:1984bd, Salopek:1988qh, Fakir:1992cg, Barvinsky:1994hx, Barvinsky:1998rn, Uzan:1999ch, Chiba:1999wt, Amendola:1999qq, Perrotta:1999am, Holden:1999hm, Bartolo:1999sq, Boisseau:2000pr, Gannouji:2006jm, Carloni:2007eu, Bezrukov:2007ep, Barvinsky:2008ia}).  The coupling constant $\\xi$ is a free parameter of the theory which should be estimated from the observational data. Recently we have shown that distant supernovae can be used to estimate the value of this parameter \\cite{Szydlowski:2008zza}.  The main advantage of dynamical system analysis is that we can visualize all the trajectories of the system admissible for all initial conditions. Therefore one can classify generic routes to the accelerating phase (the de Sitter attractor where $p_{\\psi}=-\\rho_{\\psi}$). This attractor corresponds to the model with the cosmological constant.  The paper has a following organization: in section \\ref{sec:2} we reduce dynamics to the form of an autonomous dynamical system which describes both canonical and phantom scalar field models. Section \\ref{sec:3} is devoted to a detailed analysis of the phase portraits for different values of the parameter $\\xi$. In this section we also discuss the change of evolutionary scenarios upon the value of parameter $\\xi$. ",
        "conclusions": "We study the dynamics of a scalar field with a simple quadratic potential function and non-minimally coupled to the gravity via $\\xi R \\psi^{2}$ term, where $R$ is the Ricci scalar of the Robertson-Walker spacetime. We reduce dynamical problem to the autonomous dynamical system on the phase plane in the variables $\\psi$ and its derivative with respect to the natural logarithm of the scale factor. The constraint condition is solved in such a way that the final dynamical system is free from the constraint and is defined on the plane. We investigate the whole dynamics at the finite domain and at infinity. All the trajectories for all admissible initial conditions are classified, and critical points representing the asymptotic states (stationary solutions) are found. We explore generic evolutionary paths to find the stable de Sitter state as a global attractor and classify typical routes to this point. We study the effects of the canonical scalar field as well as the phantom scalar field. The following conclusions, as the results of our studies, can be drawn: \\begin{itemize} \\item[1.]{The cosmological models with the quadratic potential function and non-minimal coupling term $\\xi R \\psi^{2}$ are represented by a 2-dimensional autonomous dynamical system which is studied in details on the phase plane $(\\psi,\\psi')$ (see Table \\ref{tab:1} for the critical points at the finite domain). The shape of the physical region $H^{2}\\ge0$ does not depend on the form of the potential function, but only on the value of the coupling constant $\\xi$;} \\item[2.]{We investigated the fixed points of the dynamical system and their stability to find the generic evolutional scenarios. We have shown the existence of a finite scale factor singular point (both future and past) where the Hubble function as well as its first cosmological time derivative diverge $H^{2}=\\infty$ and $\\dot{H}=\\infty$;} \\item[3.]{For the phantom scalar field $\\ve=-1$ we found existence of a sink type critical point (i.e. a stable node or a focus, depending on the value of the parameter $\\xi$) which represents the de Sitter solution. Only for $\\xi<1/6$ the evolutional paths avoid a past finite scale factor singularity (see for example Figs~\\ref{fig:6} and \\ref{fig:8});} \\item[4.]{For the canonical scalar field $\\ve=1$ we found that for $0<\\xi<1/6$ exists the critical point which corresponds to a past finite scale factor singularity for models with $V(\\psi)=\\frac{1}{2}m^{2}\\psi^{2}>0$. If $m^{2}<0$ then this point corresponds to a future finite scale factor singularity (see Fig.~\\ref{fig:2}).} \\end{itemize}"
    },
    "0812/0812.2814_arXiv.txt": {
        "abstract": "We report the discovery of possible overdensities of galaxies at $z \\sim 3.7$ in $Chandra$ Deep Field South (CDF-S). These overdensities are identified from a photometric redshift-selected sample, and the $BVz$-selected sample. One overdensity is identified in the proximity of two active galactic nuclei and Lyman break galaxies at $z=3.66$ and $z=3.70$ at 7$\\sigma$ significance level. The other overdensity is less significant. It is identified around six $z_{spec} \\simeq 3.6$ galaxies at 3$\\sigma$ significance level. The line-of-sight velocity dispersions of these overdensities are found to be $\\sigma_{v} \\simeq 500 - 800$ km s$^{-1}$, comparable to the velocity dispersions of clusters of galaxies today. Through spectral energy distribution fitting, we find $\\sim$ 15 massive galaxies with $M \\gtrsim 10^{11}\\,M_{\\odot}$ around the $z \\simeq 3.7$ overdensity. The mass of the $z \\simeq 3.7$ overdensity is found to be a few $\\times 10^{14}\\,M_{\\odot}$. Our result suggests that high-redshift overdense regions can be found in a supposedly blank field, and that the emergence of massive structures can be traced back to redshifts as high as $z \\sim 3.7$. ",
        "introduction": "In hierarchical galaxy formation, galaxy clusters grow through gravitational attraction of matter around high $\\sigma$ peaks, and they are predicted to be rare at high redshift. Therefore, constraining the number density of high-redshift clusters is an important method of testing hierarchical models (Evrard et al. 2002; Mantz et al. 2008).  Many previous studies have been carried out to search for evolved structures at high redshift. Radio galaxies and active galactic nuclei (AGNs) are considered to be signposts for high-redshift proto-clusters, and several studies have found proto-clusters at $2 \\lesssim z \\lesssim 4$ \\citep{lef96,kur00,pen00,wol03,kaj06,kod07,ven07}. Overdense regions are also identified out to $z \\sim 6$ using galaxies selected with the Lyman break technique or narrowband imaging (e.g, Ouchi et al. 2005).  When identifying overdense regions, it is necessary to compare the overdensity with a control field which lacks notable high-redshift overdense regions. For such a purpose, many studies have used the $Chandra$ Deep Field South (CDF-S) due to the wealth of the deep multiwavelength data over a moderately large field of view \\citep{kaj06,kod07}. However, a matter of concern is whether the CDF-S can really be considered as a field devoid of high-redshift overdense regions. There exist a fair number of high-redshift X-ray-detected AGNs in the CDF-S, which may harbor high-redshift clusters.  Motivated by this, we searched for signs of overdensity around two AGNs at $z=3.66$ and $3.70$ separated by $2\\arcmin.7$ in the CDF-S, and we report the discovery of possible overdense regions at $z \\sim$ $3.7$ in the CDF-S.  Throughout this Letter, we assume a cosmology with nonzero cosmological constant (Im et al. 1997), $\\Omega_m = 0.3$, $\\Omega_\\Lambda = 0.7$ and $h = 0.7$, consistent with the $WMAP$ cosmology \\citep{spe07}. All magnitudes are given in the AB system. ",
        "conclusions": "By examining the two-dimensional distribution of $3.45 \\lesssim z \\lesssim 3.85$ candidates in the CDF-S field selected from the photometric redshift method or the $BVz$ color-color space, we find plausible associations of the overdense region of such objects with AGNs and LBGs at $z_{spec} \\sim 3.7$. The significance of the overdense regions is found to be $5-7\\sigma$. The overdense area is abundant with massive galaxies, adding further support that this is a proto-cluster. The derived mass of the proto-cluster is found to be a few $\\times$ $10^{14}$ $M_{\\odot}$. The existence of the massive proto-cluster in the CDF-S shows that one must be cautious when examining the significance of high redshift overdensities in other fields with respect to the CDF-S."
    },
    "0812/0812.0811_arXiv.txt": {
        "abstract": "In the standard Friedmann--Lema\\^{\\i}tre--Robertson--Walker (FLRW) approach to model the Universe the violation of the so-called energy conditions is related to some important properties of the Universe as, for example, the current and the inflationary accelerating expansion phases. The energy conditions are also necessary in the formulation and proofs of Hawking-Penrose singularity theorems. In two recent articles  we have derived bounds from energy conditions and made confrontations of these bounds with supernovae data. Here, we extend these results in following way: first, by using our most recent statistical procedure for calculating new $q(z)$ estimates from the \\emph{gold} and \\emph{combined} type Ia supernovae samples; second, we use these estimates to obtain a new picture of the energy conditions fulfillment and violation for the recent past ($z\\leq 1 $) in the context of the standard cosmology. ",
        "introduction": "In the absence of constraints on the energy-momentum tensor $T_{\\mu\\nu}$ any metric satisfies Einstein's equations since they can be regarded as a definition of $T_{\\mu\\nu}\\,$, i.e., a set of equations determining $T_{\\mu\\nu}$ for any given metric $g_{\\mu \\nu}\\,$. However, if one wishes to explore general properties that hold for a variety of different physical sources it is convenient to impose the so-called \\emph{energy conditions} that limit the arbitrariness of $T_{\\mu\\nu}$ on physical grounds.\\cite{EC-basics_refs}  On scales relevant for cosmology, an important point in the  study of the energy conditions is the confrontation of their predictions with the observational data. By using model-independent energy-condition \\emph{integrated} bounds on the cosmological observables as, for example, the distance modulus and lookback time, this confrontation has been made in some recent articles\\cite{Santos2006}\\cdash\\cite{CattoenVisser} (see also the pioneering Refs.~\\refcite{M_Visser1997}  by  Visser). In Ref.~\\refcite{Lima2008a}, however,  it was shown that the fulfillment (or the violation) of these \\emph{integrated} bounds at a specific redshift $z$ is not a sufficient (nor a necessary) \\emph{local} condition for the fulfillment (or respectively the violation) of the associated energy condition at $z$.% \\footnote{Energy conditions constraints on the so-called $f(R)$--gravity have also been investigated in Ref.~\\refcite{Santiago2006} and more recently in Refs.~\\refcite{SARC2007} and \\refcite{SRA2008}.} In this way, the confrontation between the prediction of these \\emph{integrated} bounds and observational data cannot be used to draw conclusions on the fulfillment (or violation) of the energy conditions at $z$. In Ref.~\\refcite{Lima2008a} this problem was overcome by deriving new \\emph{non-integrated} energy-condition bounds, and  confrontations between the new bounds with type Ia supernovae (SNe Ia) data of the \\emph{gold}\\cite{Riess2007} and  \\emph{combined}\\cite{Combined} samples were performed by using the  upper and lower limits of confidence regions on $\\,E(z) - q(z)\\,$ plane. More recently, in Ref.~\\refcite{Lima2008b} a new statistical way for estimating the deceleration parameter $\\,q(z)\\,$ was carried out, and a new picture of the energy conditions fulfillment and violation for recent past ($z\\leq1 $) was calculated by using the recently compiled \\emph{Union} sample.\\cite{Kowalski2008}  In this work, we use the most recent statistical procedure introduced in Ref.~\\refcite{Lima2008b} along with the \\emph{gold}\\cite{Riess2007} and  \\emph{combined}\\cite{Combined} samples to obtain estimates of $q(z)$ in order to build up a new picture of the confrontation between the energy condition \\emph{integrated} bounds and these SNe Ia data sets, completing therefore the cycle of this type of analysis which involves these three samples and the  two statistical procedures to $q(z)$ estimates of Refs.\\refcite{Lima2008a} and \\refcite{Lima2008b}. ",
        "conclusions": "Since in the flat case the energy condition bounds given by Eqs.~\\eqref{eq:nec-q(z)}, \\eqref{eq:sec-q(z)} and \\eqref{eq:dec-q(z)} depend only on $q(z)$, we have obtained the $q(z_{\\star})$ estimates at $1\\sigma - 3\\sigma$ confidence levels from \\emph{gold} and \\emph{combined} SNe Ia samples by marginalizing over $E(z_\\star)$ and the other parameters ($q'_l$ 's) of the $q(z)$ function [Eq.(\\ref{eq:q_z})].% \\footnote{We note that this statistical approach has been previously used in Ref.~\\refcite{Lima2008b} but for the SNe Ia \\emph{Union} sample.\\cite{Kowalski2008}}  A global picture of the breakdown and fulfillment of the energy conditions in the recent past has been built up with the $q(z_\\star)$ estimates at 200 equally spaced redshifts in the interval $(0,1]$. Fig.~\\ref{qxz}(a) shows the \\textbf{NEC}, \\textbf{SEC}, and \\textbf{DEC} bounds along with the best-fit values and the $1\\sigma$, $2\\sigma$ and $3\\sigma$ limits of $q(z_\\star)$ in the $q(z) - z$ plane. We recall that \\textbf{WEC} bound [$(E^2(z) \\geq 0)$] is fulfilled identically.  \\begin{figure*}[h] \\includegraphics[scale=0.58]{Lima_Fig1a.eps} \\includegraphics[scale=0.58]{Lima_Fig1b.eps} \\caption{The best-fit, the upper and lower $1\\sigma$, $2\\sigma$ and $3\\sigma$ limits of $q(z_\\star)$ estimates, obtained with the \\emph{gold} [panel (a)] and the \\emph{combined} [panel (b)] samples, for $200$ equally spaced redshifts. The \\textbf{NEC} and \\textbf{SEC} lower bounds, and also the \\textbf{DEC} upper bound are shown. This figure shows that the \\textbf{SEC} is violated with $1\\sigma$ confidence level from $\\simeq 0$ until $z \\simeq 0.31$ [\\emph{gold} sample, panel (a)], and  until $z \\simeq 0.52$ [\\emph{combined} sample, panel (b)]. It also shows that the \\textbf{DEC} and \\textbf{NEC} is violated within $3\\sigma$ confidence level for high redshifts for both supernovae samples, and that the \\textbf{NEC} is violated for $z \\lesssim 0.105$ [panel (a)] and $z \\lesssim 0.085$ [panel (b)]. \\label{qxz}} \\end{figure*}  In Fig.~\\ref{qxz} it is showed that the \\textbf{SEC} bound is violated with $1\\sigma$ confidence level until $z = 0.31$ for \\emph{gold} and $z = 0.52$ for \\emph{combined} sample, while in the redshift intervals $(0.09, 0.17)$ [panel (a)] and $(0.08, 0.18)$ [panel (b)] this violation occurs with more than $3\\sigma$ confidence level, where the highest evidence is found at $z = 0.135$ with $\\simeq 3.86\\sigma$ [\\emph{gold}, panel (a)]  and  $\\simeq 4.28\\sigma$ [\\emph{combined}, panel (b)]. Unlike the result of Ref.~\\refcite{Lima2008a}, wherein these analyses have been performed by  computing the confidence regions on the $E(z_{\\star}) - q(z_{\\star})$ plane, revealing no redshift value for the \\textbf{SEC} fulfillment with at least $1\\sigma$, we note here the \\textbf{SEC} is fulfilled with more than $1\\sigma$ for $z \\gtrsim 0.615$ [\\emph{gold}, panel (a)] and $z \\gtrsim 0.855$ [\\emph{combined}, panel (b)]. According to the present \\textbf{SEC} analysis, with $1\\sigma$ confidence level, the universe crosses over from a decelerated expansion phase to an accelerated expansion during the redshift interval $(\\simeq 0.31, \\simeq 0.615)$ for \\emph{gold} and $(\\simeq 0.52, \\simeq 0.855)$ for \\emph{combined} sample.% \\footnote{We recall that in a similar \\textbf{SEC} analysis of Ref.~\\refcite{Lima2008b} performed by using the \\emph{Union} sample, the deceleration to acceleration transition expansion phase of the universe took place in the redshift interval ($\\simeq 0.4, \\simeq 0.64\\,$).}  Regarding the \\textbf{NEC}, Fig.~\\ref{qxz} indicates its breakdown within $3\\sigma$ confidence level for low redshift, $z \\lesssim 0.105$ [panel (a)] and $z \\lesssim 0.085$ [panel (b)]. For higher values of redshift, \\textbf{NEC} is violated within $3\\sigma$ at $z \\gtrsim 0.94$ for \\emph{gold} and $z \\gtrsim 0.96$ for \\emph{combined} sample.  Concerning the \\textbf{DEC}, Fig.~\\ref{qxz} indicates that it is violated within $3\\sigma$ for $z \\gtrsim 0.795$ [\\emph{gold}, panel (a)] and $z \\gtrsim 0.83$ [\\emph{combined}, panel (b)], which are intervals where the error in the estimates of $q(z)$ grow significantly, though. Finally, we note that the \\textbf{DEC} violation of the present analysis is weaker than that obtained in Ref.~\\refcite{Lima2008a} in the sense that, differently from that analysis, now the \\textbf{DEC} is fulfilled with $1\\sigma$ confidence level in the entire redshift interval for both samples."
    },
    "0812/0812.2025_arXiv.txt": {
        "abstract": "We present a pair of high-resolution smoothed particle hydrodynamics (SPH) simulations that explore the evolution and cooling behavior of hot gas around Milky-Way size galaxies.  The simulations contain the same total baryonic mass and are identical other than their initial gas density distributions.  The first is initialised with a {\\em low entropy} hot gas halo that traces the cuspy profile of the dark matter, and the second is initialised with a {\\em high-entropy} hot halo with a cored density profile as might be expected in models with pre-heating feedback.  Galaxy formation proceeds in dramatically different fashion depending on the initial setup. While the low-entropy halo cools rapidly, primarily from the central region, the high-entropy halo is quasi-stable for $\\sim 4$ Gyr and eventually cools via the fragmentation and infall of clouds from $\\sim 100$ kpc distances. The low-entropy halo's X-ray surface brightness is $\\sim 100$ times brighter than current limits and the resultant disc galaxy contains more than half of the system's baryons.  The high-entropy halo has an X-ray brightness that is in line with observations, an extended distribution of pressure-confined clouds reminiscent of observed populations, and a final disc galaxy that has half the mass and $\\sim 50\\%$ more specific angular momentum than the disc formed in the low-entropy simulation.  The final high-entropy system retains the majority of its baryons in a low-density hot halo.  The hot halo harbours a trace population of cool, mostly ionised, pressure-confined clouds that contain $\\sim 10 \\%$ of the halo's baryons after 10 Gyr of cooling.  The covering fraction for \\hi and \\mg absorption clouds in the high-entropy halo is $\\sim 0.4$ and $\\sim 0.6$, respectively, although most of the mass that fuels disc growth is ionised, and hence would be under counted in \\hi surveys. ",
        "introduction": "\\label{sec:intro} It is well known that in the absence of feedback the majority of baryons in galaxy-size dark matter haloes ($M \\sim 10^{12}$ M$_\\odot$) should have cooled into halo centres over a Hubble time (e.g. White \\& Rees 1978; Katz 1992; Benson et al. 2003).  In contrast, only $\\sim 20 \\%$ of the associated baryons in Milky-Way size haloes are observed to be in a cold, collapsed form (Maller \\& Bullock 2004 (MB04); Mo et al. 2005; Fukugita \\& Peebles 2006; Nicastro et al. 2008).  An understanding of the feedback processes that act to solve this galaxy overcooling problem is a major goal of galaxy formation today.  It is not known if the undiscovered galactic baryons exist primarily in hot gaseous halos around normal galaxies (MB04; Fukugita \\& Peebles 2006; Sommer-Larsen 2006) or if they have been mostly expelled as a result of energetic blow-out (e.g., Dekel \\& Silk 1986; Oppenheimer \\& Dav{\\'e} 2006).  There are several reasons to take seriously the possibility that a large fraction of the missing galactic baryons reside in the halos of normal galaxies. In the Milky Way, X-ray absorption lines produced by local hot gas are detected in the spectra of several bright AGN (e.g., Williams et al., 2005; Fang et al., 2006).  Many argue that this absorption corresponds to local gas ( $\\sim 50$ kpc; e.g., Wang et al., 2005;  Fang et al., 2006; Bregman \\& Lloyd-Davies, 2007), but the true origin of these features, including whether it is associated with a hot component of the Milky Way disc or with an extended hot halo, remains open to debate. Interestingly, Sembach et al. (2003) and Tripp et al. (2003 ) have argued that high-velocity features observed by FUSE highlight the boundaries between HI clouds (High-Velocity Clouds, HVCs) and an extended, hot gaseous corona around the Galaxy. Further indirect evidence for a hot Galactic halo comes from ram-pressure stripping models the Magellanic stream (e.g., Mastropietro et al. (2005); Kaufmann et al. in preparation) and the lack of neutral gas in low-luminosity satellite galaxies of the Milky Way (Grcevich et al., 2008).   The HVCs themselves might be the neutral cores of larger, pressure-supported clouds embedded within this hot gas halo (MB04; Collins et al. 2005; Thom et al. 2006; Peek et al. 2007).  Quasar absorption line studies suggest that normal Milky-Way-type galaxies at intermediate redshift are surrounded by extended, $\\sim$ 100 kpc, haloes of cool gas clouds ($T \\simeq  10^4 - 10^5$ K) with high covering factors ($f \\sim 0.6 - 0.8$). Less massive galaxies tend to have smaller, less pronounced gaseous haloes (Steidel, 1995;  Chen et al., 2001;  Tumlinson \\& Fang, 2005; Kacprzak et al., 2008;  Chen \\& Tinker, 2008) and the probability that cool halo gas is present may correlate with the color of the galaxy (Barton \\& Cooke  in preparation). Interestingly, however, X-ray observations place tight constraints on the nature of hot gas around nearby disc galaxies.  Specifically, it must be of a relatively low density in order to evade X-ray emission bounds (recent limits include $S_x < 10^{-14}$ erg cm$^{-2}$ s$^{-1}$ arcmin$^{-2}$, Pedersen et al. 2006; $L_X < 3.8 \\times 10^{41}$ erg s$^{-1}$, Benson et al. 2000; $L_X < 3 \\times 10^{39}$ erg s$^{-1}$, Li et al. 2007).  These results, together with the fairly high covering factors in cool clouds implied by absorption line studies may suggest a picture where normal galaxies are surrounded by extended, low-density hot ($\\sim 10^6$ K) haloes that are filled with fragmented, pressure supported cool ($\\sim 10^4$ K) clouds (MB04; Mo \\& Miralda-Escude 1996).   Independently, models aimed at explaining the optical properties of galaxies have relied increasingly on the idea that extended, quasi-stable hot gas haloes develop around massive galaxies (Birnboim \\& Dekel 2003; Kere{\\v s} et al. 2005; Bower et al. 2006, Croton et al. 2006; Dekel \\& Birnboim 2006). It is suggested that these hot haloes may be quite susceptible to feedback mechanisms, which could stabilise the systems to cooling and help explain the observed bimodality in galaxy properties (Dekel \\& Birnboim 2006). It is possible that in massive galactic haloes (a few times $10^{12}$ M$_\\odot$) gravitational quenching by clumpy (fragmentary) accretion could provide the source of energy (Dekel \\& Birnboim 2008).   \\begin{figure} \\includegraphics[scale=0.4]{densIC_compIIb.ps} \\caption{Initial gas density profiles for our low and high entropy models.  Note that both models have the same total gas mass within their virial radii.  The profiles shown have been evolved with an adiabatic equation of state in order to achieve full relaxation.\\label{IC}} \\end{figure}   Observational probes of the gaseous haloes of galaxies provide a potential means of testing these ideas.  Entropy injected from feedback mechanisms will alter the density distribution of halo gas and affect associated cooling rates (and thus X-ray emission) and the distribution of cooling clouds fragmenting within the hot haloes. Similarly, early feedback or {\\em pre-heating} before the halo collapses can affect halo gas profiles in a related manner, with positive consequences for galaxy properties at $z=0$ (Mo \\& Mao 2002; Oh \\& Benson 2003; Lu \\& Mo 2007).  This type of feedback has the potential for solving many of the major problems in galaxy formation. Not only will it help with overcooling, but it may also result in larger disc galaxies, relieving the so-called angular-momentum problem (Navarro \\& Steinmetz 2000, Maller \\& Dekel 2002).  Moreover, the X-ray emission in pre-heated haloes is expected to stay within observational bounds (Mo \\& Mao 2002).  The degree to which this type of pre-heating may affect the fraction of material accreted directly in the form cold flows that do not shock-heat (Birnboim \\& Dekel 2003; Kere{\\v s} et al. 2005) has yet to be investigated.    \\begin{table*} \\begin{center} \\caption{Simulated galaxies. The fiducial models are shown in the upper part, the runs investigating resolution and temperature dependence are shown in the lower parts.  \\label{table-ic}} \\begin{tabular}{||l|ccccc|lll||} \\hline  \\hline \\, \\, \\, \\, \\, (1) & (2) & (3) & (4) & (5) & (6) & (7)\\\\ \\, \\, \\, Name   & \\# gas particles &  $T_{F}$  & Softening Length & Gas Particle & Dark Matter &  $E_{tot}(t=0)$  \\\\ \\, &       \\,   &  [$10^4$ K]  & [kpc] & Mass [$M_{\\odot}$] & Particle Mass [$M_{\\odot}$] & [erg]\\\\ \\hline low entropy             &$5\\times 10^5$     &       3.0   &  0.514 & $2.8 \\times 10^5$ &$2.5 \\times 10^6$  & $-3.09\\times 10^{59}$ \\\\ high entropy           &$5\\times10^5$     &         3.0   &  0.514 & $3.6 \\times 10^5$ & $2.5 \\times 10^6$ & $-2.91\\times 10^{59}$ \\\\ \\hline low entropy low-resolution       &$1\\times 10^5$     &      3.0  & 0.514 &  $1.4\\times 10^6$ & $1.3\\times 10^7$\\\\ high entropy  low-resolution             &$1\\times10^5$     &         3.0  & 0.514 &  $1.8\\times 10^6$ & $1.3\\times 10^7$\\\\ high entropy  high-resolution             &$2\\times 10^6$     &        3.0  & 0.257 & $9.0\\times 10^4$ & $6.3\\times 10^5$\\\\ \\hline high entropy  LT             &$5\\times10^5$     &        1.5  & 0.514 & $3.6 \\times 10^5$ & $2.5 \\times 10^6$\\\\ \\hline \\end{tabular} \\end{center} {\\small (3) $T_F$ is the imposed temperature floor, LT in the name indicates a low temperature floor. \\\\ (7) The total energy content of the halo after the evolution with the adiabatic EOS. } \\end{table*}    In this paper we begin to address some of these questions by simulating the formation of a Milky Way sized spiral galaxies starting with two different initial conditions.  We focus specifically on the cooling of hot halo baryons -- a process that is expected to be important in fueling galaxy assembly at $z < 1$ in dark matter haloes of mass $\\sim 10^{12} M_\\odot$ (after the epoch of cold-mode accretion has ended for these objects, Kere{\\v s} et al. 2009; Brooks et al. 2009). Note that $10^{12} \\Msun$ dark matter haloes grow rapidly at high redshift, but then grow slowly after $z\\sim 1$ (Wechsler et al. 2002; Zhao et al. 2008) and usually do not experience major ($>1/3$) mergers since that time (Maller et at. 2006, Stewart et al. 2008).  Therefore while our simulations do not model the early growth of the galaxy well, they likely provide a reasonable estimate of late-time cooling flow behaviour.  The first case we consider has a hot gas profile that traces the density profile of the dark matter.  This, our {\\em low-entropy} halo case, has the lowest central entropy in hot gas that could realistically be expected.  Our second case explores cooling from a {\\em high-entropy} hot halo, which has a significant core in its density profile at small radius.  This hot halo has a central entropy value of $S \\sim 30$ keV cm$^2$, which is well within range expected (and seen) in cosmological simulations that include substantial non-gravitational pre-heating sources (e.g. Dav{\\'e} et al. 2008 find central entropy values of $S \\sim 100$ keV cm$^2$ in small group haloes of mass $\\sim 10^{12.7} M_{\\odot}$ as a result of feedback from outflows).  Another possibility is that dynamical processes could be responsible for creating hot gas halo profiles of this kind (Kere{\\v s} et al. 2009; Dekel \\& Birnboim 2008; Conroy \\& Ostriker 2008; Khochfar \\& Ostriker 2008).  We evolve these two haloes in isolation without any (additional) form of feedback for 10 Gyr to explore three questions. One, can large, low-density, hot haloes of the kind suggested by observations exist in a quasi-stable state for cosmological time-scales?  Two, can changes to circum-galactic hot gas at high redshifts persist until today? And three, do different initial conditions and cooling histories create features that are observable in galaxies and their gaseous halos today?    In Section 2, we present our initial conditions and the numerical techniques. In Section 3, we discuss the resulting gas halo properties after 10 Gyr of cooling and discuss these results in the context of gas halo observations. The time evolution of the hot haloes as they cool and form central galactic discs are studied in section 4. Section 5 presents a discussion of numerical convergence. We conclude and summarise in Section 6. ",
        "conclusions": "  \\begin{itemize}  \\item Changing the distribution of hot gas around galaxies, as might be expected in some pre-heating feedback schemes, dramatically alters many aspects of the subsequent cooling behaviour, as well as the energy required to stabilise the systems to subsequent cooling (see also Mo \\& Mao 2002; Oh \\& Benson 2003).  These differences are not only manifest in the resultant optical characteristics of the central disc galaxy, but they may be tested via X-ray observations and other non-optical probes of the haloes themselves.  \\item Galaxy formation that proceeds via the thermal instability and cloud infall can dramatically affect  the properties of the resultant galaxy system. In our high entropy simulations, the  resultant disc is not only less massive, but it is larger, with more specific angular momentum than in the low-entropy case.  Such behaviour is simply not possible if the cooling occurs mainly near the disc in a central cooling flow.   \\item It is possible to maintain a high-mass ($M \\sim 5 \\times 10^{10} \\Msun$), low-density, quasi-stable gaseous halo around a Milky Way size galaxy for a significant fraction of the Hubble time without a continuous input of feedback energy.  Such an extended halo (once in place) is a much better match to observations than low-entropy (cuspy/high-density) hot gas haloes that are commonly adopted in first-order analytic explorations.  \\end{itemize}  These simulations show that the large core initial gas density profile without any feedback can produce a galaxy that hardly suffers from many of the classical problems in galaxy formation.  Kere{\\v s} et al. (2009) have seen (smaller) cores in the density profiles in some of their cosmological haloes and speculate that they could be produced during the early chaotic, rapid assembly phase of haloes (without the need of AGN heating or other preventive feedback). Another possibility, explored by Tang et al. (2008) using one-dimensional simulations, is that AGN pre-heating is triggered by the initial galaxy collapse itself.   This latter collapse could be associated with the end of the cold-mode phase (Kere{\\v s} et al. 2009) or (perhaps equivalently) the end of the rapid accretion epoch (Wechsler et al. 2002).  As we have shown, once the hot gas density profile is suitably rearranged, it basically shuts down associated hot halo cooling for several Gigayears without any further input of energy after that early phase.   Tests of these ideas using high resolution numerical simulations in a full cosmological set-up are underway."
    },
    "0812/0812.3023_arXiv.txt": {
        "abstract": "\\begin{center} We investigate the oscillations of slowly rotating superfluid stars, taking into account the vortex mediated mutual friction force that is expected to be the main damping mechanism in mature neutron star cores. Working to linear order in the rotation of the star, we consider both the fundamental f-modes and the inertial r-modes. In the case of the (polar) f-modes, we work out an analytic approximation of the mode which allows us to write down a closed expression for the mutual friction damping timescale. The analytic result is in good agreement with previous numerical results obtained using an energy integral argument. We  extend previous work by considering the full range of permissible values for the vortex drag, e.g. the friction between each individual vortex and the electron fluid. This leads to the first ever results for the f-mode in the strong drag regime. Our estimates provide useful insight into the dependence on, and relevance of, various equation of state parameters. In the case of the (axial) r-modes, we confirm the existence of two classes of modes. However, we demonstrate that only one of these sets remains purely axial in more realistic neutron star models. Our analysis lays the foundation for companion studies of the mutual friction damping of the r-modes at second order in the slow-rotation approximation, the first time evolutions for superfluid neutron star perturbations and also the first detailed attempt at studying the dynamics of superfluid neutron stars with both a relative rotation between the components and mutual friction. \\end{center} ",
        "introduction": "Neutron stars have a complex interior structure. With core densities reaching several times the nuclear saturation density, these objects require an understanding of physics that cannot be gained from laboratory experiments. This makes the modelling of neutron star dynamics an interesting challenge. On the one hand, one has to consider exotic physics that is, at best, poorly constrained. On the other hand, one may ask to what extent observations can distinguish between different possible models. An excellent example of this interplay concerns the possibility that the quarks may deconfine in the high density region. If this is the case, it will have a considerable effect on transport properties associated with viscosity and heat conductivity. In fact, such a quark core is expected to be a colour superconductor \\cite{alford}. The dynamics of this exotic state of matter, and the relevance of its different possible phases, is not yet certain. In order to improve our understanding of this problem, we need to build more precise stellar models and study, for example, their oscillation properties in detail. In this context, considerable attention has been focused on the inertial r-modes of a rotating star. The r-modes are interesting because they can be driven unstable by the emission of gravitational radiation, see \\cite{nk,nareview} for literature reviews. The r-mode instability window is, however, sensitive to the physics of the neutron star interior. Since the bulk and shear viscosities are quite different in a quark core,  compared to ``normal'' npe matter, one may hope to use observations to constrain the theory, see \\cite{alford} for a discussion of the relevant literature. In absence of a direct detection of an r-mode gravitational-wave signal, this analysis would have to be based on the nature of the instability window. The idea would be that, if an observed neutron star spins at a rate that would place it inside a predicted instability region, one may be able to rule out this particular theoretical model. Of course, this argument comes with a number of caveats. It could, for example, be that  additional  physics places a stronger constraint on the r-modes than the considered mechanisms. Inevitably, this becomes a ``work in progress'' where improved theoretical models are tested against better observational data.  In order to consider ``realistic'' neutron stars, it is important to appreciate the relevance of superfluidity. A neutron star is expected to contain a number of superfluid/superconducting components \\cite{haensel}, and it is crucial to understand to what extent this affects the stars oscillation properties. It is well established that the behaviour of a superfluid system can differ significantly from standard hydrodynamics. The most familiar low-temperature system is, perhaps, He$_4$, which exhibits superfluidity below a critical temperature near 2~K. Experimentally, it has been demonstrated that this system is very well described by the Navier-Stokes equations above the critical temperature. Below the critical temperature the behaviour is very different, and a ``two-fluid'' model is generally required (see \\cite{helium} for a very recent discussion). Superfluid neutron stars are, to some extent, similar. The second sound in Helium is analogous to a set of, more or less distinct, ``superfluid'' oscillation modes \\cite{epstein,mendell1,ac01} in a neutron star. These additional modes arise because the different components of a superfluid system are allowed to move ``through'' each other. The dissipation channels in a superfluid star are also quite different. Basically, the superfluid flows without friction. In the outer core of a neutron star, which is dominated by npe matter, one expects the neutrons to be superfluid while the protons form a superconductor. As a  result, the shear viscosity is dominated by e-e scattering \\cite{cutler,npa}. The bulk viscosity, which is due to the fluid motion driving the system away from chemical equilibrium and the resultant energy loss due to nuclear reactions, is also expected to be (exponentially) suppressed in a superfluid \\cite{haensel}. These effects have direct implications for the damping of neutron star oscillations, and play a key role in determining the r-mode instability window for a mature neutron star. This is, however, not the end of the story.  A superfluid exhibits an additional dissipation mechanism, usually referred to as  ``mutual friction\". The mutual friction is due the presence of vortices in a rotating superfluid. In a neutron star core, the electrons can scatter dissipatively off of the (local) magnetic field of each vortex (see \\cite{als,mendell,trev1} for discussions and references). This effect may dominate the damping of realistic neutron star oscillation modes.  The basic requirements of a rudimentary model for superfluid neutron star oscillations should now be clear. One must account for the additional dynamical degree(s) of freedom, and also account for the mutual friction damping. This is obviously only a starting point, but the problem is sufficiently complicated that one may want to proceed with care. There has already been a number of studies of dissipative superfluid oscillations. The area was pioneered by Lindblom and Mendell who, in particular, demonstrated that the gravitational-wave instability of the fundamental f-modes would be suppressed in a superfluid star \\cite{lm95}. Following the discovery of the r-mode instability, they also provided the first accurate estimates of the relevance of the mutual friction for these modes \\cite{lm00}. Similar results were subsequently obtained by Lee and Yoshida \\cite{yl1}.  These studies provide important assessments of the relevance of the mutual friction damping. There are, however, a number of reasons why we need to return to this problem. Most importantly, we want to consider more realistic neutron star models, including finite temperature effects, magnetic fields and the possible presence of exotic (hyperon and/or quark) cores. The additional physics brings additional complications, like additional fluid degrees of freedom, boundary layers at phase-transition interfaces and fundamental issues concerning dissipative multifluid systems \\cite{helium,monster}. We also need to move away from the assumption that the vortex drag, which leads to the mutual friction, is weak. Strong arguments suggest that this is not going to be the case when the protons form a type~II superconductor and there are magnetic fluxtubes present in the system \\cite{sauls,ruderman,link03}. The neutron vortices may be ``pinned'' to the fluxtubes leading to a strong drag regime. The strong drag problem has only  been considered recently \\cite{precess,kostas}, and the first results demonstrate the presence of a new instability in systems where the two components rotate at different rates. This instability, which may be relevant for the understanding of pulsar glitches \\citep{kostas}, provides a direct demonstration that the dynamics in the strong drag regime may be both complicated and interesting. The present investigation lays the foundation for future work in this direction by allowing for strong drag. In particular, we retain the dynamic contribution to the mutual ``friction'' that has previously been neglected as a matter of course. ",
        "conclusions": "The aim of this paper was to lay the foundation for a renewed assault on the problem of dissipative superfluid neutron star oscillations. We have discussed the oscillations of slowly rotating superfluid stars, taking into account the mutual friction force at linear order in the (presumed) slow rotation of the star. We have considered both the fundamental f-modes and the inertial r-modes. Our analysis differs from previous studies in that we do not assume weak mutual friction from the outset, the final results are also valid in the strong drag regime.  In the case of the  f-modes, we worked out an analytic approximation for the mode which allowed us to write down a closed expression for the mutual friction damping timescale. This result, which is in good agreement with previous numerical results of Lindblom and Mendell \\cite{lm95}, provides useful insight into the dependence on, and relevance of, various equation of state parameters. The scaling with the harmonic index $l$ is also obvious from our final formula. The analysis is readily extended to stars with superfluid cores that do not extend all the way to the surface (as assumed in our analysis), although the result is then less transparent.  In the case of the  r-modes, we have confirmed the existence of two classes of modes. However, we demonstrated that only one of these sets will remain purely axial in more realistic situations. This agrees with previous results of Lee and Yoshida \\cite{yl1}. We discussed some peculiarities of the counter-moving r-modes. In particular, the fact that they may be unstable for some parameter values. Even though we do not expect this instability to be relevant for realistic superfluid stars, its existence is of conceptual interest.  Building on the formalism and the results presented in this paper, we are currently carrying out a detailed study of the mutual friction damping of the r-modes at second order in the slow-rotation approximation \\cite{bryn}. At the same time we are considering neutron stars with exotic hyperon and/or quark cores. Since the multifluid aspects of those problems have never been considered in detail, these are exciting developments. They are, in fact, necessary if we want to understand the dynamics of realistic models of mature neutron stars."
    },
    "0812/0812.1210_arXiv.txt": {
        "abstract": "This paper addresses alternative procedures to the ESO supplied pipeline procedures for the reduction of UVES spectra of two quasar spectra to determine the value of the fundamental constant $\\mu = M_p/M_e$ at early times in the universe.  The procedures utilize intermediate product images and spectra produced by the pipeline with alternative wavelength calibration and spectrum addition methods. Spectroscopic studies that require extreme wavelength precision need customized wavelength calibration procedures beyond that usually supplied by the standard data reduction pipelines. An example of such studies is the measurement of the values of the fundamental constants at early times in the universe. This article describes a wavelength calibration procedure for the UV-Visual Echelle Spectrometer on the Very Large Telescope, however, it can be extended to other spectrometers as well. The procedure described here provides relative wavelength precision of better than $3\\times 10^{-7}$ for the long-slit Thorium-Argon calibration lamp exposures.  The gain in precision over the pipeline wavelength calibration is almost entirely due to a more exclusive selection of Th/Ar calibration lines. ",
        "introduction": "\\label{s-int}  This paper serves two main purposes.  The first is to establish a detailed record of the UVES wavelength calibration and spectrum addition procedures used in the measurement of the fundamental constant $\\mu \\equiv M_p / M_e$ \\citep{thm08} based on spectra taken with UVES on the VLT.  The second purpose is a description of an alternative to the pipeline method of establishing a wavelength calibration of Very Large Telescope (VLT) UV-visual Echelle Spectrometer (UVES) spectra.  In a belief that a scientific result on the value of a fundamental constant in the early universe should be described well enough that it can be repeated by other researchers, we have provided more detail than is usually present in the description of a data reduction method.  The standard data reduction pipeline provided by the European Southern Observatory (ESO) for UVES does an excellent job of providing well calibrated spectra for use in most scientific investigations.  Certain investigations, however, require a more precise wavelength calibration than is provided by the pipeline. \\citet{mur07} point out that the UVES pipeline wavelength calibration has systematic deviations in accuracy that can adversely affect precise measurements.  In addition, proper validation of the accuracy of results obtained from a pipeline data reduction requires an exact knowledge of the pipeline data reduction procedures. This is often not easily available or discernible for a standard pipeline.  One example of such observations is the measurement of the values of fundamental constants in the early universe such as $\\alpha$, the fine structure constant and $\\mu$, the ratio of the proton to electron mass, through precise spectroscopy of distant objects.  The faintness of such objects often makes the use of techniques such as iodine absorption cells impractical.  This article describes a set of observing and data reduction procedures that can improve upon the wavelength accuracy relative to that available with the standard UVES pipeline.  The procedures achieve a wavelength accuracy of $\\Delta \\lambda / \\lambda \\approx 2.5 \\times 10^{-7}$ for Thorium-Argon lines in the long-slit calibration spectra for UVES with a slit width of 0.8 arc seconds.  The main component in achieving this accuracy is the proper selection of appropriate calibration lines, similar to that of \\citet{mur07}.  The examples used in this analysis are VLT UVES spectra of Q0347-383 and Q0405-443 taken in January of 2002 and 2003 respectively\\footnote{Based on observations made with ESO telescopes at the Paranal Observatory under program IDs 68.A-0106 and 70.A-0017}.  These particular spectra have been used to determine the value of $\\mu$ at early times in the universe (\\citet{rei06} and references therein, \\citet{kin08}).  The measurements of $\\mu$ are based on the absorption line wavelengths of molecular hydrogen produced in damped Lyman alpha clouds along the line of sight to the QSO.  The redshifts of the H$_2$ lines are z=3.0249007 for Q0347-383 and z=2.5947361 for Q0405-443 \\citep{thm08}.  In this work, however, we concentrate on the long-slit calibration line spectra that accompany these observations and the proper addition of the individual 2x2 binned pixel object spectra that are an intermediate product of the UVES pipeline data reduction.  In all of the following the word pixel refers to the 2x2 binned pixels produced during readout and maintained during the pipeline processing. The goal is to find, as accurately as possible, the true wavelength solution for each order of the spectrum and how that solution varies with time during the course of the observations.  All solutions are for vacuum wavelengths. The wavelength calibration techniques are independent of the objects being observed. The best results occur when a rigorous sequence of object and calibration observations is performed.  This sequence was not strictly followed in this case but the sequence is close enough to be useful. ",
        "conclusions": "The measurement of fundamental constants through astronomical spectra requires the best possible wavelength calibration that can be achieved.  This requires both observing and data reduction procedures beyond what is practical for most standard observation and data reduction procedures.  The procedures described in this paper have achieved a $1\\sigma$ fractional $\\Delta \\lambda / \\lambda$ wavelength accuracy of better than $3 \\times 10^{-7}$ for UVES spectra which is a factor of 2 better than the accuracy of the UVES pipeline reduction.  More importantly there is no detectable systematic residual wander as was discovered by \\citet{mur07} in the pipeline wavelength calibration.  The reduction methods are described in detail so that they can be repeated by other researchers and a suggested observing sequence was presented that can improve the accuracy of both the wavelength calibration and its transfer to the object spectra.  This paper serves as a detailed description of the calibration procedures used in the determination of the fundamental constant $\\mu$ through the molecular hydrogen absorption lines in the spectra of two QSOs \\citet{thm08}."
    },
    "0812/0812.1342_arXiv.txt": {
        "abstract": "{ Many problems of cold dark matter models  such as the cusp problem and the missing satellite problem can be alleviated, if galactic halo dark matter particles are  ultra-light scalar particles and in  Bose-Einstein condensate (BEC), thanks to a characteristic length scale of the  particles. We show that this finite length scale of the dark matter can also  explain  the recently observed common central mass of the Milky Way satellites ($\\sim 10^7 M_\\odot$) independent of their luminosity, if the mass of the dark matter particle is about $10^{-22} eV$. } ",
        "introduction": "Dark matter (DM) is one of the most important puzzles in modern physics and cosmology. Since the presence of DM was  inferred from gravitational effects on visible matter in galaxies~\\cite{DMreview}, a good DM model should explain the structure and evolution of galaxies. The dwarf galaxies seem to be the smallest dark matter dominated astronomical objects and, hence, are ideal for studying the nature of  DM~\\cite{gilmore}. Observational data suggest that a typical dwarf galaxy has a size never less than $O(10^2)~ pc$. This could pose a challenge to the structure formation theory based on a simple cold dark matter (CDM) model~\\cite{gilmore}, because there is no scale that can be easily accessed observationally in the CDM theory. (However, there is surely  a scale in the structure formation of baryons in DM halos. Understanding theoretically the scale is a major topic of present-day research.)  Although CDM such as WIMP( Weakly Interacting Massive Particles) with the cosmological constant ($\\Lambda$CDM) model is popular and remarkably successful in explaining large scale structures of the universe, it seems to encounter many other problems such as the cusp problem and the missing satellite problem on a scale of galactic or sub-galactic structures~\\cite{DMmodels}. For example, numerical studies with the $\\Lambda$CDM model usually predict a diverging central halo DM density $\\rho_{DM}(r)\\propto 1/r^\\alpha,\\alpha\\st{>}{\\sim}1$, where $r$ is a distance from the center of a galaxy, while  observations indicate that some low surface brightness galaxies have an almost flat core density. (However, some dwarf spheroidal  satellites of the Milky Way have steeper DM density profiles.) The CDM theory also predicts many  satellite galaxies $(O(10^3))$ around the Milky Way, however, only $O(10^2)$ satellite galaxies are observed at most ~\\cite{Salucci:2002nc,navarro-1996-462,deblok-2002,crisis}.   Recently, it was also found that the mass enclosed within a radius of $300~pc$ in these galaxies is approximately constant ($\\sim 10^7 M_\\odot$) regardless of their luminosity between $10^3$ and $10^7 L_\\odot$ ~\\cite{Strigari:2008ib}. This result implies the existence of a  minimum mass scale for the dwarf  galaxies~\\cite{Mateo:1998wg,gilmore} independent of their  baryon matter fraction, in addition to the minimum length scale. This observational fact could be also problematic with the $\\Lambda$CDM model, which usually predicts many dark matter dominated structures smaller than the dwarf galaxies down to $10^{-6}M_\\odot$~\\cite{1475-7516-2005-08-003}. ( Strigari et al ~\\cite{Strigari:2008ib} suggested that this difficulty can be overcome if we consider roles of baryonic matter.)  In this paper, we suggest that a cold dark matter model based on  Bose-Einstein condensate (BEC) ~\\cite{sin1} or scalar field dark matter (SFDM)~\\cite{myhalo} can  explain this minimum mass as well as the minimum length scale observed. Considering observational and theoretical uncertainties, we are here concerned with the scales correct  within an order of magnitude.    Let us briefly review the history of the DM theory based on  BEC. In 1992, to explain the observed  galactic rotation curves, Sin \\cite{sin1,sin2} suggested that galactic halos are   astronomical objects in BEC  of ultra-light (with mass $m\\simeq  O(10^{-24})eV$) DM particles such as pseudo Nambu-Goldstone boson. In this model the halos are like gigantic atoms, where the ultra-light boson DM particles are condensated in a single macroscopic quantum state $\\psi$. The quantum mechanical uncertainty principle prevents halos from a self-gravitational collapse.  Lee and Koh ~\\cite{kps,myhalo} suggested that  this  condensated halo DM can be described as a coherent scalar field \\cite{myreview} (dubbed as the boson halo model or, more generally,  SFDM model later ~\\cite{0264-9381-20-20-201} ). On the other hand, in usual CDM models, the wave functions of DM particles do not overlap much and the particles move incoherently. Thus, the key difference between our model and the usual CDM models is the state of DM particles rather than DM itself. Later,  similar ideas were  suggested  by many authors in terms of  the fuzzy DM, the fluid DM, or the repulsive DM ~\\cite {Schunck:1998nq,PhysRevLett.84.3037,PhysRevD.64.123528,PhysRevD.65.083514, repulsive,Fuzzy,corePeebles,Nontopological,PhysRevD.62.103517,Alcubierre:2001ea, Fuchs:2004xe,Matos:2001ps,0264-9381-18-17-101,PhysRevD.63.125016,Julien,moffat-2006,Akhoury:2008dv}. (See ~\\cite{DMmodels} for a  review.) Many related models like monopole failing models~\\cite{Matos:2000ki,Nucamendi:2000is}, unstable massless scalar field models~\\cite{Schunck:1998nq}, and quintessence models~\\cite{PhysRevD.64.123528} are suggested too. There are also extensive literature on, for example,  the collision of BEC-dark matter structures~\\cite{Bernal:2006ci,mycollision}, non-spherical collapse~\\cite{Bernal:2006it} and virialization processes~\\cite{Bernal:2006it}, and the collapse of SFDM fluctuations in an expanding universe~\\cite{Fuzzy,Guzman:2004wj,Fukuyama:2007sx,Matos:2009hf,guzman:024023}. (For a review see \\cite{myreview,Schunck:1998nq,Matos:2001ps} and references therein.)  In these models (BEC/SFDM model hereafter) the formation of DM structures smaller than the Compton wavelength ($O(1)~ pc$) of DM particles is suppressed by the quantum uncertainty principle. (We will see that the actual length scale is somewhat larger.) It was shown that this property could  alleviate the problems of the $\\Lambda$CDM model ~\\cite{corePeebles,PhysRevD.62.103517,0264-9381-17-13-101,PhysRevD.63.063506}, such as the cusp problem~\\cite{Klein:2008mn} and missing satellite problem. It is also suggested that this model can well explain the observed rotation curves~\\cite{Boehmer:2007um}, collisions of galaxy clusters~\\cite{mycollision} and the size evolution of the massive galaxies~\\cite{mysize}. Note that the BEC/SFDM particles behave like CDM ~\\cite{Matos:2003pe} at the scale larger than galaxies during the cosmological structure formation due to a small velocity dispersion of the Bose-Einstein  condensated state. Thus, this model safely satisfies the criteria of the well established structure formation theory above the galactic scales.  In this paper, we show that the minimum mass of galaxies can be explained if DM is in BEC. In Sec. 2 we show the relation between the minimum mass and the minimum length scale of galaxies. In Sec. 3 we review the proposed origin of the minimum length scale in the BEC/SFDM model. In Sec. 4 we present a result of a simple numerical study supporting our theory. Section 5 contains discussion. ",
        "conclusions": "We have shown analytically and numerically that the BEC/SFDM theory can explain the three observed properties of dwarf galaxies, i.e., the minimum length scale, the minimum mass scale, and their independence from the brightness. On the other hand, it is difficult to explain all these properties in a $single$ scenario in the CDM context, although there are proposed solutions to some of these problems~\\cite{Li:2008fx,Maccio':2008qt} relying on   roles of baryon matter. Another merit of our approach is that the stable DM configurations are not much dependent on the complicated galaxy merging history, because the equilibrium condition forces the halos to rearrange their DM distribution so that they  have a universal form regardless of their history. This could explain the observed universality of the mass and density profile independent of  possible merging history.    In conclusion, the BEC/SFDM  could be a compelling alternative to the standard CDM, because it  could not only solve some known problems of CDM model such as the cusp problem and the satellite problem, but also has a possibility to explain  recent observational mysteries of galaxy evolution in a simple way, thanks to its wave nature and the characteristic length scale. Further tests and future observations may reveal whether CDM is in a BEC."
    },
    "0812/0812.4751_arXiv.txt": {
        "abstract": "{ Several models have been proposed to explain the twin kilohertz quasi-periodic oscillations (kHz QPOs) detected in neutron-star low-mass X-ray binaries; but when confronting theory with observations, not much attention has been paid so far to the relation between their amplitudes.} {For six neutron-star atoll sources (namely 4U~1608-52, 4U~1636-53, 4U~0614+09, 4U~1728-34, 4U~1820-30, and 4U~1735-44,) we investigate the relationship between the observed fractional rms amplitudes of the twin kHz QPOs. We discuss whether this relationship displays features that could have a physical meaning in terms of the proposed QPO models.} {We consider the difference in rms amplitude between the upper and lower kHz QPOs as a function of the frequency ratio $R=\\nuU/\\nuL$. We compared two data sets. Set I is a collection taken from published data. Set II has values for the rms amplitudes obtained by automatic fitting of all RXTE-PCA observations available up to the end of 2004, corresponding to continuous segments of observation.} {\\emph{For each of the six sources, we find that there is a point in the R domain around which the amplitudes of the two twin kilohertz QPOs are the same.} We find such a point located inside a narrow interval $R=1.5\\pm\\,3\\%$. Further investigation is needed in the case of 4U~1820-30 and 4U~1735-44 to explore this finding, since we have not determined this point in Set II. There is evidence of a similar point close to R = 1.33 or R = 1.25 in the four sources \\mbox{4U~1820-30}, \\mbox{4U~1735-44}, \\mbox{4U~1608-52}, and \\mbox{4U~1636-53}. We suggest that some of these special points may correspond to the documented clustering of the twin kHz QPO frequency ratios.} {For the sources studied, the rms amplitudes of the two twin peaks become equal when the frequencies of the oscillations pass through a certain ratio $R$, which is roughly the same for each of the sources. In terms of the orbital QPO models (of both the hot-spot and disc-oscillation types), with some assumptions concerning the QPO modulation mechanism, this finding implies the existence of a specific orbit at a particular common value of the dimensionless radius, at which the oscillations corresponding to the two twin peaks come into balance. In a more general context, the amplitude difference behaviour suggests a possible energy interchange between the upper and lower QPO modes.}   \\authorrunning{G. T\\\"or\\\"ok} \\titlerunning{Reversal of the amplitude difference of kHz QPOs in six atoll sources} ",
        "introduction": "\\label{section:introduction}  A number of black-hole and neutron-star sources in low-mass X-ray binaries (LMXBs) show quasi-periodic oscillations (QPOs) in their observed X-ray fluxes, i.e., narrow features (peaks) in their power density spectra (hereafter PDS). The frequencies of these QPOs range from \\mbox{$\\sim\\!10^{-2}$\\,Hz} to \\mbox{$\\sim\\!10^{3}$\\,Hz}. Here we restrict our attention to the so-called \\emph{kHz} (or high-frequency; HF) QPOs, which have frequencies in the range $200\\fromto1300\\,\\mathrm{Hz}$, i.e., of the same order as the frequencies of orbital motion close to the compact object \\citep[see][~for a review]{Kli:2006:CompStelX-Ray:}.  In black-hole systems HF QPO peaks are typically detected at constant frequencies characteristic of a given source. When two or more HF QPO frequencies are detected, they usually come in small-number ratios, typically in a \\mbox{$3\\!:\\!2$} ratio \\citep[][]{abr-klu:2001, mc-rem:2003}. For neutron-star sources, on the other hand, kHz QPOs often arise as two simultaneously observed\\footnote{The term \\qm{simultaneously observed} is here applied to continuous observations (or their subsegments) displaying both modes. The integration time needed for these detections ranges from a few seconds to a few dozen minutes.} peaks with frequencies that change over time. This paper focuses on these kHz PDS features, referred to as {twin kHz QPOs}.  From here on, we adopt the convention of calling the two peaks forming twin kHz QPOs the \\emph{lower and upper QPO} and denote their frequencies as $\\nuL\\!<\\!\\nuU$. The twin kHz QPOs span a wide frequency range and follow a nearly linear $\\nuU$--$\\nuL$ relation for each source~\\citep[][]{Bel-Men-Hom:2005:ASTRA:,Abr-etal:2005:RAGtime6and7:CrossRef,Bur:2006:PhD:}.  In this paper we use the twin kHz QPO data obtained with the Proportional Counting Array \\citep[PCA;][]{jah-etal:1996} on board the Rossi X-ray Timing Explorer \\citep[RXTE;][]{bra-etal:1993}.  \\subsection{Orbital models of kHz QPOs and frequency ratio}  Several models have been proposed to explain kHz QPOs, and most of them involve orbital motion in the inner regions of the accretion disc \\citep[see][for a recent review]{Kli:2006:CompStelX-Ray:,Lam-Bou:2007:ASSL:ShrtPerBS}.  Among others, two frequently discussed models are based on strong-gravity properties. \\citet{Ste-Vie:1999} introduced the ``relativistic precession model'' (which we will refer to as the \\emph{RP model}) in which the QPOs are associated with the motion of blobs of matter in the inner parts of the accretion disc. The upper and lower QPO frequencies $\\nuU,~\\nuL$ are identified with the Keplerian frequency and the relativistic periastron precession frequency, respectively. Slightly after this, \\citet{Klu-Abr:2000:PHYRL:} suggested a model in which the twin kHz QPOs arise from non-linear resonance between two modes of \\emph{accretion disc oscillation}. We will refer to this as the \\emph{resonance model}. In the basic version of the resonance model, the upper and lower QPO frequencies, when in a resonant ratio, coincide with the resonant frequencies given by specific combinations of epicyclic frequencies of geodesic test particle orbital motion \\citep[see][]{abr-klu:2001,abr-etal:2003,reb:2004,Hor-Kar:2006:ASTRA:TOQPOIntRes,abr-etal:2006,swe:2008}.  Frequencies of geodesic motion at a given orbit, scale inversely with the mass~$M$ of the central compact object when the dimensionless forms of the neutron star angular momentum~$j$ and quadrupole moment~$q$ are held fixed. Therefore, within the framework of the above models, if the absolute differences between the squares of angular momenta are small among the sources being considered, the ratio $R=\\nuU/\\nuL$ between the twin kHz QPO frequencies represents a rough measure of the radial position of the QPO excitation, which is independent of the mass of the neutron star \\citep[][]{tor-etal:2008:aca:stella}. The frequency ratio~$R$ also has genuine importance for the resonance~model.  \\subsection{Strength of the signal}\\label{strength}  Because of the expected links to the orbital motion, most discussions of neutron star twin kHz~QPOs have for a long time been concentrated mainly on the frequencies $\\nuL,~\\nuU$, and on their relations and evolution.  The other QPO properties, namely the \\emph{quality factor}~$Q$ and \\emph{fractional root-mean-squared (rms) amplitude}~$r$ have also been studied, but have not attracted such wide attention. We recall that the quality factor~$Q$ characterises the coherence time of a QPO, being defined as the QPO centroid frequency over the full-width of the peak at its half-maximum, while the {fractional root-mean-squared amplitude}~$r$ represents a measure of the signal strength, which is proportional to the square root of the peak power contribution to the PDS. In the past few years, both the quality factor and the rms amplitude have been studied systematically for several sources, and possible consequences for various QPO models have been outlined~\\citep[see, e.g.,][~for further information and references]{Men:2006:MONNR:371,Bar-Oli-Mil:2006:MONNR:QPO-NS}.  However, most attention has been focussed on the quality factor and the rms amplitude as separate functions of frequency. Little attention has been paid to the \\emph{mutual relations} between the (correlated) QPO amplitudes.   In the present work we study the difference in strength between the twin kHz QPOs, $\\Delta r= r_{\\L}-r_{\\U}$, defined as the difference between the fractional rms amplitudes of the lower and upper kHz QPOs ($r_\\L, r_\\U$). ",
        "conclusions": "\\label{section:discussion}%  In five of the six atoll sources that we have considered (\\mbox{4U~1728$-$34}, \\mbox{4U~1608$-$52}, \\mbox{4U~1636$-$53}, \\mbox{4U~0614$+$09}, and \\mbox{4U~1820$-$30}), the upper kHz QPO dominates the lower one for high values of the twin kHz QPO ratio ($R\\!>\\!1.5$) corresponding to low QPO frequencies. The issue of the upper QPO dominance at low QPO frequencies in \\mbox{4U~1735$-$44} remains an open question due to the lack of the relevant datapoints in both Sets I and II. The four sources \\mbox{4U~1728$-$34}, \\mbox{4U~1608$-$52}, \\mbox{4U~1636$-$53}, and \\mbox{4U~0614$+$09} clearly exhibit the $R_1$ root close to $R\\!=\\!1.5$. Because of the lack of datapoints in Set II, however, more investigation is needed to fully confirm a similar statement about $R_1$ for the sources 4U~1820-30 and 4U~1735-44. For ratios decreasing from the value $R\\simeq1.5$, the difference between the QPO amplitudes reaches its maximum within a narrow interval of $R\\!\\in\\![1.3,~1.45]$.\\footnote{{\\cite{Bar-Oli-Mil:2005:MONNR:,Bar-Oli-Mil:2005:AN:,Bar-Oli-Mil:2006:MONNR:QPO-NS} determined the frequencies corresponding to maxima of the lower QPO quality factor in the discussed sources. Interestingly, the relevant ratios are close or coincide with those related to the maxima of $\\Delta r$  (see Fig.~\\ref{figure:rc2}).}} After this, as the ratio decreases further (but not for 4U~0614+09 where the available data terminates), the lower QPO starts to lose its dominance again. Notably, all of the determined values of the root $R_1$ lie inside the $\\sim\\!3\\%$ interval around $R\\!=\\!1.5$, which corresponds to less than $10\\%$ of the range $R\\in(1.2,\\,2.5)$. For the four sources (4U~1608-52, 4U~1636-53, 4U~1820-30, 4U~1735-44), and marginally also for 4U~1728-34, there is evidence for of another root (close to~$R= 1.33$ or~$1.25$) where the available data terminate. {Most of datapoints in Fig. \\ref{figure:rc2} appear close to the roots $R_1$ or\\,/\\,and $R_2$. We therefore suggest that there could be a link between the previously reported ratio clustering and the existence of $R_{1,2}$, which requires further investigation.}   Within the framework of the QPO resonance models, the behaviour of the amplitude difference may indicate an energy interchange between the lower and upper QPO modes, typical of non-linear resonances~\\citep[e.g.,][]{Hor-Kar:2006:ASTRA:TOQPOIntRes}. In this context, it is also rather noticeable that the roots $R_{1,\\,2}$ are \\emph{close} to the values $3/2$ and $4/3$ or $5/4$ (which we denoted by dashed vertical lines in the above figures). Nevertheless, in order to take this indication seriously, it is necessary to have a model that gives a detailed explanation of the effect.  Within the scope of the other orbital QPO models, there should also be an explanation for why the roots $R_{1,\\,2}$ appear. For instance, in the case of the relativistic precession model, one can solve equations relating the upper and lower QPO with respect to the radial coordinate $\\rho$, for the particular form of the expressions for the orbital frequencies in Schwarzschild geometry, obtaining the relation~$\\rho\\!=\\!{6MR^{2}}/{(2R-1)}$, see \\cite{tor-etal:2008:aca:stella}. According to this formula, the~$R\\!=\\!1.5$ frequency ratio would correspond to an orbital radius~$\\rho\\!=\\!6.75\\,M$, which is about~14~kilometres for a neutron star with the ``canonical'' mass~$M=1.4\\,M_{\\sun}$. The~$3\\,\\%$ scatter around $R\\!=\\!1.5$ would be projected onto an interval of only about~$\\pm1.5\\,\\%$ in~$\\rho$, i.e., onto~$\\sim\\pm0.1\\,M\\doteq\\pm200\\,\\mathrm{m}$. It is then a question not only of why the observed amplitudes corresponding to oscillations at this radius should be equal, but also why they are equal for several sources at the same radius with a scatter of less than~$250\\,\\mathrm{m}$."
    },
    "0812/0812.1174_arXiv.txt": {
        "abstract": "There is evidence for an excess in cosmic-ray electrons at about 500~GeV energy, that may be related to dark-matter annihilation. I have calculated the expected electron contributions from a pulsar and from Kaluza-Klein dark matter, based on a realistic treatment of the electron propagation in the Galaxy. Both pulsars and dark-matter clumps are quasi-pointlike and few, and therefore their electron contributions at Earth generally have spectra that deviate from the average spectrum one would calculate for a smooth source distribution. I find that pulsars younger than about $10^5$~years naturally cause a narrow peak at a few hundred GeV in the locally observed electron spectrum, similar to that observed. On the other hand, for a density $n_c = 10\\ {\\rm kpc^{-3}}$ of dark-matter clumps the sharp cut-off in the contribution from Kaluza-Klein particles is sometimes more pronounced, but often smoothed out and indistinguishable from a pulsar source, and therefore the spectral shape of the electron excess is insufficient to discriminate a dark-matter origin from more conventional astrophysical explanations. The amplitude of variations in the spectral feature caused by dark matter predominantly depends on the density of dark-matter clumps, which is not well known. ",
        "introduction": "Introduction} Only about 1 per cent of galactic cosmic rays are electrons, but their properties are of particular interest because they are very radiative at high energies and thus quickly loose their energy. The cosmic-ray electron spectrum is therefore softer than that of cosmic-ray nucleons, and high-energy electrons are few. For a long time only emulsion-chamber data of the electron flux above 100~GeV were available \\cite{elec}, which were well represented by a power law $N(E)\\propto E^{-3.2}$, but the energy resolution and statistical accuracy were limited. Recently data obtained with the ATIC balloon experiment were published, which show an excess of galactic cosmic-ray electrons at energies between 300 and 700~GeV \\cite{chang}. At TeV energies the electron flux appears to drop off rapidly \\cite{hess}. This excess is indicative of a previously unknown individual source of high-energy electrons, which could be a nearby supernova remnant \\cite{koba}, a pulsar \\cite{aha}, a microquasar \\cite{heinz}, or an annihilation site of dark-matter particles of the Kaluza-Klein type \\cite{kk}.  The dark-matter interpretation is particularly appealing, because the PAMELA collaboration has reported an increase in the cosmic-ray positron fraction above 20~GeV, suggesting the existence of a local source of both positrons and electrons \\cite{pam}. Dark-matter annihilation resulting in electron-positron pairs is possible in a number of models \\cite{cholis,berg,cirelli,arkani}, and will produce a spectrum dominated by a delta functional at the mass of the dark-matter particle \\cite{hs}. Indeed, the ATIC team finds that a Kaluza-Klein particle with mass 620~GeV fits their electron data just fine, when the expected electron source spectrum is propagated in the Galaxy using the GALPROP code \\cite{mos}. Hall and Hooper suggest that high-precision measurements of the electron spectrum be conducted with atmospheric Cherenkov observatories such as VERITAS and HESS, that would permit a discrimination between the dark-matter and pulsar hypotheses \\cite{hh}.  If dark-matter annihilation is responsible for the excess in 600-GeV electrons, then a substantial boost factor is required to match the observed electron flux which requires that the dark matter be concentrated in dense clumps. The electrons would then be injected into the Galaxy only at the location of those clumps, which introduces substantial variations in the electron flux throughout the Galaxy and can significantly modify the observed electron spectrum, as was shown in similar studies of electron propagation from supernova remnants \\cite{cl,pe98,p03}. The GALPROP code implicitely assumes a smooth source distribution on account of its using a finite-difference algorithm on a coarse grid, and therefore it will not properly describe those fluctuations.  Here we study the propagation of relativistic electrons from localized sources. ",
        "conclusions": "The ATIC collaboration has measured an excess in cosmic-ray electrons at about 500~GeV energy \\cite{chang}, which may be related to dark-matter annihilation. In this paper I have calculated the expected electron contributions from a pulsar and Kaluza-Klein dark matter. My emphasis is on a realistic treatment of the electron propagation in the Galaxy, for which I use analytical solutions to the electron transport equation. The commonly employed GALPROP code implicitely assumes a smooth distribution of the electron sources, because it uses a finite-difference algorithm on a grid.  The findings can be summarized as follows: \\begin{itemize} \\item Pulsars younger than about $10^5$~years naturally cause a narrow peak at a few hundred GeV in the locally observed electron spectrum. A single pulsar could therefore explain both the electron excess measured with ATIC and a similar excess in positrons, evidence for which at 50 to 100~GeV was obtained by the PAMELA experiment \\cite{pam}. The pulsar hypothesis does require that pulsars with ages $10^5$ to $10^6$ years leak significantly fewer electron/positron pairs, otherwise they would provide a very strong contribution in the 50 to 300~GeV band that is not observed. \\item Dark-matter annihilation occuring predominantly in dense clumps will produce a feature in the local electron spectrum, that deviates from that expected if the dark matter were smoothly distributed. The sharp cut-off in the contribution from Kaluza-Klein dark matter is often smoothed out, and the spectral feature would be indistinguishable from a pulsar source, even if the energy resolution of the electron detector were perfect. The spectral shape of the electron excess is insufficient to discriminate a dark-matter origin from more conventional astrophysical explanations, contrary to a recent claim \\cite{hh}. \\item While the mass of the dark-matter particle may be misestimated by only 20\\% or so, the amplitude of the electron excess can vary by more than a factor of 2 for a clump density $n_c = 10\\ {\\rm kpc^{-3}}$, and the required boost factors will be misestimated by the same factor. \\item All variations in the amplitude and spectral shape of the dark-matter contribution to the local electron flux depend on the density of dark-matter clumps, the variations being larger for smaller clump densities. \\end{itemize} If the clump density is the decisive parameter determining the amplitude of spectral variations in the electron excess from dark-matter annihilation, then it is of prime interest to estimate that number. Dark-matter clumpyness provides the boost factors required by the ATIC data, which may be further increased by Sommerfeld corrections \\cite{profumo,lattanzi}. Simulations of structure formation in cold-dark-matter cosmologies show clumping on a variety of scales, but the boost factors are generally very moderate \\cite{diemand,lavalle}. The clump density in those simulations is generally smaller than assumed in this paper, but that may be due to limited numerical resolution. The density of simulated particles in Via Lactea II barely exceeds $10^3\\ {\\rm kpc^{-3}}$ at the solar circle, and therefore the clump density will unavoidably be much lower. In any case, if the clump density is indeed significantly lower than $n_c = 10\\ {\\rm kpc^{-3}}$, then the dark-matter scenario can also not be distinguished from a pulsar origin by studying the high-latitude diffuse gamma-ray emission from the excess electrons, because the gamma-ray intensity distribution in the dark-matter scenario becomes similarly patchy as in the pulsar case."
    },
    "0812/0812.1945_arXiv.txt": {
        "abstract": "{  The ANTARES underwater neutrino telescope, at a depth of 2475 m in the Mediterranean Sea, near Toulon, is taking data in its final configuration of 12 detection lines. Each line is equipped with 75 photomultipliers (PMT) housed in glass pressure spheres arranged in 25 triplets at depths between 100 and 450 m above the sea floor. The PMTs look down at 45$^{\\circ}$ to have better sensitivity to the Cherenkov light from upgoing muons produced in the interactions of high energy neutrinos traversing the Earth. Such neutrinos may arrive from a variety of astrophysical sources, though the majority are atmospheric neutrinos. The data from 5 lines in operation in 2007 yielded a sufficient number of downgoing muons with which to study the detector performances, the vertical muon intensity and reconstruct the first upgoing neutrino induced muons.   } ",
        "introduction": "\\label{sec:1} The effort to build large sea water Cherenkov detectors was pioneered by the Dumand Collaboration with a prototype at great depths close to the Hawaii islands [1]; the project was eventually cancelled. Then followed the fresh water lake Baikal detector at shallow depths, which was later enlarged [2]. Considerable progress was made by the Amanda and Icecube ice telescopes in Antarctica [4, 5]. In the Mediterranean sea the NESTOR collaboration tested a deep line close to the Greek coast [5] and the NEMO Collaboration close to Sicily [6];  the ANTARES neutrino telescope at 2475 m below sea level was deployed in the Mediterranean sea, close to Toulon, France, see Figs. 1, 2, 3 [7, 8].  \\begin{figure}[t] \\begin{center} \\vspace{0.5cm} \\sidecaption \\includegraphics[scale=.28]{fig1.eps} \\caption{ Optical detection of $\\nu_{\\mu}$ neutrinos interacting in the rock below the sea and giving an upgoing muon which yields Cherenkov radiation in the detector.} \\label{fig:1}       % \\end{center} \\end{figure}  The completed ANTARES neutrino telescope, at 2475 m below sea level, is presently taking data with 12 detection lines, held vertical by buoys and anchored at the sea floor. The lines are connected to the Junction Box (JB) that distributes power and data from/to shore. The instrumented part of each line starts at 100 m above the sea floor, so that Cherenkov light can be seen also from upgoing muons coming from neutrino interactions in the rock below the sea, Fig. 1. The lines are separated by $\\sim$70 m, and are on an approximate octagonal structure. Each line has 25 floors (storeys), each with 3 Hamamatsu 10'' Photomultipliers (PMT) looking downward at 45$^{\\circ}$ from the vertical, housed inside pressure resistant glass spheres (Optical Module (OM)), Fig. 3. The storeys include frontend electronics for signal processing and digitization [9, 10, 11, 12].  \\begin{figure}[b] \\begin{center} \\sidecaption \\includegraphics[scale=.49]{fig2.eps} \\caption{ Left: Regions of sky observable by neutrino telescopes at the South Pole (Amanda, IceCube); Right: in the Mediterranean sea (Antares at 43$^\\circ$ North).} \\label{fig:2}       % \\end{center} \\end{figure}   One of the main reasons to build neutrino telescopes is to study high energy muon neutrino astronomy. Neutrinos may be produced in far away sources, where charged pions are produced and decay, like very high energy photons from $\\pi^0$ decay, and reach the earth undisturbed. Instead the photons interact with the Cosmic Microwave Background (CMB) radiation and with matter, protons are deflected by magnetic fields and neutrons are unstable. The main drawback of neutrinos is that one needs very large detectors. Neutrino telescopes may also allow dark matter searches [13], searches for exotic particles (fast magnetic monopoles, nuclearites, etc. [14, 15, 16, 17], may contribute to the study of atmospheric neutrino oscillations [18]) and test conservation laws [19].  It is worthwhile to recall that neutrino telescopes in the Northern hemisphere have access, can see, the center of our galaxy, while neutrino telescopes in Antarctica cannot see it, Fig. 2. Thus a large telescope in the Mediterranean sea (KM3) [20], and one in Antarctica (Icecube) are complementary and scientifically justified. ",
        "conclusions": "\\label{sec:4}  The Antares neutrino telescope is running in its final configuration of 12 detection lines. The preliminary results obtained with a single line and with 5 lines allowed checking and improving the reconstruction programs, and testing of MC codes.  The data taken with line 1 measured the vertical atmospheric muon intensity as a function of depth: the distribution agrees well with previous measurements.  The zenith and azimuth distributions of atmospheric downgoing muons measured with 5 lines agree with MC expectations, which have a $\\sim$30$\\%$ systematic uncertainty; the main source of uncertainty lies in the primary cosmic ray model.\\\\   I thank drs. S. Cecchini, A. Margiotta, M. Errico and prof. M. Spurio for their cooperation. I thank dr. M. Errico for typing and correcting the manuscript.   \\begin{figure}[t] \\begin{center} \\sidecaption \\includegraphics[scale=.47]{fig9.eps} \\caption{  (Preliminary) Azimuth distribution of reconstructed tracks with 5 Antares lines. Black points are data. The solid line histogram is the MC expectation with the CORSIKA code + QGSJET01 for air shower simulation plus Horandel model for CR composition. The shadowed band is an estimate of the systematics uncertanty due to environmental and geometrical parameters. } \\label{fig:9}       % \\end{center} \\end{figure}"
    },
    "0812/0812.4617_arXiv.txt": {
        "abstract": "{ Results of high-resolution simultaneous multi-frequency 8.1-15.4~GHz VLBA polarimetric observations of relativistic jets in active galactic nuclei (the MOJAVE-2 project) are analyzed. We compare characteristics of VLBI features with jet model predictions and test if adiabatic expansion is a dominating mechanism for the evolution of relativistic shocks in parsec-scale AGN jets. We also discuss magnetic field configuration, both predicted by the model and deduced from electric vector position angle measurements. }  \\FullConference{The 9th European VLBI Network Symposium on The role of VLBI in the Golden Age for Radio Astronomy and EVN Users Meeting\\\\ September 23-26, 2008\\\\ Bologna, Italy}  \\begin{document} ",
        "introduction": "In spite of the fact that our understanding of active galactic nuclei phenomenon has been significantly improved over the past decade \\cite{Lobanov06}, there are many aspects of nature of the relativistic outflows (jets), like their composition, formation, acceleration, collimation, etc, which remain to be strongly debated. VLBI observations provide us with the unique tool to explore the jets in AGN with milliarcsecond angular resolution corresponding to parsec-scale resolution in linear size. Typically, parsec-scale flows are characterized by one-sided knotty structure with sometimes pronounced curvature, rapid variations of flux density (e.g. \\cite{Wagner95,Savolainen08}), and detection of superluminal motions \\cite{Kellermann04}.  Relativistic shocks propagating down the jets are expected to be prominent on these scales, which is confirmed by detection of strong polarization \\cite{Pushkarev05,Ros00} and enhancement of the magnetic field in the brightest jet features (from the turnover frequency images) \\cite{Lobanov97}. The general evolution of a shock-induced flare in a jet is considered by \\cite{Marscher96} to pass through the Compton, synchrotron, and adiabatic stages, referring to the dominant energy loss mechanism. In the current paper we present recent results of testing adiabatic expansion of relativistic shocks in several active galactic nuclei.  Throughout the paper the $\\Lambda$CDM cosmological model with $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.3$, and $\\Omega_\\Lambda=0.7$ is adopted. Spectral index $\\alpha$ is defined as $S\\propto\\nu^\\alpha$. ",
        "conclusions": "The measured sizes and brightness temperatures of the VLBI components in parsec-scale jets in radio galaxy 1128$-$047 and quasar 2155$-$047 are found to be consistent with emission from thin relativistic shocks dominated by adiabatic energy losses. The bright distinct features in these sources may indeed be a collection of relativistic shocks developing and expanding down the flow. Applying the shock-in-jet model to the VLBI observations we predict the predominantly transverse direction of the intrinsic magnetic field associated with the central channel of the jet in 1128$-$047. Discrepancy between the model and measured brightness temperatures for the bright jet component, located at $\\sim$3.4~mas from the core in 2155$-$152, can be explained either by reacceleration of the particles due to the interaction with the ambient medium or by changing the viewing angle that becomes nearly zero."
    },
    "0812/0812.3037_arXiv.txt": {
        "abstract": "{The active galactic nucleus PG\\,1553+113 was observed by the MAGIC telescope in July 2006 during a multiwavelength campaign, in which telescopes in the optical, X-ray, and very high energies participated. Although the MAGIC data were affected by strong atmospheric absorption (calima), they were analyzed after applying a correction. In 8.5~hours, a signal was detected with a significance of 5.0\\,$\\sigma$. The integral flux above 150\\,GeV was $(2.6\\pm0.9)\\cdot10^{-7}~{\\rm ph\\,s^{-1}m^{-2}}$. By fitting the differential energy spectrum with a power law, a spectral index of $-4.1\\pm0.3$ was obtained.}   \\titlerunning{MAGIC Observations of PG\\,1553+113 during a MWL Campaign}  \\authorrunning{J.~Albert et al.}  \\date{Received 13.11.2007 / Accepted 16.10.2008 }  \\offprints{Daniela Dorner, \\email{daniela.dorner@unige.ch}} ",
        "introduction": "The Major Atmospheric Gamma-ray Imaging Cherenkov (MAGIC) telescope, located on the Canary Island of La Palma at 2200\\,m~a.s.l., is capable of extending very high energy observations to energies previously unreachable and detecting new sources at energies down to 50\\,GeV.  One of these sources is the BL Lac type object PG\\,1553+113, observed for the first time in this energy range in April and May 2005 by the MAGIC telescope and the High Energy Stereoscopic System (H.E.S.S.). From these observations a faint signal was measured, and the detection was confirmed by further observations \\citep{hess1553,magic1553}. In the subsequent years, additional VHE data were taken. From the combined data sets, strong signals were found: for H.E.S.S.\\ at 10.2\\,$\\sigma$ significance \\citep{hess1553-2} and for MAGIC at 15.0\\,$\\sigma$ significance \\citep{phd}. In addition, the VHE measurements allowed the unknown redshift of the source to be constrained. Until now the redshift of PG\\,1553+113 has been unknown, since neither emission or absorption lines could be found, nor the host galaxy resolved. Several lower limits were determined \\citep{carangelo,Sbarufatti2005,Sbarufatti2006,Scarpa2000,Treves2007}. With the MAGIC and H.E.S.S.\\ measurements, upper limits could be determined \\citep{hess1553,mazin,phd}.  To study the spectral energy distribution of a source, simultaneous data from different wavelengths are mandatory. Therefore, a multi\\-wave\\-length (MWL) campaign was performed in July 2006 to observe PG\\,1553+113. This paper concentrates on the data taken by MAGIC during this MWL campaign. ",
        "conclusions": "PG\\,1553+113 was observed in July 2006 as part of a MWL campaign with the MAGIC telescope. After correcting the data for the effect of calima, the analysis detected a gamma-ray signal of 5.0 standard deviations. Within the statistical errors, the differential energy spectrum is compatible with those derived by previous measurements including the one observed by H.E.S.S.\\ in this campaign. The inter-night light curve shows no significant variability.  The measured flux and reconstructed spectrum will be used in studies of the spectral energy distribution, which will include other data acquired during the MWL campaign \\citep{mwl}."
    },
    "0812/0812.2031_arXiv.txt": {
        "abstract": "Although quasi-periodic oscillations (QPOs) have been discovered in different X-ray sources, their origin is still a matter of debate. Analytical studies of hydrodynamic accretion disks have shown three types of trapped global modes with properties that appear to agree with the observations.  However, these studies take only linear effects into account and do not address the issues of mode excitation and decay.  Moreover, observations suggest that resonances between modes play a crucial role.  A systematic, numerical study of this problem is therefore needed.  In this paper, we use a pseudo-spectral algorithm to perform a parameter study of the inner regions of hydrodynamic disks.  By assuming $\\alpha$-viscosity, we show that steady state solutions rarely exist.  The inner edges of the disks oscillate and excite axisymmetric waves.  In addition, the flows inside the inner edges are sometimes unstable to non-axisymmetric perturbations. One-armed, or even two-armed, spirals are developed, which provides a plausible explanation for the high-frequency QPOs observed from accreting black holes.  When the Reynolds numbers are above certain critical values, the inner disks go through some transient turbulent states characterized by strong trailing spirals; while large-scale leading spirals developed in the outer disks.  We compared our numerical results with standard thin disk oscillation models. Although the non-axisymmetric features have their analytical counterparts, more careful study is needed to explain the axisymmetric oscillations. ",
        "introduction": "\\label{sec:introduction}  Quasi-periodic oscillations (QPOs) are strong coherent features in the power density spectra (PDS) that are found in different low mass X-ray binaries (LMXBs, see the review article by \\citealt{2006csxs.book...39V} and \\citealt{2006ARA&A..44...49R}). Their frequencies range from $\\sim 300$ to $\\sim 1200$ Hz, suggesting that they correspond to accretion flows very close to the central objects \\citep{2000ARA&A..38..717V, 2006csxs.book...39V}. Understanding the origin of QPOs will, therefore, lead to precise measurements of black hole spins as well as direct tests of Einstein's equivalent principle (\\citealt{2004AIPC..714...29P}, but also see \\citealt{2008PhRvL.100i1101P}).  Although several models have been proposed in the last decade to explain some aspects of the observations, the origin of QPOs is still a matter of debate.  Local analytical studies of hydrodynamic accretion disks have shown three types of trapped global modes (see the review article by \\citealt{2001PASJ...53....1K}, hereafter \\citetalias{2001PASJ...53....1K}, and references therein, or more recently \\citealt{2008NewAR..51..828W}).  The coupling between inertial oscillations and pressure variations lead to two different frequencies.  The higher frequencies correspond to inertial acoustic $p$-modes and the lower ones correspond to gravity $g$-modes \\citep{1997ApJ...476..589P, 2002ApJ...567.1043O}.  Oscillations in the vertical direction are called corrugation $c$-modes \\citep{2001ApJ...548..335S}.  Although these oscillations have properties that appear to agree with the observations, analytical studies can take only the linear effects into account and they do not address issues of mode excitation, propagation, and decay.  Moreover, observations of QPOs in black holes strongly suggest that resonances between modes play a crucial role \\citep{2001A&A...374L..19A, 2001AcPPB..32.3605K}.  This fact has led to the development of non-linear models \\citep[for example,][]{2004PASJ...56..905K, 2008PASJ...60..111K, 2008arXiv0808.1218K}.  In principle, numerical simulations can be used to test these models, including both linear and non-linear ones.  One would like to carry out simulations for many dynamical time scales to obtain accurate PDS and study their temporal variability.  However, even with the \\citet{1973A&A....24..337S} $\\alpha$-viscosity (which greatly simplifies the problem) the effective/turbulent Reynolds numbers in thin accretion disks are still beyond the value that one could perform \\emph{direct} numerical simulations with standard numerical methods.  We have developed earlier a two-dimensional, pseudo-spectral algorithm in order to model thin accretion flows around black holes \\citep{2005ApJ...628..353C}.  Spectral methods are high-order numerical methods that have a very small numerical dissipation.  For two-dimensional flows, they are at least an order of magnitude more efficient then (low order) finite difference methods.  Our algorithm solves the hydrodynamic equations in the disk plane (i.e., with $r$ and $\\phi$) by using the thin disk approximation.  In addition to the $\\alpha$-viscosity, high-order artificial viscosity is implemented by using a spectral filter, which affects only the large wavenumber modes and preserves correct shock properties \\citep[see][for details]{291525, 291526}.  We use a buffer zone to absorb outgoing waves at the outer boundary.  For the inner boundary, because the flow is always super-sonic (toward the central object), no explicit boundary condition is needed.  We apply this algorithm to study the very inner regions of two-dimensional, viscous, polytropic accretion disks around a black hole with zero spin.  We find that steady state solutions do not exist for most combinations of the parameters.  The inner edges of the disk often oscillate and excite axisymmetric waves with propagates all the way to the outer boundary.  Other than that, non-axisymmetric spirals are sometimes developed in the simulations and provide a plausible explanation of the QPOs.  Depending on the values of the parameters, these spirals can be classified as inner one-armed spirals or inner two-armed spirals; in the cases with extreme Reynolds numbers, the results are steady global one-armed spirals.  Although detailed study of all these features is beyond the scope of this paper, we will make connections between the observed numerical phenomena and the standard analytical models.  This paper is organized as following.  The assumptions and the governing equations of the problem are presented in the next section. In \\S\\ref{sec:ic}, we derive an approximate steady state solution, which is used to initialize our simulations.  We summarize the results from our simulations in \\S\\ref{sec:results} and describe our temporal analysis method in \\S\\ref{sec:fft}.  Detailed discussions of different important features are presented in the subsequent sections, namely, the axisymmetric rings in \\S\\ref{sec:rings}, non-axisymmetric spirals in \\S\\ref{sec:spirals}, and the steady global spirals in \\S\\ref{sec:global}.  We describe the limitations of our study in \\S\\ref{sec:limitations}.  Finally, we conclude with a discussion and suggest future research direction in \\S\\ref{sec:discussions}. ",
        "conclusions": "\\label{sec:discussions}  In this paper, we have performed a parameter study of two-dimensional viscous accretion disks.  We have found three types of robust features: the axisymmetric rings, the non-axisymmetric inner spirals, and the steady global one-armed spirals.  When applying a temporal analysis of the column density, a large amount of power is contained in a few modes with frequencies that are independent of the radius.  Although the properties of the axisymmetric rings do not completely agree with theoretical predictions, the frequencies $\\oring$ are always smaller than the maximum epicyclic frequency $\\kmax$.  This strongly suggests that they are axisymmetric $p$-modes.  By taking a closer look at the simulations, we find out that the local, linear assumption breaks down in the inner disk.  Moreover, based on the results of the simulations, we argue that the standard linear mode analysis is not enough to understand these rings.  The fact that these inertial acoustic modes can propagate to large radii with constant frequencies have important implications --- a frequency found in an observed oscillation can correspond to a large range of radii in the accretion disks.  Therefore, one cannot use it to estimate the size of the emission region \\citep[see][for related discussions]{2008arXiv0805.0598M}.  The two types of non-axisymmetric modes have very different properties.  For the inner spirals, their frequencies are always very close to a small multiple of the Keplerian frequency at the ISCO, $m \\OK(\\rISCO)$.  Hence, they simply are rotating patterns originating at the ISCO and corresponding to the trivial $g$-mode $\\tilde\\omega = 0$. The formation of the $m = 1$ and $m = 2$ modes is probably due to viscous overstability.  Although there are some situations in which both the axisymmetric rings and the non-axisymmetric spirals co-exist, the connection between these two features is not entirely clear.  The excitation of rings usually cleans up the inner disks and prevents the inner spirals to develop.  However, there are cases in which that the inner spirals perturb the axisymmetric rings.  Steady global spirals are then excited.  These steady patterns have properties that agree with the low frequency global $m = 1$ $p$-modes.  They are predicted as general features when the disks weakly deviated from Keplerian and have slow sound speed \\citepalias{2001PASJ...53....1K}.  The same feature can also be understood as \\emph{eccentric deformation}.  Modeling if QPOs can roughly be classified into two major schools. One conjectures that the observed oscillations correspond to the eigenfrequencies in the accretion disks.  The other proposes that an inspiral stream of matter is responsible for the modulation.  Our simulations show, depending on the parameter, both features can be excited in viscous disks.  Although the two frequencies are associated with axisymmetric rings and inner spirals are in ratio 1:3 instead of the observed 2:3, our results provide new insights to understand and model QPOs."
    },
    "0812/0812.2207_arXiv.txt": {
        "abstract": "Polarimetry of the Cosmic Microwave Background (CMB) represents one of the possible diagnostics aimed at  testing large-scale magnetism at the epoch of the photon decoupling. The propagation of electromagnetic disturbances in a magnetized plasma leads naturally to a B-mode polarization  whose angular power spectrum is hereby computed both analytically and numerically. Combined analyses of all the publicly available data on the B-mode polarization are presented, for the first time, in the light of the magnetized $\\Lambda$CDM scenario. Novel constraints on pre-equality magnetism are also derived in view of the current and expected sensitivities to the B-mode polarization. ",
        "introduction": "\\label{sec1} When (linearly) polarized electromagnetic radiation propagates in a diamagnetic material, the polarization plane of the wave changes by an angle $\\Delta\\Phi(\\lambda, B_{\\parallel})$ \\begin{equation} \\Delta\\,\\Phi(\\lambda, B_{\\parallel}) = {\\mathcal V}(\\lambda,T...)\\,\\, B_{\\parallel}\\, L, \\label{verdet} \\end{equation} where ${\\mathcal V}$ denotes the Verdet constant, $B_{\\parallel}$ characterizes the magnetic field intensity along the direction of propagation of the electromagnetic radiation of wavelength $\\lambda$; $L$ is the distance travelled by the wave within the medium. The Verdet constant ${\\mathcal V}$ not  only depends upon the wavelength but also on the specifics of the given material such as, for instance,  the temperature and the charge concentration.   In electromagnetic plasmas, the kinetic energy of the charge carriers dominates against the potential energy,  and the  Verdet constant can be computed in simple terms: \\begin{equation} \\Delta \\Phi = \\mathrm{RM}(\\tilde{n}_{\\mathrm{e}}, B_{\\parallel})\\, \\lambda^2 , \\qquad \\mathrm{RM} = \\frac{e^3}{2\\pi\\, m_{\\mathrm{e}}^2 }\\int_{0}^{L} \\tilde{n}_{\\mathrm{e}} \\,\\, B_{\\parallel} \\, ds, \\label{plasma} \\end{equation} where $m_{\\mathrm{e}}$ is the electron mass, $\\tilde{n}_{\\mathrm{e}}$ is the charge concentration of the electrons and the integral allows for the spatial variation of the magnetic field intensity as well as of the charge concentration. Equation (\\ref{plasma}) holds under two physical assumptions (see, e.g. \\cite{alfven,krall}): \\begin{itemize} \\item{} the plasma is globally neutral, i.e. the charge concentrations of the electrons and of the ions are balanced, i.e. $\\tilde{n}_{\\mathrm{e}} = \\tilde{n}_{\\mathrm{i}} = \\tilde{n}_{0}$  (where $\\tilde{n}_{0}$ denotes generically the common value of $\\tilde{n}_{\\mathrm{e}}$ and $\\tilde{n}_{\\mathrm{i}}$); \\item{} the plasma is cold, i.e. the kinetic temperature of the electrons and of the ions is always much smaller than the corresponding masses. \\end{itemize} Since the  plasma is cold (i.e. non-relativistic) and globally neutral the values of the plasma and Larmor frequencies for the electrons are always smaller than the corresponding quantities but computed for the ions: this is why, in Eq. (\\ref{plasma}) only the electron mass and the electron concentration appears. In the case of a  tokamak Eq. (\\ref{plasma}) is often used to determine the charge concentration $\\tilde{n}_{\\mathrm{e}}$. Linearly polarized electromagnetic radiation can be sent through the tokamak and the $\\Delta\\Phi(\\lambda^2)$ can be directly measured, for instance, as a function of the frequency of the incident radiation. The latter experimental information is sufficient  to determine, at least approximately, the electron concentration $\\tilde{n}_{\\mathrm{e}}$ since the geometry and intensity of the magnetic field are known from the design of the instrument.  Given the profile of the electron concentration it is relevant, in diverse circumstances, to assess the magnetic field intensity.  This is, in a nutshell, one of the objectives of cosmic polarimetry: the polarization properties of the cosmic microwave background radiation (CMB in what follows) are exploited as a diagnostic for the existence of large-scale magnetic fields. The polarization observable of the radiation field are customarily parametrized  in terms of  the Stokes parameters \\cite{jackson} whose evolution  can be written in a schematic form which is closely analog to the set of equations often employed to describe the propagation of polarized radiation through stellar atmospheres \\cite{st1,st2,st3}: \\begin{eqnarray} && \\frac{d I}{d\\tau} + \\epsilon' I = \\epsilon' {\\mathcal S}_{I}, \\label{heat1}\\\\ && \\frac{d Q}{d\\tau} + \\epsilon' Q = \\epsilon' {\\mathcal S}_{Q}  + 2 \\epsilon' \\kappa_{\\mathrm{F}} \\, U, \\label{heat2}\\\\ && \\frac{d U}{ d\\tau} + \\epsilon' U = - 2 \\epsilon' \\kappa_{\\mathrm{F}} \\, Q, \\label{heat3} \\end{eqnarray} where \\begin{itemize} \\item{} $\\epsilon'$ is the differential optical depth of the system (mainly due, in the CMB case, to photon-electron scattering) and $\\kappa_{\\mathrm{F}}$ is the Faraday rotation rate, i.e. $\\kappa_{\\mathrm{F}} = 3\\,\\lambda^2\\, ( \\hat{n} \\cdot \\vec{B})/(16\\,\\pi^2 \\, e)$; \\item{} ${\\mathcal S}_{I}$ and ${\\mathcal S}_{Q}$ are the source terms for the intensity and for the polarization; \\item{} within the present notations, the (rotationally invariant) $I$ and $V$ Stokes parameters describe, respectively, the intensity of the radiation field and the circular polarization; the $Q$ and $U$ Stokes parameters are not rotationally invariant and are both sensitive to the linear polarization. \\end{itemize} Furthermore, according to the standard conventions, we also have: \\begin{equation} \\Phi = \\frac{1}{2} \\arctan{\\biggl(\\frac{U}{Q}\\biggr)}, \\qquad P= \\sqrt{ Q^2 + U^2}. \\label{heat4} \\end{equation} Equations (\\ref{heat1})--(\\ref{heat3}) are the flat-space analog of the evolution equations studied in the present paper. In Eqs. (\\ref{heat1}), (\\ref{heat2}) and (\\ref{heat3}) the convective derivative with respect to $\\tau$ is actually a total derivative with respect to the cosmic time coordinate which receives contribution also from the fluctuations of the geometry.  The large-scale fluctuations of the geometry induce computable inhomogeneities on the Stokes parameters and they will be denoted, respectively, as\\footnote{In the spherical basis $\\hat{n} = (\\vartheta,\\, \\varphi)$ and $d\\hat{n} = \\sin{\\vartheta} \\,d\\vartheta\\, d\\varphi = - d\\cos{\\vartheta} d\\varphi$. The notation $\\mu = \\cos{\\vartheta}$ will be adopted.} $\\Delta_{\\mathrm{I}}(\\hat{n},\\tau)$, $\\Delta_{\\mathrm{Q}}(\\hat{n},\\tau)$ and $\\Delta_{\\mathrm{U}}(\\hat{n},\\tau)$. The experimental observations on the properties of the CMB polarization are expressed not in terms of $\\Delta_{\\mathrm{Q}}(\\hat{n},\\tau)$ and $\\Delta_{\\mathrm{U}}(\\hat{n},\\tau)$ but rather in terms of the corresponding E-mode and B-mode polarizations. The rationale for this common practice stems directly from the observation that the brightness perturbations of the Stokes parameters sensitive to linear polarization (i.e. $\\Delta_{\\mathrm{Q}}(\\hat{n},\\tau)$ and $\\Delta_{\\mathrm{U}}(\\hat{n},\\tau)$) are not invariant under rotations on a plane orthogonal $\\hat{n}$. Consequently, their orthogonal combinations \\begin{equation} \\Delta_{\\pm}(\\hat{n},\\tau) = \\Delta_{\\mathrm{Q}}(\\hat{n},\\tau) \\pm i \\Delta_{\\mathrm{U}}(\\hat{n},\\tau), \\label{int1} \\end{equation} transform, respectively, as fluctuations of spin-weight $\\pm 2$ \\cite{zalda} (see also \\cite{sud}). Owing to this observation,  $\\Delta_{\\pm}(\\hat{n},\\tau)$ can be expanded in terms of spin-$\\pm2$ spherical harmonics $_{\\pm 2}Y_{\\ell\\,m}(\\hat{n})$, i.e. \\begin{equation} \\Delta_{\\pm}(\\hat{n},\\tau) = \\sum_{\\ell \\, m} a_{\\pm2,\\,\\ell\\, m} \\, _{\\pm 2}Y_{\\ell\\, m}(\\hat{n}). \\label{int2} \\end{equation} The E- and B-modes are, up to a sign, the real and the imaginary parts of $a_{\\pm 2,\\ell\\,m}$, i.e. \\begin{equation} a^{(\\mathrm{E})}_{\\ell\\, m} = - \\frac{1}{2}(a_{2,\\,\\ell m} + a_{-2,\\,\\ell m}), \\qquad a^{(\\mathrm{B})}_{\\ell\\, m} =  \\frac{i}{2} (a_{2,\\,\\ell m} - a_{-2,\\,\\ell m}). \\label{int3} \\end{equation} In real space (as opposed to Fourier space), the fluctuations constructed from $a^{(\\mathrm{E})}_{\\ell\\,m}$ and $a^{(\\mathrm{B})}_{\\ell\\,m}$ have the property of being invariant under rotations on a plane orthogonal to $\\hat{n}$.  They can therefore be expanded in terms of (ordinary) spherical harmonics: \\begin{equation} \\Delta_{\\mathrm{E}}(\\hat{n},\\tau) = \\sum_{\\ell\\, m} N_{\\ell}^{-1} \\,  a^{(\\mathrm{E})}_{\\ell\\, m}  \\, Y_{\\ell\\, m}(\\hat{n}),\\qquad \\Delta_{\\mathrm{B}}(\\hat{n},\\tau) = \\sum_{\\ell\\, m} N_{\\ell}^{-1} \\,  a^{(\\mathrm{B})}_{\\ell\\, m}  \\, Y_{\\ell\\, m}(\\hat{n}), \\label{int4} \\end{equation} where $N_{\\ell} = \\sqrt{(\\ell - 2)!/(\\ell +2)!}$.  Under parity inversion $a^{(\\mathrm{E})}_{\\ell\\, m} \\to (-1)^{\\ell}  \\,a^{(\\mathrm{E})}_{\\ell\\, m}$ while $a^{(\\mathrm{B})}_{\\ell\\, m} \\to (-1)^{\\ell+1}  \\,a^{(\\mathrm{B})}_{\\ell\\, m}$. The EE and BB angular power spectra are then defined as \\begin{equation} C_{\\ell}^{(\\mathrm{EE})} = \\frac{1}{2\\ell + 1} \\sum_{m = -\\ell}^{\\ell} \\langle a^{(\\mathrm{E})*}_{\\ell m}\\,a^{(\\mathrm{E})}_{\\ell m}\\rangle,\\qquad C_{\\ell}^{(\\mathrm{BB})} = \\frac{1}{2\\ell + 1} \\sum_{m=-\\ell}^{\\ell} \\langle a^{(\\mathrm{B})*}_{\\ell m}\\,a^{(\\mathrm{B})}_{\\ell m}\\rangle, \\label{int5} \\end{equation} where $\\langle ...\\rangle$ denotes the ensemble average. In the minimal version of the $\\Lambda$CDM paradigm the adiabatic fluctuations of the scalar curvature lead to a polarization which is characterized exactly by the condition $a_{2,\\,\\ell m} = a_{-2,\\,\\ell m}$, i.e. $a_{\\ell m}^{(\\mathrm{B})} =0$.  In more general situations also the cross-correlations of the various modes can be defined in full terms. Having introduced the appropriate rotationally invariant observables, the empirical observations can now be quickly summarized, in a bulleted form, as\\footnote{Following a common practice the polarization will be described in terms of the autocorrelations and cross-correlations of the E-modes and of the B-modes (i.e., in the jargon, the EE, BB and EB angular power spectra). Relevant pieces of information are also encoded in the cross-correlation between the temperature and the polarization, i.e. the TE and TB correlations.}: \\begin{itemize} \\item{} the TE and the EE angular power spectra have been measured so far by various experiments with different accuracies; \\item{} direct measurements of the BB angular power spectra led just upper limits on the B-mode polarization since the observations are consistent with the null  result; \\item{} the polarization power spectra detected by different experiments are consistent. \\end{itemize} \\begin{table}[!ht] \\begin{center} \\begin{tabular}{|c|c|c|c|c|} \\hline Experiments & Years  & Data & Pivot frequencies &Upper limits on BB \\\\ \\hline Dasi \\cite{dasi1,dasi2,dasi3} & 2000-2003 & EE, TE, BB & $26$--$36$ GHz & ${\\mathcal G}_{\\ell}^{(\\mathrm{BB})} \\leq 50 \\,(\\mu\\,\\mathrm{K})^2$ \\\\ Boomerang \\cite{boom1}&2003  & EE, EB, BB & $145$ GHz& ${\\mathcal G}_{\\ell}^{(\\mathrm{BB})}\\leq 8.6 \\,(\\mu\\,\\mathrm{K})^2$ \\\\ Maxipol \\cite{maxip1,maxip2} & 2003 & EE, EB, BB & $140$ GHz & ${\\mathcal G}_{\\ell}^{(\\mathrm{BB})} \\leq 112.3 \\,(\\mu\\,\\mathrm{K})^2$ \\\\ Quad \\cite{quad1,quad2,quad3} & 2007 & EE,TE, TB& $100$/ $150$ GHz& ${\\mathcal G}_{\\ell}^{(\\mathrm{BB})} \\leq  10 \\,(\\mu\\,\\mathrm{K})^2$ \\\\ Cbi \\cite{cbi} & 2002-2005& EE, TE, BB  & $26$--$36$ GHz& ${\\mathcal G}_{\\ell}^{(\\mathrm{BB})} \\leq  3.76\\,(\\mu\\,\\mathrm{K})^2$ \\\\ Capmap  \\cite{capmap} & 2004-2005& EE, BB& $35$--$46$/$84$--$100$ GHz& ${\\mathcal G}_{\\ell}^{(\\mathrm{BB})} \\leq  4.8\\,(\\mu\\,\\mathrm{K})^2$ \\\\ Wmap \\cite{WMAPfirst1,WMAP51} & 2001-2006& EE, TE, BB& $23$--$94$ GHz& ${\\mathcal G}_{\\ell}^{(\\mathrm{BB})} \\leq  0.15\\,(\\mu\\,\\mathrm{K})^2$ \\\\ \\hline \\end{tabular} \\caption{The main polarization experiments, their typical frequencies, their salient observations are illustrated in the light of the polarization properties of the CMB. In the last column ${\\mathcal G}_{\\ell}^{(\\mathrm{BB})}=\\ell (\\ell +1)C_{\\ell}^{(\\mathrm{BB})}/(2\\pi)$.} \\label{TABLE1} \\end{center} \\end{table} In Tab. \\ref{TABLE1} the salient informations on different experiments measuring the polarization properties of the CMB are swiftly collected. In Tab. \\ref{TABLE1}, the last two columns at the right include specific informations on the frequency channels characteristic of each experiment and of the upper limits on the B-mode autocorrelation. The frequency of the observational channel constitutes a relevant piece of information since, as already remarked, the Faraday rate depends upon the wavelength of the propagating radiation. The main references to each experiment have been  collected in the first column of  Tab. \\ref{TABLE1}. In the case of the WMAP experiment the most recent data release can be found in \\cite{WMAP51,WMAP52,WMAP53,WMAP54,WMAP55}. The first data release corresponds to \\cite{WMAPfirst1,WMAPfirst2,WMAPfirst3}. For the specific aspects of the polarization measurements it is also interesting to bear in mind Ref. \\cite{WMAPthird2} (see also \\cite{WMAPthird1}).  The paradigm in vogue these days is the $\\Lambda$CDM scenario where $\\Lambda$ stands for the dark energy component and CDM for the cold dark matter component. In the $\\Lambda$CDM context the origin of the linear polarization of the CMB stems directly from the properties of last scattering in conjunction with the presence of a primordial spectrum of curvature perturbations which is said to be adiabatic. Adiabatic means, in the present context, that the curvature perturbations do not arise as inhomogeneities in the plasma sound speed but, on the contrary, they are related to genuine fluctuations of the energy density and of the spatial curvature. In the adiabatic case the location of the  first acoustic peak of the TT angular power spectrum determines the location of the first anticorrelation peak of the TE spectrum. This is what is observed since, roughly, $\\ell_{\\mathrm{TT}} \\simeq 220$ and $\\ell_{\\mathrm{TE}} \\simeq 3 \\ell_{\\mathrm{TT}}/4$ where $\\ell_{\\mathrm{TT}}$ denotes the location of the first acoustic peak and $\\ell_{\\mathrm{TE}}$ corresponds to the first anticorrelation peak in the TE spectrum.  Finally the success of the $\\Lambda$CDM scenario in pinning down the observed features in the TT and EE  angular power spectra gives us also, indirectly, some precious informations on the dynamics of photon decoupling and on the possibility of late-time reionization.  As indicated by Tab. \\ref{TABLE1} various experiments measured (and are now measuring) the EE and TE angular power spectra. On top of the 5-yr WMAP results \\cite{WMAP51,WMAP52,WMAP53,WMAP54,WMAP55} recently there have been rather promising results coming from the Quad collaboration \\cite{quad2} (see also \\cite{quad1,quad3}) claiming rather intriguing measurements of the EE and TE angular power spectra up to rather large harmonics\\footnote{The given measured multipole $\\ell$ is related to the angular separation. Large multipoles correspond to small angular separations, i.e., roughly $\\vartheta = \\pi/\\ell$.}, up to $\\ell \\simeq 1500$ and even $\\ell \\simeq 2000$.  In the $\\Lambda$CDM scenario with no tensors, the B-modes vanish  except for a minute component coming from the gravitational lensing of the primary anisotropies. If the $\\Lambda$CDM scenario is complemented by means of a tensor contribution, then the value of the BB autocorrelations will depend upon the ratio between the power spectrum of the tensor fluctuations and the power spectrum of the fluctuations of the scalar curvature. \\begin{figure}[!ht] \\centering \\includegraphics[height=8cm]{FIG1.eps} \\includegraphics[height=8cm]{FIG2.eps} \\caption[a]{Measurements of the BB autocorrelations as they have been reported in the various experiments listed in Tab. \\ref{TABLE1}.} \\label{Figure1} \\end{figure} The current experimental data reported in Tab. \\ref{TABLE1} are also illustrated in Fig. \\ref{Figure1}. While the possible presence of a B-mode polarization can be attributed to a stochastic background of relic gravitons, Eqs. (\\ref{heat1})--(\\ref{heat3}) suggest that it could be the signal of a pre-decoupling magnetic field.  Equation (\\ref{heat1}) describes the evolution of the intensity of the radiation field and, in a magnetized medium, the source term ${\\mathcal S}_{I}$ receives contribution from both the fluctuations of the geometry as well as from the fluctuations of the plasma which is, in turn, sensitive to the stochastic magnetic field itself. The remaining two equations determine the polarization observables and can be solved by iteration, i.e. by positing that \\begin{equation} Q = \\overline{Q} + Q^{(1)},\\qquad U= \\overline{U} + U^{(1)}. \\label{heat5} \\end{equation} The first iteration determines $\\overline{Q}$ which is a solution of Eqs. (\\ref{heat1})--(\\ref{heat2}) with $\\overline{U} =0$, i.e. \\begin{equation} \\frac{d \\overline{Q}}{d\\tau} + \\epsilon' \\overline{Q} = \\epsilon' {\\mathcal S}_{\\overline{Q}}, \\qquad \\frac{d U^{(1)}}{d\\tau} + \\epsilon' U^{(1)} =  - 2 \\epsilon' \\kappa_{\\mathrm{F}} \\overline{Q}. \\label{heat6} \\end{equation} The result of the second iteration produces $U^{(1)}\\neq 0$ and, ultimately, a B-mode polarization. In the present paper the magnetized B-mode will be computed both analytically and numerically. Furthermore various bounds on the magnetic field intensity will be derived by comparing the  Faraday-induced B-mode both with the upper limits on the BB angular power spectra (see  Tab. \\ref{TABLE1}) as well as with the published data on the B-mode angular power spectrum (see e. g. Fig. \\ref{Figure1}).  The present investigation is based on a number of previous results. The large-scale magnetic fields are included both at the level of the evolution equations as well as at the level of the initial conditions. This program has been undertaken  in \\cite{mg1} (see also \\cite{prev,maxrev1}) and subsequently developed in \\cite{mg2,mg3}. Semi-analytic results for the polarization observables which are relevant for the present investigation have been discussed in \\cite{mg2}. Recently the results of \\cite{mg1,mg2,mg3}  have been complemented by a dedicated numerical approach \\cite{gk1,gk2,gk3} which is based, indirectly, on COSMICS \\cite{cosmics1,cosmics2} and, more directly, on CMBFAST \\cite{cmbfast1,cmbfast2}.  In \\cite{gk3} a bound on pre-decoupling magnetic fields has been derived, by including, as suggested in \\cite{mg2}, the large-scale magnetic fields in the heat transfer equations as well as in the initial conditions of the full Einstein-Boltzmann hierarchy.   The obtained results showed that the bounds obtainable from the TT correlations were still more constraining than the direct Faraday bounds stemming from the upper limits on the B-mode autocorrelations. The assumptions of \\cite{gk3} will be relaxed and the finite thickness of the last scattering surface will be taken into account in the analytical and numerical estimate of the Faraday autocorrelation. While the sudden decoupling limit (i.e. infinitely thin last-scattering surface) is  justified for large angular scales (i.e. $\\ell < 100$) over small angular scales finite thickness effects are comparable with diffusive effects.  The latter statements can be understood also in analogy with the calculations of the temperature autocorrelations. While the finite thickness effects can be ignored for the estimate of the Sachs-Wolfe plateau, the same cannot be said for the region of the acoustic peaks where both the finite thickness effects compete with diffusive damping to determine the well known shape of the temperature autocorrelations. Semi-analytic  treatments of the finite thickness effects can be found in  \\cite{zal,wyse,zeld1,pav1,pav2}. Finally, finite thickness effects also determine the correlated distortions of the first three acoustic peaks  in the temperature autocorrelations when large-scale magnetic fields are consistently included in the calculation \\cite{mg3,gk2}.  In the light of some of the recent (and forthcoming) experimental data, finite thickness effects play a relevant role in the estimate of the Faraday autocorrelations as well as in the estimate of the B-mode autocorrelations. Indeed (see Tab. \\ref{TABLE1} and Fig. \\ref{Figure1}) the upper limits on the B-mode autocorrelations stem from rather small angular scales. Over those scales the Faraday induced B-mode cannot be trusted in the sudden decoupling limit.  While physically the small angular scales are the realm of diffusive effects, geometrically the small angular separations imply that the microwave sky  degenerates into a plane. On a 2-sphere the angular eigenfunctions of the Laplacian are  spherical harmonics. When the radius of the sphere is taken to infinity (i.e. the angular separation becomes small) the eigenfunctions of the Laplacian become plane waves. In the language of group representations this is one of the simplest examples of group contraction (in the In\\\"on\\\"u-Wigner sense) where the group $SO(3)$ contracts to $M(2)$, i.e. the semidirect product of the Euclidean group in two dimension and of the $SO(2)$ group \\cite{vilenkin}.  This observation can simplify the evaluation of complicated correlators and will lead, ultimately, to an analytical expression of the Faraday-induced autocorrelations at small angular scales.  The plan of the present investigation is the following. In section 2 the essentials of the interplay between large-scale magnetism and CMB physics will be reviewed. Faraday  autocorrelations will be computed in section 3. The B-mode autocorrelations will computed in section 4. In Section 5 the parameters of the magnetized $\\Lambda$CDM model will be analyzed within a frequentist approach. The concluding remarks are collected in section 6. \\renewcommand{\\theequation}{2.\\arabic{equation}} \\setcounter{equation}{0} ",
        "conclusions": "In summary, the B-mode due to Faraday rotation has been calculated both semi-analytically and numerically. The parameter space of the magnetized background has been investigated by enforcing the upper limits on the B-mode polarization and also by performing a dedicated $\\chi^2$ analysis. The two analyses led to compatible results. It has been shown that forthcoming experiments aiming at a direct detection of a B-mode polarization arising from a putative tensor power spectrum can also be able to cut through the parameter space of a magnetized background.  An interesting byproduct of the present analysis has been an analytic scrutiny of the Faraday induced B-mode at small angular scales.  In the small-scale limit the B-mode autocorrelations can be computed with reasonable accuracy by contracting (in the In\\\"on\\\"u-Wigner sense) the relevant matrix elements of the rotation group. The latter results have been cross-checked by direct numerical calculations."
    },
    "0812/0812.0491_arXiv.txt": {
        "abstract": "The DA white dwarfs in the Sloan Digital Sky Survey, as analyzed in the papers for Data Releases 1 and 4, show an increase in surface gravity towards lower effective temperatures below 11500~K. We study the various possible explanations of this effect, from a real increase of the masses to uncertainties or deficiencies of the atmospheric models. No definite answer is found but the tentative conclusion is that it is most likely the current description of convection in the framework of the mixing-length approximation, which leads to this effect. ",
        "introduction": "Almost every contribution at this white dwarf meeting mentions the Sloan Digital Sky Survey (SDSS) or is even based on its data. Because of the huge impact the survey has on this field, it is very important to critically analyze all aspects of the data collection and reduction process, to understand possible uncertainties and biases. One problem related to white dwarfs, which was obvious already in the first white dwarf catalogue publication \\citep{Kleinman.Harris.ea04}, was the apparent systematic increase of the mean surface gravity of DAs below effective temperatures of approximately 12000~K, determined from spectral fits, was again demonstrated in the data of the Data Release 4 \\citep{Eisenstein.Liebert.ea06} and confirmed by \\citet{Kepler.Kleinman.ea07} and \\citet{DeGennaro.von-Hippel.ea08}.  While the huge number of objects and the homogeneity of the survey make this feature very obvious, it is by no means a new discovery. A similar trend was first recognized with data from the Palomar-Green Survey \\citep[][and several other papers of the Montreal group]{Bergeron.Wesemael.ea90, Bergeron92}, and more recently by \\citet{Liebert.Bergeron.ea05}. A similar effect was also found in the data from the ESO Supernova Ia Progenitor Survey (SPY) by \\citet{Voss06}. A first conclusion therefore is that this effect is not a peculiarity of the SDSS data collection or reduction processes. In this paper we study the remaining possibilities \\begin{itemize} \\item the effect could be real, i.e. the DAs show a real increase of the average mass towards lower temperatures \\item it could be caused by the fitting between observations and theoretical models used to determine the atmospheric parameters \\item uncertainties in input physics or a flaw in the models used \\end{itemize} For our study we use the data of DR4 as provided from the SDSS database, with slight changes in reduction software compared to \\citet{Eisenstein.Liebert.ea06}.  We also use very nearly the same models as used in that work for our ``standard analysis'' (see below). The major difference are the fitting algorithms: whereas the SDSS uses the whole spectra, adjusted with a polynomial using information from the photometry, we here use only the spectral lines, with the continuum adjusted on both sides of the line center.  This technique minimizes the influence of uncertainties in the flux calibration.  The SDSS photometry gives an independent estimate of parameters, which is used to distinguish between the two solutions, which exist for most objects between 8000 and 18000~K.  In the main part of this study we vary some of the parameters in the input physics, in the hope of finding some hints for the origin of the effect. This concerns the treatment of convection, of the non-ideal effects in the occupation numbers, and a possible admixture of invisible helium. All these possible explanations have been proposed and studied in various papers of the Montreal group. We partially repeat this work here with a much larger and very homogeneous sample. ",
        "conclusions": "The frustrating result of this study is that we cannot propose any change of approximations in the input physics, which would clearly lead to the desired result -- a constant average mass for all white dwarfs for all effective temperatures. Variations of parameters in the Hummer-Mihalas-D\\\"appen occupation probabilities can be excluded. Helium pollution cannot be definitely excluded, but could be tested with high resolution, high S/N spectra of (apparent) DAs in the range 10000 - 12000~K. Because of the asteroseismology results discussed above we consider this solution unlikely.  That leaves by default convection as the only suspect. Given the current amount of uncertainty about the correct description, and the contradictions in the current applications, we consider it quite possible that a more physically sound description of the convective flux and gradient might solve the problem.  At the end of this study a final word of caution: fitting observed spectra with theoretical models is as much an art as an objective science. One should be very sceptical of results from spectra with S/N $<$ 30, as well as DA spectra that do not include the higher Balmer lines with high S/N. Never believe the errors so readily calculated by a $\\chi^2$ routine; these are errors which consider only statistical errors in the observed data. With modern CCD spectra from large telescopes systematic effects of reduction, the fitting process, and model deficiencies result in much larger errors. These can only be estimated in a meaningful way from comparing results for the same objects from different observations, and if possible, from different authors, with different fitting routines and model grids.    \\subsection"
    },
    "0812/0812.2185.txt": {
        "abstract": "%Abstract { This paper reviews the scientific use of the TNG from the beginning of regular observations %with the complete set of instruments (2000) %regular observations (1998) till the end of 2007. Statistics are given for the time request, use and productivity  of the telescope and its focal plane instruments. Information on the down-times and a list of the major technical works/upgrades are also included. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.1048_arXiv.txt": {
        "abstract": "The density profile of simulated dark matter structures is fairly well-established, and several explanations for its characteristics have been put forward. In contrast, the radial variation of the velocity anisotropy has still not been explained. We suggest a very simple origin, based on the shapes of the velocity distributions functions, which are shown to differ between the radial and tangential directions.  This allows us to derive a radial variation of the anisotropy profile which is in good agreement with both simulations and observations.  One of the consequences of this suggestion is that the velocity anisotropy is entirely determined once the density profile is known. We demonstrate how this explains the origin of the $\\gamma$--$\\beta$ relation, which is the connection between the slope of the density profile and the velocity anisotropy.  These findings provide us with a powerful tool, which allows us to close the Jeans equations. ",
        "introduction": "The natural outcome of cosmological structure formation theory is equilibrated dark matter (DM) structures. According to numerical simulations, the mass density profile, $\\rho(r)$, of these structures changes from something with a fairly shallow profile in the central region, $\\gamma \\equiv d{\\rm ln}\\rho/d{\\rm ln}r \\sim -1$ (or maybe zero), to something steeper in the outer region, $\\gamma \\sim -3$ (or maybe steeper) \\citep{nfw,moore,diemand} (see also \\cite{reed,stoehr,navarro2004,alister,merritt,ascasibar,stadel08,springel08}).  For the largest structures, like galaxy clusters, there appears to be fair agreement between the numerical predictions and observations concerning the central steepness \\citep{pointe,sand,buote,broadhurst,vikhlinin}, however, for smaller structures, like galaxies or dwarf galaxies, observations tend to indicate central cores~\\citep{salucci,gilmore,wilkinson} (see also~\\cite{rubin85,courteau97,palunas00,blok01,blok02,salucci01,swaters02,corb03,salucci2}). The various theoretical approaches still make different predictions \\citep{taylornavarro,hansenjeans,austin,dehnenmclaughlin,gsmh,henriksen,henriksen2}, varying from central cores to cusps.  The second natural quantity to consider (after the density profile) is the velocity anisotropy, which is defined through \\begin{equation} \\beta \\equiv 1 - \\frac{\\sigma^2_t}{\\sigma^2_r} ~, \\end{equation} where $\\sigma^2_t$ and $\\sigma^2_r$ are the 1-dimensional tangential and radial velocity dispersions~\\citep{binneytremaine}. If most dark matter particles in an equilibrated structure were purely on radial orbits, then $\\beta$ could be as large as 1, and for mainly tangential orbits $\\beta$ could be arbitrarily large and negative. Since dark matter is collision-less $\\beta$ does not have to be zero, and it could in principle even vary as a function of radius.   Numerical N-body simulations of collision-less dark matter particles show that the dark matter velocity aniso\\-tro\\-py is indeed radially varying, and that $\\beta$ goes from roughly zero in the central region, to 0.5 towards the outer region~\\citep{colelacey,carlberg}.  Only very recently has this velocity anisotropy been measured to be non-zero in galaxy clusters \\citep{hansenpiff}, and it has even been observed to be increasing as a function of radius \\citep{host2008}, in excellent agreement with the numerical predictions. For smaller structures, like our own galaxy, this has not been observed yet. In principle $\\beta$ of our Galaxy can be measured in an underground directional sensitive detector, however, it will require a large de\\-dicated experimental programme~\\citep{hosthansen}.  Very little theoretical understanding of the origin of this velocity anisotropy exists, and to my knowledge no successful derivation of it has been published (see, however, \\cite{hansenmoore,smgh07,wojtak2008}).  We will in this paper present an attempt towards deriving $\\beta$. ",
        "conclusions": "One of the consequences of the above considerations is, that the appearance of the radial variation of $\\beta$ is dictated by the density profile. That is, given any density profile, the velocity anisotropy is entirely determined, as long as the structure has had time to equilibrate. We are thus stating explicitly, that $\\beta$ is unrelated to the infall of matter in the outer region, and that the only connection $\\beta$ has to the formation process is through the radial structure of the density profile. This is naturally supported by the very non-cosmological simulations (see figs.~\\ref{fig:tangent}, \\ref{fig:tangent.log}) which also produce a $\\beta$-profile in agreement with cosmological simulations~\\citep{hansenzemp}.   \\begin{figure}[thb] \\centering \\includegraphics[angle=0,width=0.49\\textwidth]{f7.pdf} \\caption{The velocity anisotropy, $\\beta$, as function of the density slope, $\\gamma$. The blue line (triangles) is for an NFW profile, and the black solid line (crosses) is for a power-law density profile.  The orange (squares) is for the non-trivial double-bump profile, the green (stars) is for $\\rho = 1/(1+r^2)^2$ profile, and the red diamonds are for the profile suggested in \\cite{navarro2004}. The dashed straight line is the suggestion from \\cite{hansenstadel}. These results are roughly fit by $\\beta = -0.13\\gamma$.} \\label{fig:gamma.beta} \\end{figure}   Another consequence is that $\\beta$ must always be positive in equilibrated structures, since the density profile at most can develop a core (see fig.~\\ref{fig:beta.r}). This is also in good agreement with dark matter simulations \\citep{dehnenmclaughlin,barnes,faltenb,bellovary}. Virtually no numerical simulations find negative velocity anisotropy, and when they do, this is usually only in the very inner region where numerical convergence may be questioned. Few analytical treatments have predicted a negative $\\beta$ in the inner region \\citep{zait}, however, this result may be an artefact of assuming that the pseudo phases-space is a perfect power-law in radius, which is generally not correct \\citep{schmidt09}. If future high-resolution numerical simulations instead will establish that the central velocity anisotropy is negative (in agreement with the predictions of \\cite{zait}), then that would be a proof that the present analysis is flawed somehow.  It has previously been suggested that a connection between the anisotropy and the slope of the density profile should exist.  This connection appears to hold even for structures which have profiles with non-trivial radial variation in $d{\\rm log}\\rho/d{\\rm log}r$ \\citep{hansenmoore}.  We can now test this connection. In figure~\\ref{fig:gamma.beta} we show $\\beta$ as function of the density slope for the NFW profile (solid line, blue triangles), the profile suggested by \\cite{navarro2004} (red diamonds), and also for a ``double-bump'' structure (the sum of two spatially separated profiles of the form $1/(1+r)^3$, orange squares)).  This double-hump profile has a very non-trivial radial variation of $d{\\rm ln}\\rho/d{\\rm ln}r$, which cannot be well approximated by any generalized double power-law profile.  All these structures appear to land near a connection roughly given by $\\beta = -0.13\\gamma$. We also show the results for the $\\rho = 1/(1+r^2)^2$ profile (green stars), as well as for single power-law profiles (fat black line, crosses). The dashes line is $\\beta = -0.2(\\gamma-0.8)$ as suggested in \\cite{hansenstadel}, based on a set of cosmological and non-cosmological simulations. These two results differ by approximately 0.1 in $\\beta$. We indeed see that all structures land in a relatively narrow band in the $\\gamma$--$\\beta$ plane, and hence likely explaining the origin of the $\\gamma$--$\\beta$ relations.     The most important practical implication of this suggestion is, that it will allow us to close the Jeans equation. As is well-known, the Jeans equation depends on the density, dispersion, anisotropy and the total mass. Now, having demonstrated (or at least suggested strongly) that the anisotropy is uniquely determined once the density is known, we see that it is possible to close the Jeans equation for systems that are fully relaxed.     We have been making simplifying assumptions above, which all need to be tested through high resolution simulations.  First, we assume that the radial VDF is very similar to the one appearing from the Eddington formula, even in the presence of a non-zero $\\beta$.  We also assume that the $\\sigma$ entering eq.~(\\ref{eq:tsallis}) is the total one.  If the correct $\\sigma$ to use is instead closer to $\\sigma_{tan}$, then $\\beta$ will be slightly larger, but the radial variation will remain.  From these assumptions we estimate the accuracy of the present work to be about 0.1 or up to about $30\\%$ in $\\beta(r)$.        One could naturally ask why and how the radial and tangential VDF's get their shapes?  It is slightly disappointing that it is not a deep physical principle, like a generalized entropy, which is responsible. Instead it is simply the density profile (either the radially varying, or the tangentially constant) which through the Eddington formula demands that the VDF's take on these forms. These forms will therefore appear when there is sufficient amount of violent relaxation to allow enough energy exchange between the particles."
    },
    "0812/0812.3421_arXiv.txt": {
        "abstract": "In this paper, we revisit brane inflation models with the WMAP five-year results. The WMAP five-year data favor a red-tilted power spectrum of primordial fluctuations at the level of two standard deviations, which is the same as the WMAP three-year result qualitatively, but quantitatively the spectral index is slightly greater than the three-year value. This result can bring impacts on brane inflation models. According to the WMAP five-year data, we find that the KKLMMT model can survive at the level of one standard deviation, and the fine-tuning of the parameter $\\beta$ can be alleviated to a certain extent at the level of two standard deviations. ",
        "introduction": "The inflation paradigm of the early universe provides a compelling explanation for many big puzzles in the standard big bang cosmology, such as the problems of homogeneity, isotropy and flatness of the universe \\cite{Guth81}. This period of accelerated expansion in the early universe predicts a nearly scale-invariant perturbation power spectrum, which has already been supported by the measurement of temperature fluctuations in the cosmic microwave background radiation \\cite{Smoot92,Bennett03}. In spite of many phenomenological successes of the inflation based on effective field theory, there exist serious problems, concerning the origin of the scalar field driving inflation, namely the singularity problem and the trans-Plankian problem. Therefore, it is expected that inflation should be realized in a more natural way from some fundamental theory of microscopic physics. String theory is one of the most promising candidates for the fundamental microscopic physics, so it is very important to try to embed the inflation scenarios into string theory. One possible inflation scenario within the framework of string theory is the brane inflation, in which the inflation is driven by the potential between the parallel dynamic brane and anti-brane \\cite{Dvali98,Burgess01,Quevedo02}. The distance between the branes in the extra compactified dimensions plays the role of the inflaton field. However, generically, getting a sufficiently flat inflaton potential in brane inflation models is not an easy thing \\cite{Quevedo02}.  A rather realistic slow roll inflation scenario based on string theory was first proposed by Kachru et al. \\cite{kklt}. In \\cite{kklt}, by introducing some $\\overline{\\rm D}$3-branes in a warped geometry in type IIB superstring theory to break supersymmetry, the authors successfully lift the AdS vacuum to a metastable de Sitter vacuum with sufficiently long lifetime. This mechanism is often called KKLT mechanism. Furthermore, by putting an extra pair of the brane and anti-brane in this scenario, one can successfully realize a slow roll inflation model, namely the KKLMMT model \\cite{Kachru03}. In this model, the anti-brane is fixed at the bottom of a warped throat, while the brane is mobile and experiences a small attractive force towards the anti-brane. The inflation ends when the brane and the anti-brane collide and annihilate, initiating the hot big bang epoch. The annihilation of the brane and anti-brane makes the universe settle down to the string vacuum state that describes our universe. During the process of the brane collision, cosmic strings are copiously produced \\cite{Jones02,Sarangi02}. For a well-worked inflationary scenario, the production of topological defects other than cosmic strings must be suppressed by many orders of magnitude. For extensive studies on the KKLMMT model and other types of brane inflation models, see, e.g., \\cite{Tye05,Baumann06,Huang06,Huang:2006zu,Zhang06,Bean:2007hc,Lorenz:2007ze,Alabidi:2008ej}.  Recently, the WMAP team released the five-year data \\cite{Komastu08}. For the $\\Lambda $CDM model, WMAP five-year data show that the index of power spectrum satisfies\\footnote{The pivot scale is taken to be $k_{\\rm pivot}=0.002 {\\rm Mpc}^{-1}$.} \\begin{equation} n_{s}=0.963_{-0.015-0.028}^{+0.014+0.029}~~~(1\\sigma, 2\\sigma~ \\text{CL});  \\label{nsnumber1} \\end{equation} combining WMAP with SDSS and SNIa, the result is \\begin{equation} n_{s}=0.960_{-0.013-0.027}^{+0.014+0.026}~~~(1\\sigma, 2\\sigma~ \\text{CL}),  \\label{nsnumber2} \\end{equation} which is a little bluer than the WMAP three-year result, though the red power spectrum is still favored at the level of $3\\sigma $ CL. Meanwhile, the running of the spectral index is not favored anymore. With WMAP five-year data only, the running of the spectral index is\\footnote{In this case, the result of the spectral index is $n_s=1.031^{+0.054}_{-0.055}$.} \\begin{equation} \\alpha _{s}=\\frac{dn_{s}}{d\\ln k}=-0.037_{-0.028}^{+0.028}~~~(1\\sigma~\\text{CL}). \\label{runnumber1} \\end{equation} combining with SDSS and SNIa data, the result is\\footnote{In this case, the result of the spectral index is $n_s=1.022^{+0.043}_{-0.042}$.} \\begin{equation} \\alpha _{s}=\\frac{dn_{s}}{d\\ln k}=-0.032_{-0.020}^{+0.021}~~~(1\\sigma~\\text{CL}). \\label{runnumber2} \\end{equation} It is notable that the result of the spectral index in five-year WMAP is slightly larger than that in three-year WMAP. More evidently, the result of the running of the spectral index in five-year WMAP is much smaller than that in three-year WMAP. Finally, let's see the result of the amplitude of the primordial gravitational waves. When combining WMAP data with SDSS and SNIa data, the tensor-to-scalar ratio $r$ is limited under a much tighter bound:\\footnote{In this case, the result of the spectral index is $n_s=0.968\\pm 0.015$.} \\begin{equation} r<0.20~~~(95\\%~\\text{CL}).  \\label{tensornumber1} \\end{equation}  Based on the WMAP three-year results, Huang et al. \\cite{Huang06} investigated brane inflation models and showed that the KKLMMT model cannot fit WMAP+SDSS data at the level of one standard deviation and a fine-tuning (at least eight parts in a thousand) is needed at the level of two standard deviations. Now, the WMAP data are updated to the five-year ones, so it is meaningful to see how the status of brane inflation is affected by the arrival of the new WMAP data. Therefore, in this paper, we address this issue, i.e., we revisit brane inflation with the WMAP five-year results.  This paper is organized as follows: In section \\ref{sec:Dp}, we discuss a simple brane inflation model neglecting the problem of dynamic stabilization. In section \\ref{sec:kklmmt}, we focus on the KKLMMT model and investigate to what degree the model is fine-tuned under the WMAP five-year data. The conclusion is given in the last section. ",
        "conclusions": "In this paper, we have studied brane inflation with the WMAP five-year data. We first considered a general brane inflation model in which the problem of dynamic stabilization is neglected. Furthermore, we examined the KKLMMT model by using the WMAP five-year results.  For the general inflation model, we show that, according to the WMAP five-year data, the model survives at the level of 1$\\sigma$. This conclusion is different from that of a previous work \\cite{Huang06} which is based on the analysis of the WMAP three-year data. In \\cite{Huang06}, with the WMAP three-year data, it was indicated that the brane inflation model for $n=4$ (D3-brane case) is near the boundary of the WMAP3 only and cannot fit the WMAP3+SDSS data at the level of 1$\\sigma$; the model with $n=2$ (D5-brane case) is outside the range allowed by WMAP3 only or WMAP3+SDSS at the level of 1$\\sigma$. Therefore, we find that the WMAP five-year data save the brane inflation model to the level of $1\\sigma$.  For the KKLMMT model, we first discuss how the various observables can bring constraints on the parameter $\\beta$ of the model. Then, we compare the model to the WMAP five-year results. In \\cite{Huang06}, by using the WMAP three-year data, the authors find that the KKLMMT model cannot fit WMAP3+SDSS data at the level of 1$\\sigma$ and a fine-tuning, at least eight parts in a thousand, is needed at the level of 2$\\sigma$. However, in this work, we show that the KKLMMT model is consistent with both the data of WMAP5 only and of WMAP5+BAO+SN at the level of $1\\sigma$. Moreover, comparing to the WMAP five-year results, we find that the value of the parameter $\\beta$ is relaxed to ${\\cal O}(10^{-2})$ at the level of 2$\\sigma$. Therefore, the problem of fine-tuning of $\\beta$ is alleviated to a certain extent when confronting the WMAP five-year data. This is definitely a good news for the KKLMMT model.  This paper only discusses the the simplest realistic brane inflationary model, namely the KKLMMT model, in which the usual slow-roll scenario is employed, but does not involve the Dirac-Born-Infeld (DBI) scenario, in which the rolling of the inflaton is albeit slow but relativistic. In the DBI model, large tensor mode and/or non-Gaussianity may emerge, providing a possible stringy signature. The DBI model has been investigated in detail in \\cite{Alabidi:2008ej} by using the WMAP five-year data."
    },
    "0812/0812.2561_arXiv.txt": {
        "abstract": "{}{We present new results from our ongoing multiplicity study of exoplanet host stars, carried out with SofI/NTT. We provide the most recent list of confirmed binary and triple star systems that harbor exoplanets.}{We use direct imaging to identify wide stellar and substellar companions as co-moving objects to the observed exoplanet host stars, whose masses and spectral types are determined with follow-up photometry and spectroscopy.}{We found two new co-moving companions of the exoplanet host stars HD\\,125612 and HD\\,212301. HD\\,125612\\,B is a wide M4 dwarf (0.18\\,$M_\\odot$) companion of the exoplanet host star HD\\,125612, located about 1.5\\,arcmin ($\\sim$4750\\,AU of projected separation) south-east of its primary. In contrast, HD\\,212301\\,B is a close M3 dwarf (0.35\\,$M_\\odot$), which is found about 4.4\\,arcsec ($\\sim$230\\,AU of projected separation) north-west of its primary.}{The binaries HD\\,125612\\,AB and HD\\,212301\\,AB are new members in the continuously growing list of exoplanet host star systems of which 43 are presently known. Hence, the multiplicity rate of exoplanet host stars is about 17\\,\\%.} ",
        "introduction": "For more than a decade, radial-velocity and photometric planet search campaigns have indirectly identified more than three hundred exoplanets that revolve around mostly sun-like stars, in the solar neighborhood. The majority of these exoplanet host stars are isolated single stars, but during recent years more and more of them turned out to be the brighter primary component of stellar systems, identified by ongoing multiplicity studies. These studies were carried out with seeing limited near infrared imaging \\cite[see e.g.][2004b, 2005b, 2006a, 2007a, 2007b]{mugrauer2004a}\\nocite{mugrauer2004b}\\nocite{mugrauer2005b}\\nocite{mugrauer2006a}\\nocite{mugrauer2007a}\\nocite{mugrauer2007b}, high contrast diffraction limited AO observations \\citep[see e.g. Patience et al. (2002), Luhman \\& Jayawardhana (2002), Chauvin et al. 2006,][and most recently Eggenberger et al. 2007]{neuhaeuser2007}\\nocite{eggenberger2007}\\nocite{chauvin2006}\\nocite{patience2002}\\nocite{luhman2002}\\nocite{neuhaeuser2007}, as well as from space \\citep{luhman2007}. In addition, data from visible and infrared all sky surveys like POSS or 2MASS are used to identify new companions of exoplanet host stars \\citep[as reported e.g. by Bakos et al. 2006, Raghavan et al. 2006, or most recently by][]{desidera2007} \\nocite{raghavan2006}\\nocite{bakos2006}.  Most of the detected stellar companions of exoplanet host stars are low-mass main sequence stars with projected separations of between a few tens up to more than 10000\\,AU. In a few cases the companions themselves turned out to be close binaries, i.e. these systems are hierarchical triples \\citep[see][for a summary]{mugrauer2007a}. The closest of these systems presently known is HD\\,65216\\,A+BC, with a projected separation of about 250\\,AU \\citep{mugrauer2007b}. Also white dwarfs have been identified as companions of exoplanet host stars, suggesting that exoplanets can survive the post main sequence evolution of a nearby star \\citep[e.g. Gl\\,86\\,B, $\\sim$20\\,AU, see ][for more details]{mugrauer2005a}. Later on, the first directly imaged substellar companion of an exoplanet host star, the T7-T8 brown dwarf HD\\,3651\\,B, was discovered \\citep[see Mugrauer et al. 2006b, and][]{liu2007}\\nocite{mugrauer2006b}. In addition to these imaging surveys a dynamical characterization of the closest exoplanet host binaries, like Gl\\,86\\,AB \\citep[see][]{lagrange2006}, or $\\gamma$\\,Cep\\,AB \\cite[see][]{neuhaeuser2007}, is being carried out to determine the full set of their orbital elements.  All these efforts will help to reveal the true impact of stellar multiplicity on the formation process of planets and the evolution of their orbits. For recent statistical studies we refer here to e.g. \\cite{mugrauer2007c} or \\cite{bonavita2007}.  In this paper we present new results of our ongoing multiplicity study carried out at La Silla observatory with SofI/NTT. We detected two new low-mass stellar companions of the exoplanet host stars HD\\,125612 and HD\\,212301. Our SofI astro- and photometry is described in section\\,2, the follow-up spectroscopy in section\\,3. The properties of the newly found exoplanet host binaries are described in section\\,4, where the SofI detection limits also are presented and the separation space of possible, so far undetected, additional companions is discussed. A list of all presently known and confirmed exoplanet host star systems is given in the Appendix. ",
        "conclusions": "As described in detail in the last sections HD\\,125612\\,B is a wide M4 dwarf companion with a mass of $\\sim$0.18\\,$M_{\\odot}$ which is separated by about 4750\\,AU from its primary, the exoplanet host star HD\\,125612\\,A. In contrast, HD\\,212301\\,B is a close M3 dwarf (0.35\\,$M_{\\odot}$) companion, located at a separation of about 230\\,AU from its primary.  The typical SofI detection limit for both exoplanet host stars is shown in Fig.\\,\\ref{limit_hd212}. At an angular separation of less than 1.5\\,arcsec ($\\sim$80\\,AU) saturation occurs, i.e. companions cannot be detected in this region. According to the evolutionary models from \\cite{baraffe1998} all stellar companions ($mass>0.078\\,M_{\\rm Jup}$) are detectable beyond 6\\,arcsec ($\\sim$320\\,AU) around both stars, for assumed ages between 1 and 5\\,Gyr. In the background limited region beyond about 15\\,arcsec ($\\sim$ 800\\,AU) a sensitivity of H=18\\,mag ($S/N=10$) is reached, i.e. companions with absolute magnitudes down to $M_{\\rm H}=14.4$\\,mag can be imaged in this region. This allows the detection of brown dwarf companions with masses down to 37 to 65\\,$M_{\\rm Jup}$. Companions with angular separations of up to about 143\\,arcsec ($\\sim$ 7500\\,AU of projected separation) can be detected together with the exoplanet host stars in our SofI images.  \\begin{figure} [htb] \\resizebox{\\hsize}{!}{\\includegraphics{0639fig7.eps}}\\caption{The average detection limit ($S/N=10$) in our SofI images of the exoplanet host stars HD\\,125612\\,A and HD\\,212301\\,A plotted for a range of angular separations up to 25\\,arcsec. The detected co-moving companion HD\\,212301\\,B is indicated as a black square.}\\label{limit_hd212} \\end{figure}  HD\\,125612 was observed only in the 1st epoch with SofI. A 2nd epoch follow-up imaging could not be carried out in 2008 due to bad weather. For HD\\,212301 we obtained a 2nd epoch image in January 2008 but only to confirm the detection of the companion HD\\,212301\\,B and to rule out that it is a fast moving foreground object e.g. a planetoid in the solar system. The expected total motion of the companion between both observing epochs is only about 40\\,mas, too small to be clearly detected with SofI in its large field mode. Hence, follow up SofI imaging is needed for both stars with sufficient epoch difference to check the companionship of all faint companion candidates not imaged by 2MASS.  Nevertheless, we can already exclude additional companions around both exoplanet host stars based on the 2MASS detection limit. According to \\cite{skrutskie2006} the 2MASS H-band limit is 15.1\\,mag at $S/N=10$, the limit for all detected sources with accurate astrometry. As determined by us, this limit is reached in the 2MASS images beyond about 20\\,arcsec ($\\sim$1060\\,AU) around both stars. According to the evolutionary models from \\cite{baraffe1998} and an assumed age of 1 to 5\\,Gyr companions with masses down to 63 and 74\\,$M_{\\rm Jup}$ can be detected in 2MASS. All stellar companions ($mass>0.078$\\,$M_{\\rm Jup}$) can be detected beyond about 13\\,arcsec ($\\sim$690\\,AU) around both stars.  In addition, we can also rule out companions that should not be on long-term stable orbits around the exoplanet host stars in both binary systems. According to \\cite{holman1999} only companions with semi-major axes smaller than the critical semi-major axis exhibit long-term stable orbits. In the case of the HD\\,125612\\,AB system (HD\\,125612\\,A ($\\sim$1.1\\,$M_{\\odot}$), HD\\,125612\\,B ($\\sim$0.18\\,$M_{\\odot}$), 4750\\,AU of projected separation) we expect the long-term stable zone to extend up to $\\sim$650\\,AU (12.4\\,arcsec of angular separation), assuming that HD\\,125612\\,B is on a circular orbit, whose semi-major axis corresponds to the observed projected separation of the companion. This value can be considered as an upper separation limit for additional long-term stable companions. In the case of an eccentric orbit of HD\\,125612\\,B the extent of this long-term stable zone should be smaller. Indeed, there is one faint companion candidate in our SofI image detected about 10\\,arcsec west of HD\\,125612\\,A. Follow-up imaging is needed to test the companionship of this faint companion candidate.  In the case of the HD\\,212301\\,AB system (primary mass of about 1.27\\,$M_{\\odot}$, secondary mass of $\\sim$\\,0.35\\,$M_{\\odot}$, with a projected separation of $\\sim$\\,230\\,AU) we derive an extent of the long-term stable zone of only $\\sim$\\,38\\,AU, i.e about 0.7\\,arcsec of angular separation. No objects are detected within this given angular radius around HD\\,212301\\,A in our SofI images.  The binaries HD\\,125612\\,AB and HD\\,212301\\,AB are two new members of the growing list of exoplanet host multiple star systems. Today 250 exoplanet host stars are known and ongoing multiplicity studies so far have found 43 of them to be components of a multiple star system (see the Appendix for summary). Hence, the multiplicity-rate of the exoplanet host stars is at least 17\\,\\%."
    },
    "0812/0812.1964_arXiv.txt": {
        "abstract": "{ We present a quantitative and relatively model-independent way to assess the radial structure of nearby AGN tori.  These putative tori have been studied with long-baseline infrared (IR) interferometry, but the spatial scales probed are different for different objects.  They are at various distances and also have different physical sizes which apparently scale with the luminosity of the central engine.  Here we look at interferometric visibilities as a function of spatial scales normalized by the size of the inner torus radius $\\Rin$.  This approximately eliminates luminosity and distance dependence and, thus, provides a way to uniformly view the visibilities observed for various objects and at different wavelengths.  We can construct a composite visibility curve over a large range of spatial scales if different tori share a common radial structure.  The currently available observations do suggest model-independently a common radial surface brightness distribution in the mid-IR that is roughly of a power-law form $r^{-2}$ as a function of radius $r$, and extends to $\\sim$100 times $\\Rin$.  Taking into account the temperature decrease toward outer radii with a simple torus model, this corresponds to the radial surface density distribution of dusty material directly illuminated by the central engine roughly in the range between $r^0$ and $r^{-1}$.  This should be tested with further data.  } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.3199_arXiv.txt": {
        "abstract": "We demonstrate that magnetic reconnection is not necessary to initiate fast CMEs. The Aly-Sturrock conjecture states that the magnetic energy of a given force free boundary field is maximized when the field is open. This is problematic for CME initiation because it leaves little or no magnetic energy to drive the eruption, unless reconnection is present to allow some of the field to escape without opening. Thus, it has been thought that reconnection must be present to initiate CMEs. This theory has not been subject to rigorous numerical testing because conventional MHD numerical models contain numerical diffusion, which introduces uncontrolled numerical reconnection. We use a quasi-Lagrangian simulation technique to run the first controlled experiments of CME initiation in the complete lack of reconnection. We find that a flux rope confined by an arcade, when twisted beyond a critical amount, can escape to an open state, allowing some of the surrounding arcade to shrink and releasing magnetic energy from the global field. This mechanism includes a true ideal MHD instability. We conclude that reconnection is not a necessary trigger for fast CME eruptions. ",
        "introduction": "Coronal mass ejections (CMEs), are large expulsions of magnetic field and plasma from the solar corona. The kinetic and gravitational potential energy contained in a CME is around $10^{31-32}$ ergs, making these events some of the most energetic in our solar system \\citep{canfieldenergy,forbesreview,hundhausenetal1994,low1990,low2001}. It is thought that CMEs derive their energy from the magnetic field of the solar corona because this field is the only possible source for such a large reserve of energy \\citep[e.g.][]{forbesreview,klimchuckreview,lowreview1996}.  The flux rope model is one possible pre-eruptive CME configuration. A flux rope is a length of magnetic field that has been twisted along its axis, often held in place in the corona by an overlying arcade or ambient field.  It is thought that cool photospheric plasma can become be trapped in the center of a flux rope, creating a solar filament or prominence \\citep{vanBallegooijen1989,ridgeway1991,Priest1989}. The flux rope configuration easily explains the clear three-part structure seen in many CMEs, specifically those associated with prominence eruptions \\citep{hudsonetal1999}. Because these structures are present in the low corona where magnetic field is strong and plasma density is low, they are magnetic-field-dominated. Recently, movies from \\emph{Hinode} have shown interesting dynamics that are not described by the flux rope model which imply that at least some prominences are not low-$\\beta$ \\citep{berger2008}, but the flux rope model remains useful. The lower coronal environment is also frequently modeled as being force-free because flow speeds are low, and $\\mathbf{J}\\times\\mathbf{B}$ is the dominant force in the equation of motion. Gravity is also frequently ignored because it is a factor of $\\sim5$ weaker than the magnetic forces. Low-$\\beta$ flux ropes are stable when the outward magnetic pressure force is balanced by an inward-directed tension force. In models, an exterior arcade field is often added to increase the tension force and keep the flux rope from simply expanding in length and width as twist is added. The approximate energy per unit length along the axis, $U$, stored in an unconfined Gold-Hoyle flux rope \\citep{goldhoyle1960} is given by \\begin{equation} U=\\frac{1}{8\\pi^{2}}\\frac{\\Phi^{2}b^{2}}{ln\\left(1+b^{2}R^{2}\\right)}\\label{eq:flux tube energy}\\end{equation} where $\\Phi$ is the magnetic flux, $R$ is the radius of the tube, and $b$ is the twist parameter such that $\\frac{B_{\\phi}}{B_{z}}=br$ \\citep{sturrock2001}. As the twist accumulates, $b$ increases, $\\Phi$ is constant and $R$ remains approximately constant so the total energy increases. When an unconfined flux rope anchored at both ends on the photosphere accumulates twist, its equilibrium state is expanded in length relative to the untwisted state so the twist per unit length does not necessarily increase with the total twist. Thus, flux ropes that are confined by an overlying arcade contain more energy because their length changes very little as $b$ increases.  The energy stored in the magnetic field is given by the volume integral of $B^{2}/8\\pi$, up to conversion factors. The minimum energy of a magnetic system with a given photospheric boundary occurs in the potential, or vacuum field, configuration. As the field is stressed away from this configuration due to photospheric movements, the energy is increased above the evolving potential state by an amount commonly referred to as the magnetic {}`free energy'. If the field reverts to the potential configuration, this free energy is released and in the case of solar active regions, is available to drive a CME. If the reconnection is localized and helicity is conserved, the lowest accessible energy state may not be potential, so the free energy is an upper boundary on the amount of energy that can be released.  There is a global cap on the amount of energy that the magnetic field can provide. The magnetic virial theorem asserts that the total pressure force can not exceed the tension force in a stable plasma environment \\citep{priestbook}. The related Aly-Sturrock conjecture states that the global magnetic energy of a force free field is at a maximum when the field is completely {}`open'. This refers to magnetic flux that is anchored at the solar surface and extends radially outward a significant distance so that, near the Sun, the field appears open \\citep{aly1984,aly1991,sturrock1991}.  Many CME observations show prominences lifting off of the surface of the Sun, expanding to several solar radii and leaving behind long radial field lines. If this conjecture is correct, then the implication is that CMEs which open large amounts of field must derive the bulk of their kinetic energy from sources other than the magnetic field because the field energy is actually greater in the post-CME configuration.  Order of magnitude analysis has shown, however, than the magnetic field is the only source of energy that can potentially drive a $10^{32}$erg CME \\citep{forbesreview}. This poses a significant problem to ideal CME models. In an azimuthally symmetric 2.5-D case, all of the field lines originally above a prominence-like feature would have to open to release the filament (Fig. \\ref{fig:2.5d vs 3d}a). Thus, in the 2.5-D case, to have an eruption which results in a net decrease of magnetic energy, reconnection must be present. There have been studies in which a 2.5-D field is shown to have magnetic energy exceeding the open field energy when mass loading is present, but it has not been demonstrated that these fields can erupt without reconnection \\citep{lowreview1996,fonglowfan,zhanglow}. In three dimensions, the flux rope is anchored in the photosphere, and the surrounding field can move away in the direction parallel to the flux rope axis (Fig. \\ref{fig:2.5d vs 3d}b). Reconnection is not necessary in the fully three-dimensional case, as not all of the field must be open to have an eruption, only the flux rope opens, and thus the eruption is not relevant to the hypothesis of Aly and Sturrock because some of the field remains closed \\citep{low1986}.  \\begin{figure} \\begin{centering} \\includegraphics{f1.jpg} \\par\\end{centering}  \\caption{Behavior of flux rope expansion in 2.5D (a) and 3D (b).} \\label{fig:2.5d vs 3d}   \\end{figure}  The full Aly-Sturrock conjecture has yet to be disproved. However, it is not usually relevant to 3-dimensional models and analysis for the reasons stated above. A more relevant question which has been asked, is whether a configuration with \\emph{some} open force free field can contain less energy than a configuration with the same boundary conditions that is fully closed \\citep{low1990}. This question has been addressed semi-analytically by \\citet{woldsonandlow1992} and \\citet{wolfson1993}, who showed that a fully closed field can contain magnetic free energy above the partially open field threshold, but they did not demonstrate a release mechanism for this energy. A demonstration of the ideal evolution of a field whereby free energy is stored then released, resulting in a partially open state, would once and for all eliminate the problem posed by the magnetic virial theorem and the Aly-Sturrock conjecture in a fully 3D case.  \\citet{sturrock2001} describe the possibility of driving CMEs with metastable magnetic fields. They analytically describe one system in particular as a known metastable state: the previously mentioned twisted flux rope under an overlying arcade. This system is metastable because it is stable (due to the confining arcade) against small perturbations, but the energy of the erupted flux rope is lower than the contained flux rope. If the rope is tightly wound, it can open and escape the arcade by herniating through it, leaving the deflated arcade near the footpoints. The amount of twist needed is not unreasonable for the solar surface. Analytically, for a flux rope that is ten times longer than its radius, the rope need only exceed 1.5 total turns about its axis to be in this metastable state \\citep{sturrock1991}.  Numerous simulations exist which model flux rope CME initiation of the metastable configuration described above \\citep[e.g.][]{aulanier2005,g+f04,titov+demoulin,torok+kliem2004,rousesevetal2003}. Most of these simulations have found that it is possible to herniate through the arcade, but they do not agree exactly on how much twist is needed, or how unstable the resulting configuration is after the onset of writhe (helical geometry in the central field line) in the flux rope.  Typically, these codes agree that the the critical twist needed to erupt is around 1.5 turns, and that it is possible to get an eruptive event by twisting the footpoints of a flux rope.  CME initiation with these flux ropes has been modeled both with and without reconnection as the intended primary destabilizing factor. Theories that do not involve reconnection are referred to as ideal {}``loss of equilibrium'' models \\citep[e.g.][]{rousesevetal2003}. The initial structure generally undergoes an ideal instability, such as the MHD kink instability, caused by a large amount of twist. Another possibility, if mass loading is critical in keeping the structure contained, is that mass displacement could upset the force balance and start the CME \\citep{klimchuckreview,fonglowfan,zhanglow}. Other theories, such as tether cutting \\citep{mooreroum} or {}``breakout'' \\citep{breakout}, explicitly include reconnection to decrease the strength of the overlying arcade. In models such as tether cutting and breakout, there are generally two stages of reconnection. Slow -- Sweet-Parker style \\citep{parker_sweetparker1963} -- reconnection occurs early in the evolution and destabilizes the system, allowing it to expand. This is often followed by fast -- Petcheck style \\citep{petschek1964} -- reconnection, which releases large amounts of energy in a short time and is believed to be the primary driver for fast, impulsive CMEs.  Essentially all existing numerical simulations of CME onset use Eulerian methods, in which a 3-D grid of values is used. With magnetic fields (and indeed all vector fields and flows), sharp gradients are not conserved because derivatives are represented as finite differences.  For magnetic fields, this means that the field will reconnect if gradients approach the size of the grid whether the modeler wishes it or not. Techniques such as adaptive mesh refinement can reduce the rate of numerical reconnection, but can not remove it altogether. Hence it is not possible to separate the effects of ideal MHD evolution and magnetic reconnection with an Eulerian grid code. This can be problematic insofar as the reconnection destabilizes a metastable system. By switching to a Lagrangian (field-aligned) formulation, we eliminate all reconnection, allowing study of ideal MHD instabilities \\citep{fluxonpaper}.  With our model, we are able to analyze simplified systems where topology is locked in and reconnection is not present. Note that we do not hypothesize that reconnection is not present in the Sun, only that to have a controlled numerical experiment, the effect of reconnection must be isolated, and we do this by eliminating it. Our method is unique in this way, and may offer insights that grid simulations can not. In particular, we are able to demonstrate the existence of a true MHD instability that can release free energy into a CME, even with no triggering reconnection. ",
        "conclusions": "Our simulations show that reconnection is not necessary to initiate a CME and that impulsive CMEs may be possible without explosive reconnection.  This theory is not new; it was originally published by \\citet{sturrock2001}, who describe the existence of metastable states, specifically a system similar to the one we have studied. Since then, other solar physicists have studied this system computationally \\citep[etc.]{g+f04,torok+kliem2004,aulanier2005}.  The results from these studies show that a highly twisted flux rope can herniate through a confining magnetic arcade and reconnect into a plasmoid, causing an eruption. However, this is the first study of this system in the complete lack of reconnection. In general, our results agree with those of other research groups.  The fact that many of these simulations, including ours, agree that about 1.5 total turns is needed to herniate through an arcade implies that reconnection is not greatly important to the overall stability of the system. If it were, then we would expect our ideal simulation to support significantly more twist and therefore release more energy than the dissipative simulations. The exact value of the critical twist will depend on the configurations of the system: strength of the arcade, width of the flux rope, the twist profile within the flux rope, etc. But even with these variables, we conclude that highly twisted flux ropes can not easily be confined by external field, even when the reconnection rate is extremely small.  Current research on the twist available photospheric fields indicates that there may be an excess of one full turn available in many active regions \\citep{leka}. This implies that many pre-eruptive active regions may be on the brink of an ideal instability when they flare or erupt, regardless of the eventual trigger mechanism.  Our results together with the results of \\citet{woldsonandlow1992} and \\citet{wolfson1993} can finally put to rest the concerns that the Aly-Sturrock conjecture have created over the initiation of CMEs. Although it has not been proven that a closed force free field can have more energy than a fully open one, that question is not relevant in the complex three dimensional system that is our Sun. A more appropriate question is whether a closed force free field can contain more energy than a configuration with some open field given the same lower boundary and connectivity. \\citet{woldsonandlow1992} began to answer that question with semi-analytic techniques and successfully showed that it was possible. We have proven that it is possible to transition between closed and partially open field while still releasing free energy without reconnection and without the need for gravitational or other non-magnetic confinement. At the beginning of the century there was, {}``still no model which demonstrates that a partly open magnetic field can be achieved solely by a loss of ideal MHD equilibrium or stability.''  \\citep{forbesreview}. Happily, this statement is no longer true."
    },
    "0812/0812.1649_arXiv.txt": {
        "abstract": "Two charge- and current neutral plasma beams are modelled with a one-dimensional PIC simulation. The beams are uniform and unbounded. The relative speed between both beams is 0.4c. One beam is composed of electrons and protons and one out of protons and negatively charged oxygen (dust). All species have the temperature 9 keV. A Buneman instability develops between the electrons of the first beam and the protons of the second beam. The wave traps the electrons, which form plasmons. The plasmons couple energy into the ion acoustic waves, which trap the protons of the second beam. A proton phase space hole grows, which develops through its interaction with the oxygen and the heated electrons into a rarefaction pulse. This pulse drives a strong ion acoustic double layer, which accelerates a beam of electrons to about 50 MeV, which is comparable to the proton kinetic energy. The proton distribution eventually evolves into an electrostatic shock. Beams of charged particles moving at such speeds may occur in the foreshock of supernova remnant shocks. This double layer is thus potentially relevant for the electron acceleration (injection) into the diffusive shock acceleration by supernova remnants shocks. ",
        "introduction": "The blast shell of a supernova remnant (SNR) can accelerate charged particles to cosmic ray energies through diffusive shock acceleration by the SNR shocks \\citep{Uchiyama,Reynolds,Hillas}. Particles are scattered upstream and downstream of a shock and some particles cross the shock repeatedly, gaining energy each time \\citep{D1,D2,D3,Luke1}. Diffusive shock acceleration requires a seed population of particles with energies well in excess of the thermal ones, since only such particles can cross the shock repeatedly and be scattered efficiently by the magnetohydrodynamic waves on either side of the shock. The presence of shock-accelerated electrons is evidenced by the radio emissions of SNRs. It is not clear yet how the electrons can reach the energy threshold for diffusive shock acceleration; that requires that their kinetic energies are comparable to those of the ions. Some electrons must be accelerated locally, that is close to the SNR shock or close to one of its precursors \\citep{Reynolds,Hillas}.  This injection of electrons is thought to be accomplished by the interplay of energetic ions with electrons and dust \\citep{Luke2} by means of plasma instabilities \\citep{Cargill}. The instabilities occur in the upstream plasma, just ahead of the SNR shock. This region is called the foreshock. One source of the ion beams is the downstream plasma. Ions can leak through the shock boundary and outrun it \\citep{Leaking}. The shock can also reflect a substantial fraction of the upstream ions. If the reflection is specular, then the ions can cross the foreshock at twice the shock speed. Reflected and leaked ion beams can be observed in situ at the Earth bow shock \\citep{Foreshock}. Particle-in-cell (PIC) simulations of fast plasma shocks suggest, that they may also occur at SNR shocks when the upstream magnetic field is orthogonal to the shock normal \\citep{Hoshino,Lem2,Chapman,Amano}. Precursor shocks move ahead of the main SNR shock. They can reach mildly relativistic speeds for the most extreme supernova explosions \\citep{Kulkarni}. Such shocks can also reflect ions in the presence of a quasi-parallel magnetic field \\citep{NewApJ}.  We focus in this work on ion beam instabilities in the foreshock, but we do not model the shock itself. We rely on a foreshock model that has been widely used, e.g. in the works by \\citet{H1,H2} and by \\citet{TwoBeam,Die}. This model considers an ion beam, which moves through the upstream plasma, consisting of cool electrons and protons. The system is spatially uniform. Often a second, counter-propagating ion beam is introduced, which cancels the current of the first one. This second beam can be due to ions that have been reflected by a perpendicular shock and which return to the shock, after they have been rotated by the upstream magnetic field \\citep{McClements}. Here a beam of co-moving negatively charged oxygen ions cancels the current of a proton beam.  This foreshock model can be used if all processes develop on scales well below an ion Larmor radius and an inverse ion cyclotron frequency in the case of a perpendicular shock. The foreshock model may be representative over larger scales otherwise, because ions can expand much farther if they move along the magnetic field, as it is observed at the Earth's bow shock. This foreshock model is not self-consistent, because it decouples the foreshock physics from that of the shock. It also assumes a plasma that is spatially uniform, while the shock reflected ion beams vary in their mean speed and temperature. However, we can examine with this model the acceleration of the upstream electrons as a function of the bulk parameters, such as the beam temperatures and mean speeds and the background magnetic field. We can also separate individual processes from the global shock dynamics, that involves coupling over many scales \\citep{h3}. This simplifies their identification and interpretation. The periodic boundary conditions also allow us to reduce the box size, making PIC simulations computationally efficient even for realistic ion-to-electron mass ratios. The insight gained from such simulations can then guide us in the selection of plasma parameters for the more realistic and demanding PIC simulations of full shocks.  Simulation studies employing this foreshock model have revealed that proton beam-driven instabilites are not efficient electron accelerators. An obstacle is the mass difference between the electrons and the protons. Most plasma instabilities saturate by accelerating the electrons. The fields are too weak to accelerate the ions and the heated electrons quench most of the known plasma instabilities, e.g. through Landau damping. Often only a few per cent of the ion energy is transfered to the electrons \\citep{H1,H2,TwoBeam,Die}. Initially, this is also the case for ion beams that expand into the upstream plasma \\citep{DM} prior to the formation of a shock. Certain mechanisms \\citep{VV,Kat} can, in principle, accelerate the electrons to relativistic speeds, but this remains yet to be demonstrated with self-consistent PIC simulations. However, the electrons, that have been accelerated by an instability like the Buneman instability (BI) \\citep{Buneman}, are often spatially bunched, for example in form of phase space holes \\citep{Roberts,Schamel1,Schamel2,TwoBeam}. Electron bunches or plasmons have a much higher inertia than electrons and they modulate the ions to give ion acoustic waves. The ion acoustic waves can trigger further processes, one of which we consider in this work.  We investigate a plasma consisting of two spatially uniform beams. One beam is composed of protons and electrons and the second one of protons and oxygen carrying a single negative charge. Each of the beams is charge and current neutral. Both beams move at a relative speed of 0.4c. All species have the temperature 9 keV. These initial conditions are the simplest ones that take into account three particle species and they permit a first assessment of the impact of heavier ions. Section 2 discusses in more detail these initial conditions and their potential relevance for a foreshock plasma. Section 3 describes the simulation results, which can be summarized as follows. A BI develops between the electrons and the protons of the second beam. The BI saturates by forming electron phase space holes and these plasmons feed wave energy into the ion acoustic wave \\citep{Mendonca,Mendonca2}. Proton phase space holes develop, once the ion acoustic waves are strong enough to trap the beam protons. The electrons eventually thermalize and they reach a mildly relativistic thermal speed, which is stabilizing the system for some time. The BI plays no further role after the electron thermalization. We examine in detail a proton phase space hole that grows into being a proton density pulse with a negative density \\citep{Infeld1,Infeld2}. The growth of this proton phase space hole may be caused by its interaction with the negative oxygen \\citep{Bengt} or by the disruption of the electron current \\citep{Smith}. The pulse triggers a strong ion acoustic double layer and it evolves subsequently into an electrostatic shock. The electron phase space structure and evolution of this double layer resembles that, found by a recent simulation study by \\citet{Newman}. That study addressed the electron acceleration in the Earth's auroral ionosphere. However, our simulation yields the much stronger electron acceleration that is necessary for injecting the electrons into the diffusive shock acceleration at SNR shocks. The proton phase space distribution also differs from the standard picture of double layers outlined, for example, in the review papers by \\citet{Smith} and by \\citet{RaaduRas}. The trapped and free streaming protons in our simulation are well-separated in velocity, which is not generally the case for double layers. The substantial free energy stored in the difference of the mean velocities of both beams causes the extreme electron acceleration up to 50 MeV. We discuss our simulation results in section 4, we relate it to previous investigations of double layers and we discuss the necessary future research. ",
        "conclusions": "This work has examined secondary instabilities triggered by the Buneman instability (BI) \\citep{Buneman} with a PIC code \\citep{Eastwood}. One spatial x-direction has been resolved, which implies that the wave spectrum driven by the BI is monodirectional. This limits the realism of the simulation but it firstly allows us to resolve all relevant spatio-temporal scales and, secondly, it facilitates the identification of the relevant physical processes. The large box has ensured that the periodic boundary conditions have not affected the results. We have considered two unmagnetized, interpenetrating, uniform and unbounded beams that have been composed of three particle species. The beam 1 has contained protons and electrons and the beam 2 has been formed by protons and by oxygen with a single negative charge. Each beam has been charge- and current neutral. The relative speed between both beams has been set to 0.4c and the initial temperature of all species has been selected to be 9.1 keV.  The BI has saturated by the trapping of electrons \\citep{Roberts}, which groups the electrons into phase space holes that give rise to charge density modulations. The latter can behave as quasi-particles and are thus referred to as plasmons. We could match the mean (group) speed of these plasmons with that of the protons of beam 2 by selecting negatively charged oxygen rather than electrons. Such a resonance facilitates the coupling of the plasmon energy into the ion acoustic wave mode \\citep{Mendonca}.  Initially the equilibrium between the electric fields and the electrons has been preserved \\citep{Schamel1,Schamel2} and the ion beams had lost only a small fraction of their energy; the electrons moved consequently only at mildly relativistic speeds. The thermal anisotropy of the electron distribution would trigger an instability similar to that proposed by \\citet{Weibel}, resulting in the growth of magnetic fields and multi-dimensional field structures that cannot be represented by a 1D PIC simulation. The Weibel instability may, however, not be important. The proton phase space holes form simultaneously with the collapse of the electron phase space holes that gives rise to the thermal anisotropy. The electromagnetic fields of the Weibel instability may not be strong enough to affect the interplay between the already present proton phase space holes and the oxygen and might be negligible.  As the time progressed, two large density rarefaction pulses \\citep{Infeld1,Infeld2} have developed in the protons of the beam 2. We have examined the larger one in more detail. The long time between the formation of the proton phase space hole and the pulse has suggested that it has been caused by the oxygen response. We have found that the proton density pulse grows, once the electric field of the proton phase space hole has displaced a substantial fraction of the oxygen ions. A related effect has been observed by \\citet{Bengt}. There it was shown that an electron phase space hole was ejected by a density depletion in the background ions, which its electric fields excavated. Here, the consequence of the oxygen depletion is not the acceleration of the phase space hole but its growth. The electrons may also contribute to the growth of the proton phase space hole \\citep{Dupree}.  The growth of the proton density pulse has resulted in an electron double layer that accelerated the electrons to ultrarelativistic speeds. To the best of our knowledge this is the first time that such an electron acceleration could be observed in PIC simulations of nonrelativistic beam instabilities. The electrons have reached a peak energy of 50 MeV. This immense electron acceleration implies, that the 1D PIC simulation geometry is not appropriate any more. Firstly, the electrons moving at such speeds drive the oblique mode instability \\citep{BretRel}, rather than the two-stream instability. The oblique modes are excluded by the simulation geometry. Secondly, the extreme electron temperature anisotropy resulting from the thermalization of the ultrarelativistic electron beam would drive a very strong Weibel-type instability. The subsequent magnetic field growth would imply that electrons emit synchrotron radiation on top of the bremsstrahlung \\citep{Brems1,Brems2}. The balance of these processes \\citep{Schlickeiser} requires a three-dimensional geometry and a method to deal with the radiation, which is not possible with our PIC code.  At the same time, this electron double layer constitutes a plasma structure that can easily accelerate the electrons to a speed, which allows them to enter the diffusive shock acceleration mechanism. This could be achieved with beam speeds that are not much higher than the flow speeds of supernova blast shells \\citep{Vink,Fransson}. The oxygen ions we have modelled can also be found in such plasmas, although not in the high density and in the ionisation number we have assumed here \\citep{Oxygen}. Our negative charge is more representative of dust, which we also find in supernova blast shells \\citep{Meikle}. Our findings can thus not be applied directly to a supernova scenario. However, the formation of the electron double layer in the presence of heavier ions clearly demonstrates the need to introduce these into PIC simulations, which so far only consider one ion species \\citep{Amano,Chapman}.  Future simulation work must address the minimum speed, at which the double layer can form. This has to be done in form of parametric simulation studies, because it is not clear for which combinations of the beam speeds and temperatures the double layers can form. Existing analytic studies \\citep{Verheest} do not cover the case we consider here. The effects of the obliquely propagating modes driven by the initial beam instability have to be examined, as well as those due to the Weibel instability driven by the thermal anisotropy of the electron distribution after the initial BI has saturated. This requires at least two-dimensional PIC simulations. Finally, other ion species and plasma compositions have to be examined to assess how easily double layers can form. This is important, because the ion composition is not easily obtained \\citep{Luke2} and not unique.  The modelling of double layers and their consequences is also of interest for other astrophysical scenarios. Ion double layers, which are characterized by ion distributions that resemble the electron distribution here, may regulate the mass inflow onto neutron stars \\citep{Layer1} and accelerate particles in accretion discs \\citep{Layer2}. They may also contribute to solar coronal ejections \\citep{Alfven}."
    },
    "0812/0812.3827_arXiv.txt": {
        "abstract": "{The IUE database provides a large number of UV high and low-resolution spectra of RS CVn-type stars from 1978 to 1996. In particular, many of these stars were monitored continuously during several seasons by IUE. }{ Our main purpose is to study the short and long-term chromospheric activity of the RS CVn systems most observed by IUE: HD 22468 (V711 Tau, HR 1099, K1IV+G5V), HD 21242 (UX Ari, K0IV+G5V) and HD 224085 (II Peg, K2IV).}{We first obtain the Mount Wilson index $S$ from the IUE high and low-resolution spectra. Secondly, we analyse with the Lomb-Scargle periodogram the mean annual index $\\langle S\\rangle$ and the amplitude of the rotational modulation of the index $S$.}{For HD 22468 (V711 Tau, HR 1099), we found a possible chromospheric cycle with a period of $\\sim$18 years and a shorter cycle with a period of $\\sim$3 years, which could be associated to a chromospheric ``flip-flop'' cycle. The data of HD 224085 (II Peg) also suggest a chromospheric cycle of $\\sim$21 years and a flip-flop cycle of $\\sim$9 years.  Finally, we obtained a possible chromospheric cycle of $\\sim$7 years for HD 21242 (UX Ari).}{} ",
        "introduction": "One of the principal diagnostics for solar and stellar chromospheric activity is the emission in the Ca \\scriptsize{II} \\normalsize H and K resonance lines (at 3968 and 3934 \\AA). In particular, the largest dataset of activity measurements available at present, comprising observations of over 2000 stars, is the one obtained with the Mount Wilson HK spectrophotometers, which have been used since 1966 to measure high-precision Ca \\scriptsize{II} \\normalsize H+K fluxes. As an indicator of stellar activity, an index $S$ has been defined as the ratio between the line-core fluxes and the flux in the continuum nearby.  Since the Mg \\II  h and k lines (at 2803 and 2796 \\AA\\-)  are formed in a similar way to the Ca \\scriptsize{II} \\normalsize lines, they are  also good indicators of the thermal structure of stellar atmospheres, specially from the high photosphere to the upper part of the chromospheric plateau.  The \\emph{International Ultraviolet Explorer (IUE)} provides a large database of low and high-resolution UV spectra in the band 1150-3350 \\AA\\-, from 1978 to 1996, which can be used to study long-term stellar activity.  In a previous work (\\citealt{2008A&A...483..903B}, hereafter Paper I) we obtained a colour dependent relation between the Mount Wilson index $S$ and Mg \\II line-core surface fluxes derived from IUE high-resolution spectra for F5 to K3 main sequence stars.  As an application of this calibration, we computed the Mount Wilson index for all the dF to dK stars which have high-resolution IUE spectra. In that paper we also combined IUE observations with those obtained by \\cite{1996AJ....111..439H}, and by our group (see \\citealt{2004A&A...414..699C} and \\citealt{2007astro.ph..3511C}). For some of the most frequently observed main sequence stars, we combined the Mount Wilson index $S$ from the IUE spectra, together with the ones derived from visible spectra, covering the period between 1978 and 2005. Analysing the data with the Lomb-Scargle periodogram, we were able to confirm the chromospheric activity cycle of $\\epsilon$ Eri (HD 22049) and $\\beta$ Hydri (HD 2151) with periods of $\\sim$5 and $\\sim$12 years respectively, and we found a magnetic cycle in $\\alpha$ Cen B (HD 128621) of $\\sim$8 years.    Similarly,  the purpose of the present study is to calibrate IUE low-resolution observations of Mg \\II h+k lines with the Mount Wilson index $S$, as a complement of the work presented in Paper I, and to apply this calibration to study the activity in RS CVn stars.   RS CVn-type stars are well known due to their strong chromospheric plages, coronal X-ray, and microwave emissions, as well as strong flares in the optical, UV, radio, and X-ray.  These stars belong to a class of close detached binaries where the more massive primary component is a G-K giant or subgiant and the secondary is a subgiant or dwarf of spectral classes G to M.  The magnetic activity signatures are generally dominated by the primary star of the system. Since these stars are fast rotators enforced by tidal synchronization, they are, in general, more active than single stars of similar mass and age.  The remarkable activity and high luminosity of these stars make them favourite targets for light curve modeling \\citep{1986A&A...165..135R,1995A&A...293..107D,1999A&A...350..626B,2007ApJ...659L.157B}, Doppler imaging \\citep{1992A&A...265..682D,1997MNRAS.291..658D,1999MNRAS.302..457B,2000MNRAS.318.1171K} and spectral line analysis \\citep{1987A&A...186..241W,1996ApJ...470.1172D,2000A&AS..146..103M,2001A&A...365..128P}. In general, most long-term stellar activity studies of RS CVn-type stars are based on easily detected optical photometric variations produced by their long-lived large spots.  Recently, \\cite{2008A&A...480..495M} presented the results of a long-term photometric monitoring project, where he carefully  studied the magnetic activity of 14 late-type active components of close binary systems and its evolution on different time scales. \\cite{1998A&A...332..541L,2006A&A...455..595L} and \\cite{2007ApJ...659L.157B} presented similar long-term studies for individual RS CVn stars.  In general, these studies concluded that the mean activity level of RS CVn-type stars presents cyclic variations similar to the 11-year solar cycle. Moreover, they found that large starspots tend to appear always at particular longitudes, called ``active longitudes'' (ALs), which are generally separated by 180$^\\circ$ on average and differ in their level of activity.  \\cite{1998A&A...336L..25B} found that the dominant activity switches periodically from one active longitude to the other in most RS CVn stars. This type of stellar activity cycle, known as ``flip-flop cycle'', is mainly related to the non-axisymmetric redistribution of the spotted area.    The IUE database provides a large number of UV high and low-resolution spectra of RS CVn-type stars. Furthermore, IUE monitored these stars continuously during several seasons. In most cases, the observations were aimed at uniform coverage with respect to orbital phase over more than one rotation period in each season (e.g. \\citealt{1996ApJ...470.1172D}, \\citealt{1999A&A...350..571B}). In this work we analyse the short and long-term chromospheric activity of the non-eclipsing RS CVn stars most observed by IUE and we compare the chromospheric and photospheric patterns of variability.     In particular, in Section \\S\\ref{sec.calibbaja}, we derive a relation to determine the Mount Wilson index from the Mg \\II line-core fluxes measured on IUE low-resolution spectra. As an application of this calibration, in Section \\S\\ref{sec.rscvn} we study the Mount Wilson indices we derived from both IUE high and low-resolution spectra of the RS CVn stars HD 22468 (V711 Tau, HR 1099),  HD 21242 (UX Ari) and HD 224085 (II Peg). ",
        "conclusions": "The main purpose of this paper is to study the short and long-term chromospheric activity of the RS CVn-type stars most observed by IUE.  In the first part of this work we calibrate the Mg \\II fluxes observed in  low-resolution spectroscopic observations available in the IUE archives with the Mount Wilson index $S$, which allows us to use these observations for systematic studies of magnetic activity in late-type stars. In particular, we found that the Mg \\II line-core fluxes derived from IUE low-resolution spectra could be only considered as an activity indicator in stars cooler than G3.   To complement the calibration between the Mount Wilson index $S$ and the IUE high-resolution Mg \\II line-core fluxes we obtained in Paper I, we analysed a set of 96 nearly simultaneous observations, which belong to 11 G3 to K3 main sequence stars, of the index $S$ and the Mg \\II fluxes derived from IUE high-resolution spectra re-sampled into the low-resolution wavelength domain.  We obtained a $B-V$ colour dependent relation between the index $S$ and the IUE low-resolution Mg \\II line-core fluxes.  In the second part of this work, we studied the long and short-term chromospheric activity of the non-eclipsing RS CVn stars most observed by IUE: HD 22468 (V711 Tau, HR 1099, K1IV+G5V), HD 21242 (UX Ari, K0IV+G5V) and HD 224085 (II Peg, K2IV).  To do so, we analysed with the Lomb-Scargle periodogram the mean annual index $\\langle S\\rangle$ and the amplitude of the rotational modulation of $S$ of each star.  We obtained a possible chromospheric activity cycle in HD 22468 with a period of $\\sim$18 years and a chromospheric flip-flop cycle of period $\\sim$3 years. We also found a flip-flop like cycle with a period of $\\sim$9 years in the HD 224085 data.  These results are in agreement with the ones reported in the literature.  On the other hand, we also detected a cyclic chromospheric pattern in the mean annual index of HD 224085 and HD 21242 with periods of $\\sim$21 and $\\sim$7 years respectively. %"
    },
    "0812/0812.4400_arXiv.txt": {
        "abstract": "% This paper reviews some of the developments that over the last 10 years have allowed us to go from deciphering the physical origin of several of the enigmatic features of the second solar spectrum to discovering unknown aspects of the Sun's hidden magnetism via sophisticated radiative transfer modeling. The second solar spectrum is the observational signature of radiatively induced quantum coherences in the atoms and molecules of the solar atmosphere. Magnetic fields produce partial decoherence via the Hanle effect, giving rise to fascinating observable effects in the emergent spectral line polarization. Interestingly, these effects allow us to ``see\" magnetic fields to which the Zeeman effect is blind within the limitations of the available instrumentation. In the coming years, the physical interpretation of observations of the spectral line polarization resulting from the joint action of the Hanle and Zeeman effects might lead to a new revolution in our empirical understanding of solar magnetic fields. ",
        "introduction": "\\begin{quote} {\\small ``{\\em Recent accurate measurements of the wavelength dependence of linear polarization of the solar limb radiation has resulted in the discovery of what may be called the second solar spectrum, with dozens of polarization ($Q/I$) features seen both `in absorption' and `in emission', i.e., with correspondingly smaller and larger polarization than in the adjacent continuum.}\"  (Ivanov 1991)}. \\end{quote}  Before 1980 ``non-magnetic\" scattering line polarization had been detected only in a few spectral lines. Except perhaps for the Ca {\\sc i} line at 4227 \\AA, the shapes of the fractional linear polarization ($Q/I$) profiles were practically unknown \\citep[e.g.,][]{trujillo_2_bruckner63,trujillo_2_stenflo74,trujillo_2_wiehr75}. In 1980 \\cite*{trujillo_2_stenflos80} published the results of a first systematic exploration of the linearly polarized spectrum produced by scattering processes in the solar atmosphere. Of particular interest was finding  that the $Q/I$ pattern across the D-lines of Na {\\sc i} changes sign between the D$_2$ and D$_1$ lines\\footnote{This paper follows the usual convention of choosing the reference direction for $Q>0$ along the parallel to the observed solar limb.}. Even more interesting was  discovering the first hints of a similar sign reversal in the $Q/I$ pattern of the K and H lines of Ca {\\sc ii}, which are 35 \\AA\\ apart. In a theoretical paper \\cite{trujillo_2_stenfloHK80} could show that this is due to quantum mechanical interference between the two upper atomic states with total angular momentum $J=3/2$ and $J=1/2$.  It was however in 1983 when the first surveys of scattering line polarization were published. The observations of the UV spectral region between 3165 \\AA\\ and 4230 \\AA\\ were obtained using the vertical spectrograph of the Kitt Peak McMath Telescope \\citep*[]{trujillo_2_stenflo83a}, while the survey between 4200 and 9950 \\AA\\ was achieved with the McMath Telescope using the Fast Fourier Spectrometer (FTS) as a polarimeter \\citep[]{trujillo_2_stenflo83b}. Although the typical spectral resolution was 0.1 \\AA\\ (for the UV region) and the polarimetric sensitivity ${\\sim}10^{-3}$, such observational surveys revealed a number of interesting linear polarization signals, most of which were considered as enigmatic. In particular, it was found that many spectral lines for which $\\epsilon_Q{\\approx}0$ (with $\\epsilon_Q$ the contribution of scattering processes to the emissivity in Stokes $Q$) show instead conspicuous linear polarization peaks (e.g., the 8662 \\AA\\ line of Ca {\\sc ii} whose $\\epsilon_Q=0$ showed a mysterious, positive polarization profile, which was considered to provide a challenge for the theoretical efforts). Another unexpected finding was that the CN molecule shows significant linear polarization, increasing to a maximum at each band head. The impossibility of explaining the observed enigmatic $Q/I$ features via the standard theory of scattering line polarization led \\cite{trujillo_2_stenflo83b} to write, ``clearly, the theoretical developments lag far behind in providing answers to the questions posed by our data\".  It is interesting to mention that in the same year \\cite{trujillo_2_landi83} published his paper on `Polarization in Spectral Lines:  A Unifying Theoretical Approach', where the scattering line polarization phenomenon is described as the temporal succession of 1st-order absorption and re-emission processes, interpreted as statistically independent events (complete redistribution in frequency). Actually, the phenomenon of scattering polarization in a spectral line is intrinsically a 2nd-order process \\citep[e.g., the review by][]{{trujillo_2_casini-landi07}}, where frequency correlations between the incoming and outgoing photons can occur (partial redistribution in frequency). However, it is still possible to treat consistently the phenomenon of scattering to 1st order if the atomic system is illuminated by a spectrally flat radiation field\\footnote{For the flat-spectrum approximation to hold, the incident  radiation field must be flat over a frequency interval $\\Delta{\\nu}$ larger than the natural width of the atomic levels, and, when coherences between non-degenerate levels are involved, $\\Delta{\\nu}$ must then be larger than the corresponding Bohr frequencies \\citep[]{trujillo_2_landibook}.}. The key point to keep in mind is that within the framework of the 1st-order theory of spectral line polarization, the solution of the statistical equilibrium equations for the multipolar components of the atomic density matrix determines the excitation state of the atomic or molecular system, from which the emergent radiation can then be calculated by solving the radiative transfer equations \\citep[see][for the numerical methods and computers programs with which we could solve this {\\em Non-LTE Problem of the Second Kind}]{trujillo_2_trujillo99,trujillo_2_trujillo-tubin,trujillo_2_manso-bueno-spw3}.  The 1983 exploratory surveys provided ``a glimpse of the wealth of information accessible to us in this new field of observational solar physics\" and made obvious the interest in pursuing the development of the Z\\\"urich Imaging Polarimeter (ZIMPOL), an instrument capable of measuring the polarization of the solar spectrum with a sensitivity limited only by photon statistics. ZIMPOL was invented by \\cite{trujillo_2_povel95}, a physicist working on the development of instrumentation at the ETH. Since then, it has been used by Stenflo and collaborators in many observing campaigns in combination with several solar telescopes, mainly with the 1.5m McMath Telescope at Kitt Peak. It was in 1997 when we could see in detail the beauty of the second solar spectrum at a polarimetric sensitivity of $10^{-5}$ in $Q/I$ \\citep[see][]{trujillo_2_stenflo-keller97}. A variety of fascinating $Q/I$ signals were discovered, many of which were considered (again) as ``enigmatic\" by Stenflo and collaborators \\citep*[see also][]{trujillo_2_stenflos00}. A suitable way to appreciate the improvement achieved with ZIMPOL is to compare Wiehr's (1975) observation of the $Q/I$ pattern across the sodium doublet (see his Fig. 1c) with Fig. 8 of \\cite{trujillo_2_stenflos80}, and both of them with Fig. 2$a$ of \\cite{trujillo_2_stenflo-keller97}.  The $Q/I$ observations of Stenflo and coworkers \\citep[see also][]{trujillo_2_gandorfer00,trujillo_2_gandorfer02} were soon confirmed and extended to the full Stokes vector by several groups using the Canary Islands telescopes. For example, Fig. 4 of  \\cite{trujillo_2_trujillos01} shows an {\\em on-disk} observation of the O {\\sc i} triplet around 7774 \\AA, which led to the discovery of {\\em negative} $Q/I$ polarization for the weakest line at 7776 \\AA. Another interesting observational result was that the D-lines of Na {\\sc i} show anomalous Stokes $V/I_{\\rm c}$ profiles in ``quiet\" regions close to the North solar limb, with the red lobe much more enhanced than the blue one \\citep[see Figs. 2 and 3 of][]{trujillo_2_trujillos01}. ",
        "conclusions": "Observing the second solar spectrum with high spatial resolution would allow us to discover hitherto unknown aspects of the Sun's hidden magnetism. For this reason, the design of the European Solar Telescope (EST), and of any future space telescope (e.g., SOLAR-C), should incorporate the scientific case of the spectral line polarization produced by radiatively induced quantum coherences in atomic and molecular systems. You can be sure that the magnetic fields of the extended solar atmosphere are continuously giving rise to amazing signatures in the emergent spectral line polarization, whose observation 100 years after Hale's (1908) discovery would lead to a new revolution in our empirical understanding of solar magnetism."
    },
    "0812/0812.1225_arXiv.txt": {
        "abstract": "The spectro-photometric properties of galaxies in galaxy formation models are obtained by combining the predicted history of star formation and mass accretion with the physics of stellar evolution through stellar population models. In the recent literature, significant differences have emerged regarding the implementation of the Thermally-Pulsing Asymptotic Giant Branch phase of stellar evolution. The emission in the TP-AGB phase dominates the bolometric and near-IR spectrum of intermediate-age ($\\sim 1$ Gyr) stellar populations, hence it is crucial for the correct modeling of the galaxy luminosities and colours. In this paper for the first time, we incorporate a full prescription of the TP-AGB phase in a semi-analytic model of galaxy formation. We find that the inclusion of the TP-AGB in the model spectra dramatically alters the predicted colour-magnitude relation and its evolution with redshift. When the TP-AGB phase is active, the rest-frame $V-K$ galaxy colours are redder by almost $2$ magnitudes in the redshift range $z \\sim 2-3$ and by $1$ magnitude at $z \\sim 1$. Very red colours are produced in disk galaxies, so that the $V-K$ colour distributions of disk and spheroids are virtually undistinguishable at low redshifts. We also find that the galaxy K-band emission is more than $1$ magnitude higher in the range $z \\sim 1-3$. This may alleviate the difficulties met by the hierarchical clustering scenario in predicting the red galaxy population at high redshifts. The comparison between simulations and observations have to be revisited in the light of our results. ",
        "introduction": "The modern theory of the formation and evolution of galaxies is built on the CDM hierarchical structure formation paradigm, and is supported by two powerful tools, semi-analytic models and smoothed particle hydrodynamics (SPH) simulations. In both cases, they combine the results of N-body simulations, producing the assembly of dark matter structures, with the baryonic physics describing cooling, star formation, feedback, chemical enrichment and dynamical interactions. The interface between models and observations is mainly constituted by the galaxy spectra, which can be modeled and interpreted to infer mass, age, chemical composition, star formation rate and mass assembly history of the galaxy populations (see Baugh, 2006, for a recent review).  The spectra of galaxies are obtained using stellar population models, that predict the spectral energy distribution (SED) of ensembles of stars as a function of age, metallicity and the Initial Mass Function (IMF). At any given galaxy age, the predicted emitted spectrum is a superposition of the synthetic spectra of the Single Stellar Populations (SSPs) that compose the total stellar mass, assembled according to the model's star formation rate (SFR) as a function of time.  Stellar population models currently used in published galaxy formation models and SPH simulations include Bruzual \\& Charlot (2003; adopted by Bower et al. 2006, Croton et al. 2006, De Lucia et al. 2006, Khochfar et al. 2007, Menci et al. 2006, Naab et al. 2007, Nagamine et al. 2005, Somerville et al. 2008), GRASIL by Silva et al. (1998; for instance Granato et al. 2004, Monaco et al. 2007), STARDUST by Devriendt et al. (1999; see Hatton et al. 2003) and PEGASE by Fioc \\& Rocca-Volmerange (1997; see Cattaneo et al. 2005, Mori \\& Umemura 2006). These stellar population models make use of the Padova (1994) stellar evolutionary tracks, which do not include the full contribution of the Thermally Pulsing - Asymptotic Giant Branch phase, as pointed out by Maraston (1998, 2005 hereafter M05) and more recently by other authors (e.g. Bruzual 2007, Conroy, Gunn \\& White 2008, Eminian et al. 2008, Marigo \\& Girardi 2007). As shown in Maraston (1998, 2005), the TP-AGB can account for up to $40 \\%$ of the bolometric light and $80 \\%$ of the light in the K band. Hence, the neglection of this phase in the galaxy formation models leads to an underestimation of the flux redwards of $5000 \\ \\mathring{A}$, with a major effect in the near-IR.  The emission in this phase is significant for the stars in the mass range $2-3 \\ M_{\\odot}$, and therefore peaks in a very short period of the life of the SSP, around the age of $\\sim 1$ Gyr. For this reason, the TP-AGB light is a powerful marker of an intermediate-age stellar population and dominates the K-band luminosity in the post-starburst phases of star formation. Since the rest of the emission in the K band is contributed by old and/or metal rich low-mass stars, the inclusion of the TP-AGB stars on top of the K-band background leads to dramatic differences in the colours and the mass-luminosity relation of model galaxies, especially in the early phases of their evolution.  To quantify the effect of the TP-AGB emission on the predicted galaxy colours at different redshifts, we run a semi-analytical code with different input SSPs, the M05 models (with TP-AGB) and the PEGASE models (without TP-AGB), and study the evolution of the colour-magnitude relation.  For the scope of this paper, the PEGASE and the Bruzual \\& Charlot model SSPs are equivalent, since they are based on the same input physics. In particular, they make use of the same stellar evolutionary tracks (Padova tracks) and they employ the isochrone synthesis modeling, producing very similar SSP spectra (see M05 for an extended discussion on the properties of these models).  We use GalICS as our reference semi-analytical model, set with standard LCDM cosmology and default astrophysical parameters. The model includes chemical enrichment and metallicity evolution based on the instantaneous recycling approximation (for details on the GalICS recipe see Hatton et al. 2003). We adopt no reddening by dust, to better isolate the TP-AGB effect.  Between the runs, we do not vary the input physics except for the different SSP models. The results presented here do not depend, in essence, on the particular galaxy formation model, and can be applied to any other model based on CDM hierarchical clustering (see the Discussion Section). ",
        "conclusions": "The offset in the predicted color-magnitude relation, the galaxy colour distribuition and their evolution with redshift described in the previous Section is dramatic, and is entirely produced by different prescriptions for the input stellar population models. The test galaxy formation model GalICS was run with the same input physics in both cases, ruling out that any feature in the model is responsible for the offset. The results presented here are solid against different semi-analytical model or SPH simulation implementations, as long as the framework in use is the hierarchical growth of structures. The relative contribution of the TP-AGB to the overall galaxy emission mainly depends on the average age of the stellar populations, and therefore is only sensitive to the mass assembly and star formation history. For this reason we expect similar results with any galaxy formation model based on CDM hierarchical clustering. The details of this offset will depend on individual model recipes, through second-order effects that are expected to be produced by different implementations of cooling and feedback.  The implications of these results have a great impact on the models of structure formation. The redshift-dependent, non-linear offset in the predicted K-band luminosities shifts the rest-frame $V-K$ colours redwards and the $M/L$ ratio in the red/infrared downwards, especially for massive galaxies. The key difference lies in producing extremely red $V-K$ colours in not-so-massive, young or intermediate-age galaxies, still actively star forming and possibly very blue in $B-V$, rather than only to old and dead, passively evolving systems.  It has been pointed out that hierarchical clustering is unable to account for the observed red galaxy population at high redshifts (Cirasuolo et al. 2008, Yan et al. 2004). The comparison between observations and simulations needs to be reconsidered in the light of the results presented here. We produce very red optical to near IR colours in young galaxies, without the need of large stellar masses. Simulations would be relieved of the constraint to produce high star formation rates at high redshifts, hard to achieve because the dark matter halo assembly process and the baryonic cooling strain to put together much stellar mass very fast at such early epochs.  We are now set on testing our results to see in how far the new model alleviates the difficulties in the convergence between theoretical predictions and observations. In follow-up papers we will analyse the impact of our results on key observational tests such as luminosity function, Tully-Fisher relation and the properties of the population of extremely red objects."
    },
    "0812/0812.2360_arXiv.txt": {
        "abstract": "The Magellanic Clouds offer unique opportunities to study star formation both on the global scales of an interacting system of gas-rich galaxies, as well as on the scales of individual star-forming clouds. The interstellar media of the Small and Large Magellanic Clouds and their connecting bridge, span a range in (low) metallicities and gas density. This allows us to study star formation near the critical density and gain an understanding of how tidal dwarfs might form; the low metallicity of the SMC in particular is typical of galaxies during the early phases of their assembly, and studies of star formation in the SMC provide a stepping stone to understand star formation at high redshift where these processes can not be directly observed. In this review, I introduce the different environments encountered in the Magellanic System and compare these with the Schmidt-Kennicutt law and the predicted efficiencies of various chemo-physical processes. I then concentrate on three aspects that are of particular importance: the chemistry of the embedded stages of star formation, the Initial Mass Function, and feedback effects from massive stars and its ability to trigger further star formation. ",
        "introduction": "The Magellanic Clouds (MCs) are our closest gas-rich galaxy neighbours. New instrumental advances mean we can now study their resolved stellar populations in great detail. At the same time, in the MCs we can study these populations from outside the galaxies themselves, allowing us a unique view of these populations and how they relate to the gas and dust distributions and the galaxies' structure.  The formation of stars in the early Universe took place in a metal-poor environment, however most of what is currently known about the star formation process is derived from observations in the Milky Way. A great advantage of the MCs is that their Interstellar Mediums (ISM) are characterised by metallicity significantly lower than that of the Milky Way. Thus the MCs are ideal templates to test whether metallicity significantly influences star formation and thus are unique probes for the environmental conditions more typical of the high-redshift Universe.  Only since relatively recently are we able to study the details of the star formation in the MCs. Firstly, I describe the star formation environment in the Magellanic Clouds. I will then concentrate on three particular facets of the star formation process. I will start by discussing the effects of a lower metal content on the chemistry in molecular clouds. I will then discuss the stellar Initial Mass Function (IMF), as well as massive star feedback and triggered star formation at low metallicity. ",
        "conclusions": ""
    },
    "0812/0812.0365_arXiv.txt": {
        "abstract": "We examine the evolution and influence of viscosity-induced diskoseismic modes in simulated black hole accretion disks. Understanding the origin and behavior of such oscillations will help us to evaluate their potential role in producing astronomically observed high-frequency quasi-periodic oscillations in accreting black hole binary systems.  Our simulated disks are geometrically-thin with a constant half-thickness of five percent the radius of the innermost stable circular orbit. A pseudo-Newtonian potential reproduces the relevant effects of general relativity, and an alpha-model viscosity achieves angular momentum transport and the coupling of orthogonal velocity components in an otherwise ideal hydrodynamic numerical treatment.  We find that our simulated viscous disks characteristically develop and maintain trapped global mode oscillations with properties similar to those expected of trapped g-modes and inner p-modes in a narrow range of frequencies just below the maximum radial epicyclic frequency. Although the modes are driven in the inner portion of the disk, they generate waves that propagate at the trapped mode frequencies out to larger disk radii. This finding is contrasted with the results of global magnetohydrodynamic disk simulations, in which such oscillations are not easily identified. Such examples underscore fundamental physical differences between accretion systems driven by the magneto-rotational instability and those for which alpha viscosity serves as a proxy for the physical processes that drive accretion, and we explore potential approaches to the search for diskoseismic modes in full magnetohydrodynamic disks. ",
        "introduction": "\\label{sec:introduction}  Since the detection of the first high-frequency quasi-periodic oscillations (HFQPOs) from black hole candidate GRS 1915+105 \\citep{1997ApJ...482..993M}, much effort has been made to relate such oscillations to natural accretion disk frequencies. Some early analysis by \\citet{1997ApJ...477L..91N} suggested that these HFQPOs were manifestations of global oscillation modes in galactic black hole binaries (GBHBs), the theory of which had been explored extensively by, for example, \\citet{1987PASJ...39..457O} and \\citet{1991ApJ...378..656N,1992ApJ...393..697N,1993ApJ...418..187N}. The discovery by \\citet{2001ApJ...552L..49S} of a pair of HFQPOs in GRO J1655-40 with an approximate 3:2 frequency ratio, however, lent support to an alternative parametric resonance model \\citep{2001A&A...374L..19A} in which HFQPOs result from resonance between orbital and radial epicyclic motion of disk material. This and similar resonance models have the advantage of naturally generating the small-integer frequency ratios seen in GRO J1655-40 and some subsequently observed sources \\citep[for a summary of current HFQPO observations, see][]{2006ARA&A..44...49R}. Still, parametric resonance models have yet to incorporate convincing physical mechanisms by which to excite HFQPOs \\citep{2008NewAR..51..855R}, and it has been noted by \\citet{2008GApFD.102...75O} that multiple global oscillation modes of differing mode number produce a 3:2 frequency ratio equally well. A detailed physical understanding of such oscillations is crucial since their observed frequencies ($\\sim 100$ Hz) are comparable to orbital frequencies near the innermost stable circular orbit (ISCO) of stellar-mass black holes. Interpreted correctly, HFQPOs thus have the potential to tell us much about the inner portions of accretion disks and the black holes they orbit.  While numerical simulations have great promise to elucidate the nature of disk oscillations, they unfortunately have failed thus far to produce convincing, identifiable HFQPOs. Despite some preliminary claims of HFQPO generation in relatively low-resolution simulations by \\citet{2004PASJ...56..931K}, subsequent numerical magnetohydrodynamic (MHD) studies have shown that simulated HFQPOs typically are transient \\citep{2006ApJ...651.1031S}, require external driving \\citep{2006astro.ph.11269C}, or, in the case of the first paper in this series, remain completely undetected \\citep[hereafter \\citetalias{2008arXiv0805.2950R}]{2008arXiv0805.2950R}. Interestingly, the simulations of ideal {\\it hydrodynamic} disks in \\citetalias{2008arXiv0805.2950R} did generate trapped gravity-driven (\\ie g-mode) global disk oscillations such as those described in \\citet{1992ApJ...393..697N}, but these oscillations were not seen in their otherwise comparable MHD disks. In fact, oscillations of the amplitude seen in X-ray observations or in their hydrodynamic disks would have fallen below the level of turbulent noise generated in the MHD case, so they could not determine whether the modes were hidden or actively damped by turbulence in the manner discussed by \\citet{2006ApJ...645L..65A}. Regardless, it is clear that the prototypical MHD disks of \\citetalias{2008arXiv0805.2950R} failed to excite to a detectable level either the global diskoseismic oscillations of \\citet{1993ApJ...418..187N} or the parametric resonance instability of \\citet{2001A&A...374L..19A}.  In this paper, we describe complementary work to \\citetalias{2008arXiv0805.2950R} in the form of simulations with oscillations induced by the viscous tapping of orbital energy. While the standard physical model for black hole accretion is predicated upon the magneto-rotational instability \\citep[MRI,][]{1991ApJ...376..214B}, an intrinsically MHD process that naturally generates turbulence and the transport of angular momentum, the traditional alternative to full MHD simulations has been to mimic the influence of magnetic fields through the introduction of an ``alpha-model'' viscosity \\citep{1973A&A....24..337S}. This approach subsumes all physical details of accretion into a single dimensionless viscous parameter $\\alpha$ designed to achieve the appropriate global level of angular momentum transport. This approach is by no means completely equivalent to a full MHD treatment, as \\citet{1998RvMP...70....1B} and \\citet{2008MNRAS.383..683P} note, and there even remain some basic order-of-magnitude discrepancies between values of $\\alpha$ inferred from observation and those derived from MHD simulations \\citep{2007MNRAS.376.1740K}. Still, studying alpha-disks provides us with a method by which to evaluate the influence of viscosity independent of MRI-driven turbulence and other typically complex behaviors of fully MHD disks.  \\citet{2000ApJ...537..922O} have provided a linear analysis of normal modes in viscous, rotating, Newtonian fluids that is applicable to our simulated viscous accretion disks. In particular, they show that the presence of viscosity should cause the fundamental g-modes in rotating disks to grow at a rate that scales with the characteristic orbital frequency in the system. In relativistic or pseudo-Newtonian gravitational potentials, these g-modes are predicted to be non-evanescent at radii where $|\\omega| < \\kappa$, where $\\omega$ is the wave frequency, and $\\kappa$ is the radial epicyclic frequency [see \\citet{1991ApJ...378..656N,1992ApJ...393..697N,2000ApJ...537..922O}, or Section 2.2 of \\citetalias{2008arXiv0805.2950R} for the appropriate dispersion relations and derivations]. In black hole accretion disks, this means that the modes are trapped just under the maximum radial epicyclic frequency $\\kappa_{\\rm max}$. Similarly, \\citet{2000ApJ...537..922O} find that viscosity also causes the inner pressure-driven (p-mode) oscillations to grow for $\\kappa_{\\rm max} > |\\omega| > \\kappa$, although the successful trapping of such modes depends strongly upon on the nature of the inner boundary of the disk.  While the viscosity-induced trapped g-mode of \\citet{2000ApJ...537..922O} has never been identified explicitly in simulations, numerical models of viscous hydrodynamic disks have generated identifiable waves at frequencies comparable to this mode. In an early numerical analysis of axisymmetric, vertically integrated (\\ie 1D) disks, \\citet{1992PASJ...44..529H} found that viscosity above a critical value $\\alpha \\sim 0.1$ caused global disk oscillations near $\\kappa_{\\rm max}$. Likewise, \\citet{1995ApJ...441..354C} and \\citet{1996MNRAS.283..919M} identified global oscillations near $\\kappa_{\\rm max}$ in vertically integrated disk simulations for a range of moderate accretion rates. This work was followed by the 2D simulations of \\citet{1997MNRAS.286..358M}, which focused on convection in optically thick disks but also found oscillations near $\\kappa_{\\rm max}$, particularly for low accretion rates and large viscosities. More recently, \\citet{2008arXiv0805.0598M} revisited the vertically integrated models of \\citet{1996MNRAS.283..919M}, pointing out that waves propagating from the inner portions of the disk could easily be locally super-Keplerian. All of these studies associate the observed signals with radial inertial-acoustic oscillations corresponding to the previously mentioned inner p-modes. This is a particularly valid interpretation in the case of vertically integrated disks where motion is constrained to the radial dimension, but distinguishing between trapped g- and inner p-modes in 2D viscous disks is not as straightforward, as we shall discuss in this work.  Our simulations of viscous accretion disks are intended to complement this previous work by exploring in more detail how viscosity can induce diskoseismic modes in accretion disks and how these modes affect the body of the disk. First, we seek to discover whether we can produce and identify in our models any of the viscosity-induced modes of \\citet{2000ApJ...537..922O}. In particular, we are interested in the trapped global g-modes since they exist in a narrow frequency range near $\\kappa_{\\rm max}$, the value of which in principle can be used as a diagnostic of the fundamental physical properties of the black hole. Since these g-modes are trapped well away from the inner boundary of the disk, we also expect them to be less susceptible than inner p-modes to leakage across the ISCO. We further examine how such modes generate waves that propagate through the entire body of the disk, far beyond the formal mode trapping region. Since trapped g-modes were identified in the hydrodynamic simulations -- but not the MHD disks -- of \\citetalias{2008arXiv0805.2950R}, we further seek to understand and evaluate the observed differences between viscous alpha-disks and full MHD models.  In $\\S$ \\ref{sec:modeling}, we outline the computational framework used and describe our simulated disk models. In $\\S$ \\ref{sec:results}, we present the results of our simulations and discuss our identification of trapped diskoseismic modes and the effects of these modes on viscous disks. We place our findings in a broader context in $\\S$ \\ref{sec:discussion}, comparing our results to previous work, and present our conclusions in $\\S$ \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We have conducted an ensemble of axisymmetric simulations of black hole viscous accretion disks to explore the generation of diskoseismic modes and their influence on disks. While we are still far from a definitive identification of the origin of HFQPOs, we have uncovered and explored several interesting facets of viscous disk evolution, and we summarize our findings here:  1) For viscous disks with $\\alpha \\geq 0.05$, we see indications of the trapped diskoseismic modes of \\citet{2000ApJ...537..922O}. These modes have all of the expected properties of trapped g-modes or inner p-modes, and are located at $r \\lesssim r_{\\rm max}$ with frequencies $\\nu \\sim \\nu_{\\rm max}$. This confirms that modes similar to those seen in earlier simulations of vertically integrated models are present for 2D optically- and geometrically-thin accretion disks.  2) We note that viscous disk models with trapped diskoseismic modes also develop related waves that pervade much of the body of the disk. These outward-propagating waves are continuous extensions in frequency and power of the trapped modes, despite extending beyond the region of formal mode trapping. This too is similar to the 1D result and suggests that diskoseismic modes can effectively communicate their characteristic frequencies to portions of the disk in which the modes themselves would be strongly damped.  3) By comparing our viscous disks to a full 3D MHD simulation of \\citetalias{2008arXiv0805.2950R}, we have further shown that the trapped mode signal for the corresponding alpha disk would fall far below the current noise level of the MHD simulation. This suggests that, to produce detectable trapped modes, MHD simulations may need to feature larger effective viscosities, possibly through the natural incorporation of net magnetic flux. Alternately, larger trapped mode signal-to-noise ratios should be achievable by extending the time domain of these simulations."
    },
    "0812/0812.0479_arXiv.txt": {
        "abstract": "\\noindent We perform a statistical analysis with the prospective results of future experiments on neutrino-less double beta decay, direct searches for neutrino mass (KATRIN) and cosmological observations. Realistic errors are used and the nuclear matrix element uncertainty for neutrino-less double beta decay is also taken into account. Three benchmark scenarios are introduced, corresponding to quasi-degenerate, inverse hierarchical neutrinos, and an intermediate case. We investigate to what extend these scenarios can be reconstructed. Furthermore, we check the compatibility of the scenarios with the claimed evidence of neutrino-less double beta decay. ",
        "introduction": "Introduction} Neutrino mass and lepton mixing represent an unambiguous proof that the Standard Model (SM) of elementary particles is incomplete. Various experiments with solar \\cite{sol}, atmospheric \\cite{atm} and man-made \\cite{kl,minos} neutrino sources imply non-trivial lepton mixing angles, as well as non-zero and non-degenerate neutrino masses. Their values are extremely suppressed with respect to the masses of the other (electrically charged) fermions of the SM. The most prominent and often studied mechanism to explain the smallness of neutrino masses is the see-saw mechanism \\cite{seesaw}. The neutrino mass scale is here inversely proportional to the scale of its origin. In addition, lepton number violation is predicted: neutrinos are Majorana particles. Searching for this property will be a crucial test of the see-saw mechanism, but also of other mechanisms leading to small Majorana neutrino masses. Possible phenomenological consequences of lepton number violation are the generation of the baryon asymmetry of the Universe \\cite{lepto} or, at low energies, \\onbb~(\\obb) \\cite{APS}. This decay of certain nuclei, $(A,Z) \\rightarrow (A,Z+2) + 2 \\, e^-$, which has not yet been observed, clearly violates lepton number by two units, and is intensively searched for \\cite{APS}. We will assume here that light Majorana neutrinos are exchanged in the diagram responsible for \\obb. In this case, the amplitude for this process is proportional to the coherent sum \\be \\label{eq:meff} m_{ee}  \\equiv \\sum\\limits_{i=1}^3 U_{ei}^2 \\, m_i \\,, \\ee where $m_i$ are the individual neutrino masses and $U$ is the leptonic mixing, or Pontecorvo-Maki-Nakagawa-Sakata (PMNS), matrix. The absolute value of $m_{ee}$ is called the effective mass. The entries $U_{ei}$ can be written as $U_{e1} = \\cos \\theta_{12} \\, \\cos \\theta_{13}$, $U_{e2} = \\sin \\theta_{12} \\, \\cos \\theta_{13} \\, e^{i \\alpha}$ and $U_{e3} = \\sin \\theta_{13} \\, e^{i \\beta}$, where $\\alpha$ and $\\beta$ are two currently unknown ``Majorana phases'' and $\\theta_{12, 13}$ are mixing angles. While $\\theta_{13}$ is constrained mainly by short-baseline reactor experiments, $\\theta_{12}$ is probed by solar and long-baseline reactor neutrino experiments. Their current best-fit values as well as $1\\sigma$ and $3\\sigma$ ranges can be obtained from three-flavor fits, the result being \\cite{data} \\be \\label{eq:data} \\sin^2 \\theta_{12} = 0.32\\,(\\pm 0.02)^{+0.08}_{-0.06} ~,~ \\sin^2 \\theta_{13} = 0^{+0.019, \\, 0.050}\\,. \\ee In what regards the neutrino masses, for a normal ordering one has $m_3 > m_2 > m_1$ with $m_2^2 = m_1^{2} +\\dms$ and $m_3^2 = m_1^{2}+\\dma$. In case of an inverted ordering we have $m_2 > m_1 > m_3$ with $m_2^2 = m_3^{2}+\\dms+\\dma$ and $m_1^2 = m_3^{2} + \\dma$. Here $\\dms$ and $\\dma$ are mass-squared differences with best-fit values and $3\\sigma$ ranges $(7.9^{+1.1}_{-0.9}) \\cdot 10^{-5}$ eV$^2$ and $(2.6^{+0.6}_{-0.6}) \\cdot 10^{-3}$ eV$^2$, respectively \\cite{data}. Quasi-degenerate neutrino masses occur when $m_{1,2,3}^2 \\gg \\dma, \\dms$. If neutrinos are Majorana particles, all low energy neutrino phenomenology can be described by the neutrino mass matrix $m_\\nu = U^\\ast \\, m_\\nu^{\\rm diag} \\, U^\\dagger$. It contains nine physical parameters. Seven out of the nine parameters of the neutrino mass matrix appear in \\meff. Therefore, it contains a large amount of information, in particular if complementary measurements of some of the other parameters exist. We also note that all parameters of $m_\\nu$ which do {\\it not} influence neutrino oscillations show up in the effective mass. Those are the the Majorana phases and, in particular, the individual neutrino masses (neutrino oscillations are only sensitive to mass-squared differences). For a review on the dependence of \\meff~on the various neutrino parameters see refs.~\\cite{APS,Lindner:2005kr,0vbb_rev} and references therein. In the present paper, in contrast to other works statistically analyzing future neutrino mass measurements including \\obb~\\cite{PPS,WR_old,HP,AdG,steen2,Fogli:2006yq,nme1}, we focus on the neutrino mass scale, i.e.\\ the value of the smallest neutrino mass. To this end we define three natural benchmark scenarios and investigate how future experiments may be able to constrain them. Our goal here is to combine as much mass-related information as possible. ",
        "conclusions": ""
    },
    "0812/0812.3067.txt": {
        "abstract": " ",
        "introduction": "\\hspace{1.1cm} The recombination of primordial plasma is a process that ultimately leads to the formation of neutral atoms from ions and free electrons due to the decrease in temperature through cosmological expansion. This process has three distinct epochs at which the fraction of free electrons changes significantly: (1) HeIII$\\rightarrow$HeII recombination ($z\\simeq 5000 - 7000$), (2) HeII$\\rightarrow$HeI recombination ($z\\simeq 1500 - 3000$), and (3) HII$\\rightarrow$HI recombination ($z\\simeq 900 - 1600$), where z is the cosmological redshift. Since other nuclides (D, $^3$He, Li, B, etc.) in the primordial plasma are much fewer in number than $^1$H and $^4$He ($<10^{-4}$), the recombination of hydrogen-helium plasma is usually considered (Zeldovich et al. 1968; Peebles 1968; Matsuda et al. 1969). The recombination of other elements is considered in isolated cases for special problems, such as the effect of lithium recombination on the cosmic microwave background (CMB) anisotropy (Stancil et al. (2002) and references therein), the formation of primordial molecules (Galli and Palla (2002) and references therein), etc.  The recombination of primordial plasma affects significantly the growth of gravitational instability and the formation of CMB spectral distortions and anisotropy (Peebles 1965; Dubrovich 1975). The appearance of the first experimental data on CMB anisotropy (Relikt 1, COBE) rekindled interest in the recombination of primordial plasma in the mid- 1980s. A number of improvements in the model of hydrogen plasma recombination were suggested (Jones and Wyse 1985; Grachev and Dubrovich 1991).  Significant progress in CMB anisotropy observations achieved in the second half of the 1990s (BOOMERANG, WMAP) necessitated including a number of subtle effects that could affect the recombination of primordial hydrogen and helium at a level of 0.1 - 1\\% (Leung et al. 2004; Dubrovich and Grachev 2005; Novosyadlyj 2006; Burgin et al. 2006; Kholupenko and Ivanchik 2006; Wong and Scott 2007; Chluba and Sunyaev 2006, 2007, 2008a, 2008b; Hirata and Switzer 2008; Sunyaev and Chluba 2008; Hirata 2008; Grachev and Dubrovich 2008).  One of the most important (for the primordial plasma recombination kinetics) effects considered in recent years is the absorption of HeI resonance photons by neutral hydrogen, which leads to an acceleration of the HeII$\\rightarrow$HeI recombination (Kholupenko et al. 2007 [hereinafter KhIV07]; Switzer and Hirata 2008; Rubino-Martin et al. 2007). Recently, this effect was taken into account by Wong et al. (2008) in the {\\bf recfast} computational code developed by Seager et al. (1999). This code is most widely used to compute the primordial plasma recombination kinetics when the CMB anisotropy is analyzed. To take into account the effect of neutral hydrogen on the HeII$\\rightarrow$HeI recombination kinetics, Wong et al. (2008) used a simple approximation formula with adjustable parameters. The {\\bf recfast} modified in this way allows the results of computations with the multilevel code\\footnote {By the multilevel code we mean a computational program that uses a multilevelmodel atom.} by Switzer and Hirata (2008) to be quickly and accurately reproduced for any reasonable values of the cosmological parameters. Nevertheless, the approach by Wong et al. (2008) is inapplicable for more detailed studies of the helium recombination, since their formula is not universal for all resonance transitions in HeI, but can be used only in describing the absorption of HeI $2^1P\\rightarrow 1^1S$ resonance photons by neutral hydrogen. When using the multilevel codes by Switzer and Hirata (2008) and Rubino-Martin et al. (2007), which allow the absorption of HeI $nP\\rightarrow 1S$ ($n\\ge 2$) resonance photons by neutral hydrogen to be taken into account, much of the radiative transfer calculations in the HeI resonance lines are performed numerically, which is computationally demanding and time-consuming. These circumstances forced us to seek for an analytic approach to the problem of including the effect of neutral hydrogen on the HeII$\\rightarrow$HeI recombination kinetics that, on the one hand, would be more universal than the approach of Wong et al. (2008) (i.e., would allow the effect of the absorption of HeI $nP\\rightarrow 1S$ resonance photons (where $n\\ge 2$, and not only $n=2$) on the HeII$\\rightarrow$HeI recombination to be estimated) and, on the other hand, would not reduce the speed and accuracy of the primordial plasma recombination computations for various parameters of the cosmological model. This approach was implemented on the basis of the papers by Chugai (1987) and Grachev (1988), who analytically investigated the diffusion of resonance radiation in the presence of continuum absorption. In this paper, we present extended and more detailed justifications of the key suggestions made in KhIV07. We take into account the fact that the scattering in the HeI resonance lines (for transitions in the singlet structure of the HeI atom) occurs with a partial redistribution in frequency, which turns out to be important for the results (Switzer and Hirata 2008; Rubino-Martin et al. 2007). Thus, our goal is to numerically compute the HeII$\\rightarrow$HeI recombination kinetics using analytic formulas (Chugai 1987; Grachev 1988; this paper) to allow for the peculiarities of the radiative transfer in the HeI resonance lines (a partial redistribution of HeI resonance photons in frequency and their absorption in the neutral hydrogen continuum). ",
        "conclusions": "We calculated the HeII$\\rightarrow$HeI recombination kinetics by including the effect of neutral hydrogen. We additionally took into account the partial redistribution of HeI resonance photons in frequency upon their scattering in the HeI $2^1P\\rightarrow 1^1S$ line (Eqs. \\ref{P_H_int3} - \\ref{f_limit_final}).  It is shown that the calculated relative numbers of free electrons $N_e/N_H$ could be satisfactorily reconciled with the results of recent studies of the HeII$\\rightarrow$HeI recombination kinetics (Rubino-Martin et al. 2007; Wong et al. 2008) using the formulas derived here. The achieved accuracy is high enough for the HeII$\\rightarrow$HeI recombination model suggested here to be used in analyzing the CMB anisotropy data from future experiments (Planck and others). This accuracy is also high enough to calculate the intensities, frequencies, and profiles of the HeI recombination lines formed during the cosmological HeII$\\rightarrow$HeI recombination (Dubrovich and Stolyarov 1997; KhIV07; Rubino-Martin et al. 2007).  It should also be noted that using Eqs. (\\ref{P_H_D_approx}), (\\ref{approx_P_H}), (\\ref{P_H_W}), (\\ref{P_H_int3}), (\\ref{P_H_R2} - \\ref{f_limit_final}) derived here, along with Eq. (\\ref{P_H_G}) derived by Grachev (1988), we can find not only the dependence of $P^{H}$ on $\\gamma$, but also the dependences of $P^{H}$ on the Voigt parameter $a$ and the fraction of coherent scatterings $\\lambda$ (and, accordingly, the fraction of incoherent scatterings $(1-\\lambda)$, if the atomic transitions produced by electron collisions may be neglected). This will be important in further detailed theoretical studies of the HeII$\\rightarrow$HeI recombination."
    },
    "0812/0812.4016_arXiv.txt": {
        "abstract": "We present a complete analysis of the cosmological constraints on decaying dark matter. Previous analyses have used the cosmic microwave background and Type Ia supernova. We have updated them  with the latest data as well as extended the analysis with the inclusion of Lyman-$\\alpha$ forest, large scale structure and weak lensing observations. Astrophysical constraints are not considered in the present paper. The bounds on the lifetime of decaying dark matter are dominated by either the late-time integrated Sachs-Wolfe effect for the scenario with weak reionization, or CMB polarization observations when there is significant reionization. For the respective scenarios, the lifetimes for decaying dark matter are $\\Gamma^{-1} \\gtrsim 100$ Gyr  and $ (f \\Gamma) ^{-1} \\gtrsim 5.3 \\times 10^8$ Gyr (at 95.4\\% confidence level), where the phenomenological parameter $f$ is the fraction of the decay energy deposited in baryonic gas. This allows us to constrain particle physics models with dark matter candidates through investigation of dark matter decays into Standard Model particles via effective operators. For decaying dark matter of $\\sim 100$ GeV mass, we found that the size of the coupling constant in the effective dimension-4 operators responsible for dark matter decay has to generically be $ \\lesssim 10^{-22}$. We have also explored the implications of our analysis for representative models in theories of gauge-mediated supersymmetry breaking, minimal supergravity and little Higgs. ",
        "introduction": "The past decade and a half has seen tremendous progress in the field of cosmology. The numerous experiments and the plethora of accumulated data have elevated cosmology into a precision science. With the knowledge we have gained, a remarkably consistent consensus known as the standard model of cosmology has emerged from these attempts to understand the nature of the universe. The picture that we have presently is that of a universe that is composed of 74\\% dark energy, 22\\% dark matter (DM) and 4\\% baryonic matter \\cite{Hinshaw2008, Komatsu2008}. Despite its extraordinary success in explaining a variety of diverse observations, fundamental questions do remain. Arguably the most vexatious is the question, ``What are these dark components of the universe?\" Despite the fact that it makes up 96\\% of the universe, we have so far been unable to say definitively what they are. Perhaps the most compelling ideas to resolve this question arise from particle physics. Dark energy is commonly attributed to the vacuum energy while the identity of dark matter is hypothesized to be one of the new particles in theories that  extend the Standard Model of particle physics. This confluence and cross-fertilization of ideas from two major fields of scientific endeavor promises to herald an exciting new era in understanding of the universe. With the resumption of operation of the Large Hadron Collider at CERN, we are possibly months away from collider detection, albeit indirectly through missing energy signatures, of dark matter particles.  Yet this avenue too is plagued with troubling questions that need to be addressed if we are to take the idea seriously that some particle from an extension of the Standard Model is indeed the elusive yet ubiquitous dark matter of the universe. Most worryingly, to ensure the presence of a dark matter candidate in a number of beyond the Standard Model theories of particle physics, it is often necessary to impose global symmetries. For instance, we have the T-parity in Little Higgs \\cite{Schmaltz:2005ky} and R-parity in supersymmetry \\cite{Martin:1997ns}. In the limit where the global symmetries are exact, the lightest particle carrying such a global charge would be stable from decay to lighter particles that do not possess such a charge. It is this point that could potentially destroy this promising marriage of ideas from cosmology and particle physics, for it is well known that global symmetries are never exact.  The presence of anomalies, as in the case of T-parity \\cite{Hill:2007zv}, or R-parity violating terms in supersymmetry \\cite{Barbier:2004ez} would often mean that the dark matter candidates arising from these theories are neither stable nor long-lived in the cosmological sense.  Even if this had not been the case, the presence of gravity would necessarily induce the violation of global symmetries as was first revealed in studies of black holes \\cite{b1, b2, b3, b4}. So the lightest particle charged under a particular global symmetry would have, at best, a very long lifetime. Indeed, it has even been conjectured that discrete global symmetries are violated maximally by gravity \\cite{Banks:1989ag, Krauss:1988zc}.  Additional motivation can be found in numerical simulations of the universe (based on the conventional $\\Lambda$CDM cosmology) which predict an overabundance of substructures as compared to actual observations. Models with decaying dark matter \\cite{Cen:2000xv, Borzumati:2008zz} provide an extremely compelling and natural mechanism for suppressing the power spectrum at small scales thus resolving the discrepancy.  Continuing on the line of thought leading from particle physics to cosmology, the question that then naturally springs to mind is ``What can cosmology say about decaying dark matter and the particle physics theories that contain them?\" It is this intriguing prospect that we explore in this paper.  There have been a few papers \\cite{Ichiki:2004vi, Gong:2008gi, Lattanzi:2008zz} in recent years analyzing DM decay into electromagnetically non-interacting particles using just the cosmic microwave background data (Ref.\\cite{Gong:2008gi} also includes supernova data). In this paper, we revisit this scenario using a Markov Chain Monte Carlo (MCMC) analysis employing all available datasets from the cosmic microwave background (CMB), Type Ia supernova (SN), Lyman-$\\alpha$ forest (Ly$\\alpha$), large scale structure (LSS) and weak lensing (WL) observations. We find that the lifetime of decaying DM is constrained predominantly by the late time Integrated Sachs Wolfe (ISW) effect to be $\\Gamma^{-1} \\gtrsim 100$ Gyr. In the main body of the paper, we will comment on the discrepancies between the results of Refs.\\cite{Ichiki:2004vi, Gong:2008gi, Lattanzi:2008zz}.  The studies in the preceding paragraph considered only the case where there was negligible reionization of the universe due to DM decay. In an attempt to address this, Ref.\\cite{Zhang:2007zzh} analyzed the scenario of DM decaying into only electromagnetically interacting products, that get partially absorbed by the baryonic gas, using a subset of the available CMB datasets. Our paper extends their analysis by using all the available CMB datasets, and also the SN, Ly$\\alpha$, LSS and WL datasets. Besides the smaller selection of datasets, their analysis also ignores the impact of DM decay on cosmological perturbations which renders it ineffectual in the parameter space where there is negligible reionization. Our treatment allows the decay products to not only reionize the universe but also takes into account the effect of DM decay on cosmological perturbation. This allows us to generate many other observables, particularly, the late time ISW effect that is crucial to constrain the lifetimes at low reionization. Another key difference between the analyses is that we use a combined reionization parameter for both DM reionization and phenomenological star formation reionization, rather than just treating them separately as was done in Ref.\\cite{Zhang:2007zzh}, because current observations cannot distinguish which contribution to reionization is the dominant one. Doing the MCMC analysis, we find that the lifetime of decaying DM in this scenario constrained to be $ (f \\Gamma) ^{-1} \\gtrsim 5.3 \\times 10^8$ Gyr, where $f$ is a phenomenological parameter introduced by Ref.\\cite{Zhang:2007zzh} and related to the degree of reionization.  Astrophysical constraints together with additional assumptions  have also been used \\cite{PalomaresRuiz:2007ry, Yuksel:2007dr}  to give even tighter bounds than ours on the lifetime of the decaying DM. While interesting and complementary, these lie outside the scope of our present paper.  Having obtained the bounds on the lifetime of decaying dark matter, we will then explore the implications of our cosmological analysis on particle physics models beyond the Standard Model. We will present a complete list of cross-sections for spin-0, spin-1/2 and spin-1 dark matter to decay into Standard Model degrees of freedom via effective operators. Obviously, this can be easily extended to other models with additional light degrees of freedom (for instance, hidden valley models \\cite{Strassler:2006im}) by appropriate substitution of the parameters. Applying the bound on the lifetime of the decaying DM, we can then place limits on the size of the parameters of theories. For generic theories with a decaying dark matter of $\\sim 100$ GeV mass, the coupling constant in the effective dimension-4 operators responsible for dark matter decay will be shown to be $\\lesssim 10^{-22}$. We will also look at specific representative cases of theories beyond the Standard Model physics and investigate the possibility of viable dark matter candidates: the spin-0 messenger DM in the context of gauge mediation messenger number violation, the spin-1/2 bino DM in the scenario with R-parity violation and the spin-1 ``massive photon partner''DM in the framework of T-parity violation.     The rest of this paper is organized as follows. In Section 2, we discuss the physics of decaying dark matter cosmology as well as introduce the datasets that we will be using. Section 3 contains our Markov Chain Monte Carlo results and discussions of the cosmological implications. In Section 4, we explore the consequences of these results for particle physics theories by enumerating the decay channels and partial widths. Representative models from theories of gauge-mediated supersymmetry breaking, minimal supergravity and little Higgs were also investigated using the results of our analysis. We conclude and briefly comment on future prospects in Section 5. ",
        "conclusions": "We have performed a full cosmological analysis using the available datasets from cosmic microwave background, Type Ia supernova, Lyman-$\\alpha$ forest, galaxy clustering and weak lensing observations to determine the extent by which we can constrain decaying dark matter models which are very typical in most extensions of the Standard Model of particle physics.  In the scenario where there is negligible reionization of the baryonic gas by the decaying dark matter, we have found that the late-time Integrated Sachs-Wolfe  effect gives the strongest constraint. The lifetime of a decaying dark matter has the bound $\\Gamma^{-1}\\gtrsim 100$Gyr (at 95.4\\% confidence level). Because of cosmic variance, the results are not likely to improve significantly with the WMAP-9yr data.  When there is significant reionization of the baryonic gas due to the decaying dark matter, the bounds become more restrictive as the CMB polarization is well measured. In this scenario, the lifetime of a decaying dark matter is $(f \\, \\Gamma) ^{-1}\\gtrsim 5.3 \\times 10^8$ Gyr (at 95.4\\% confidence level) where $f$ is a phenomenological factor related to the degree of reionization. With even more CMB polarization data, one could conceivably distinguish the reionization due to decaying dark matter from reionization due to star formation, thereby giving us even better bounds on the lifetime of the dark matter. We expect that the the 21cm cosmological observation in the future would give us even greater precision as it is expected to probe the reionization history at redshifts $6<z<30$.  Having obtained the cosmological constraints, we turned our attentions to the particle physics aspects of it. For completeness and motivated by the utility of such an exercise, we systematically tabulated the decay cross-sections for a spin-0, spin-1/2 and spin-1 dark matter candidate into the Standard Model degrees of freedom. This enabled us to simply sum up all the relevant contributions for a particular model of particle physics and arrive at the functional form of the lifetime of the decaying dark matter. We repeated this process for a variety of representative models from the following classes of theories: generic supersysmmetric scenario, gauge-mediated supersymmetry breaking models and the little Higgs theories. Imposing the limits from our cosmological analysis, we find that generically for most models we have looked at, the dimensionless coupling for a decaying dark matter to Standard Model fields should be smaller than $10^{-22}$.  This restriction can be slightly relaxed if the dark matter decays solely into light particles via helicity suppressed interaction terms, in which case, the small mass of the decay products suppresses the decay rate. If, for instance, the dark matter decays purely via helicity suppressed terms into $\\nu\\bar{\\nu}$ with Dirac mass of $\\sim 2 \\textrm{eV}$, then the dimensionless coupling can be as large as $10^{-11}$. In addition to constraining the coupling, one can assume it to be of ${\\cal O}(1)$ and estimate the scale of new physics which suppresses the decay rate. In all cases, either the coupling attains a small value or the new physics come from a huge scale, both of which would need interesting and exotic physics to realize if indeed the dark matter does decay via dimension-4 or dimension-5 terms. In the case of exclusive helicity suppressed decays, moreover, one has to explain why other interaction terms are absent in the model. A more promising avenue, which we briefly mentioned in the previous section, is to look at models where the dark matter decays via dimension-6 operators. The Large Hadron Collider might provide us with the identity for dark matter in the very near future, but on the basis of our analysis, there will still be much to understand about physics of the dark matter sector and how it interacts with the Standard Model.   In a future paper, we hope to address some of the astrophysical issues of decaying dark matter. The recent spate of results from astrophysical experiments \\cite{Strong:2005zx,Aguilar:2007yf,Bernabei:2008yi,Adriani:2008zr} has given us much to ponder. The immediate goal would of course be to combine all the astrophysical datasets with the cosmological ones that we have considered in the present paper and arrive at a set of characteristics that a phenomenologically viable decaying dark matter must possess. However, we expect considerable tension between the two classes of constraints. The astrophysical ones require that the decays to be significant enough to account for the as-yet-unexplained phenomena, while the cosmological ones need decays to be small enough because of the late time ISW effect and the CMB polarization observations. To reconcile and resolve these two seemingly conflicting classes of observations could be the defining challenge of dark matter physics in the next decade."
    },
    "0812/0812.4827_arXiv.txt": {
        "abstract": "The properties of quintessence are examined through the study of the variation of the electromagnetic coupling. We consider two simple quintessence models with a modified exponential potential and study the parameter space constraints derived from the existing observational bounds on the variation of the fine structure constant and the most recent Wilkinson Microwave Anisotropy Probe observations.  \\vskip 0.5cm ",
        "introduction": "Over the last decade, the temporal and spatial variation of fundamental constants has become a very popular subject. The motivation partially comes from theories unifying gravity and other interactions, which suggest that fundamental constants could have indeed varied during the evolution of the universe~\\cite{fund-th}. It is therefore particularly relevant to search for these variations and try to establish correlations, if any, with other striking properties of the universe, as for instance with dark energy.  From the observational point of view, the time variation of the fine structure constant $\\alpha$ has been widely discussed in several contexts. In particular, from the spectra of quasars (QSO), one obtains from the KecK/HIRES instrument~\\cite{quasars1} \\begin{align} \\frac{\\Delta\\alpha}{\\alpha}=(-0.57 \\pm 0.11)\\times 10^{-5}~,~\\text{for}~ 0.2 < z < 4.2~, \\label{murphy} \\end{align} while the Ultraviolet and Visual Echelle Spectrograph (UVES) instrument~\\cite{quasars2, quasars3} implies \\begin{align} \\frac{\\Delta\\alpha}{\\alpha}=(-0.64 \\pm 0.36)\\times 10^{-5}~,~\\text{for} ~0.4< z < 2.3~. \\label{chand} \\end{align} The Oklo natural reactor also provides a bound, \\begin{align} -0.9\\times 10^{-7}<\\frac{\\Delta\\alpha}{\\alpha}<1.2\\times 10^{-7}~,~\\text{for}~z<0.14~, \\label{oklo} \\end{align} at $95 \\%$ C.L.~\\cite{oklo1,oklo2,oklo3}. Furthermore, estimates of the age of iron meteorites, corresponding to $z \\simeq 0.45$, when combined with the measurement of the Os/Re ratio resulting from the radioactive decay $^{187}$Re $\\to\\, ^{187}$Os, yield \\cite{met1,met2,met3}: \\begin{align} \\frac{\\Delta\\alpha}{\\alpha}=(-8 \\pm 8) \\times 10^{-7}~, \\label{meteorites1} \\end{align} at $1\\sigma$, and \\begin{align} -24\\times 10^{-7}<\\frac{\\Delta\\alpha}{\\alpha}<8\\times 10^{-7}~, \\label{meteorites2} \\end{align} at $2\\sigma$.  Big bang nucleosynthesis (BBN) also places bounds on the variation of $\\alpha$~\\cite{BBN}: \\begin{align} \\label{BBN} -0.05<\\frac{\\Delta\\alpha}{\\alpha}<0.01~,~\\text{for}~10^9<z<10^{10}~. \\end{align} Finally, the 5-year data from the Wilkinson Microwave Anisotropy Probe (WMAP) with Hubble Space Telescope (HST) prior provides the following bound at 95\\% C.L.~\\cite{Nakashima:2008cb}: \\begin{align} \\label{WMAP} -0.028 <\\frac{\\Delta\\alpha}{\\alpha}< 0.026~, ~\\text{for}~z \\sim 10^3. \\end{align} Without HST prior this bound is less restrictive~\\cite{Nakashima:2008cb}: \\begin{align} \\label{WMAP2} -0.050 <\\frac{\\Delta\\alpha}{\\alpha}< 0.042~, \\end{align} at 95\\% C.L.  On the other hand, recent observations of high redshift type Ia supernova and, more indirectly, of the CMB and galaxy clustering, indicate that the universe is undergoing a period of accelerated expansion~\\cite{expansion}. This suggests that the universe is dominated by a form of energy density with negative pressure (dark energy). An obvious candidate for dark energy could be an uncanceled cosmological constant~\\cite{BentoUncanceledCC}, which however would require an extremely high fine-tuning. Quintessence-type models~\\cite{quint} with one~\\cite{quint1} or two~\\cite{quint2} scalar fields, k-essence~\\cite{k-ess} and the Chaplygin gas with an exotic equation of state~\\cite{GCG3,bil} are among other possibilities.  In most of the theoretical approaches, scalar fields are present in the theory. Thus, one could expect that the two observational facts, namely the variation of $\\alpha$ and the recent acceleration of the universe, are somehow related. Indeed, the coupling of such a scalar field to electromagnetism would lead to a variation of the fine structure constant~\\cite{Bekenstein}. In several contexts~\\cite{carroll,supergrav,mohammad,BM}, the above question has already been addressed.  In this work we shall consider two simple quintessence models with a modified exponential potential, proposed by Albrecht and Skordis~\\cite{skordis1,skordis2}. Our aim is to restrict the parameter space of these models by using the observational constraints on the variation of the fine structure constant and by imposing consistency with the 5-year data from WMAP. ",
        "conclusions": "The time variation of fundamental constants is a common prediction of theories that attempt to unify the four fundamental interactions. The experimental bounds obtained on the variation of these constants are therefore a useful tool for testing the validity of these theories. In this work we have studied the implications of the coupling of electromagnetism to quintessence fields. In our analysis we have considered two simple quintessence models with a modified exponential potential~\\cite{skordis1,skordis2}, with the aim to constrain the parameter space using the observational data on the variation of the fine structure constant $\\alpha$ in combination with the 5-year data from WMAP. Our results are summarized in Eqs.~(\\ref{eq:boundperm1})-(\\ref{eq:boundperm2}).  We should point out that, in constraining the models, the observational data on the variation of $\\alpha$, obtained from the quasar absorption systems [cf. Eqs.(\\ref{murphy}) and (\\ref{chand})], turned out to be crucial when establishing limits on the parameters of the AS1 and AS2 models. This is due to the fact that the QSO bounds imply a non-vanishing value of $\\Delta\\alpha/\\alpha$ at redshift $0.2 \\lesssim z \\lesssim 4.2$, which in turn requires a non-vanishing coupling $\\xi$ between the scalar and electromagnetic fields. On the other hand, it was not possible to use in our approach the Oklo, meteorite, BBN and CMB bounds to put any tight constraint on the model parameters, since these observational bounds are all consistent with no variation of $\\alpha$ at their corresponding redshift scales. \\vspace{0.5cm}"
    },
    "0812/0812.4220_arXiv.txt": {
        "abstract": "Recent radioastronomical observations of Faraday rotation in the solar corona can be interpreted as evidence for coronal currents, with values as large as $2.5 \\times 10^9$ Amperes \\citep{Spangler07}. These estimates of currents are used to develop a model for Joule heating in the corona. It is assumed that the currents are concentrated in thin current sheets, as suggested by theories of two dimensional magnetohydrodynamic turbulence. The Spitzer result for the resistivity is adopted as a lower limit to the true resistivity.  The calculated volumetric heating rate is compared with an independent theoretical estimate by \\cite{Cranmer07}. This latter estimate accounts for the dynamic and thermodynamic properties of the corona at a heliocentric distance of several solar radii.  Our calculated Joule heating rate is less than the Cranmer et al estimate by at least a factor of $3 \\times 10^5$. The currents inferred from the observations of \\cite{Spangler07} are not relevant to coronal heating unless the true resistivity is enormously increased relative to the Spitzer value. However, the same model for turbulent current sheets used to calculate the heating rate also gives an electron drift speed which can be comparable to the electron thermal speed, and larger than the ion acoustic speed. It is therefore possible that the coronal current sheets are unstable to current-driven instabilities which produce high levels of waves, enhance the resistivity and thus the heating rate. ",
        "introduction": "In a recent paper, \\cite{Spangler07} reported radioastronomical observations which were consistent with the presence of coronal currents in the range of hundreds of MegaAmperes to a few GigaAmperes.  This measurement was made using Faraday rotation observations of a radio source occulted by the corona, and the coronal plasma probed was at heliocentric distances of 5.2 to 6.7 $R_{\\odot}$.  In the present paper, I discuss the implications of these observations for the process of coronal heating by Joule heating.  As discussed in \\cite{Spangler07} the currents reported (and summarized in Section 2 below)  correspond to the {\\em net} current within an Amperian loop defined by the two, closely-spaced lines of sight through the corona to the different parts of the radio source.  The measured net current could be, and probably is, a residual due to numerous current filaments with alternate positive and negative current density within the Amperian loop.  This topic is of interest because Joule heating has been identified as the primary mechanism for heating the closed-field part of the corona \\citep{Gudiksen05,Peter06}.  The purpose of this paper is to make model-dependent estimates of the heating rate due to Joule dissipation of these currents.  As expected, the calculation involves introduction of several ``imponderables'', i.e. physical characteristics of the turbulence in the corona which are poorly constrained by observations, but which play an important role in coronal heating.  I feel this exercise is worthwhile in identifying coronal parameters which are important in coronal heating, so that they can be targeted for future observational investigations.  The outline of this paper is as follows.  In Section 2, I briefly summarize the observational results of \\cite{Spangler07} which were the basis of the estimates of the coronal current.  Section 3 is the most important part of the paper; it introduces a model for current-carrying coronal turbulence, and identifies the most important characteristics of this turbulence. A glossary of the variables and parameters introduced in this discussion is given in Table 1. This model is used to obtain an estimate of the volumetric heating rate due to Joule heating. Section 4 briefly considers the possibility that current densities in these sheets could be large enough to generate turbulence via current-driven instabilities, and thus produce high wave levels which enhance the resistivity and thus the Joule heating rate.  Section 5 summarizes what has been learned from this exercise and presents conclusions. ",
        "conclusions": ""
    },
    "0812/0812.0545_arXiv.txt": {
        "abstract": "\\baselineskip 11pt The statistical properties of dark matter halos, the building blocks of cosmological observables associated with structure in the universe, offer many opportunities to test models for cosmic acceleration, especially those that seek to modify gravitational forces.   We study the abundance, bias and profiles of halos in cosmological simulations for one such model: the modified action $f(R)$ theory. The effects of $f(R)$ modified gravity can be separated into a large- and small-field limit. In the large field limit, which is accessible to current observations, enhanced gravitational forces raise the abundance of rare massive halos and decrease their bias but leave their (lensing) mass profiles largely unchanged.  This regime is well described by scaling relations based on a modification of spherical collapse calculations.   In the small field limit, the enhancement of the gravitational force is suppressed inside halos and the effects on halo properties are substantially reduced for the most massive halos.  Nonetheless, the scaling relations still retain limited applicability for the purpose of establishing conservative upper limits on the modification to gravity. ",
        "introduction": "\\label{sec:intro} In the so-called $f(R)$ class of models (see \\cite{Sotiriou:2008rp,Nojiri:2008nt} and references therein) cosmic acceleration arises not from an exotic form of energy with negative pressure but from a modification of gravity that replaces the Einstein-Hilbert action by a function of the Ricci or curvature scalar $R$  \\cite{Caretal03,NojOdi03,Capozziello:2003tk}.   Cosmological simulations are crucial for exposing the phenomenology of $f(R)$ models. In order to satisfy local tests of gravity, $f(R)$ models exhibit a non-linear process, called the chameleon mechanism, to suppress force modifications in the deep potential wells of cosmological structure \\cite{khoury04a,Cembranos:2005fi,Navarro:2006mw,Faulkner:2006ub,HuSaw07a}. Upcoming tests of cosmic acceleration from gravitational lensing, galaxy and cluster surveys have most of their statistical weight in the weakly to fully non-linear regime. Stringent constraints on modified gravity can be expected from current and future surveys, once the impact on observables in the non-linear regime is understood.  In the previous papers  in this series, we have established the methodology for cosmological $f(R)$ simulations \\cite{oyaizu08b} and conducted a suite of simulations that uncover the chameleon mechanism and its effect on the matter power spectrum \\cite{Pkpaper}. In this paper, we continue our exploration of the non-linear aspects of the $f(R)$ model by examining the properties of the basic building blocks of cosmological structure: dark matter halos.   Specifically, we quantify their abundance, i.e. the halo mass function, clustering properties, i.e. the linear bias, and their density profiles, to see how each are modified from the standard cosmological constant, cold dark matter model $\\Lambda$CDM.  We begin in \\S \\ref{sec:meth} with a brief review of the important properties of $f(R)$ models and a discussion of the simulation and analysis methodology. We present our results on halo statistics in \\S \\ref{sec:results} and discuss them in \\S \\ref{sec:discussion}. Throughout we place a special emphasis on exploring the impact of the chameleon mechanism and highlighting differences between the simulations and conventional scaling relations based on linear theory and  $\\Lambda$CDM. These differences expose crucial distinctions that must be considered when observationally testing modified gravity theories. ",
        "conclusions": "\\label{sec:discussion}  Dark matter halos are the building blocks of cosmological observables associated with structure in the universe. Their statistical properties provide many interesting tests of cosmic acceleration, especially of those that seek to modify gravitational forces.  Here we have examined the abundance, clustering and profiles of dark matter halos in $f(R)$ modified gravity models.   In these models, gravitational forces are enhanced below the {\\it local} Compton scale of an extra scalar degree of freedom $f_R$. Generically, this extra force leads to an enhanced abundance of massive halos and a decrease in the bias of such halos, but relatively little change to the density profile or mass correlation around halos of fixed mass. The extent of these effects on halo statistics depends strongly on whether the background scalar field is in the large field ($|f_{R0}|\\gtrsim 10^{-5}$) or small field limit ($|f_{R0}|\\lesssim 10^{-5}$).  In the \\textit{large field limit}, forces are modified everywhere below the {\\it background} Compton scale ($\\lambda_C\\simgt 10\\;\\Mpch$ today \\cite{Pkpaper}).  The modifications in this regime are relatively well described by scaling relations for halo statistics. By modifying spherical collapse parameters to include the enhanced forces, we have shown that the mass function and linear halo bias can be described well by the Sheth-Tormen prescription.   The halo-mass correlation and average density profiles are little changed from $\\Lambda$CDM due to a cancellation of effects from the enhanced forces and decreased bias.  Together these provide a description of the enhanced matter power spectrum that corresponds to a relatively small overestimate of $|f_{R0}|$ by $\\sim 50\\%$ or less.  This level of accuracy more than suffices for an order of magnitude constraint on field values. Moreover, the overestimate depends only weakly on $|f_{R0}|$ and can largely be corrected. In this prescription, concentration uncertainties which are unresolved in our simulations should be marginalized. Concentration uncertainties also arise from baryonic effects in $\\Lambda$CDM \\cite{RuddEtal} and marginalization over these leaves only the more unique intermediate scale deviations to distinguish modifications of gravity \\cite{ZentnerEtal}.  In the \\textit{small field limit}, potential wells of dark matter halos are comparable to or larger than the background $f_R$ field, so that the local Compton wavelength decreases substantially from the background value. Modifications to gravitational forces then decrease in the interior of halos by the so-called chameleon mechanism. This decrease has the effect of bringing deviations in all of the halo statistics down at the high mass end.    At intermediate masses, the excess in the halo abundance can actually increase further due to a pile up of halos which also suppresses the change in the bias.  Scaling relations are not  as easily modified to include the chameleon effect but do still have limited applicability.  Due to a fortuitous cancellation of problems associated with a small background Compton wavelength and the chameleon mechanism, the modified Sheth-Tormen mass function can still be used to provide upper limits on the field values that err only on the conservative side. Likewise the bias description is reasonably accurate for intermediate to high mass halos. We caution that this fortuitous cancellation does not apply to all quantities that can be built out of halo statistics.  For example the halo model for the power spectrum overpredicts the enhancement in the weakly non-linear regime.  To summarize, in the large field limit which encompasses the range that current cosmological observations can test, the scaling relations presented here should already enable strong tests of the model. However, more work in calibrating the effects of $f(R)$ gravity will be required when cosmological observations reach the $\\sim$10\\% percent level precision required to test the small field limit of $f(R)$ modified gravity.  \\bigskip  \\noindent {\\it Acknowledgments}:  We thank N. Dalal, B. Jain, A. Kravtsov, U. Seljak, A. Upadhye, and A. Vikhlinin for useful conversations. This work was supported in part by the Kavli Institute for Cosmological Physics (KICP) 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. WH and ML  were additionally supported by U.S.~Dept.\\ of Energy contract DE-FG02-90ER-40560 and WH by the David and Lucile Packard Foundation. Computational resources were provided by the KICP and by the KICP-Fermilab computer cluster.      \\appendix"
    },
    "0812/0812.4528_arXiv.txt": {
        "abstract": "The Kozai mechanism often destabilises high inclination orbits. It couples changes in the eccentricity and inclination, and drives high inclination, circular orbits to low inclination, eccentric orbits. In a recent study of the dynamics of planetesimals in the quadruple star system HD98800 \\citep{VE08}, there were significant numbers of stable particles in circumbinary polar orbits about the inner binary pair which are apparently able to evade the Kozai instability.  Here, we isolate this feature and investigate the dynamics through numerical and analytical models. The results show that the Kozai mechanism of the outer star is disrupted by a nodal libration induced by the inner binary pair on a shorter timescale. By empirically modelling the period of the libration, a criteria for determining the high inclination stability limits in general triple systems is derived. The nodal libration feature is interesting and, although effecting inclination and node only, shows many parallels to the Kozai mechanism. This raises the possibility that high inclination planets and asteroids may be able to survive in multistellar systems. ",
        "introduction": "\\label{sec:intro}  The Kozai instability is a well-known destabilizing mechanism of high inclination satellites. Arnold (1990), in his book {\\it Huygens \\& Barrow, Newton \\& Hooke}, reports a neat illustration due to Lidov (1963).  Lidov discovered that if the orbit of the Moon is turned through $90^\\circ$, its eccentricity increases so rapidly under the action of the tidal forces of the Sun that it collides with the Earth in four years! The effect was described and analysed by Kozai (1962, 1980) in his studies of the survival of high inclination comets and asteroids.  The Kozai mechanism is a secular effect whereby, above a critical inclination $\\icrit$ of around $40^\\circ$ (depending on the system), the inclination $i$ is coupled to eccentricity $e$ though the Kozai constant~\\citep[see e.g.,][]{Th96} \\begin{equation} H_{\\rm Koz} = \\sqrt{a(1-e^2)} \\cos i \\end{equation} and is driven to vary in cycles between its initial value and the critical value. The argument of pericentre also librates about $\\pm90^\\circ$ \\citep{Ko62,Th96,Ta08}.  In other words, nearly circular, high inclination orbits are driven to high eccentricities in exchange for lower inclination, so that in the Solar system this produces the population of Sun-grazing comets (e.g., Stagg \\& Bailey 1989).  At first glance, this suggests that planets or asteroids in high inclination orbits in multistellar systems cannot survive for long times.  A planet in a circumstellar orbit in a binary star system is known to be subject to Kozai cycles on a timescale \\citep{Ki98,Ta08} \\begin{equation} \\tau_{\\rm{Koz}} \\simeq \\frac{2}{3\\pi}\\frac{P^2_{\\rm bin}}{P_p}(1-e_{\\rm bin}^2)^{3/2}\\frac{m_A+m_B+m_p}{m_B} \\label{eq:koz} \\end{equation} where $P_{\\rm bin}$ is the period of the binary, $P_p$ that of the planet, $e_{\\rm bin}$ is the binary's eccentricity, $m_A$ and $m_B$ the stellar masses and $m_p$ the planet's mass. Surprisingly, then, a recent study of planetesimals in the stellar system of HD 98800~\\citep{VE08} found a long-lived, high inclination circumbinary population.  HD 98800 has four stars, in two close binaries A and B, in highly inclined and eccentric orbits and hosts a circumbinary debris disc around the B binary pair. The system is well approximated as a hierarchical triple star system, as the secondary star in the A binary is fairly small and in a close orbit. High inclination particles are found to be in long-term stable orbits inclined by $55^\\circ$ to $135^\\circ$ to the inner binary B. This is well above the critical inclination relative to both stellar orbits, yet the orbits are not undergoing Kozai cycles.  It is known that the Kozai effect can be disrupted if another mechanism, such as general relativity, tidal forces or other interactions between planets, causes orbital precession on a shorter timescale \\citep{KN91,Wu03, TR05, Ta08}. However, the simulations in \\citet{VE08} deal with the Newtonian dynamics of test particles, so the only additional factor present is the mutual gravitational perturbations of the stars.  This paper explains the origin of the surprising stability.  The inner binary causes a nodal precession, instead of Kozai cycles. This stabilises the test particles against any Kozai instability driven by the outer star. We first present the evidence in favour of this explanation in Section~\\ref{sec:numexpt}, and then provide a more detailed study of the nodal precession, together with the dynamics behind it, in Section~\\ref{sec:analytic}.  Finally, we summarize our conclusions in Section~\\ref{sec:conclusion}.   \\begin{table} \\caption[Initial orbital parameters for the simulations of HD 98800]{ \\label{tab:orbit} The initial orbital parameters for the simulations of system HD 98800 used in \\citet{VE08}. These are the orbital elements used to evolve the system from 1 Myr in the past to the present state, so the starting parameters for the inner binary are also different for each wide orbit case. The longitudes given are relative to the arbitrary reference direction of the simulations.} \\begin{center} \\begin{tabular}{lcccccc} \\hline Orbital               & \\multicolumn{3}{c}{Wide orbit A-B} & \\multicolumn{3}{c}{Inner orbit Ba-Bb}      \\\\ Parameter             &  I     & II     & III                       & I & II & III                       \\\\ \\hline $m_1$ ($M_\\odot$)     & \\multicolumn{3}{c}{1.281}                  & \\multicolumn{3}{c}{0.699}          \\\\ $m_2$ ($M_\\odot$)     & \\multicolumn{3}{c}{1.3}                    &  \\multicolumn{3}{c}{0.582}         \\\\ $a$ (au)              & 61.9   & 67.6  & 78.6                      & 0.983 & 0.983 & 0.983              \\\\ $e$                   & 0.3    & 0.5   & 0.6                       & 0.640 & 0.790 & 0.570              \\\\ $i$ (${}^\\circ$)      & 130.7  & 144.7 & 131.7                     & 0.0   & 0.0   & 0.0                \\\\ $\\omega$ (${}^\\circ$) & 214.1  & 323.4 &  65.9                     & 82.0  & 11.3  & 34.89              \\\\ $\\Omega$ (${}^\\circ$) & 224.5  & 309.2 &  66.6                     & 0.0   & 0.0   & 0.0                \\\\ $M$ (${}^\\circ$)      & 166.3  & 90.8  &  16.2                     & 225.9 & 32.6  & 259.7              \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "\\label{sec:conclusion}  This paper presents a foray into an area of dynamics that has not been extensively studied, namely small particles around a highly eccentric binary star in a hierarchical triple system. At outset, such particles might be expected to be disrupted by the Kozai instability.  Here, we have demonstrated the existence of stable, high inclination circumbinary test particles. They owe their stability to the high eccentricity of the inner binary. This is somewhat surprising, as it goes against every expectation of planetary dynamics. The inner binary, instead of inducing Kozai cycles, causes smooth inclination variations and nodal precession for certain initial longitudes. This suppresses Kozai cycles that would otherwise occur due to the outer star in the hierarchical triple.  An analytical theory to explain these variations is desirable. However, the stars are very eccentric, their mass ratio is high, and the test particle's inclination is also large, leaving no obvious small parameters for the standard perturbation expansions of celestial mechanics. The semimajor axis ratios are not very small, but do not explain why existing secular theories (e.g.  \\citealt{Ko62,FKR00}) are unable to describe the dynamics. The libration islands, inclination variations and critical initial inclination are surprisingly familiar to, but contrasting with, the dynamics of the Kozai mechanism. The Kozai mechanism occurs if $\\dot\\omega \\approx 0$, whereas the nodal libration appears to be a similar process occurring when $\\dot\\Omega \\approx0$.  There exists an approximate integral of motion -- namely, the test particle's component of angular momentum along the direction of the line of apses of the binary star's orbit. This also is analogous to the Kozai mechanism, which has an approximate integral of motion.  The high inclination libration feature may have important consequences for planetary stability in circumbinary orbits. Many stellar members of triples have exchanged into the system, so high mutual planet-star inclinations are very likely. If there are regions of stability then the outlook for planetary systems in these environments is more promising than previously thought."
    },
    "0812/0812.3414_arXiv.txt": {
        "abstract": "Baryon Acoustic Oscillations (BAO) in the radial direction offer a method to directly measure the Universe expansion history, and to set limits to space curvature when combined to the angular BAO signal. In addition to spectroscopic surveys, radial BAO might be measured from accurate enough photometric redshifts obtained with narrow-band filters. We explore the requirements for a photometric survey using Luminous Red Galaxies (LRG) to competitively measure the radial BAO signal and discuss the possible systematic errors of this approach. If LRG were a highly homogeneous population, we show that the photo-z accuracy would not substantially improve by increasing the number of filters beyond $\\sim 10$, except for a small fraction of the sources detected at high signal-to-noise, and broad-band filters would suffice to achieve the target $\\sigma_z = 0.003 (1+z)$ for measuring radial BAO. Using the LRG spectra obtained from SDSS, we find that the spectral variability of LRG substantially worsens the achievable photometric redshift errors, and that the optimal system consists of $\\sim$ 30 filters of width $\\Delta \\lambda / \\lambda \\sim 0.02$. A $S/N > 20$ is generally necessary at the filters on the red side of the $H\\alpha$ break to reach the target photometric accuracy. We estimate that a 5-year survey in a dedicated telescope with etendue in excess of 60 ${\\rm m}^2\\, {\\rm deg}^2$ would be necessary to obtain a high enough density of galaxies to measure radial BAO with sufficiently low shot noise up to $z= 0.85$. We conclude that spectroscopic surveys have a superior performance than photometric ones for measuring BAO in the radial direction. ",
        "introduction": "An important observable to constrain the nature of dark energy is the Hubble parameter $H(z)$ \\cite{SVJ05}, since it constitutes a more direct probe to the dark energy equation of state than the angular diameter distance $d_a(z)$ or the luminosity distance $d_L(z)$, which depend on an integral of $H(z)$. Einstein's equations imply that a homogeneous and isotropic universe, which is described by the FRW metric, that is composed of matter and dark energy with equation of state $p_Q = w_Q(z) \\rho_Q$ expands according to \\begin{eqnarray} \\frac{H(z)}{H_0} = \\left[\\rho_T(z)\\over \\rho_T(0) \\right]^{1/2} = \\\\ \\nonumber \\left [ \\Omega_M (1+z)^3+ \\Omega_k(1+z)^2+\\Omega_Q\\exp\\! \\left ( 3 \\int_0^z \\frac{1+w_Q(z')}{1+z'} dz' \\right ) \\right ]^{1/2}\\!, \\end{eqnarray} where the subscripts $Q$, $k$, $M$ and $T$ refer to dark energy, space curvature, matter, and total energy density, respectively. In the absence of space curvature, the quantities $d_A(z)$ and $d_L(z)$ are related to $H(z)$ via $d_A(z)(1+z)=d_L(z)/(1+z) = \\int_0^zdz'/H(z')$.  The acoustic oscillations in the photon-baryon plasma, observed as acoustic peaks in the CMB power spectrum, are also imprinted in the matter distribution at the scale of the sound horizon at the radiation drag epoch, when baryons were released from the photon pressure. The mass distribution is traced by the distribution of galaxies, and the Baryon Acoustic Oscillations (BAO) can be observed as a peak in the galaxy correlation function or as a series of harmonic oscillations in the galaxy power spectrum. The sound horizon at the radiation drag epoch can be computed very accurately from CMB observations (e.g., $r_s=153.2\\pm 2.0$ Mpc from WMAP5, \\cite{wmap5}), so it provides a natural standard ruler. In fact, the galaxy power spectrum can be used to measure both the angular diameter distance through the clustering perpendicular to the line-of-sight, and the expansion rate $H(z)$ through the clustering along the line-of-sight. Therefore, BAO measurements test the relation between $H(z)$ and $d_A(z)$, providing constraints on dark energy and a limit to space curvature (e.g., \\cite{Polarsky05,Huangetal06,Clarcksonetal07,Barenboimetal08} ). The BAO technique is being considered a powerful probe to the nature of dark energy \\cite{sdss,2dF,Percival,Percival07} because of its potential to provide a standard ruler at different redshifts and its robustness to systematic effects.  For an ideal galaxy survey, \\cite{SeoEisenstein07} have shown that a volume of $1$ (Gpc/h)$^3$ at low redshift can constrain $H_0$ to the 7 \\% level. Forthcoming surveys with larger volume are expected to reach the statistical power to constrain $H(z)$ at the \\% level. There are a number of requirements that these surveys need to satisfy for measuring BAO: covering a large survey volume, modeling the effects of galaxy bias and non-linearity, characterizing the covariance between different modes, evaluating the galaxy selection function to sufficient accuracy, reducing photometric calibration errors to low enough levels to avoid contamination of the BAO signal, etc. Measuring the BAO scale in the radial direction demands in addition that galaxy redshifts are measured to a sufficiently high accuracy, $\\sigma_z$, to avoid an excessive smoothing of the BAO peak, which has an intrinsic width $\\delta r \\sim 10$ Mpc, i.e., $\\delta z \\leq \\delta r H(z)/c$. As shown by \\cite{SeoEisen03}, a redshift accuracy $\\sigma_z < 0.003(1+z)$ is required to avoid substantial loss of accuracy of the $H(z)$ measurement from a given survey volume. Spectroscopic surveys usually yield a redshift accuracy much higher than this minimum requirement.  An alternative approach to measure the large number of redshifts required for BAO detection are photometric redshifts from imaging surveys. Broad-band photometry with $\\sim 6$ filters usually reaches only to $\\sigma_ z / (1+z) \\sim 0.03$ for the general population (with red galaxies having slightly smaller photo-z errors than blue ones), insufficient for measuring radial BAO. However, as the Combo-17 survey \\cite{combo17} has demonstrated, the galaxy photo-z accuracy can be improved by using a larger number of narrower filters. Photometric surveys can cover a large area of the sky faster than a spectroscopic survey and reach a higher number density of observed objects. We are therefore motivated to investigate the requirements for a photometric survey with medium to narrow bands to deliver interesting BAO measurements, and the optimization of the number of filters. Previous work has already explored this issue (\\cite{Bridle, SeoEisen03,Dahlen,Benitez}). Here we concentrate specifically on the impact of the number of bands, signal-to-noise, and non-uniformity of the galaxy sample. We conclude with several considerations on the systematic effects that the photometric approach entails.  In our investigation we use both synthetic stellar population models and Sloan Digital Sky Survey (SDSS)-DR6 spectra of Luminous Red Galaxies. We concentrate on LRG because they have several properties that make them particularly useful for BAO surveys: they are a fairly homogeneous population, the form of their spectra makes them particularly suitable for good photo-z determinations, and their high luminosity facilitates reaching a high enough signal-to-noise up to high redshifts. For these reasons, LRG are the target of choice for $z<2$ BAO surveys (at higher redshift the Ly$\\alpha$ forest probably provides the best method for measuring BAO; see \\cite{McDonaldEisenstein}). The main conclusion we reach is that a spectroscopic survey is superior to narrow-band photometric surveys for measuring the radial BAO signal. We also show that if it were possible to find a very homogeneous population of LRG (with spectra closely matched by a single spectral template), then the photometric redshift accuracy would not substantially increase with the number of filters used beyond a total number of $\\sim 10$ at a fixed total exposure time. This is not the case for galaxies measured at high signal-to-noise, for which a larger number of filters is optimal, but this high signal-to-noise cannot be achieved for a large enough number of objects in a way that is competitive with the spectroscopic approach. In reality, however, the variability of realistic galaxy spectra worsens the photometric redshift accuracy, making it optimal to increase the number of filters to $\\sim 30$ and requiring a higher signal-to-noise per filter to reach the desired redshift accuracy. These results are generally in good agreement with those of \\cite{Dahlen}.  This paper is organized as follows: in \\S~2 we review the requirements for a survey to measure radial BAO. The modeling of the LRG population is described in \\S~3, and in \\S~4 we describe our fiducial survey model. The results are presented in \\S~5, where we analyze in detail the photo-z accuracy as a function of the number of filters, galaxy luminosity and redshift, first for ideal galaxies that match the templates precisely and then for real galaxies with SDSS spectra. Discussion and conclusions are presented in \\S~6 and \\S~7. Throughout this paper, we use a cosmological model with $H_0 = 70$ km s$^{-1}$  Mpc$^{-1}$, $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$. Readers who want to quickly see the main conclusions or our study may wish to go directly to Figure 14. This shows the number density of LRG in several redshift bins with photometric redshift better than $0.003(1+z)$ as a function of the etendue times the exposure time of a survey. It also shows the number density required to reach $nP=1$, necessary to make shot noise subdominant in the Fourier modes near the line-of-sight useful for measuring BAO. ",
        "conclusions": "The principal objective of this paper is to examine the potential for narrow-band photometric surveys to measure the radial  BAO signal  in the LRG  power spectrum. We have revisited this issue after \\cite{SeoEisen03,SeoEisenstein07,Bridle}, addressing the optimization of the number of filters and filter width, discussing the dependence on the spectra of the target population of LRG, and considering the signal-to-noise that could be realistically achieved in a survey. We have found it to be particularly important to take into account the spectral variability of a realistic galaxy population for properly evaluating the photometric redshift accuracy that can be achieved. Our conclusions can be summarized as follows:  {\\it a)} In agreement with \\cite{SeoEisen03, SeoEisenstein07, Bridle} the photometric approach, even if it can achieve the target photo-z error of 0.3\\%, will be advantegeous only if it can cover a much larger fraction of the sky, with a higher galaxy density and reach higher redshifts than  the spectroscopic surveys currently under way. Not only one needs higher source density in a photometric survey to measure the oscillations up to $k \\sim 0.2$, but also the distribution of redshift errors needs to be known accurately to correct for their effect on the power spectrum shape and be able to measure the BAO shape. This is equivalent to  the requirement of obtaining photo-z  with better than 0.3\\% accuracy to fainter magnitudes than spectroscopy. We have  concentrated on the LRG galaxies  because they are a fairly homogeneous population,  the form of their spectra makes them particularly suitable for  good photo-z determinations and because they represent the target of choice for $z<2$  BAO surveys.  {\\it b)} If LRG were a perfectly homogeneous population (with spectra closely matched by a single spectral template), the photometric redshift accuracy would not improve very much beyond a total number of filters $N_f \\simeq 10$, except for galaxies at high signal-to-noise which would have a low number density in a survey with $Et\\simeq 1.5 \\times 10^5\\, {\\rm m}^2\\, {\\rm deg}^2$ hr. In reality, galaxy spectra are variable and the ideal number of filters is $N_f\\simeq 30$ over a wavelength range $530\\, {\\rm nm} < \\lambda < 830\\, {\\rm nm}$, focusing on the redshift range $0.5 < z < 0.9$. Medium-to-narrow band filters photometric  redshift errors show periodic oscillations with a spacing $\\delta z$ that corresponds to the filter widths. For a  filter width of $\\sim 100 \\AA$,  the scale of the oscillation corresponds to the BAO scale. The presence of these oscillations may introduce significant biases in the measurement of the BAO scale (see Section 5.2).  {\\it c)} The variability of realistic galaxy spectra demand high resolution (or more filters in a photometric survey) and high signal-to-noise.  We have quantified, with realistic simulations of a photometric survey, the  minimum galaxy luminosity and the corresponding number density for which 0.3\\% photo-z error can be obtained as a function of the survey Etendue$\\times$  exposure time $Et$ (see Figure \\ref{fig:etendue}). We estimate the minimum value of $Et$ necessary to measure the full second baryonic acoustic oscillation in the power spectrum, at $k<0.164$ $h$/Mpc  near the line-of-sight, with shot noise $nP > 1$ and up to $z=0.85$, at $Et\\simeq 3\\times 10^5\\, {\\rm m}^2\\, {\\rm deg}^2$ hr. The actual required value of $Et$ would likely be larger when taking into account that our assumptions for the seeing and the overall throughput of the telescope-camera system are optimistic, and that the presence of outliers, optimistically clipped and ignored, will degrade the estimated performance. Moreover, the true level of galaxy variability should likely be higher than in the SDSS LRG sample we have used (which is for $z \\lesssim 0.4$ and $L \\gtrsim 3 L_*$), because the variability of LRG spectra is expected to increase with redshift and decrease with luminosity.  Current spectroscopic surveys (e.g.,  SDSSIII--BOSS) are obtaining spectra of the massive ($ > 2 L^{*}$) population of LRG at $z<0.75$ to measure radial and tangential BAO scale. Here we have investigated the conditions necessary to obtain the required photo-z accuracy of $0.3$\\% from a similar or  fainter population using photometry, which in principle is less expensive in terms of observing. We have concluded that  an etendue of $30$ ${\\rm m}^2\\, {\\rm deg}^2$ is required for tracing the same population as BOSS ($z<0.75$, $L > 2 L^{*}$)  with radial BAO signal degraded  only by $\\sim 60$\\% due to photo-z errors. But to improve on this i.e.,  to reach $L_*$ galaxies   at slightly higher redshifts,  the minimum etendue required is $60$ ${\\rm m}^2\\, {\\rm deg}^2$ for a 5-year survey (corresponding to a 4m telescope with the largest fields of view that have been made), and probably needs to be larger given our optimistic assumptions we have made for the seeing, throughput, observing time, removal of outliers and systematic errors.   In this paper we have restricted our attention to the objective of measuring radial BAO with LRG. However, a large-area imaging optical survey with a large number of narrower bands than existing surveys would have many other applications, and might lead to a large number of interesting astronomical discoveries. We have shown in this  paper that, for realistically achievable values of $Et$ at the present time, the density  of LRG sources that could be measured with a narrow-band photometric survey with the target redshift accuracy for radial BAO will not be large enough to make it competitive with a spectroscopic survey. Therefore, a photometric survey with a large number of optical narrow-bands needs to find its scientific justification in other astronomical applications."
    },
    "0812/0812.3849_arXiv.txt": {
        "abstract": "{Grains in disks around young stars grow from interstellar submicron sizes to planetesimals, up to thousands of kilometres in size, over the course of several Myr. Thermal emission of large grains or pebbles can be best observed at centimetre wavelengths. However, other emission mechanisms can contribute, most notably free-free emission from stellar winds and chromospheric activity.} {We aim to determine the mechanisms of centimetre emission for three T~Tauri stars. WW~Cha and RU~Lup were recently found to have grain growth at least up to millimetre sizes in their circumstellar disks, based on millimetre data up to 3.3~mm. CS~Cha has   similar indications for grain growth in its circumbinary disk.} {The T~Tauri stars WW~Cha and RU~Lup were monitored over the course of several years at millimetre and centimetre wavelengths, using the Australia Telescope Compact Array (ATCA). The new ATCA 7~mm system was also used to observe CS~Cha at 7~mm.} { WW~Cha was detected on several occasions at 7 and 16~mm.  We obtained one detection of WW~Cha at 3.5~cm and upper limits only for 6.3~cm.  The emission at 16~mm was stable over periods of days, months and years, whereas the emission at 3.5~cm is found to be variable. A second young stellar object, Ced~112~IRS~4, was found in the field of WW~Cha at 16~mm. RU~Lup was detected at 7~mm. It was observed at 16~mm three times and at 3 and 6~cm four times and found to be variable in all three wavebands. CS~Cha was detected at 7~mm, but the signal-to-noise was not high enough to resolve the gap in the circumbinary disk. The typical resolution of the 7 and 16~mm observations were 5--10 arcsec with rms $\\sim0.2$ mJy. } {The emission at 3, 7 and 16~mm for WW~Cha is well explained by thermal emission from millimetre and centimetre-sized ``pebbles.'' The cm spectral index between 3.5 and 6.3~cm is consistent with the emission from an optically-thick ionised wind, although the high variability of the cm emission points to a non-thermal contribution. The spectral energy  distributions of both RU~Lup and CS~Cha from 1 to 7~mm are consistent with thermal emission from mm-sized grains.  The variability of the longer-wavelength emission for RU~Lup and the negative spectral index suggests non-thermal emission, arising from an optically-thin plasma.} ",
        "introduction": "\\label{sect: introduction}  Planet formation takes place in the disks around young stellar objects (YSOs), where submicron-sized grains have to grow to planet size in about 10~Myr or less \\citep{haisch:2001, carpenter:2005, setiawan:2008}. The very first step, where the grains grow from submicron size to several microns, may be traced by the 10-$\\mu$m silicate feature. Silicates of a few microns in size give considerably flatter features than those of submicron sizes \\citep[e.g.][]{przygodda:2003}. However, the strength and shape of the 10-$\\mu$m feature may also be influenced by the crystallinity of the grains \\citep[e.g.,][]{meeus:2003}, or by their porosity \\citep[][]{voshchinnikov:2008}. Furthermore, the 10-$\\mu$m silicate feature only probes the surface layers of the circumstellar disks.  Growth up to millimetre (mm) sizes can be more readily observed by looking at the mm slope in the spectral energy distribution (SED). From the slope $\\alpha$, where $F_\\nu \\propto \\nu^{\\, \\alpha}$, the opacity index $\\beta$ in the mm regime, where $\\kappa_\\nu \\propto \\nu^{\\, \\beta}$, can be obtained through \\begin{equation} \\label{eq: beta} \\beta \\approx (\\alpha - 2)(1 + \\Delta), \\end{equation} where $\\Delta$ is the ratio of optically thick to optically thin emission from the disk \\citep{beckwith:1990, beckwith:1991, rodmann:2006}. An opacity index $\\beta \\approx 1$ indicates that grains have grown at least up to mm sizes, irrespective of their exact physical structure \\citep{draine:2006}. Indeed, about 30 sources with grains up to at least mm sizes have been found over the past several years \\citep[][]{natta:2004,rodmann:2006,andrews:2007,lommen:2007}.  The next step, growth to centimetre (cm) sizes and beyond, is much harder to observe directly. The reason for this is that at cm wavelengths, where these ``pebbles'' can be observed, the emission is several orders of magnitudes weaker than at mm wavelengths. Furthermore, other processes such as chromospheric activity \\citep[][]{forbrich:2007}, gyrosynchrotron emission \\citep[][]{andre_etal:1988}, electron-cyclotron maser emission \\citep[][]{dulk:1985,smith:2003} or an ionised wind \\citep[][]{girart:2004} may also contribute significantly to the cm emission.  A particularly interesting source of variable cm emission may be that from (interacting) helmet streamers \\citep{massi:2006,massi:2008}. To rule out these other cm emission mechanisms, the sources have to be monitored over days, months, and years, to ascertain that the cm emission is optically thin and not variable, and ideally be spatially resolved. So far, TW~Hya is the only T~Tauri star that has been monitored over sufficient time periods to safely characterise the cm emission as thermal emission from pebbles \\citep{wilner:2005}. [\\citet{natta:2004} and \\citet{alonso-albi:2008} also claim thermal emission from cm-sized grains though their observations have not yet been monitored over an extended period of time.]  \\citet{lommen:2007} observed T~Tauri stars in the southern constellations Lupus and Chamaeleon with the Australia Telescope Compact Array (ATCA\\footnote{The Australia Telescope Compact Array is part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.}) at 3.3~mm, and found three sources that were resolved and had an opacity index $\\beta < 1.0$. We chose the two brightest of these, WW~Cha and RU~Lup, for follow-up observations with the ATCA at 7~mm through 6~cm, the results of which are presented in this paper. The source CS~Cha was not resolved at 3~mm but found to have an opacity index of $\\beta = 1.0 \\pm 0.6$ by \\citet{lommen:2007} and was added for high-resolution observations at 7~mm.  WW~Cha is located in a reflection nebula in the Ced~112 region of the Chamaeleon~I molecular cloud, and is thought to drive the highly collimated jets HH~915 and HH~931 \\citep{bally:2006, wang:2006}. \\citet{wang:2006} claim that two near-infrared molecular-hydrogen emission knots detected by \\citet{gomez:2004} on the opposite side of WW~Cha may be the counterparts of HH~915. WW~Cha has a relatively weak 10-$\\mu$m amorphous-silicate feature \\citep{przygodda:2003}, indicating that the surface layers are dominated by micron-sized grains. No clear crystalline features are detected in the 10-$\\mu$m region. \\citet{reipurth:1996} placed the bright mm source Cha-MMS2 about $9\\arcsec$ from IRAS~11083-7618, which is the infared counterpart to the T~Tauri star WW~Cha. However, given the similar fluxes found from SEST observations by both \\citet{henning:1993} and \\citet{reipurth:1996}, the lack of other mm sources within several tens of arcseconds, and the fairly large beam used for these and the IRAS observations, we claim that Cha-MMS2 and WW~Cha are one and the same source. Another T~Tauri star, Ced~112~IRS~4, is located about $40\\arcsec$ to the north from WW~Cha.  RU~Lup is a very active and well-studied T~Tauri star, showing variability in the optical, UV and X-ray bands \\citep[e.g.][]{lamzin:1996, stempels:2002, herczeg:2005, robrade:2007}. The mass accretion rate onto the central star is found to be relatively large, at $\\sim 10^{-7}$~M$_\\odot$~yr$^{-1}$ \\citep{lamzin:1996, podio:2007}. \\citet{stempels:2002} studied the variations in the radial velocity in RU~Lup, which they found to be periodic with a period of 3.71~days. \\citeauthor{stempels:2002} attribute this periodic variability to long-lived star spots, and find the solution of a substellar companion unlikely. Olofsson et al. (2008, in prep.) show a very strong and boxy 10-$\\mu$m feature for RU~Lup, indicating the presence of submicron-sized particles in the disk photosphere. Contributions from crystalline silicates are detected between 20 and 35~$\\mu$m \\citep{kessler-silacci:2006}.  CS~Cha was classified as a so-called transitional disk by  \\citet{espaillat:2007a}.  The SEDs of transitional disks show a lack of infrared emission, indicating a deficit of warm dust close to the star.  SEDs of transitional disks are well fit by models that include an inner hole, suggesting that the disks are in a  transitional stage. The loss of  \t\twarm dust can be explained by photo-evaporation, by dust growing to larger sizes and effectively moving the flux to longer wavelengths in the SED, and/or by the presence of an unseen planet that sweeps up material in the inner disk.  However, CS~Cha was recently found to be a binary \\citep{gunther:2007}, and so is now classified as a circumbinary disk rather than a transitional disk. Basic parameters of the sources WW~Cha, RU~Lup, and CS~Cha are presented in Table~\\ref{tab: source list}. \\begin{table*} \\caption[]{Source list of sources observed with the ATCA.} \\label{tab: source list} \\centering \\begin{tabular}{lllcclll} \\hline \\hline Source          \t& Cloud\t\t& Distance$^\\mathrm{a}$ & Age$^\\mathrm{b}$ \t& Luminosity$^\\mathrm{b}$\t& Mass$^\\mathrm{b}$\t& Spectral \t\t& Position$^\\mathrm{d}$\t\t\\\\ &\t\t& (pc)\t\t\t& (Myr)\t\t\t& (L$_\\odot$)\t\t\t& (M$_\\odot$)\t\t& type$^\\mathrm{c}$\t&\t\t\t\t\\\\ \\hline WW Cha      & Cha I         & $160\\pm15$\t       & 0.4--0.8\t       & 2.2\t\t\t       & 0.6--0.8\t       & K5\t\t       & (11:10:00.7, -76:34:59)       \\\\ RU Lup      & Lup II        & $140\\pm20$\t       & 0.04--0.5\t       & 2.1\t\t\t       & 2.0--2.8\t       & K7-M0  \t       & (15:56:42.3, -37:49:15.47)    \\\\ CS Cha      & Cha I         & $160\\pm15$\t       & 2--3\t\t       & 1.3\t\t\t       & 0.9--1.2\t       & K4\t\t       & (11:02:25.1, -77:33:35.95)    \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^\\mathrm{a}$]     Distances from \\citet{whittet:1997} (WW~Cha and CS~Cha) and \\citet{hughes:1993} (RU~Lup). \\item[$^\\mathrm{b}$]     Ages, luminosities and masses from \\citet{hughes:1994} (RU~Lup) and \\citet{lawson:1996} (WW~Cha and CS~Cha). \\item[$^\\mathrm{c}$]     Spectral types from \\citet{gauvin:1992} (WW~Cha and CS~Cha) and \\citet{hughes:1994} (RU~Lup). \\item[$^\\mathrm{d}$]     RA and dec positions (J2000) from SIMBAD. \\end{list} \\end{table*}  We here present observations of WW~Cha and RU~Lup, taken with the ATCA at wavelengths ranging from 7~mm to 6.3~cm, to determine their cm emission mechanisms.  We also present ATCA 7~mm observations of the binary CS~Cha, with the aim to obtain a longer-wavelength flux point and resolve the hole in the dust disk, which has a diameter of $\\sim 85$~AU \\citep{espaillat:2007a}, at mm wavelengths. The observations are described in Sect.~\\ref{sect: observations}, with the basic results presented in Sect.~\\ref{sect: results} and further discussed in Sect.~\\ref{sect: discussion}.  Conclusions are presented in Sect.~\\ref{sec:concs}. ",
        "conclusions": "\\label{sec:concs}  We have been monitoring the mm and cm emission of two T~Tauri stars, WW~Cha and RU~Lup, over the course of several years with the ATCA, and more recently performed 7~mm observations of the third young T~Tauri star CS~Cha. We find that emission up to 7~mm for all three sources is well explained by thermal dust emission from mm-sized grains.  The stability of the 16~mm flux in WW~Cha, along with the low peak brightness temperature at this wavelength, indicates that this emission is dominated by even larger, cm-size ``pebbles\", making it the second protoplanetary disk known to contain such large grains.  The 16~mm emission of RU~Lup may also include dust emission from pebbles, but other emission mechanisms appear to dominate at this wavelength.  This work underlines the necessity to observe young stellar objects at multiple wavelengths and to monitor them over extended periods of time, in order to disentangle the various candidate emission mechanisms at mm and cm wavelengths. The ATCA is well suited to do this and with the upgrade of the correlator that is currently underway, increasing the bandwidth of the telescope by a factor of 16, extended surveys of southern protoplanetary disks  will soon be within reach. This will allow us to put more stringent constraints on the processes involved in the first steps of planet formation, telling us where and when they take place."
    },
    "0812/0812.0976_arXiv.txt": {
        "abstract": "We study the formation of galaxies in a $\\Lambda$ cold dark matter ($\\Lambda$CDM) universe using high resolution  hydrodynamical simulations with a multiphase treatment of gas, cooling and feedback, focusing on the formation of discs. Our simulations follow eight  isolated haloes similar in mass to the Milky Way and extracted from a large cosmological simulation without restriction on spin parameter or merger history. This  allows us  to  investigate  how the final properties of the simulated galaxies correlate with the formation histories of their haloes. We find that, at $z=0$, none of our galaxies contain a disc with more than $20$ per cent of its total stellar mass. Four of the eight  galaxies nevertheless have well-formed disc components, three have dominant spheroids and very small discs, and one is a spheroidal galaxy with no disc at all. The $z=0$ spheroids are made of old stars, while discs are younger and formed from the inside-out. Neither the existence of a disc at $z=0$ nor the final disc-to-total mass ratio seems to depend on the spin parameter of the halo. Discs are formed in haloes with spin parameters as low as $0.01$ and as high as $0.05$;  galaxies with little or no disc component span the same range in spin parameter. Except for one of the simulated galaxies, all have significant discs at $z\\gtrsim 2$, regardless of their $z=0$ morphologies. Major mergers and instabilities, which arise when accreting cold gas is misaligned with the  stellar disc, trigger a transfer of mass from the discs to the spheroids. In some cases, discs are destroyed, while in others, they survive or reform. This suggests that the survival probability of discs depends on the particular formation history of each galaxy. A realistic $\\Lambda$CDM model will clearly require weaker star formation at high redshift and later disc assembly than occurs in our models. ",
        "introduction": "According to the standard picture of cosmic structure formation, dark matter and baryons acquire similar specific angular momentum  during their collapse through external torques from  neighbouring structures (Hoyle 1953; Peebles 1969; White 1984). Discs are formed as the gas cools and condenses within the dark matter haloes conserving its specific angular momentum (White \\& Rees 1978; Fall \\& Efstathiou 1980; Mo, Mao \\& White 1998). In these models, the final properties of discs are closely linked to those of their parent haloes: for example,  disc sizes and stability properties depend on the spin parameter of the dark matter haloes in which they reside.     However, discs are  unstable against rapid changes in the gravitational potential, such as those produced by the accretion of satellites (Toth \\& Ostriker 1992; Quinn, Hernquist \\& Fullagar 1993; Velazquez \\& White 1999). Mergers of gas-poor disc galaxies are expected to produce  early-type galaxies as remnants (Toomre \\& Toomre 1972; Toomre 1977). Isolated galaxy mergers have been extensively studied with hydrodynamical simulations, confirming that major mergers can indeed transform disc galaxies into elliptical-like systems (Barnes 1992; Barnes \\& Hernquist 1992, 1996; Mihos \\& Hernquist 1994, 1996). Given the fragility of discs and the fact that hierarchical models predict mergers to occur quite frequently, the existence of present-day spiral galaxies has been interpreted as an indirect proof of late disc formation in systems with a quiescent recent past.     Numerical hydrodynamical simulations are a powerful tool to study galaxy formation and evolution in its full cosmological context, since mergers are then naturally accounted for. State-of-the-art simulations currently include treatments of star formation, cooling, chemical enrichment and supernova (SN) feedback, and in some cases other processes such as black hole formation, active galactic nuclei (AGN) feedback and/or cosmic ray generation. These have been  successful in reproducing some observed properties of galaxies, such as overall scaling relations, star formation histories and chemical abundances (e.g. Springel \\& Hernquist 2003; Brook et al. 2004; Okamoto et al. 2005; Oppenheimer \\& Dav\\'e 2006; Stinson et al. 2006; Governato et al. 2007; Dalla Vecchia \\& Schaye 2008; Dubois \\& Teyssier 2008; Finlator \\& Dav\\'e 2008; Scannapieco et al. 2008 (hereafter S08); among recent examples). However, such simulations have generally had difficulty reproducing disc-dominated galaxies in typical dark matter haloes, when taking into account  the cosmological setting. A major problem is known as the ``angular momentum catastrophe''. This occurs when baryons condense early and then transfer a significant fraction of their angular momentum to the dark matter as the final galaxy is assembled (Navarro \\& Benz 1991; Navarro \\& White 1994). As a result, discs contain too small a fraction of the stellar mass in comparison to  observed spirals.       In a limited number of  cases, simulations have been able to produce individual examples of realistic disc galaxies in $\\Lambda$ cold dark matter ($\\Lambda$CDM)  universes (e.g. Abadi et al. 2003; Governato et al. 2004; Robertson et al. 2004; Okamoto et al. 2005; Governato et al. 2007; S08; see also Croft et al. 2008). Though, none has produced a viable late-type spiral. These simulations use the  `zoom' technique which allows  an individual galaxy to be simulated at high resolution in its proper cosmological context. The relative success of these models appears to reflect the inclusion of efficient supernova feedback and improved numerical resolution. Typically, the simulated haloes are chosen so as to improve the chances of disc formation: masses similar to the Milky Way, relatively high spin parameters, and quiet merger histories. Such haloes are believed to be the most likely to produce disc galaxies, but given our incomplete knowledge of galaxy formation and evolution, and our repeated failure to reproduce late-type spirals, it is  uncertain whether these assumptions are correct.     In a recent paper (S08), we studied the effects of SN  feedback on the formation of galaxy discs by using cosmological simulations of a Milky Way-type galaxy. This study  used an extension of the Tree-PM SPH code {\\small GADGET-2} (Springel 2005) which includes star formation, chemical enrichment and SN feedback, metal-dependent cooling and a multiphase model for the gas component (Scannapieco et al. 2005, 2006; hereafter S05, S06 respectively). Unlike other schemes, this treatment does not require ad hoc elements such as suppression of cooling  after SN explosions or the explicit insertion of winds. Galactic winds are {\\it naturally} generated when the multiphase gas model and SN feedback are included. In S08, we  showed that this model can produce realistic discs from cosmological initial conditions. This is a consequence of self-regulation of the star formation and the generation of galactic winds. Most combinations of the input parameters yield  disc-like components, although with different sizes and thicknesses, indicating that the code can form discs without fine-tuning the implemented physics. In all cases tested, however, the dominant stellar component by mass was the spheroid.    In this paper, we investigate these issues further by following the evolution of eight different  galaxy-halo systems in a $\\Lambda$CDM universe at higher resolution than in our previous work. We use an updated simulation code {\\small GADGET-3} (Springel et al. 2008)  which includes our multiphase gas and feedback algorithms. Target haloes, similar in mass as the Milky Way and mildly isolated,  were chosen  from a  dark matter only simulation of a cosmological volume, with  no restriction on their spin parameter or merger history. These haloes were resimulated including baryons and at improved mass and force resolution. Our main goal here  is to test whether realistic discs form in typical $\\Lambda$CDM haloes and whether  spin parameter or merger history influences their properties as usually assumed.    Our paper is organized as follows. In Section~\\ref{sims}, we describe the main characteristics of the initial conditions and the implemented physics. In Section~\\ref{results}, we present the  properties of our simulated galaxies at $z=0$, while in Section~\\ref{halo_prop}, we investigate the relation between the present-day morphology of the galaxies and the properties of their parent dark matter haloes. In Section~\\ref{disk_survival}, we study the evolution of discs with time and investigate the processes leading to a morphological transformation. Finally, in Section~\\ref{conclu}, we summarize our main results. ",
        "conclusions": "\\label{conclu}  We have used a series of simulations to study the formation of galaxies in  a $\\Lambda$CDM cosmology. Our simulations have enough resolution to follow the internal structure of galaxies and, at the same time, to keep track of larger scale processes such as mergers, infall and tidal stripping. We have simulated the evolution of eight galaxies with virial masses in the range $7-16\\times 10^{11}$M$_\\odot$. These were selected purely on the basis of their mass ($\\sim 10^{12}$ M$_\\odot$) and in isolation (no companion exceeding half of their mass within $1.4$ Mpc), with no additional restriction on spin parameter or  merger history. The only condition  was the absence of  close massive companions at $z=0$. As a result, our sample, although still small, should be representative for the formation of isolated galaxies similar to the Milky Way.   Dark matter only simulations of these haloes with much higher resolution than those studied here have been carried out for the Aquarius Project of the Virgo Consortium (Springel et al. 2008). In this paper, we have included gas dynamics using a specialized extension of the code {\\small GADGET-3} which includes star formation and supernova feedback (both from Type II and from Type Ia SNe), metal-dependent cooling, a multiphase treatment of the gas component, and photoheating by the UV background (S05; S06). This hydrodynamical code includes no  scale-dependent parameters, and is among the most detailed and least ad hoc attempts to describe the interstellar medium  physics and  SN feedback in galaxy formation. One of its main  advantages is that galactic winds are naturally generated, without the ad hoc prescriptions required in most alternative simulation codes.   With the parameters adopted here, none of our eight haloes grows a $z=0$ disc with more than about $20$ per cent of the total stellar mass. Four of the simulated galaxies do have  well-formed and substantial discs in rotational support, three have dominant spheroids and little discs, and one has no disc at all. The star formation histories of all galaxies are similar, with an early starburst followed by a more quiescent phase and little star formation at $z=0$. Old stars populate the spheroids, while younger stars are found in the discs. As a result, typical half-mass formation times are $\\tau_{\\rm s}\\sim 1-3$ Gyr for spheroids and  $\\tau_{\\rm d}\\gtrsim 4$ Gyr for discs. We find that, regardless of  the presence or absence of a disc, all spheroids have  similar mass profiles and stellar age distributions. While all discs formed from the inside-out, their mass distributions show greater variety.    Even the most massive  discs formed in these simulations produce disc-to-total mass ratios of the order of $0.2$, much lower than those of late-type spirals or of the Milky Way. This is primarily because of the small amount of cold gas available for star formation at late times. Presumably, this  is either due to an excess of star formation at early times which leaves the systems with too little gas to make a large disc, or due to  too efficient feedback which expels the gas which would otherwise make the disc. We will investigate this problem in more detail in a forthcoming paper.    We find  that disc growth in dark matter haloes does not depend on spin parameter. We have significant discs in haloes with spin parameter as low as $0.01$ and as high as $0.05$, and galaxies with little or no disc component in haloes with the same range of spin parameters. We also find  no relation between disc mass fraction and halo spin parameter.   As expected, our  discs have high specific angular momentum, even higher than that of their dark matter haloes. In fact, the discs in our simulations all have  similar specific angular momenta, independently of their mass. Spheroids, on the contrary, span a wide range of specific angular momentum, although in all cases their values are smaller than those of the parent dark matter haloes. Those spheroids with the highest specific angular momenta are found in haloes which have prominent discs at $z=0$.    The variation of disc-to-total ratio with redshift reveals that all our galaxies were able to form discs at some time during their evolution. At high redshift,  some discs  are comparable in mass to their spheroids. However,  discs often fail to survive until the present time. We find that both major mergers and misalignment between the discs and newly accreted gas  can completely or partially destroy them. Only one of our galaxies experiences a major merger which completely  destroys its disc and, at $z=0$, it only has a spheroid. Misalignment between stellar discs and newly accreted cold gas is more frequent and induces instabilities which transfer a significant amount of mass from the discs to the spheroids. Among the galaxies which do have significant discs at $z=0$, there are substantial differences in disc evolution. Discs can be stable since $z=3$; they can be completely destroyed and can regrow; or they can lose part of their mass but survive until $z=0$.   Our results suggest that  discs can form easily in $\\Lambda$CDM scenarios, regardless of the spin parameter of their parent  haloes. However, their survival  is strongly affected by major mergers and by disc instabilities. It is clear that the particular implementation of star formation and feedback physics in our current models, while it produces individual objects which resemble real galaxies, cannot produce a galaxy population which matches observation. In particular, the disc mass fractions of our simulated galaxies are typically much lower than those of real galaxies of similar stellar mass. We produce nothing which looks like the Milky Way or other late-type spirals. Apparently, real galaxies formed fewer stars at early times than our models and substantially more stars in a disc at late times."
    },
    "0812/0812.4308_arXiv.txt": {
        "abstract": "We discuss the possibility to suppress downward atmospheric neutrinos in a high energy neutrino telescope. This can be achieved by vetoing the muon which is produced by the same parent meson decaying in the atmosphere. In principle, atmospheric neutrinos with energies $E_\\nu > 10$~TeV and zenith angle up to $60^\\circ$ can be vetoed with an  efficiency of  $>99$\\%. Practical realization will depend on the depth of the neutrino telescope, on the muon veto efficiency and on the ability to identify downward moving neutrinos with a good energy estimation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.3234_arXiv.txt": {
        "abstract": "This work deals with bifurcation and pattern changing in models described by two real scalar fields. We consider generic models with quartic potentials and show that the number of independent polynomial coefficients affecting the ratios between the various domain wall tensions can be reduced to $4$ if the model has a superpotential. We then study specific one-parameter families of models and compute the wall tensions associated with both BPS and non-BPS sectors. We show how bifurcation can be associated to modification of the patterns of domain wall networks and illustrate this with some examples which may be relevant to describe realistic situations of current interest in high energy physics. In particular, we discuss a simple solution to the cosmological domain wall problem. ",
        "introduction": "Bifurcation is a well known phenomenon, which occurs in several distinct areas of nonlinear science \\cite{R1}. Usually, it can be controlled by some forcing parameter, which drives the system to change from a given behavior to another. Interesting examples can be found, for instance, in a competition between the roll and hexagon structures which appear in an experiment of thermal convection \\cite{C}, in nematic liquid crystals subject to rotating magnetic fields where spiral waves were observed \\cite{LC}, and in periodically forced nonlinear oscillatory systems, where patterns are seen to spring induced by bifurcation, as reported for instance in \\cite{E1,Fr1}.  A distinguishing feature of these studies is that they are all non-relativistic. However, bifurcation also appears in various high energy physics models where static domain wall solutions are shown to bifurcate under the variation of model parameters \\cite{AT}. Other examples are given in \\cite{O1}, relevant for the cosmological evolution in brane-world scenarios, in \\cite{O2}, associated with critical gravitational collapse, and in \\cite{D,O3}, in the context of hybrid inflation.  In the present work, we take a specific direction by studying bifurcation in a class of relativistic models described by two real scalar fields, $\\phi$ and $\\chi$. The basic motivation is to study the possibility that a network of domain walls may bifurcate into different configurations, giving rise to modifications of the network pattern. If we choose two minima arbitrarily, it may be possible that they are connected with two or more distinct orbits. When this happens the system develops a bifurcation, since one can go from a given minimum to another one, following two or more distinct paths. In one dimension, given three minima $j$, $k$, and $l$, the direct $jl$ path is energetically favored over the path $jk+kl$ if the energy density or tension of the $(jl)$ structure is lower than the sum of the other two tensions, that is $\\tau_{jl}<\\tau_{jk} +\\tau_{kl}$. A bifurcation occurs when the inequality becomes an equality, since we can go from $j$ to $l$ following the direct $jl$ path, or alternatively visiting $k$ through the path $jk+kl$, at no extra energy cost. In more than one dimension the analysis needs to take into consideration the intersection angle of the domain walls.  In this paper we consider generic models with quartic potentials with or without a superpotential. In the former case the analysis is simpler since the tension of the domain walls associated with a given BPS sector is fully determined by the values of the superpotential, $W(\\phi,\\chi)$, at the minima of the potential. We take $W(\\phi,\\chi)$ to be a polynomial function of degree $3$ which gives rise to a quartic potential, $V(\\phi,\\chi)$ (note that for superpotentials of order lower than $3$ there are no domain wall solutions).  The situation is more complicated when the two-field model is not the bosonic portion of some supersymmetric theory. In this case, there is no superpotential and the determination of the domain wall tensions is not so straightforward. However, an interesting  procedure was followed  in \\cite{BBL} where a model (the so-called BBL model) with two real scalar fields, $\\phi$ and $\\chi$, was constructed under the assumption that its potential could be written as the sum of two quartic potentials, $V(\\phi)$ and $V(\\chi)$, plus a small interaction term.  Although the model as a whole has no superpotential, the importance of the interaction term can be controlled by adjusting a small parameter, $\\varepsilon$. The domain wall tensions could then be determined analytically for $|\\varepsilon| \\ll 1$.  The evolution of the macroscopic properties of domain wall networks is strongly dependent on whether or not junctions are present and, in the former case, on the dimensionality of the junctions. Although the presence of junctions is not enough to make domain walls a viable dark energy candidate, their impact on the dynamics of the network needs to be taken into account in any quantitative study of cosmological consequences of domain wall evolution \\cite{ammmo1}. In this paper we perform domain wall network simulations of a new class of one-parameter multi-tension domain wall models, showing that it has a number of interesting properties. In particular, we discuss a simple solution to the cosmological domain wall problem which arises naturally in this class of models.  In this paper we study bifurcation and pattern changing in models described by two real scalar fields, with or without a superpotential. The outline of this paper is as follows. We start by reviewing the BBL model in Section II, analytically and numerically calculating  the domain wall tensions and comparing our results with analytical results, valid for $|\\varepsilon| \\ll 1$, obtained in  \\cite{BBL} (see also  \\cite{ammmo}). We discuss how bifurcation arises in this model as well as the corresponding implications for the dimensionality of the junctions which may be present when we consider more than one spatial dimension. In section III we study the other possibility, that is, we consider models which support bifurcation and are described by a superpotential. We investigate this case in detail because it is expected to be of more general interest, since it can be seen as the bosonic portion of a larger, supersymmetric theory. We also discuss some implications of our findings in the light of the cosmological domain wall problem. We end the paper in section IV with some comments and conclusions. ",
        "conclusions": "In this work we have investigated specific models described by two real scalar fields and studied the role of bifurcation in changing domain wall network geometry and topology.  First we have considered a class of models without a superpotential controlled by a small parameter, $\\varepsilon$. In this case a perturbative expansion up to first-order in $\\varepsilon$ or the computation of the tensions associated with straight line orbits allows for the analytical determination of the domain wall tensions in the $|\\varepsilon |\\ll 1$ limit. We tested the analytical results by comparing with numerical simulations and found a very good agreement for $|\\varepsilon | \\lsim 0.2$. We also confirmed that the model bifurcates which, in more than one spatial dimension, results in a change from a phase where $Y$-type junctions are favored ($\\varepsilon > 0$) to another one with only $X$-type junctions ($\\varepsilon < 0$).  Note that there are significant differences on the evolution of the macroscopic properties of domain wall networks which depend on whether or not junctions are present and, in the former case, on the dimensionality of the junctions. Although the presence of junctions does not seem to help much from the point of view of making domain walls a viable dark energy candidate \\cite{ammmo1}, their impact on the dynamics of the network needs to be taken into account in any quantitative study of cosmological consequences of domain wall evolution. A good example are models in which the domain walls separate domains with different values of fundamental couplings.  In this paper we have also considered generic models with two real scalar fields with a quartic potential governed by a cubic superpotential. In this case we have shown that the number of parameters affecting the ratios between the various domain wall tensions can be reduced to $4$. We then studied, in detail, a particular subset of models controlled by a single parameter  which, despite its simplicity, has revealed a very rich structure, with several regions of distinct behavior. We have also shown that these models incorporate a natural solution to the cosmological domain wall problem.  If the $Z_4$ symmetry of the scalar field potential is broken into $Z_2$, by moving $b$ away from unity, a bias is created which tends to favor one minimum of the potential over the others. This type of mechanism might operate in realistic domain wall models, leading to a suppression of the domain walls before they become cosmologically relevant."
    },
    "0812/0812.1231_arXiv.txt": {
        "abstract": "We prove that the anisotropic inflationary background of the Ackerman-Carroll-Wise model, characterized by a fixed-norm vector field, is unstable.  We found the instability by explicitly solving the linearized equations for the most general set of perturbations around this background, and by noticing that the solutions diverge close to horizon crossing. This happens because one perturbation becomes a ghost at that moment. A simplified computation,  with only the perturbations of the vector field included, shows the same instability, clarifying the origin of the problem. We then discuss several other models, with a particular emphasis on the case of a nonminimal coupling to the  curvature, in which vector fields are used either to support an anisotropic expansion, or to generate cosmological perturbations on an isotropic background. In many cases, the mass term of the vector needs to have the ``wrong'' sign; we show that, as a consequence, the longitudinal vector mode is a ghost (a field with negative kinetic term, and negative energy; not simply a tachyon). We comment on problems that arise at the quantum level. In particular, the presence of a ghost can be a serious difficulty for the UV completion that such models require in the sub-horizon regime. ",
        "introduction": "\\label{sect-intro}   Cosmological observations over the past decade have ushered in the long heralded era of precision cosmology with theoretical models of the universe now being compared quantitatively to the data. In particular, Cosmic Microwave Background (CMB) anisotropies measurements from both orbital and sub-orbital experiments and Large Scale Structure (LSS) observations have revealed a universe in overall agreement with the standard inflationary paradigm. This picture assumes an epoch preceding the standard, radiation dominated era when the scale factor increased exponentially with a nearly constant Hubble rate. This quasi deSitter expansion not only provides a solution to the classical cosmological problems (homogeneity, isotropy, monopole and flatness) but also results in a nearly scale-invariant spectrum of super-horizon curvature perturbations which seed the structure formation process.  However, certain features of the full sky CMB maps recently observed by the WMAP satellite \\cite{WMAP} seem to be in conflict with the standard picture. These so called `anomalies' include the low power in the quadrupole moment~\\cite{cobe,wmap1,lowl}, the alignment of the lowest multipoles, also known as the `axis-of-evil'~\\cite{axis}, and an asymmetry in power between the northern and southern ecliptic hemispheres~\\cite{asym}. The statistical significance of these effects has been debated extensively in the literature. For instance, the significance of the anomalous lack of large-angle correlations, together with the alignment of power has also grown in strength with the latest data~\\cite{Copi:2008hw}, with only one in 4000 realizations of the concordance model in agreement with the observations.  Recent studies on properly masked data have shown that an anisotropic covariance matrix fits the WMAP low-$\\ell$ data at the 3.8$\\sigma$ level~\\cite{Groeneboom:2008fz}. \\footnote{An upper bound on the anisotropy has instead been obtained in Ref.~\\cite{ArmendarizPicon:2008yr}.} The fit was done using the full anisotropic covariance matrix and found a best-fit anisotropy at the $10\\%$ level. The increase in significance arises from the fact that much information is contained in the off-diagonal terms of the covariance matrix in these models \\cite{gcp,gcp2} and this is missed by simple fits to the diagonal power spectrum.   Such violations of statistical isotropy are considered at odds with the standard phase of early inflation. However, an albeit more plausible explanation of these anomalies arising from a systematic effect or foreground signal affecting the analysis is not forthcoming.  This has led to a number of attempts at reconciling some of the anomalies with the standard inflationary picture through various modifications.   It has been suggested in~\\cite{gcp,gcp2} that the alignment of lowest multipoles could be could be due to an early anisotropic expansion. Indeed, an anisotropic expansion results in a nonvanishing correlation between different $a_{\\ell m}$ coefficients of the CMB multipole expansion (the formalism for dealing with the cosmological perturbations on an anisotropic background was also developed in~\\cite{uzan1,uzan2}). To realize this with a minimal departure from the standard inflationary picture, the anisotropy was simply set as an initial condition at the onset of inflation. It was then shown in \\cite{gkp} that one of the gravity waves polarization experiences a large growth during the anisotropic era (this is intimately related to the instability of Kasner spaces \\cite{gkp}), which may result in a large $B$ signal in the CMB. Inflation however rapidly removes the anisotropy.  The modes that leave the horizon well after the universe has isotropized were deep inside the horizon while the universe was anisotropic, and one recovers a standard power spectrum at the corresponding scales. The signature of the earlier anisotropic stage are therefore present only on relatively large scales, which were comparable to the horizon when the universe was anisotropic. If the duration of inflation was too long, such scales get inflated well beyond the present horizon, and are irrelevant for phenomenology. Therefore, the above suggestions require a tuned duration of inflation.   To avoid this tuning, one can obtain prolonged anisotropic inflationary solutions by introducing some ingredients that violate the premises of Wald's theorem~\\cite{wald} on the rapid isotropization of Bianchi universes. This has been realized through the addition of quadratic curvature invariants to the gravity action~\\cite{barrow}, with the use of the Kalb-Ramond axion~\\cite{nemanja}, or of vector fields~\\cite{ford}. \\footnote{The possibility of a late time anisotropic expansion has instead been considered in~\\cite{dark}. Although the present statistics of the observed Supernovae does not show any evidence for the anisotropy~\\cite{SN}, such studies are motivated by the large increase of data that is expected in the next few years.} The basic idea underlying these models is that one or more vector fields acquire a spatial expectation value (vev). The background solutions are homogeneous, so that the vev is only time dependent, and aligned along one direction; for definiteness, we take it to coincide with the $x-$axis. The models differ from the way in which the vector field(s) acquires its vev. Most of this work we focus on one of the above models: namely the model by Ackerman, Carroll, and Wise (ACW) \\cite{acw}, in which the vector acquires a fixed norm due to the constraint equation enforced by a lagrange multiplier. There is also a cosmological constant in the model, which is responsible for the accelerated expansion. The vev of the vector field controls the difference between the expansion rate of the $x$-direction and that of the two orthogonal directions.  The original paper studied the power spectrum generated for a test field $\\delta \\chi$ in this background. It was then argued that this spectrum gets transferred to the spectrum of primordial density perturbations through the mechanism of modulated perturbations \\cite{modulated}. The basic idea is as follows: imagine that the effective cosmological constant in the model is replaced by an inflaton field, and that the decay rate of the inflaton is controlled by the vacuum expectation value of the test field $\\chi$. Then, fluctuations $\\delta \\chi$ will result in fluctuations of the decay rate. This generates fluctuations of the energy density of the decay products (since the energy density of the inflaton and of the decay products scale differently with time). The resulting power spectrum was the starting point of the analysis of Ref.~\\cite{Groeneboom:2008fz} cited above.  The claim of~\\cite{Groeneboom:2008fz} renders the ACW model particularly interesting. However, the mechanism of modulated perturbations requires that the perturbations of the metric and of the energy densities which are unavoidably generated during inflation are subdominant; otherwise the resulting density perturbations will be a superposition of the original ones, generated without the test field, plus those imprinted by $\\delta \\chi$ at reheating. The original ACW work computed the perturbations of $\\chi$ on an unperturbed background. The underlying assumption is that the original perturbations in the metric and in the energy densities can be neglected with respect to $\\delta \\chi \\,$. In this work we show that, unfortunately, this assumption does not hold. On the contrary, each mode (in Fourier space) of the linearized perturbations (which are unavoidably generated irrespectively of the test field $\\chi$) becomes nonlinear while it leaves the horizon during inflation. As a consequence, the background solution of \\cite{acw} is unstable, and any phenomenological study based on it is unreliable.  The computation of the instability is rather complicated, due to the large number of perturbations involved. An important step in this computation was performed in Ref.~\\cite{dgw},~\\footnote{Another interesting analysis can be found in Ref.~\\cite{dg}, where the late time stability of the ACW model was studied, assuming that the cosmological constant is replaced by a fluid.} where the linearized equations for the complete set of perturbations of the ACW model were presented. These equations were solved in~\\cite{dgw} in the approximation of very large or very small wavelength of the perturbations. Since the Hubble expansion rates for the ACW model are constant, the wavelength of any mode increases exponentially with time. Therefore, at sufficiently early times during inflation, the wavelength is much smaller than the inverse hubble rates; this is the small wavelength regime. On the contrary, at sufficiently late times during inflation, the wavelength is much greater than the inverse hubble rates; this is the large wavelength regime. The moment at which those length become equal is known as horizon crossing~\\footnote{The language is slightly loose, since there are two different expansion rates for the $x-$direction, and for the perpendicular $y-z$ plane. However, we have in mind a small anisotropy (since a large one is forbidden by observations), so that the two rate are parametrically identical.}. We show here that the linearized perturbations diverge close to horizon crossing, where the equations were not solved in~\\cite{dgw} (apart from this, our computation agrees with that of~\\cite{dgw}). This computation is presented in Section \\ref{sect-linearized}. It is just a conceptually straightforward (although technically involved) computation and explicit solution of the equations for the complete linearized system of perturbations for the ACW model. Each mode of the perturbations becomes unstable at a time $t_*$ close to horizon crossing (since different modes have different wavelengths, they all become unstable at different times). This analysis is all one needs to prove that the ACW background solution is unstable.  To understand the reason for the instability, in Section~\\ref{sect-ghost} we compute the quadratic action for the perturbations. After the action is diagonalized, we find that the kinetic term for one of the modes goes from positive to negative precisely at the time $t_*$. This has two consequences for the perturbations: 1) the equations of motion for the perturbations diverge at this point, as we have seen; 2) the mode is a ghost for $t>t_* \\,$ (we note that this second point alone renders the model unstable, due to the decay of the vacuum in ghost-nonghost pairs, with a rate that increases with the momentum of the produced particles). This last point is very interesting in the light of the findings of Ref.~\\cite{clayton}. It was shown there that, in models in which the norm of the vector field is fixed as in the ACW model, it is possible to construct field configurations with arbitrary large negative energy. Those configurations were obtained ignoring gravity. However, combining the result of Ref.~\\cite{clayton} and ours, we are inclined to conclude that such models are generally unstable due to ghosts.  In Section~\\ref{sect-simplified} we present a simplified computation, with only the perturbations of the vector field included (no metric perturbations). This is by itself an inconsistent computation, since metric perturbations are unavoidably sourced by the perturbations of the vector field. However, we show it for illustrative purposes, since it is much simpler than the complete computations presented in the first two parts. The qualitative results of this simplified computation are in perfect agreement with those of the complete one (the solutions diverge close to horizon crossing, when one mode becomes a ghost). This can be seen with extremely simple algebra: the mode that becomes a ghost is the longitudinal perturbation of the vector field (``longitudinal'' in the $y-z$ plane; see Section~\\ref{sect-simplified} for the technical detail). Including the perturbations of the metric (as was done in the first two parts of this paper) simply ``dresses up'' the problem with technical complication, but does not remove the instability of the model, which is ultimately related to the vector field.  The instability found for the ACW model motivated us to study also different models in which vector fields with broken U(1) invariance play a role during inflation. In Section \\ref{sect-other}, we specifically focus on models in which the mass term is due to a nonminimal coupling of the vector to the curvature \\cite{Dimopoulos1,mukhvect,soda1}. These models also present a ghost. This is not related to an anisotropic expansion (in fact, the background is isotropic for \\cite{Dimopoulos1}, in which the vector fields do not contribute to the expansion, and and can be made isotropic for \\cite{mukhvect} by using $3N$ mutually orthogonal vectors with equal vev), but rather to the specific sign that the mass term needs to have in these models. This is due to the fact that, also in this case, the ill-behaved mode is the longitudinal polarization of the vector, which is present (and, therefore, acquires a kinetic term) only when the mass is nonvanishing. Since this is not a trivial point, as since this has not yet been pointed out for the above models, we present three different proofs of this statement. In addition to the ghost, we have verified that also for the models  \\cite{mukhvect,soda1} the linearized equations for the perturbations diverge close to the horizon crossing, precisely as in the ACW model. The detailed computations will be presented elsewhere.  In Section \\ref{sect-conclusion} we present some general discussions of these models. We do not provide quantitative computations here, as in the rest of the paper, but we present some comments and some worries that we believe deserve consideration. Particularly, we comment on some of the implications of having a ghost state in the theory. In the above models, the mass term provides a hard breaking of the gauge invariance. The interactions of the longitudinal vector polarization become strong at an energy scale parametrically set by the mass of the vector. This will occur for all the interactions involving the longitudinal vector, or mediated by it, renormalizing the values of the couplings between all the fields in the theory. Since the mass in these models is the hubble rate, or below, this may invalidate the predictions that rely on the initial quantization in the short wavelength regime (depending on the precise point in which the theory goes out of control). The most immediate solution to this problem is to provide a UV completion of the theory, in which the gauge symmetry is broken spontaneously from the condensation of a field that becomes dynamical above that mass scale. However, providing a UV completion appears to be a harder task when the longitudinal vector is a ghost rather than a normal field (for instance, in the higgs mechanism, the field responsible for the mass would need to be itself a ghost). We believe that such issues deserve further study.  \\smallskip  Most of the computations presented here we summarized in our previous work \\cite{hcp1}. ",
        "conclusions": "\\label{sect-conclusion}  We have shown by explicit computations that the ACW background solution is unstable. We did so in Section \\ref{sect-linearized} by writing down and solving the linearized equations for the most general set of perturbations of the model. These equations, and their solutions, diverge at some point $t_*$ close to horizon crossing (this occurs for any mode in the Fourier decomposition; we stress that all the modes are decoupled from each other at the linearized level). The computations of Section \\ref{sect-linearized} prove by themselves that the background is unstable (it is possible that the divergency does not occur at the full nonlinear level. However, when the nonlinear interactions become important, the solution will be substantially different from the ACW background solution). The computations the we have presented next clarify the origin of the instability. Specifically, we have shown in Section \\ref{sect-ghost} that one dynamical variable of the system becomes a ghost at $t_*$. The kinetic term for this mode vanishes at $t_* \\,$, which explains why the corresponding equation of motion, and the solutions, diverge at this point. We stress that the presence of a ghost indicates on its own an instability of the vacuum (see below). Our analysis is conceptually straightforward, but algebraically very involved. For this reason, we have also shown that our expressions pass several nontrivial checks. To gain a further understanding on the instability, we have also given, in Section~\\ref{sect-simplified}, a simplified stability study, including only the perturbations of the vector field. The algebra involved is now considerably simpler than in the previous case. We verified that this system presents the same instability of the complete computation: the solutions diverge when one mode becomes a ghost. As we mentioned in the Introduction, the metric perturbations complicate the computation, but do not remove the intrinsic instability of the model, which is due to the vector field.  In a loose sense, the ACW model is ``doubly unstable'', since the linearized equations for the perturbations diverge - we call this ``instability I'' - and since it has a ghost - ``instability II'' (in a generic model, either instability may appear without the other one). Motivated by this finding, we investigated whether other models with vector fields playing some role during inflation are also unstable. To our knowledge, the fist model where anisotropic inflation is due to a vector is the one of Ford~\\cite{ford}. The original proposal studied a subset of perturbations in the infinite wavelength limit. A general computation, performed along the line of the one presented here, shows that this model contains a ghost in the short wavelength regime (it has the instability II mentioned above, but not the instability I). The models \\cite{mukhvect,soda1,Dimopoulos1} are characterized by a nonminimal coupling of the vector to the curvature, with a precise factor. We have showed in Secton \\ref{sect-other} why this leads to a ghost (the presence of a ghost is confirmed by the complete computation of the perturbations; we stress that the instability is due to the specific sign mass term, and not to the anisotropic expansion). The models \\cite{mukhvect,soda1} present also the instability of the type I, as defined above, although this is not manifest. The detailed computations will be presented elsewhere.  An instability of the type I is a clear problem of the linearized theory. The implications of having a ghost instead require an understanding of the full theory. Such a study is beyond the goal of the present work. Nonetheless, we would like to provide some general comments. Although, contrary to what we have presented so far, these final considerations are not based on explicit computations, we believe that they are worth mention, and, possibly, more specific future investigations.  The most obvious problem associated to a ghost is the instability of the vacuum. If a ghost is coupled to a normal field (and, in all the above theories, there are at least gravitational couplings), the vacuum will decay in ghost-nonghost excitations, with a rate which is UV divergent. This is simply due to phase space considerations. In normal situations, a source of any given energy can produce quanta up to some given momentum; this cuts-off the phase space available to the decay products. However, the quanta of a ghost field have negative energy: the higher their momentum, the more negative the energy is. Therefore, even a zero energy source decays into ghost-nonghost excitations conserving both energy and momentum. One may hope that the nonlinear interactions may somewhat arrange to cancel the total decay rate. However, this does not appear likely. The usual approach is to regard theories with ghosts as effective theories, valid only below some energy scale. Inflationary predictions heavily rely on the initial conditions; for instance, a slowly rolling inflaton field results in a nearly scale invariant spectrum of cosmological perturbations because the initial modes of the inflaton have amplitude $\\vert \\delta \\phi \\vert \\propto k^{-1/2} \\,$. This is due to the choice of an initial adiabatic vacuum in the quantized theory for the perturbations, which is performed in the deep UV regime (energy $\\gg H^{-1}$) \\cite{mfb}. In presence of a cut-off, the initial adiabatic vacuum cannot be imposed at arbitrarily early times, and, depending on the cut-off, it may not be possible to impose it at all. This casts doubts on the phenomenological predictions obtained for such models.  In fact, all these theories require a cut-off which makes them invalid at high energies, irrespectively of the sign of the mass term.~\\footnote{Several of the considerations presented here emerged from discussions with N. Kaloper, who we thank for either pointing them out, or for stressing their relevance. We also acknowledge very useful discussions with D.H. Lyth on these issues.} We can see this based on the behavior of massive vector fields at high energies. The models studied here have a gauge invariance that is broken in a hard way by the explicit mass term $M^2 \\, A^2 \\,$ for the vector. It is well known that, in such cases, the interactions of the longitudinal bosons violate unitarity at a scale which is parametrically set by $M$, leading to a quantum theory out of control. For the present models, $M$ is the Hubble rate or below, so that the entire sub-horizon regime may be ill-defined. Although we are aware of explicit computations of this problem only in Minkowski spacetime, we believe that it applies also to the inflationary case, if one has unbroken Lorentz invariance and transitions to a locally flat frame (moreover, during inflation the mass term for the vector is nearly constant, and the momentum is adiabatically varying in the sub horizon regime); this issue is less transparent for the specific ACW model, in which the constraint enforced by $\\lambda$ will modify the interactions of the vector; however, we believe that this concern is legitimate also for that model. One may then adopt the approach of ignoring the ill behaved longitudinal vector mode, and concentrate on the other degrees of freedom in the model. We do not believe that this approach is justified. At high energies, the longitudinal mode will also interact strongly with the other fields in the theory (this will renormalize the coupling constants). Then, depending on exactly when this happens, the quantum theory of the perturbations may be out of control throughout the entire short wavelength regime. If this is the case, all initial conditions would become unjustified, and the theory would lose its predictive power.  We are well aware that these problems, if indeed they are present, will affect also the finding of the present work. However, in our opinion, the computations presented above would not lose their importance. One reason for this is that the instabilities that we have found have not emerged so far in the literature. Different studies have adopted different approximation schemes, which are very useful in the case of scalar fields, as for instance studying separately the short or long wavelength regimes \\cite{dgw}, approximating the source with fluids \\cite{mota1} \\footnote{Clearly, the microscopic description of the fluids considered in \\cite{mota1} does not need to involve vector fields, although this has been considered as one of the possibilities in that work.}, or using the $\\Delta N$ formalism \\cite{Dimopoulos1} (which allows to follow the evolution of the perturbations only in the long wavelength regime). Although such studies are perfectly legitimate in the regime of their validity, and can also give very useful results in the present case, one should be aware that these models require extra checks. We believe that our work offers some nontrivial technical advance in this respect.  A more important reason is that ``curing'' a theory which has a hard vector mass and a ghost is more problematic than curing a theory with only a hard vector mass. The most immediate UV completion of a theory with a hard mass is through a higgs mechanism. The mass would be then due to the vev of a scalar field that becomes dynamical above the scale $M$. In this way the theory remains under control also in the short wavelength regime, and one can apply all the standard computations valid for scalar fields during inflation. However, if $M^2$ needs to have the wrong sign, the scalar field in this UV completed theory needs to be a ghost. UV completing a theory with a condensed ghosts may present analogous problems to the UV completion~\\cite{Cheng:2006us} of the ghost condensation model of \\cite{ArkaniHamed:2003uy}. We hope to return to this issue in some future investigation."
    },
    "0812/0812.3372_arXiv.txt": {
        "abstract": "We present high angular resolution observations of HC$_3$N J=5--4 line and 7 mm continumm emission from the extreme carbon star CIT 6. We find that the 7 mm continuum emission is unresolved and has a flux consistent with black-body thermal radiation from the central star. The HC$_3$N J=5--4 line emission originates from an asymmetric and clumpy expanding envelope comprising two separate shells of HC$_3$N J=5--4 emission: (i) a faint outer shell that is nearly spherical which has a radius of 8\\arcsec; and (ii) a thick and incomplete inner shell that resembles a one-arm spiral starting at or close to the central star and extending out to a radius of about 5\\arcsec. Our observations therefore suggest that the mass loss from CIT 6 is strongly modulated with time and highly anisotropic. Furthermore, a comparison between the data and our excitation modelling results suggests an unusually high abundance of HC$_3$N in its envelope. We discuss the possibility that the envelope might be shaped by the presence of a previously suggested possible binary companion. The abundance of HC$_3$N may be enhanced in spiral shocks produced by the interaction between the circumstellar envelope of CIT 6 and its companion star. ",
        "introduction": "The circumstellar envelope around carbon-rich asymptotic giant branch (AGB) stars is well known to be the site of very active photochemistry and hence a rich source of molecular emission lines. The prime example is the massive carbon rich circumstellar envelope of the nearby AGB star IRC+10216 (CW Leo), where more than 50 molecular species have been detected. Many spectral line surveys in the millimeter wavelength range have been conducted toward IRC+10216 (He et al. 2008, Cernicharo \\& Gu\\'{e}lin 2000, Kawaguchi et al. 1995). These surveys show that the emission lines from cyanopolyyne molecules such as HC$_3$N and HC$_5$N are very prominent in the millimeter wavelengths. Interferometric observations of HC$_3$N lines in the 3-mm band by Bieging \\& Tafalla 1994 with BIMA and by Gu\\'{e}lin et al. (2000) with the IRAM PdBI show that the cyanopolyyne molecules have a hollow shell distribution, consistent with the predictions of chemical models for a carbon rich circumstellar envelope (Cherchneff et al. 1993, Millar \\& Herbt 1994, Millar et al. 2000, Brown \\& Millar 2003). Even higher angular resolution observations of HC$_3$N J=5--4 and HC$_5$N J=16--15 lines in the 7 mm band by  Dinh-V-Trung \\& Lim (2008) with the Very Large Array (VLA), which are detectable over a more radially extended region, revealed the presence of a number of incomplete shells representing different episodes of mass loss enhancement from the central AGB star. These observations show that the emission of cyanopolyyne molecules can be used to trace the small scale structures of the envelope.  The spatial distribution of cyanoplyyne molecules in the circumstellar envelope around other carbon-rich stars has been much less studied at high angular resolution. Here, we present observations of the HC$_3$N J=5--4 line around CIT 6 (RW LMi, GL 1403 or IRAS 10131+3049) with the VLA. CIT 6, a semiregular variable, is and extreme carbon star enshrouded in a massive molecular envelope, very similar to the archetypical carbon star IRC+10216. It has a pulsation period of 640 days, which, based on the period-luminosity relation for evolved stars, translates to a distance of $\\sim$400 pc (Cohen 1980 and Cohen \\& Hitchon 1996). For easy comparison with previous work, we will adopt a distance of 440 pc for CIT 6, which is similar to that used by Schoeier et al. (2002). High resolution optical imaging of Schmidth et al. (2002) shows a series of arcs in the optical reflection nebula around CIT 6, within 2\\arcsec\\, of the star. These arcs comprise the density enhancement in the envelope and indicate variations of mass loss on time scale of several hundred years. Such dusty arcs have previously been seen in a number of other circumstellar envelopes and AGB and post-AGB stars such as IRC+10216 (Mauron \\& Huggins 1999), the Egg nebula (Sahai et al. 1998) and some proto-planetary nebulae (Hrivnak et al. 2001).  Interferometric observations of CO J=1--0 line by Neri et al. (1998) and Meixner et al. (1998) at angular resolution of 3\\arcsec--6\\arcsec\\, reveal a bright central core surrounded by a more diffuse but roughly spherical envelope, expanding at a velocity $\\sim$18 kms$^{-1}$. Using sophisticated radiative transfer models, Schoeier \\& Olofsson (2001) and Schoeier et al. (2002) estimated a mass loss rate of 5 10$^{-6}$ M$_\\odot$ yr$^{-1}$ for CIT 6 by fitting the the strength of CO rotational lines. They also noted that a more satisfactory match between model predictions and observations could be obtained if mass loss variation is introduced into the model. Teyssier et al. (2006) also reached similar conclusion from modelling the single dish observations of CO rotational lines in CIT 6. Apart from CO, several other molecules have been detected in CIT 6, including cyanopolyynes (Zhang et al. 2008). The overall chemical makeup of its envelope is similar to that of the archetypical carbon rich envelope IRC+10216, with the exception of enhanced SiC$_2$ and cyanopolyyne emissions. The distributions of other molecules such as the cyanopolyynes (HC$_3$N and HC$_5$N) and HCN and CN have been mapped at an angular resolution of about 3\\arcsec\\, by Lindqvist et al. (2000). Only the emission of cyanogen radical CN was well resolved into an incomplete and elongated expanding shell.  In this paper we present our high angular resolution observations of HC$_3$N J=5--4 emission from CIT 6 obtained with the Very Large Array (VLA\\footnote{The VLA is a facility of the National Radio Astronomy Observatory, which is operated by Associated Universities, Inc., under contract with the National Science Foundation}). The observations allow us to probe the structure of the envelope at arcsec scale or $\\sim$6 10$^{15}$ cm. ",
        "conclusions": "We have imaged at high angular resolution the distribution of HC$_3$N J=5--4 emission from the envelope around carbon star CIT 6. We found that the emission of HC$_3$N J=5--4 traces (1) a faint outer spherical shell located at a radial distance of 8 arcsec and centered at the position of the AGB star CIT 6 revealed by the detection of 7 mm continuum emission. (2) a thick and incomplete inner shell resembling a one-arm spiral. The presence of multiple shells in CIT 6 suggests that the mass loss from this star is highly anisotropic and episodic. From excitation modelling of HC$_3$N molecules we inferred that the abundance of HC$_3$N is unusually high in CIT 6 in comparison to the well known carbon star IRC+10216. We suggest that the observed spatial distribution of the emission and the inferred high abundance of HC$_3$N might be caused by the presence of a binary companion in a wide orbit around CIT 6."
    },
    "0812/0812.3658_arXiv.txt": {
        "abstract": " ",
        "introduction": "Non-Gaussianities in the Cosmic Microwave Background (CMB) have become very important in the last few years. On the observational side, there has been a huge improvement in the sensitivity with experiments such as WMAP. Upcoming experiments such as Planck will be sensitive to even smaller non-Gaussian signals. On the theoretical side, standard slow-roll inflation predicts an extremely small amount of non-Gaussianities \\cite{Maldacena:2002vr}, undetectable by next generation experiments. However, several large modifications to the standard inflationary picture have been proposed both in the case of single field inflation \\cite{Arkani-Hamed:2003uz,Alishahiha:2004eh,Shandera:2006ax,Chen:2006nt}, where all the models have been unified in an effective field theory description in \\cite{Cheung:2007st}, and in the case of multi-field inflation \\cite{Lyth:2002my,Zaldarriaga:2003my}. Recently even a consistent bouncing cosmology has been proposed \\cite{Lehners:2007ac,Buchbinder:2007ad,Creminelli:2007aq}, which, though significantly less compelling than inflation, does predict a possibly detectable level of non-Gaussianities \\cite{Creminelli:2007aq}. A detection  of primordial non-Gaussianities would therefore imply a radical departure from the standard and familiar slow-roll inflation scenario. Non-Gaussianities are probes of the interactions of the inflaton, and therefore contain an unprecedented  amount of information on inflation. If detected, we would have to abandon the standard slow-roll inflation picture, but would be left with new information to understand the dynamics of the inflaton.  Of course, non-Gaussianities in the CMB are not generated only by the inflaton, there are many contributions coming from the post-inflationary evolution. These can be roughly subdivided into effects which are important around or before recombination, and foreground contamination or secondary anisotropies generated after recombination. The latter include scattering secondaries, such as the thermal Sunyaev-Zel'dovich (SZ) effect \\cite{SZ}  and the kinetic SZ or Ostriker-Vishniac (OV) effect \\cite{OV}; and gravitational secondaries, such as weak lensing and the Rees-Sciama (R-S) effect \\cite{RS}. In order to be able to fully exploit upcoming CMB experiments, it is necessary to understand quantitatively the contributions from all these post-inflationary mechanisms.  The effect of the secondaries on the CMB bispectrum has been already extensively analyzed in the literature. The thermal SZ effect and its frequency dependence are investigated in \\cite{0002238}. The R-S effect is studied in \\cite{9811252v2}. Gravitational lensing and the ISW effect are considered in \\cite{1999PhRvD..60d3504S} and \\cite{9811251v2}. The OV effect is analyzed in e.g. \\cite{astro-ph/0511060v1} and \\cite{0610007v2}. The signal-to-noise for upcoming experiment generated by these secondaries is very large and it will have to be taken into account when looking for primordial non-Gaussianities.  Let us now consider non-Gaussianities generated from around recombination and concentrate on the 3-point function or bispectrum. Due to translational invariance, the bispectrum in Fourier space is a function of three wavevectors which must add up to zero forming a closed triangle. For the squeezed triangle limit, where one of the modes is well outside the horizon, the bispectrum generated around recombination has been obtained exactly  in \\cite{2004PhRvD..70h3532C}. Oversimplifying, we can  imagine that the non-Gaussianities are given by the following relation between the observed curvature perturbation $\\zeta$ and a gaussian random variable $\\zeta_g$: \\be \\zeta(\\vec x)=\\zeta_g(\\vec x)-\\frac{3}{5}f_{\\rm NL}^{\\rm loc.}\\left(\\zeta_g(\\vec{x})^2-\\langle\\zeta_g(\\vec{x})^2\\rangle\\right)\\ , \\ee where $f_{\\rm NL}^{\\rm loc.}$ is a measure of the level of non-Gaussianity, and the superscript loc., which stands for local, refers to the above local-in-space structure of the non-Gaussiantiy. Then, in \\cite{2004PhRvD..70h3532C} it was found that the non-Gaussianity in these squeezed triangles  corresponds to an $f_{\\rm NL}^{\\rm loc.}<1$, which is undetectable even by Planck.  The purpose of this paper is to go beyond this simplified limit, and take a step further in the computation of the non-Gaussianities generated around recombination for all triangles. The full calculation of the bispecturm is a very hard task. In order to obtain it, one is forced to solve the second order Boltzmann and Einstein equations, which are very complicated. We give those in a companion paper \\cite{inprep1}, following and  correcting the equations given in \\cite{bartolo} and \\cite{bartolo2}, and including the effects of the free electron number density perturbation. This we calculated at first order in the approximation of the Peebles' effective 3-level atom \\cite{peebles}. In \\cite{inprep1} we showed that around recombination, the amplitude of the perturbations to the free electron density $\\delta_e$ is enhanced by a factor of $\\sim5$ relative to the baryon density perturbations due to the relatively small timescale of recombination~\\footnote{As usual, we denote the change in a given quantity $q$ by $\\delta q$, while the fractional change we denote by $\\delta_q\\equiv \\delta q/q$}.  Without accounting for the enhancement of $\\delta_e$, one very naively expects that the second order evolution will lead to a CMB bispectrum with an $f_{\\mathrm{NL}}^{\\rm loc.}\\sim 1$, which for an experiment like Planck corresponds to a signal-to-noise ratio of $\\sim0.3$ (see for example \\cite{2004PhRvD..70h3005B}). Consequently, we may expect that the enhanced $\\delta_e$ will lead to an observable three-point signature in the CMB corresponding to an $f_{\\rm NL}^{\\rm loc.}\\sim 5$. In this paper, we therefore  concentrate only on those non-linear perturbations which are induced by $\\delta_e$.  Notice that, as long as the tight-coupling between photons and baryons holds, the CMB temperature does not depend on the electron number density $n_e$, and the contribution from $\\delta_e$ to the second order CMB anisotropies vanishes. This means that the effect of $\\delta_e$ is most important during recombination.  At first order, the solution for the CMB temperature anisotropies $\\Theta\\equiv \\delta T/T$, once expanded in multipoles, is given as an integral over the first order source $S^{(1)}$ and the photon visibility function $g(\\eta)$ \\cite{cmbfast}. For a mode of wave vector $\\vec k$, the solution reads: \\be\\label{firstorder} \\Theta_l^{(1)}(k,\\eta_0)\\sim\\int_0^{\\eta_0} d\\eta g(\\eta) S^{(1)}(k,\\eta)j_l[k(\\eta_0-\\eta)]\\ , \\ee where $j_l$ is the spherical Bessel function, $\\eta$ is conformal time and $\\eta_0$ corresponds to $\\eta$ of today. The order in perturbation theory is given as a superscript in parenthesis and will be dropped whenever that does not cause confusion. If $S^{(1)}$ is expressed using the photon monopole, dipole, and quadrupole without accounting for photon diffusion, then including the effects of photon diffusion approximately amounts to replacing $S^{(1)}\\to S^{(1)} \\exp[-k^2/k_D^2]$, where $k_D$ is the photon diffusion scale (see for example Section 8.5 of \\cite{Dodelson}).  The electron density $n_e$ enters only in the visibility function and the diffusion scale, which means that $\\delta_e$ multiplies the first order source. Thus, any perturbations to $n_e$ will affect $\\Theta$ at second order only. This can be understood also intuitively.  In a homogeneous universe, before, during and after recombination the radiation temperature decreases as $a^{-1}$, $a$ being the scale factor, irrespective of the electron density. This means that a perturbation to the electron density changes the position of the last scaterring surface and the mean free path of the photons before recombination, but not the observed radiation temperature. At first order we perturb $n_e$ and keep the other quantities unperturbed, and therefore the CMB anisotropies are not affected by $\\delta_e$ at this order.  From (\\ref{firstorder}), we expect that the expression for the second order temperature anisotropies after including the perturbation to $n_e$ will be schematically given by~\\footnote{In the paper we expand the perturbations as $$f(\\vec x,p^i,\\eta)=f^{(0)}(p^i,\\eta)+f^{(1)}(\\vec x,p^i,\\eta)+\\frac{1}{2}f^{(2)}(\\vec x,p^i,\\eta)\\ .$$} \\be\\label{guess} \\Theta^{(2)}_{lm}(k,\\eta_0)\\sim 2 \\int_0^{\\eta_0} d\\eta g(\\eta) \\left[{\\delta_g}+2 \\frac{k^2}{k_D^2}\\delta_{k_D}\\right] S^{(1)}(k,\\eta)j_l[k(\\eta_0-\\eta)]\\ , \\ee where $\\delta_g$ and $\\delta_{k_D}$ are the fractional perturbations to the visibility function and diffusion scale. When later in the paper we perform the full calculation, we will see that the above schematic equation is not far from the right expression.  The paper is organized as follows. In sec.~{\\ref{sec:estimates}} we elaborate on eq.~(\\ref{guess}) and provide simple estimates of the induced bispectrum. Then, in sections~\\ref{sec:second_order} and~\\ref{sec: 3-point function}, we go on and present the full calculation of the bispectrum concentrating on modes inside the horizon so that we expect to be able to neglect second order metric perturbations. We summarize the result of the analytical calculation in sec.~\\ref{sec:summary of analytical}. In sec.~\\ref{sec:resultsand discussion}, we show the results of the numerical integration and describe the induced signal. We conclude in sec.~\\ref{sec:summary}. In App.~\\ref{appendix:A} we summarize some useful formulas about rotation invariance. ",
        "conclusions": "}  Let us summarize the results of our analytical treatment. It is easy to realize that our second order expression for $\\Delta^{(2)}$ in eq.~(\\ref{deltalm2main}) exactly reproduces what we expected from eq.~(\\ref{guess}) in the Introduction, which implies that also the form of the bispectrum is  very similar to the one given in sec.~\\ref{sec:estimates}. In fact, if we start from part $(b)$ of the source, we can see that $\\tilde S^{b,2}$ is nothing but the first order source, while $\\tilde S^{b,1}$ is just equal to $\\delta_g$:  part $(b)$ of the source reproduces exactly what we expected from the perturbation to the visibility function. Concerning the perturbation to the damping scale, this is given by part $(a)$ of the source. Also in this case,  $\\tilde S^{a,2}/k^2$ turns out to reproduce nothing but the first order source (where we have not kept control of the metric perturbations, that therefore have disappeared). $\\tilde S^{a,1}$ is just equal to $2\\delta_{k_D}/k_D^2$. Putting all these pieces together, and obviously neglecting the complications coming from rotation invariance, we see that expression~(\\ref{deltalm2main}) for $\\Delta^{(2)}_{lm}$ reproduces what we had anticipated in eq.~(\\ref{guess}).     The final expression we obtained for the bispectrum in  (\\ref{eq:B_total}) is of course quite more complicated. A useful simplification is obtained if we express $\\Delta_l^T$ given in (\\ref{firstT}) using the standard photon temperature transfer function, $\\Theta_l$ (given in e.g.~\\cite{Dodelson} and which is returned from CMBFAST for example \\cite{cmbfast}). Combining \\be\\label{transferfunction} \\Delta_l^{T,(1)}(k,\\eta_0)=(-1)^l 4\\sqrt{4\\pi(2l+1)}\\Theta_l\\ , \\ee where $T$ stands for transfer function, with (\\ref{eq: legendre_to_spherical}), we see that $\\Delta_l^{(1)}=(-i)^l 4 \\Theta_l$. Using these transformations we can check that the multipoles of the CMB anisotropies defined in (\\ref{alm}) reduce to the standard form  (up to an irrelevant phase): \\be\\label{almdef} a_{lm}=(-i)^{-l}\\int\\frac{d^3k}{(2\\pi)^3}\\int d\\vec{n}Y_{lm}^*(\\vec{n})\\Theta(\\vec{k},\\vec{n},\\eta_0)\\ . \\ee  Further, since the source contains multipoles only up to the quadrupole, we have the following inequalities:  $0\\le l_1''\\le 2\\ $, $l_1-2\\le l_1'\\le l_1+2\\ $, $l_3-2\\le \\tilde l_3\\le l_3+2\\ $. Combining these with the triangle inequalities for $l_1,l_2,l_3$, we can explicitely do the sums over $l_1',l_1'',\\tilde l_3$.   Using (\\ref{transferfunction}) we find~\\footnote{One may wonder why our expression for the bispectrum contains two $k$-integrals, while the bispectrum in the local model contains 3 $k$-integrals. This is due to the fact that the perturbation $\\delta_e$ affects the CMB only locally in real space: this is true for the effect on the visibility function along the line of sight, and approximately in the squeezed limit for the effect due to the damping. Technically, this boils down to the fact that the second order source factorizes as in eq.~(\\ref{eq:source_redefinition}), which allows to eliminate one of the $k$ integrals using the two delta functions coming from (\\ref{deltaf}). } \\begin{eqnarray}\\label{eq:B_total_simple} &&\\langle a_{l_1m_1}a_{l_2m_2}a_{l_3m_3} \\rangle=\\\\ \\nonumber &&\\ \\ \\  \\ \\ \\ \\ \\ \\ \\mathcal{G}_{l_1l_2l_3}^{m_1m_2m_3}i^{l_1+l_2+l_3}\\int_0^{\\eta_0} dx g(\\eta_0-x) \\left[\\alpha^{a}_{l_2}(x)\\beta^{a}_{l_3}(x)+(a\\to b)\\right]+ {\\rm 5 \\ symm.}\\ ,\\\\ &&\\alpha^{i}_{l_2}\\equiv\\frac{2}{\\pi}\\int dk_2 k_2^2 \\Theta_{l_2}(k_2,\\eta_0)P(k_2) j_{l_2}(k_2 x_{\\tilde\\eta}) \\tilde S^{i,1}(k_2,\\eta_0-x)\\ ,\\\\ &&\\beta^{i}_{l_3}\\equiv\\frac{1}{8\\sqrt{\\pi}}\\frac{2}{\\pi}\\int dk_3 k_3^2 \\Theta_{l_3}(k_3,\\eta_0)P(k_3)\\times\\\\ &&\\left( j_{l_3}(k_3 x) \\tilde S^{i,2}_0(k_3,\\eta_0-x) +\\frac{\\sqrt{3}}{2l_3+1} \\left[l_3j_{l_3-1}(k_3 x)-(l_3+1)j_{l_3+1}(k_3 x)\\right] \\tilde S^{i,2}_1(k_3,\\eta_0-x)+\\right.\\nonumber\\\\ &&\\left.+\\frac{\\sqrt{5}}{2}\\left[3j''_{l_3}(k_3 x)+j_{l_3}(k_3x)\\right]\\tilde S^{i,2}_2(k_3,\\eta_0-x) \\right)\\ . \\nonumber \\end{eqnarray} where $x_{\\tilde\\eta}=x$ for all terms in the source, except the term containing the integral over $\\delta_e$ in $\\eta'$, in which case $x_{\\tilde\\eta}=\\eta_0-\\eta'$, and $j_{l_2}$ must be brought inside the integral over $\\eta'$ in $\\tilde S^{b,1}$. This is our final expression for the bispectrum generated during recombination, which we will evaluate in the next section. Notice that in the limit in which the time dependence of $\\delta_e$ is slow compared to the width of the visibility function, $\\alpha_l^{a}$ and  $\\alpha_l^{b}$ can be evaluated at $x=\\eta_r$ and approximately brought out of the $x$ integral. What remains is proportional to the variation of $C_{l_3}$ with respect to $\\log X$, where $X$ are the quantities perturbed by $\\delta_e$, reproducing the approximate expression for the bispectrum  found in (\\ref{eq:approx_bispectrum}).         }  \\subsection{Signal-to-noise of the bispectrum}  We evaluate numerically eq.~(\\ref{eq:B_total_simple}), working in synchronous gauge, using the following cosmological parameters: $$(\\Omega_b,\\,\\Omega_\\Lambda,\\,h,\\,T_{\\rm cmb}\\,,Y_p,\\,n_s,\\tau)=(0.0441,\\,0.742,\\,0.719,\\,2.725,\\,0.24,\\,0.963,\\,0.087)\\, ,$$ and we are now ready to compute the signal-to-noise of the bispectrum induced by perturbations to the recombination history.  We concentrate on a full sky cosmic variance limited experiment between $l_{\\rm min}$ and $l_{\\rm max}$, without accounting for lensing~\\footnote{This is a good approximation up to $l$ of order 3000}. The Fisher matrix between two bispectra $i$ and $j$ is given by (see for example \\cite{komatsu}): \\be\\label{Fish} F_{ij}\\equiv \\sum\\limits_{l_{\\rm min}\\leq l_1 \\leq l_2 \\leq l_3 \\leq l_{\\rm max}}\\frac{B^{(i)}_{l_1l_2l_3}B^{(j)}_{l_1l_2l_3}}{C_{l_1}C_{l_2}C_{l_3}\\Delta_{l_1l_2l_3}}\\ , \\ee where  $\\Delta_{l_1l_2l_3}=1,2,6$ for triangles with no, two or three equal sides; and $C_l$ is the CMB angular power spectrum. We decide to  concentrate only on the bispectra from recombination and from the local model (with $f_{\\rm NL}^{\\rm loc.}=1$). The local model is a good representative for the non-gaussianities one can expect from inflation (in this case multi-field inflation \\cite{Creminelli:2004yq, Cheung:2007sv} ), and therefore it offers a qualitative measure of the degradation between the primordial and the cosmological signals~\\footnote{Since we will find that the degradation on the local model due to recombination is relatively small, we can avoid to compute explicitly the degradation induced on the equilateral model, which can come from single field inflation.}.  The signal-to-noise ratio $S/N$ for the recombination bispectrum is given by \\be\\label{SoN} \\left(S/N\\right)^2= F_{rec,rec}\\ . \\ee This assumes a negligible degradation from the local model~\\footnote{If the correlation with the local model is included in calculating the $S/N$, the result for $l_{\\rm min}=100$ and $l_{max}=3000$ is changed only by 2\\%.}.  We plot the total $S/N(<l_{\\rm max})$ as a function of $l_{\\rm max}$ in Fig.~\\ref{fig:sn}. The minimum $l$ we include in the calculation is $l_{\\rm min}=100$. This means that, consistently with our approximation, we sum over modes inside the horizon at recombination, which corresponds to $l\\approx70$. We obtain $S/N\\approx 0.4$ for $l_{max}=3000$. This means that such a signal should not be detectable by a satellite like Planck from the analysis of the temperature signal alone. However, the information contained in the polarization signal is expected to be comparable to the one in the temperature  \\cite{2004PhRvD..70h3005B}, and therefore we expect that if our calculation were extended to  include polarization, the resulting signal may be Planck detectable. The $f_{\\rm NL}^{\\rm loc.}$ which in the local model generates the same signal-to-noise for the given $l_{\\rm min}$ and $l_{\\rm max}$ is usually referred to as $f_{\\rm NL}^{\\rm eff.}$ and thus our result corresponds to \\be f_{\\rm NL}^{\\rm eff.}\\approx -3.5\\pm\\mathcal{O}(1)\\ , \\ee where the correction $\\mathcal{O}(1)$ comes from what we naively expect from the terms not enhanced by $\\delta_e$ that we dropped. The effective $f_{\\rm NL}^{\\rm loc.}$ is a convenient way to parametrize the total signal-to-noise but one should keep in mind that it carries no information about the shape of the bispectrum. The sign of $f_{\\rm NL}^{\\rm eff}$ is chosen according to the convention discussed below.  We also compute the correlation coefficient between the bispectra induced by recombination and by the local model, $r_{rec,loc}\\equiv F^{-1}_{rec,loc}/\\sqrt{F^{-1}_{rec,rec}F^{-1}_{loc,loc}}$ as a function of $l_{\\rm max}$. This is shown in Fig.~\\ref{fig:sn}, where one can see that for low $l_{\\rm max}$ it is equal to -1, and it decreases rapidly as we include higher $l$'s. The Fisher matrix with the local model, for $l_{\\rm max}=3000$, is given by: \\bea F=\\left( \\begin{array}{cc} 0.19 & 0.011 \\\\ 0.011 & 0.016 \\end{array} \\right)\\ , \\eea where index 1 corresponds to the non-Gaussianities from recombination, while index 2 corresponds to  the local model.  From this we can deduce that the bias of the non-Gaussianities from perturbing recombination into the estimate of $f_{\\rm NL}^{\\rm loc.}$ is about 0.7 at this resolution.    In Fig.~\\ref{fig:sn} we also plot the contribution to $S/N$ from each $l_{max}$--, and each $l_{min}$--bin by plotting~$d(S/N)^2/dl_{\\rm min}$ and $d(S/N)^2/dl_{\\rm max}$. It is easy to see that the biggest contribution to the total $S/N$ comes from squeezed triangles with \\be\\label{config} l_2 \\approx l_3\\gtrsim2000,\\ l_1\\sim 200\\ , \\ee i.e. with the small $l$ around the first acoustic peak, and the two high $l$s in the diffusion damping tail. This result is consistent with our expectations stated at the beginning of sec.~\\ref{sec:estimates} based on considering the timescales at recombination. The fact that the signal is not dominated by triangles with the smallest of the $l$'s we consider (it decays for $l\\lesssim 200$) is a confirmation of the viability of our approximation of neglecting metric perturbations. This result matches what was found explicitly in \\cite{2004PhRvD..70h3532C} for triangles  with the smallest side out of the horizon.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=15cm]{SNs.pdf} \\end{center} \\caption{\\small The plot shows the signal-to-noise for the recombination and local model bispectra; $d(S/N)^2/d l_{max}$ and $d(S/N)^2/d l_{min}$ (for triangles with $l_{min}=100$ and $l_{max}=3000$, respectively) of the bispectrum generated around recombination; and the correlation coefficient between the local model and recombination vs $l_{\\rm max}$ with $l_{\\rm min}=100$. The computation was done for an ideal (cosmic variance limited) experiment without accounting for lensing. Clearly triangles with two high $l$s, and a low-$l$ around 200 dominate the recombination signal-to-noise. The local model and recombination bispectra start decorrelating only after high-$l $ modes are included. } \\label{fig:sn} \\end{figure}  \\subsection{The shape of the non-Gaussianities}  Let us discuss the shape of the non-gaussian signal produced by fluctuations in $\\delta_e$. The signal-to-noise per triangle, which is given by  $|B_{l_1l_2l_3}|/\\sqrt{C_{l_1}C_{l_2}C_{l_3}}$, can be read off from eq.~(\\ref{SoN}). It includes the additional weighting from the 3-$j$ symbols which takes into account the phase space distribution of the triangles. The physical information in the bispectrum is contained in the $S/N$ per triangle divided by the phase space weight of that triangle shape, which gives $b_{l_1l_2l_3}/\\sqrt{C_{l_1}C_{l_2}C_{l_3}}$ (we prefer to preserve the information encoded in the sign of $b$). $b_{l_1l_2l_3}$ is the reduced bispectrum and is defined as $B_{l_1l_2l_3}^{m_1m_2m_3}\\equiv\\mathcal{G}_{l_1l_2l_3}^{m_1m_2m_3} i^{l_1+l_2+l_3} b_{l_1l_2l_3}$. It can be read off from (\\ref{eq:B_total_simple})~\\footnote{ The bispectrum is defined up to a phase. Our definition of the bispectrum contains an extra factor of $i^{l_1+l_2+l_3}$ (which is real since $l_1+l_2+l_3$ is even) compared to \\cite{komatsu} due to the extra factor of $i^l$ in our definition of $a_{lm}$ (\\ref{almdef}).} . The reduced bispectrum reduces to the bispectrum in the flat-sky approximation \\cite{komatsu}, and thus it contains the physical information of the three-point function.  Let us comment briefly on the sign of the bispectrum. With our definition, the signs of $B_{l_1l_2l_3}$ and $b_{l_1l_2l_3}$ are the same. These are chosen such that they coincide with the one of the reduced bispectrum in \\cite{komatsu} in the ``local'' model, in which a positive $f_{\\rm NL}^{\\rm loc.}$ parameter corresponds to a negative skewness of the one-point distribution function of the temperature anisotropies, and to a negative reduced bispectrum in the Sachs-Wolfe regime (see eq.~4.24 of~\\cite{komatsu}). With our definition, a generally positive bispectrum (negative $f_{\\rm NL}^{\\rm eff.}$) corresponds to high-$l$ modes having enhanced amplitude in hot long-wavelength patches of the sky.   Let us now analyze the shape of the signal. In Fig.~\\ref{fig:l1700a} we plot $l_3^{3/2}\\times b_{l_1l_2l_3}/(C_{l_1}C_{l_2}C_{l_3})^{1/2}$, for $l_1\\leq l_2\\leq l_3$ (to avoid redundancy) subject to the triangle inequalities for a typical high-$S/N$ $l_3$. In the plot we choose $l_3=1700$. For comparison, we plot the analogous quantity in the case of the local model. The factor of $l_3^{3/2}$ is introduced so that after normalization of the $l_1$ and $l_2$ axes by dividing by $l_3$, the integral over the square of the plotted function, multiplied by the phase space weight (see eq.~(4.54) of \\cite{komatsu}), is equal to the $(S/N)^2$  per logarithmic $l_{\\rm max}=l_3$ bin. The most important feature to notice is that, as anticipated, the signal is peaked on squeezed triangles, but the shape is different from the one of the local model.  \\begin{figure*}[h!] \\begin{minipage}{155mm} (a) \\hskip 68mm (b) \\begin{center} \\epsfig{file=L1700p.pdf, width=7.2cm,bbllx=0,bblly=0,bburx=360,bbury=270} \\epsfig{file=3dLOC.pdf, width=7.2cm,bbllx=-20,bblly=0,bburx=340,bbury=270} \\end{center} \\caption{\\small In panel (a) we show $b_{l_1l_2l_3}^{rec.}/(C_{l_1}C_{l_2}C_{l_3})^{1/2}$ for an high-$S/N$ $l_3$ (we choose $l_3=1700$). For comparison, in panel (b) we plot the analogous quantity from the local model with $f_{\\rm NL}^{\\rm loc.}=1$. The peak in the lower-right corner of panel (a) ($l_1\\approx 200$, $l_2\\approx l_3=1700$) corresponds to the high-$S/N$ squeezed triangles given by (\\ref{config}). These triangles are such that they sample both the Silk damped tail, and the first acoustic peak. As for the local model, the signal is dominated by squeezed triangles, though the shape is clearly different. The acoustic oscillations can be clearly seen. The color scheme of the plots is chosen to highlight $b_{l_1l_2l_3}=0$. }\\label{fig:l1700a} \\end{minipage} \\end{figure*}  For purposes of comparison with the literature (see for example \\cite{komatsu}), we plot in Fig.~\\ref{fig:blll} the reduced bispectrum as a function of the smallest $l$, $l_1$, for isosceles configurations with the higher $l$'s, $l_2=l_3$, ranging from $\\sim2200$ to $\\sim 2300$. As $l_1$ changes, we clearly see the acoustic oscillations in the bispectrum. The change in value as $l_2$ varies is also due to the acoustic oscillation.   \\begin{figure}[h!] \\begin{center} \\includegraphics[width=15cm]{blll.pdf} \\end{center} \\caption{\\small The reduced bispectrum for triangles with $l_2=l_3$. The different curves correspond to equally spaced values of $l_2=l_3$. The bottommost (light-gray) curve is for $l_2=l_3=2200$, while the topmost (dark-gray) curve is for $l_2=l_3=2300$. The similarity between the bispectrum for these triangles and the temperature anisotropy power spectrum is apparent.} \\label{fig:blll} \\end{figure}  In order to get some intuition on how much the signal is scale invariant, in Fig.~\\ref{fig:shapes} we plot the $S/N$ per triangle for isosceles triangles ($l_2=l_3$). Each curve represents a triangle with a certain level of squeezing,  characterized by the ratio $l_1/l_2$, with $l_1\\leq l_2$, as we vary the overall size of the triangle characterized by $l_2$.  The topmost curve corresponds to the most squeezed triangles, which confirms that most of the signal is in that kind of triangles. However, we see that as the size of the triangle varies, the amplitude of the signal fluctuates a lot. This is a violation of scale invariance. In fact, the thick curve represents the analogous quantity for a local shape with $f_{\\rm NL}^{\\rm loc.}=1$ for the maximum squeezing, whose signal is generated by primordial scale invariant perturbations. We find that the equilateral and obtuse triangles have a completely negligible contribution. Only squeezed triangles with ratio $l_1/l_2\\lesssim 0.1$ have a signal larger than the ones corresponding to $f_{\\rm NL}^{\\rm loc.}=1$.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=15cm]{shapes7.pdf} \\end{center} \\caption{\\small The $S/N$ per triangle for triangles of various shapes with $l_2=l_3$. The different curves correspond to different ratio of $l_1/l_2$ which varies between 0.05 and 1 in equal logarithmic steps. Each curve terminates on the left with a dot, where $l_1=100$, which approximately is the minimum $l$ below which we can not trust our solution anymore. On the rightmost side of the plot, there are 40 curves, and this number is reduced to the left. The bottommost (light-gray) curve corresponds to equilateral triangles, while the topmost (dark-gray) curve is for highly squeezed triangles with $l_1/l_2=0.05$. The thick curve represents the analogous quantity for the local model with $f_{\\rm NL}^{\\rm loc.}=1$ and $l_1/l_2=0.05$. From the plot we can see that only extremely squeezed triangles with ratio $l_1/l_2\\lesssim 0.1$ have enhanced $S/N$. Obtuse triangles have negligible contribution, though we do not plot them explicitly. } \\label{fig:shapes} \\end{figure}                          \\subsection{Physics of the bispectrum}     In sec.~\\ref{sec:estimates}, we provided simple estimates for the origin and the size of the effects of the fluctuations of $\\delta_e$ on the bispectrum. After the full calculation, we are now able to check those results. From Fig.~\\ref{fig:sn}, we saw that most of the signal comes from squeezed triangles with $l_1\\ll l_2\\simeq l_3$. In that limit, the leading contributions to $b_{l_1l_2l_3}$ come from the terms $\\alpha^{a}_{l_1}(\\beta^{a}_{l_{2}}+\\beta^{a}_{l_{3}})$ and $\\alpha^{b}_{l_1}(\\beta^{b}_{l_{2}}+\\beta^{b}_{l_{3}})$  of eq.~(\\ref{eq:B_total_simple}). These terms respectively correspond to the following combinations: low-$k$ $\\delta_{k_D}$ and high-$k$ first order source $S^{(1)}$; and low-$k$ $\\delta_g$ and high-$k$ $S^{(1)}$. The other interesting term, high-$k$ $\\delta_g$ and low-$k$ $S^{(1)}$, gives a negligible contribution as we can expect from our estimates. This is exactly the same structure we found in sec.~\\ref{sec:estimates}. In particular, the fact that the highest signal from perturbations to the diffusion scale due to $\\delta n_e$, which we call $\\delta_{k_D}$, comes from squeezed triangles with a slowly varying $\\delta_e$, is a consistency check of our approximate treatment of the diffusion damping at second order in sec.~\\ref{sec: part a}. The contribution from the diffusion damping perturbation to the total $S/N(<l_{max})$ from triangles for which that approximation breaks down is $\\sim10\\%$, which is equivalent to a contribution of $\\Delta f_{\\rm NL}^{\\rm eff.}\\sim 1/2$ to the bispectrum. These order one corrections to $f_{\\rm NL}^{\\rm eff.}$ are of the order of the ones we expect from terms we neglect such as second order corrections from metric perturbations, and are therefore beyond our control.   \\begin{figure}[h!] \\begin{center} \\includegraphics[width=15cm]{sq.pdf} \\end{center} \\caption{\\small The $S/N$ per triangle (keeping the sign of $b$) for typical high-$S/N$ squeezed triangles with $l_2=l_3$ and $l_1=200$. We also show separately the three important contributions to $S/N$. Notice the similarity between this plot and the plot obtained using the approximate treatment of the bispectrum in the squeezed triangle limit, Fig.~\\ref{fig:btot}. We reproduce all qualitative features of that plot, which serves as a check of our code, and allows us to pinpoint the main physical effects generating the bispectrum (see the text). } \\label{fig:sq} \\end{figure}   In Fig.~\\ref{fig:sq} we plot the three contributions discussed above to the $S/N$ per triangle for  high-$S/N$ squeezed triangles.  The bispectrum generated from $\\delta_{k_D}$ is generally positive and it dominates the one from $\\delta_g$ which is generally negative. Thus, the total bispectrum from recombination is generally positive, i.e. the amplitude of the high-$l$ modes is enhanced in hot long-wavelength patches of the sky corresponding to the first acoustic peak scale. Notice that $\\delta_g$ receives contributions from perturbations  to $k_D$ due to the shift in the position of the last scattering surface ($\\delta_{k_g}$ in the notation of sec.~\\ref{sec:estimates}), from perturbations to the phase of the modes at the last scattering surface, and from perturbations to  the probability for the CMB photons to originate from different times within the recombination era ($\\delta {\\rm Area}$ still in the notation of sec.~\\ref{sec:estimates}). As we anticipated, the leading contributions from $\\delta_{k_g}$ and $\\delta_{k_D}$ cancel in the limit in which $\\delta_e$ is much slower than the timescale of recombination, which explains the partial cancellation between the two terms seen in Fig.~\\ref{fig:sq}. This can be seen also in Fig.~\\ref{fig:l1700b}, where we plot the analogous of Fig.~\\ref{fig:l1700a} for the same three contributions.    \\begin{figure*}[h!] \\begin{minipage}{155mm} (a) \\hskip 68mm (b) \\begin{center} \\epsfig{file=LA1700p.pdf, width=7.2cm,bbllx=-20,bblly=0,bburx=340,bbury=270} \\epsfig{file=LBl1700p.pdf, width=7.2cm,bbllx=-20,bblly=0,bburx=340,bbury=270} \\end{center} \\hskip 40mm  (c) \\begin{center} \\epsfig{file=LBh1700p.pdf, width=7.2cm,bbllx=0,bblly=0,bburx=360,bbury=270} \\end{center} \\caption{\\small The same plot as in Fig.~\\ref{fig:l1700a} for the three interesting contributions to the recombination bispectrum. a) Bispectrum generated by low-$k$ $\\delta_{k_D}$; b) Bispectrum generated by low-$k$ $\\delta_{k_g}$; c) Bispectrum generated by high-$k$ $\\delta {\\rm Area}$. Each of the two main contributions is peaked in the squeezed limit. Note that, as anticipated, the contribution to the bispectrum from the perturbations to the visibility function is opposite in sign to the contribution from diffusion damping (see vertical-axis label). }\\label{fig:l1700b} \\end{minipage} \\end{figure*}  If we compare Fig.~\\ref{fig:btot} from our approximate treatment of sec.~\\ref{sec:estimates} with the same one from the full calculation, Fig.~\\ref{fig:sq}, we notice that the semianalytical treatment reproduces with good accuracy the general trends and the oscillation frequencies, phases and relative amplitudes of the various terms. The additional averaging coming from the integration in time of the visibility function in the full calculation suppresses the bispectrum generated by $\\delta$-phase in the full calculation with respect to the semianalytical result, as expected. Very importantly, this tells us that we can trust both our code and the semianalytical analysis, and pinpoint the main physical effects generating the recombination bispectrum, as described in sec.~\\ref{sec:estimates}. In particular, the low-$k$ $\\delta_g$ has a larger oscillation amplitude than the contribution from the low-$k$ $\\delta_{k_D}$ that is due to the perturbation to the phase of the first order source (see  Fig.~\\ref{fig:bL}) entering in (\\ref{blllsecond}). The general upward trends (in magnitude) of the low-$k$ $\\delta_{k_D}$ and $\\delta_g$ are due to perturbations to the diffusion damping generated by $\\delta n_e$ and $\\delta\\eta_r$, respectively. By comparing the shape of the signal induced by a high-$k$ $\\delta_e$ with what we obtained in Fig.~\\ref{fig:bH}, we can confirm that the high-$k$ $\\delta_g$ bispectrum is negligible. Overall, the sum of all these contributions gives a bispectrum which has an $f_{\\rm NL}^{\\rm eff.}\\simeq -3.5$, which is enhanced from the first naive expectations of $f_{\\rm NL}^{\\rm eff.}\\sim 1$.                    The calculation of the bispectrum induced on the CMB by the standard cosmological evolution has become necessary to fully exploit the signal measured by the next generation CMB experiments such as Planck.  In this paper we have computed in an approximate way the bispectrum due to perturbations in the recombination history. In a companion paper \\cite{inprep1} we computed the first order perturbations to the free electron density $\\delta_e$, which is one of the relevant quantities for studying the perturbed recombination history. This quantity does not enter in the first order calculation of the CMB temperature fluctuations which allows us to predict the power spectrum, but instead it enters at non-linear level, and therefore it can induce a bispectrum. What we found is that the perturbation to $n_e$ is about a factor of 5 larger than the other perturbations, which implies the possibility for these fluctuations to generate a bispectrum with $f_{\\rm NL}^{\\rm eff}\\sim 5$, about a factor of 5 larger than naive expectations, and possibly detectable by next generation experiments.  We have therefore set up to make this computation. The full second order calculation of the CMB bispectrum is a very hard task, with many terms that source the second order matter, radiation, and metric perturbations. By concentrating on modes well inside the horizon at recombination, we expect to be able to neglect the second order metric perturbations. This means that our results will be wrong by an effective $f_{\\rm NL}^{\\rm eff.}\\sim{\\cal{O}}(1)$, which is the signal expected to be induced by metric perturbations~\\footnote{Recent estimates in \\cite{Pitrou:2008ak} and \\cite{Bartolo:2008sg} gives an $f_{\\rm NL}^{\\rm equil.}\\sim{\\cal{O}}(10)$ for equilateral configurations, which corresponds to roughly   $f_{\\rm NL}^{\\rm eff.}\\sim{\\cal{O}}(3)$.}. Because of the large value of $\\delta_e$, we have concentrated on those second order source terms that are proportional to $\\delta_e$, and that therefore are expected to give an enhanced effect with respect to the rest.  What we find is that the induced bispectrum on the temperature signal corresponds to an $f_{\\rm NL}^{\\rm eff.}\\simeq-3.5\\pm{\\cal{O}}(1)$. The signal is peaked on squeezed triangles with the lowest $l\\sim200$ and the highest ones at about $l\\gtrsim 2000$,  and it is very far from being scale invariant. Such an $f_{\\rm NL}^{\\rm eff.}$ corresponds to a signal-to-noise for an experiment like Planck of order one half. It is expected that the polarization signal will contain an amount of information comparable to the one in the temperature, and therefore such a signal may be detectable by such an experiment.  We find that the physical origin of the signal can be understood quite simply in terms of three physical effects each one giving an $f_{\\rm NL}^{\\rm eff.}$ of order a few. The first effect is due to the time shift $\\delta \\eta_r$ induced on the time of recombination by a low-$k$ $\\delta_e$ mode. This induces  a perturbation to the phase of the high-$k$ first order mode at the last scattering surface which is proportional to the wavenumber itself, and that therefore grows for more and more squeezed triangles. This growth however does not persist until very high $k$, because there is an averaging over the width of the last scattering surface, and we are left with an $f_{\\rm NL}^{\\rm eff.}$ of order a couple.  A second effect due to the time delay is that it changes the amount of time during which photons can diffuse. This second effect goes together with another effect, the perturbation to the diffusion damping scale $k_D^2\\propto n_e$ due to a low-$k$ $\\delta_e$ mode. In the limit in which the $\\delta_e$ perturbation is much slower that the timescale of recombination (which is approximately the case for low-$k$ modes), the enhanced effect from  the perturbation to $n_e$ is cancelled by the change in the  time during which the photons diffuse, as expected. However, what is left of the perturbations is boosted in the squeezed limit by the square of ratio of the $k$ of the high-$k$ mode to the damping scale $k_D$, and again we are left with an effective $f_{\\rm NL}^{\\rm eff}$ of a few.  By summing all of these effects, we get our final signal corresponding to $f_{\\rm NL}^{\\rm eff.}\\simeq-3.5\\pm{\\cal{O}}(1)$. Notice that this sign of $f_{\\rm NL}^{\\rm eff}$ corresponds to enhancing the short scale power in hotter long scale regions. Though physically well defined, experimentally quite large, and possibly detectable, the effect we find is nevertheless not much larger than the naive expectations for the effects coming from the terms we have neglected, which are expected to give rise to $f_{\\rm NL}^{\\rm eff.}$ of order one. We can not therefore exclude that other sources might give comparable signal. In particular it is not clear to us that other first order quantities that are not exponentially suppressed at high $k$ do not give rise to a similar effect. Therefore, in addition to providing a clear calculation of some of the second order effects, we think that our result further motivates the full calculation of the CMB anisotropies at second order."
    },
    "0812/0812.1970_arXiv.txt": {
        "abstract": "We show that vector theories on cosmological scales are excellent candidates for dark energy. We consider two different examples, both are theories with no dimensional parameters nor potential terms, with natural initial conditions in the early universe and the same number of free parameters as $\\Lambda$CDM. The first one exhibits  scaling behaviour during radiation and a strong phantom phase today, ending in a \"big-freeze\" singularity. This model provides the best fit to date for the SNIa Gold dataset. The second theory we consider is standard electromagnetism. We show that a temporal electromagnetic field on cosmological scales generates an effective cosmological constant and that primordial electromagnetic quantum fluctuations produced during electroweak scale inflation could naturally explain, not only the presence of this field, but also the measured value of the dark energy density. The theory is compatible with all the local gravity tests, and is free from classical or quantum instabilities. Thus,  not only the true nature of dark energy could be established without resorting to new physics, but also the value of the cosmological constant would find a natural explanation in the context of standard inflationary cosmology. ",
        "introduction": "The fact that today matter and dark energy have comparable contributions to the energy density, $\\rho_\\Lambda\\sim\\rho_M\\sim (2 \\times 10^{-3}$ eV)$^4$ in natural units, poses one of the most important problems for models of dark energy. Thus, if dark energy is a cosmological constant, its energy density would remain  constant throughout the history of the universe, whereas those of the rest of components (matter and radiation) grow as we go back in time. Then the question arises as to whether it is a {\\it coincidence} (or not) that they have comparable values today when they have differed by many orders of magnitude in  the past. Notice also that if $\\Lambda$ is a fundamental constant of nature, its scale (around $10^{-3}$ eV) is more than 30 orders of magnitude smaller than the natural scale of  gravitation, $G=M_P^{-2}$ with $M_P\\sim 10^{19}$ GeV. On the other hand, alternative models in which dark energy is a dynamical component rather than a cosmological constant also require the introduction of unnatural scales in their Lagrangians or initial conditions in order to account for the present phase of accelerated expansion. Such models are usually based on new physics, either in the form of new  cosmological fields or modifications of Einstein's gravity \\cite{quintessence1,quintessence2,quintessence3,quintessence4, quintessence5} and they are generically plagued by additional problems such as classical or quantum instabilities, or inconsistencies with local gravity constraints.   Therefore, we would like to find a description for dark energy without dimensional scales (apart from Newton's constant $G$), with the same number of free parameters as $\\Lambda$CDM, with natural initial conditions, with good fits to observations and no consistency problems. In this work we will show that vector theories can do the job. With that purpose, we present two examples of such theories which have been recently proposed \\cite{BM,BMEM} (for other vector models see references in those works). ",
        "conclusions": "We have shown that  vector theories offer a simple and accurate description of dark energy in which the coincidence problem could be easily avoided. In our first example, the scaling behaviour during radiation and the natural initial conditions for the vector field offer a neat way around the problem. Moreover, in our second example, the presence of a cosmological electromagnetic field generated during inflation provides a natural explanation for the cosmic acceleration. This result not only offers a solution to the problem of establishing the true nature of dark energy, but also explains the value of the cosmological constant without resorting to new physics. In this scenario the fact that matter and dark energy densities {\\it coincide} today is just a consequence of inflation taking place at the electroweak scale. Present and forthcoming astrophysical and cosmological observations will be able to discriminate these proposals from the standard $\\Lambda$CDM cosmology.  \\vspace{0.5cm}  {\\em Acknowledgments:} This work has been  supported by DGICYT (Spain) project numbers FPA 2004-02602 and FPA 2005-02327, UCM-Santander PR34/07-15875 and by CAM/UCM 910309. J.B. aknowledges support from MEC grant BES-2006-12059. \\vspace{0.2cm}"
    },
    "0812/0812.1638_arXiv.txt": {
        "abstract": "{} {The nearby galaxy clusters Abell~496 and Abell~85 are studied in the very high energy (VHE, E$>$ 100 GeV) band to investigate VHE cosmic rays (CRs) in this class of objects which are the largest gravitationally bound systems in the Universe.} {H.E.S.S., an array of four Imaging Atmospheric Cherenkov Telescopes (IACT), is used to observe the targets in the range of VHE gamma rays.} {No significant gamma-ray signal is found at the respective position of the two clusters with several different source size assumptions for each target. In particular, emission regions corresponding to the high density core, to the extension of the entire X-ray emission in these clusters, and to the very extended region where the accretion shock is expected, are investigated. Upper limits are derived for the gamma-ray flux at energies E$>570$ GeV for Abell~496 and E$>460$ GeV for Abell~85.} {From the non-detection in VHE gamma rays, upper limits on the total energy of hadronic CRs in the clusters are calculated. If the cosmic-ray energy density follows the large scale gas density profile, the limit on the fraction of energy in these non-thermal particles with respect to the total thermal energy of the intra-cluster medium (ICM) is 51\\% for Abell~496 and only 8\\% for Abell~85 due to its larger mass and higher gas density. These upper limits are compared with theoretical estimates. They predict about $\\sim$10\\% of the thermal energy of the ICM in non-thermal particles. The observations presented here can constrain these predictions especially for the case of the Abell~85 cluster.} ",
        "introduction": "Clusters of galaxies, the most extended gravitationally bound systems in the Universe, are expected to contain a significant population of non-thermal particles. The mechanisms for producing cosmic rays (CR) in galaxy clusters can be divided into external processes and internal processes. In external processes, the particle acceleration is driven by the assembly of the cluster. An efficient production of high energy particles is expected especially at the strong accretion shock at the outskirt of the cluster, where cold in-falling material plunges into the already existing hot intra-cluster medium (ICM). Therefore, large-scale shock waves caused by the cosmological structure formation may populate these objects with a non-thermal component of particles (see e.g. Colafrancesco \\& Blasi \\cite{colafrancesco00}, Loeb \\& Waxman \\cite{loeb00}, Ryu et al. \\cite{ryu03}, Miniati \\cite{miniati03}). Particles can also be accelerated by turbulence in the intra-cluster medium(ICM) generated by major sub-cluster merger events (e.g. Brunetti et al. \\cite{brunetti04}). In contrast in the internal mechanisms, the CRs are accelerated by cluster galaxies and injected into the whole cluster volume afterwards. Internal sources of CRs in clusters can be supernova-driven galactic winds (V\\\"olk et al. \\cite{voelk96}) or AGNs (e.g. En\\ss lin et al. \\cite{ensslin97}, Aharonian \\cite{aharonian02}, Hinton et al. \\cite{hinton07}).  Accelerated, highly energetic CR particles can subsequently produce gamma rays through two main mechanisms in galaxy clusters. VHE gamma rays can be generated by hadronic processes, through inelastic collisions between high-energy protons and target protons provided by the thermal hot ICM and subsequent $\\pi^0$-decay (Dennison \\cite{dennison80}). In the leptonic scenario, very high energy (VHE, E~$>$~100 GeV) electrons up-scatter cosmic microwave background (CMB) photons to gamma-ray energies via the inverse Compton process (e.g. Atoyan \\& V\\\"olk \\cite{atoyan00}, Gabici \\& Blasi \\cite{gabici03}, \\cite{gabici04}). VHE electrons which up-scatter low energy photons to the gamma-ray range can also be of secondary origin either due to photon - photon interactions (Timokhin et al. \\cite{timokhin04}) or generated by ultra-high energy protons (E $>$ 10$^{18}$ eV) interacting with CMB photons in the Bethe-Heitler process (Inoue et al. \\cite{inoue05}).  Galaxy clusters are huge structures with a linear scale of several Mpc. This length scale together with typical cluster magnetic fields in the $\\mu$G range (Carilli \\& Taylor \\cite{carilli02}, Govoni \\& Feretti \\cite{feretti04}) leads to a large diffusion timescale. This diffusion time is longer than the Hubble time for CR protons with energies less than $\\sim$10$^{15}$ eV. Therefore no losses of such particles will occur in these systems (V\\\"olk et al. \\cite{voelk96}, Berezinsky et al. \\cite{berezinsky97}). Gamma rays produced through the hadronic scenario can thus be used to probe particle acceleration in clusters of galaxies over the entire formation history.  In contrast to hadronic CRs, electrons which are accelerated to VHE energies in clusters will suffer from substantial energy losses due to synchrotron and inverse Compton radiation. Therefore, an electron with an energy of $\\sim$1 TeV has only a limited lifetime of $\\sim$10$^6$ years in a typical cluster magnetic field of 1 $\\mu$G (e.g. Atoyan \\& V\\\"olk \\cite{atoyan00}). Consequently VHE electrons do not accumulate in clusters and it is expected that the electron spectrum will cut off at an energy that is below the VHE regime outside the acceleration sites. Hence leptonic VHE gamma-ray production is likely less relevant in galaxy clusters (Atoyan \\& V\\\"olk \\cite{atoyan00}).  Observationally, a significant population of non-thermal electrons is found in several clusters that can be detected in radio waves (Giovannini \\& Feretti \\cite{giovannini00}, Feretti et al. \\cite{feretti04}, Bagchi et al. \\cite{bagchi06}) and possibly hard X-rays (Rephaeli \\& Gruber \\cite{rephaeli02}, Fusco-Femiano et al \\cite{fusco04}). Given the evidence/hints of accelerated particles in galaxy clusters, these objects are potential sources for observation of gamma-rays (see Blasi et al. \\cite{blasi07} for a recent review). However, despite these indications, no cluster has been established as a gamma-ray emitter so far (Reimer et al. \\cite{reimer03}). In the VHE gamma-ray range Perkins et al. (\\cite{perkins06}) have reported upper limits for two nearby massive clusters (Perseus and Abell~2029) with the \\textit{Whipple} telescope. For the Perseus cluster, these results imply that the cosmic-ray proton energy density is less than 8\\% of the thermal energy density.  Throughout the paper a standard cosmology with $\\Omega_{\\Lambda}$~=~0.7, $\\Omega_{m}$~=~0.3 and $H_0$~=~70~km~s$^{-1}$~Mpc$^{-1}$ is used.  In this paper, the observations of the galaxy clusters Abell~496 and Abell~85 with the H.E.S.S. array of telescopes in the VHE gamma-ray regime are presented. ",
        "conclusions": "No significant point-like or extended gamma-ray flux F($>570$ GeV (Abell~496) and F($>460$ GeV) (Abell~85) has been found in the H.E.S.S. observations. Upper limits on the VHE gamma-ray flux from both clusters are derived. With the presented results, it is possible to constrain the energy fraction in hadronic CRs. For the case where the energy density of the CRs follows the density of the thermal gas and where the CRs have a spectral index of 2.1 this is not more than 51\\% (Abell~496) and 8\\% (Abell~85) of the thermal energy. Especially for the case of Abell~85 the upper limits obtained with H.E.S.S. can already constrain model predictions on $E_{\\rm nonth}/E_{\\rm th}$. These limits are the best determined so far without a dependency on the magnetic field. However, for a uniform energy density of CRs or a softer spectrum of the particles, the limit on the fraction of energy in hadronic CRs versus thermal energy of the ICM is not constraining. Our theoretical considerations also suggest that the most promising clusters for VHE gamma-ray observations are those with the largest mass (M$_{\\rm tot}$ $>$ 10$^{15}$ M$_{\\odot}$) and highest temperatures (T $>$ 7 keV) and, as far as supernova activity is concerned, those with the largest iron mass (M$_{\\rm Fe}$ $>$ 5$\\times$10$^{10}$ M$_{\\odot}$). Observations of this kind of galaxy cluster at a distance of $\\sim$100 Mpc with the present generation of IACTs to a flux of $F(>1\\ \\mathrm{TeV}) < 10^{-12}$ ph. s$^{-1}$ cm$^{-2}$ within a radius of 1$^{\\circ}$ will provide the ability to test models of the non-thermal hadronic component in galaxy clusters, which predict a considerable fraction ($\\sim$10\\%) of energy in CRs."
    },
    "0812/0812.4537_arXiv.txt": {
        "abstract": " ",
        "introduction": "Great progress has been made in understanding the the large scale structure of the universe in past decades. Not only observations of the major galaxy surveys such as  SDSS  \\cite{tegmark03,zehavi02,zehavi05}, 2dF \\cite{colless,hawkins,madgwick,percival}, APM \\cite{maddox,padilla}, and REFLEX \\cite{collins,schueker} etc, have revealed the cosmic structures of increasingly large sky dimension with new detailed features being found, theoretical studies have also achieved important results, through numerical simulations \\cite{White,springel}, perturbation method \\cite{peebles,davis,fry,goroff,bernardeau,valageas}, and thermodynamics \\cite{saslaw}. From  view point of dynamics, the Universe filled with galaxies and clusters is a many-body self-gravitating system in an asymptotic relaxed state, since the cosmic time scale $1/H_0$ is longer than the local crossing time scale \\cite{saslaw}. A systematic  approach to  statistical mechanics of many-body systems is to convert the degrees of freedom of discrete particles into a  continuous field. Thereby, the fully-fledged techniques of field theory can apply to study the systems  \\cite{zustin}. The Landau-Ginzburg theory is a known example in this regard. We have formulated such a density field theory of self-gravitating systems and applied it to the large scales structure of the Universe \\cite{zhang}. Under the Gaussian approximation, the field equation of the 2-point correlation function and the solution have been derived explicitly. The result qualitatively interprets some observational features, but it suffers from insufficient clustering on small scales. In this paper, we will go beyond the Gaussian level and include nonlinear terms of density fluctuations up to the second order, yielding a more satisfying description of the large scale structure. ",
        "conclusions": "Overall, the solution $\\xi(r)$ of Eq.(\\ref{eq3}) improves the Gaussian one considerably, as shown FIG.\\ref{landau-nonlinear}, in which the 2dFGRS active galaxies are taken for demonstration purpose. It is found that the nonlinear term $-b(\\xi') ^2$ has the effects of strongly enhancing $\\xi(r)$ on small scales, and making  $P(k)$ flatter in $k<k_0$. The friction term $a\\xi'$ has the following effects: slightly increasing the height of $\\xi(r)$ on small scales; moving the zeros of $\\xi(r)$ to larger $r$; strongly damping the amplitude of oscillations of $\\xi$ on large scales \\cite{broadhurst,Tucker,Einasto1997a,Einasto1997b,Einasto2002,Tago}, as seen in FIG.\\ref{oscillation}; smoothing out the sharp peak of $P(k)$ at $k_0$ and turning $P(k)$ positive for $k<k_0$.  The main conclusion of this paper is the following. The universe containing galaxies or clusters is viewed as a many-particle system in asymptotic relaxed state and can be described by an effective density field, whereby the techniques of field theory applies, yielding a perspective on the large scale structure of the universe other than the conventional methods. There appears  the Jeans scale $k_0\\simeq (0.04\\sim 0.06)$ hMpc$^{-1}$, which is the unique scale underlying the large scale structure as a gravitational system. Up to nonlinear terms $(\\delta \\psi)^2$ beyond the Gaussian approximation, the nonlinear field equation of the 2-pt correlation function of the density fluctuations has been derived and solved analytically. This analytic result of field theory interprets several observed features of large scale structures. With fixed values of the parameters $(a,b, k_0)$, the solution matches the observed $\\xi(r)$ and $P(k)$ of both galaxies and of clusters, simultaneously.     Although our results match  the observational data on large scales very well, our model is still preliminary  at the present stage, there are several problems that need to be further addressed in the future as in the following.     Firstly the calculated $\\xi_{gg}(r)$ for galaxies increases too fast on very small scales $r\\leq 0.2 h^{-1}$Mpc. This indicates that the model may break down on such small scales close to a galaxy size. This may suggest that either higher order terms of perturbations should be taken into account, or the effects of galaxy formation and local virilization need to be included.  Next,  there is a limitation in applying the Groth-Peebles ansatz in Eq.(\\ref{groth-peebles}). As is known, for descriptions of any many-particle dynamic system based on Gibbs-Boltzmann equation, a procedure is usually taken, which decomposes the complicated equation into a set of differential equations, so that each one in the set is possibly manageable. However, thereby, a hierarchy of BBKGY type, or the like, is inevitably arises. Different treatments of the hierarchy are employed for different systems, and a cut-off of the hierarchy is usually is used. For instance, in the case of the photon gas of CMB, the multipole decomposition is involved for the temperature anisotropies and polarization, and the common practice is to cut off the hierarchy of multipoles by letting higher multipole components to zero, yielding a very accurate description of CMB spectra \\cite{ZhaoZhang,Zhang et al,XiaZhang}). But for the calculation of $G^{(2)}_c({\\bf r})$ in our context, one can not set $G^{(3)}_c({\\bf r}_1,{\\bf r}_2,{\\bf r}_3)$  to zero, since these are important and give rise to nonlinear effects, due to the  long range nature of gravity force. The Groth-Peebles ansatz has been used in our analytic treatment, because it is a simple one and  an analytic solution can be derived. We would like to mention that the ansatz only approximately reflects the actual distribution since the forms of factor, $Q$,  is in fact a function of $r$. In this case Eq.(\\ref{eq2}) for $G^{(2)}_c({\\bf r})$ would be more complicated and  analytic solutions for the general case would be difficult to derive explicitly.   Thirdly, in fitting the observation of galaxies, we have not separated the dark matter from galaxies in our present model. Therefore, no bias is introduced and the baryon acoustic oscillations are not incorporated in our model. A comprehensive treatment of two components, dark matter and baryons, in our theory would require substantial extensions to the model discussed the present paper.    Finally, it should be mentioned that our theoretical model deals with only the quasi-relaxed state of the large scale structure of the universe, which, as an assumption, is qualitatively good approximation since the overall expansion rate $H_0$ is smaller than the particle collision rate. It would be much desired to have an extension of the present model to take into account of an evolutionary description. These issues  will to be addressed in our further studies.          ACKNOWLEDGMENT: Y. Zhang's research work was supported by the CNSF No.10773009, SRFDP, and CAS.    \\baselineskip=12truept"
    },
    "0812/0812.3647_arXiv.txt": {
        "abstract": "Azimuthal color (age) gradients across spiral arms are one of the main predictions of density wave theory; gradients are the result of star formation triggering by the spiral waves. In a sample of 13 spiral galaxies of types A and AB, we find that 10 of them present regions that match the theoretical predictions. By comparing the observed gradients with stellar population synthesis models, the pattern speed and the location of major resonances have been determined. The resonance positions inferred from this analysis indicate that 9 of the objects have spiral arms that extend to the outer Lindblad resonance (OLR); for one of the galaxies, the spiral arms reach the corotation radius. The effects of dust, and of stellar densities, velocities, and metallicities on the color gradients are also discussed. ",
        "introduction": "Density wave phenomena have been proposed to explain the spiral structure seen in disk galaxies \\citep{lind63,lin64,too77,ber89a,ber89b}. Observationally, these phenomena are best studied in the near infrared (near-IR), especially the $K$-band, that mostly traces the old stars in the disk \\citep[e.g.,][]{rix93}. The old stellar disk, however, has a disordered optical counterpart of young stars, gas, and dust \\citep{zwi55,blo91}, so the question arises whether the spiral structure, seen in near-IR bands, and the star formation, seen in the optical bands, are coupled or not. If disk dynamics and star formation are indeed related, the star formation rate per unit gas mass should be affected by the presence of the density wave. Unfortunately, H$_2$ cannot be directly observed, and the relation between detected CO and total H$_2$ mass is a whole controversial topic in itself \\citep[e.g.,][ and references therein]{allen96}.  The alternative approach of comparing star formation rates, past (as traced by optical and near-IR surface photometry) and present (as probed by $H\\alpha$ emission), in galaxies with different Hubble types has been carried out; as a result, both a positive correlation between disk dynamics and star formation \\citep[e.g.,][]{sei02}, and the absence of such a correlation \\citep[e.g.,][]{ryd94,elm86} have been claimed.  In this work, we will focus on the relation between star formation and density wave phenomena, as proposed in the large scale shock scenario \\citep{rob69,shu72}. Evidence has been gathered, from observations of dust, gas compression \\citep[e.g.,][]{mat72,vis80}, and molecular clouds \\citep{vog88,sch04} near the concave regions of spiral arms, that suggests that star formation is triggered there. Through a simple interpretation of these ideas and observations, the existence of azimuthal color-age gradients has been predicted. As seen in Fig.~\\ref{colorgrads}, if we assume that the spiral pattern rotates with constant angular speed, and that the gas and stars have differential rotation, a corotation region exists where these two angular speeds are equal. At smaller radii, bursts of star formation take place where the differentially rotating gas overtakes the spiral wave, and show up as brillant HII regions \\citep[e.g.,][]{mor52,elm83a}. As young stars age, they drift away from their birth site, thus creating a color gradient. Star formation can occur also beyond the corotation radius, when the spiral pattern overtakes the gas. The assumption that the pattern rotates with constant angular speed can be corroborated with numerical simulations \\citep{tho90,don94,zha98}, but the only observational evidence of this premise is the apparent persistence of the spiral structure for up to a Hubble time in nearby galaxies \\citep{elm83b}, in avoidance of the winding dilemma.  Another way to test the assumption of constant angular speed of the spiral pattern is by the dynamical consequences it has on the disk environment and material. In this regard, the prediction of azimuthal color gradients can be seen as another test of the spiral density wave theory. So far, the search for these color gradients has in general yielded inconclusive results \\citep{sch76,tal79,cep90,hod90,delRio98}, and a few exceptional positive cases like \\citet{efr85} for M31, \\citet{sit89} for the Milky Way, and \\citet{gon96} for the spiral galaxy M~99.  \\begin{figure} \\epsscale{1.00} \\plotone{f1.eps} \\caption{Stellar age gradients across the spiral arms are indicated by arrows that go from blue to red. The azimuthal age gradients are produced by stars born in the spiral shock, where the shocked interstellar medium forms a dust lane, that later drift away as they age. The direction of the gradients changes at the corotation radius, $R_{\\rm CR}$. Inside this radius, the disk material overtakes the spiral wave, and beyond it the spiral wave catches up with the material [see also figure 1 in \\citet{pue97}]. \\label{colorgrads}} \\end{figure}  \\begin{figure} \\epsscale{1.00} \\plotone{f2.eps} \\caption{Theoretical $Q(rJgi)$ vs.\\ time, CB07 models. The duration of the burst is $2 \\times 10^7$ years, with a Salpeter IMF; the fraction of young stars ranges from 1\\% to 5\\% by mass. Lower and upper mass limits are 0.1 and 10 $M_{\\sun}$, respectively.   \\label{Qpercent}} \\end{figure}  \\begin{figure} \\epsscale{1.00} \\plotone{f3.eps} \\caption{ Model Johnson filter $Q(RJVI)$ vs.\\ time $(t)$ for CB07 models, reddened as per the ``starburst galaxy\" model of \\citet{wit92}. The duration of the burst is $2 \\times 10^7$ years, with a Salpeter IMF and 2\\% by mass of young stars. Lower and upper mass limits are 0.1 and 100 $M_{\\sun}$, respectively. \\label{Qdust}} \\end{figure}  \\citet[][ GG96 hereafter]{gon96} used for the first time a supergiant sensitive and reddening-free photometric index\\footnote{ See Binney \\& Merrifield (1998), for other reddening-free indices.} to trace star formation, defined as:  \\begin{equation} Q(rJgi) = (r-J) - \\frac{E(r-J)}{E(g-i)}(g-i). \\end{equation}  \\noindent Using the extinction curves of \\citet{sch83}, and \\citet{rie85} for a foreground screen, the color excess term is $E(r-J)/E(g-i) = 0.82$. According to population synthesis models, following a star formation burst the $Q(rJgi)$ photometric index increases its value for $\\sim 2.6 \\times 10^7$ years, and then starts to decline, as shown in Fig.~\\ref{Qpercent} for most recent Charlot \\& Bruzual (2007, in preparation, CB07 hereafter) models with different stellar mixtures in which young stars constitute from 1\\% to 5\\% by mass. Dust lanes arising from spiral shocks are expected to precede the azimuthal color gradients, hence the importance of reddening-free diagnostics. With the stellar population synthesis (SPS) models of \\citet{bru93}, and the radiative transfer models of \\citet{bru88} and \\citet{wit92}, GG96 investigated the behavior of $Q(RJVI)$ for mixtures of dust and stars with different relative spatial distributions. The main effect of a mixture of dust and stars on the $Q$ index, illustrated in Fig.~\\ref{Qdust} for the CB07 population synthesis models, is to increase its value relative to the $\\tau = 0$ and foreground dust screen cases, and $Q(RJVI)$ is no longer reddening-free when $\\tau_V > 2$.\\footnote{The same conclusion can be applied to $Q(rJgi)$, since the optical bands $R$, $V$, and $I$ have approximately the same effective wavelengths, respectively, as the $r$, $g$, and $i$ filters; the main difference between both sets is that the latter passbands are narrower.} However, the requirement that $\\tau_V < 2$ is likely fulfilled by the disks of nearly face-on galaxies \\citep{pel95,kuc98,xil99}; moreover, as stars age, dust dissipates, and reddening diminishes \\citep[e.g.,][]{char00}.  Based on their study of M~99, GG96 propose that real data are best matched by SPS models with 0.5\\% to 2\\%, by mass, of young stars. These values are in accordance with the amount of $B$ light in the arms contributed by young stars, as estimated by \\citet{sch76}. GG96 also hypothesize there is an inverse correlation between high-mass star-forming regions and detectable azimuthal color gradients. The color gradient in M~99 lies where no HII regions are identified in $H\\alpha$ images of the galaxy. This counterintuitive finding could help explain the dearth of positive detections to date, if contamination from bright emission lines produced in HII regions around the most massive stars masks the color gradients associated with star formation in spiral shocks \\citep{shu97}. On the other hand, evidence has been found recently of star forming regions with very few massive stars that generate only scant $H\\alpha$ emission \\citep{inde08}. In this investigation we apply the GG96 method, using the $Q(rJgi)$ photometric index, to a new sample of galaxies in order to search for and analize color-age gradients near spiral arms.  \\subsection{Color gradients: the link between star formation and spiral dynamics.} \\label{linkdyn}  According to various studies, spiral density waves must propagate between orbital resonances \\citep{lin70,mar76,lin79,too81,con86}. Of these resonances, the most important ones are the inner Lindblad resonance (ILR), the 4:1 resonance, corotation (CR), and the outer Lindblad resonance (OLR).\\footnote{ At the ILR (OLR), the epicyclic frequency $\\kappa = \\pm 2(\\Omega - \\Omega_p)$; $\\kappa = 4(\\Omega - \\Omega_p)$ at the 4:1 resonance. }  \\noindent Azimuthal color-age gradients retain some information about the stellar drift relative to the spiral shock, allowing us to obtain $\\Omega_p$, the angular velocity of the spiral pattern. In order to find $\\Omega_p$ from the gradient information, we can use:  \\begin{eqnarray} \\label{eqOMEGA_I} \\Omega_{p} = \\frac{1}{t} \\left(\\int_{0}^{t} \\frac{\\vec v(t') \\cdot \\hat{\\varphi}(t')}{R(t')} \\mathrm{d}t' - \\left(\\theta_{\\rm shock} + \\Delta\\theta \\right)\\right), \\\\ \\Delta\\theta = \\cot(-i) \\ln \\left(\\frac{R(t)}{R(0)}\\right), \\end{eqnarray}  \\noindent where $t$ is the age of the young stellar population at an angle $\\theta_{\\rm shock}$ away from the shock position ($t'$ is a variable of integration), $\\vec v(t')$ is the velocity vector of the young stellar population in an inertial reference system; $\\hat{\\varphi}(t')$ is the unit vector in plane polar coordinates $\\rho$, $\\varphi$, in a non inertial reference system; $R(t')$ is the orbital radius of the studied region, measured from the center of the disk to the center of mass of the young stellar population;\\footnote{$R(t)$ corresponds to $\\theta_{\\rm shock}$.} and $i$ is the arm pitch angle,\\footnote{The angle between a tangent to the spiral arm at a certain point and a circle, whose center coincides with the galaxy's, crossing the same point.} if the azimuthal angle increases in the direction of rotation and the spiral arms trail. The angular quantity $\\Delta\\theta$ accounts for the logarithmic spiral shape of the shock. Assuming the departures from circular motion are small, we have:  \\begin{equation} \\Omega_{p} = \\frac{1}{R_{mean}} \\left(\\frac{\\int_{0}^{t} v_{rot}(t') \\mathrm{d}t' }{t} - \\frac{d}{t} \\right), \\end{equation}  \\noindent where $R_{mean}$ is the mean orbital radius of the studied region, $v_{rot}$ is the circular orbital velocity in an inertial reference system, and $t$ is the age of the young stellar population at the azimuthal distance $d$ from the shock. Expanding the term $\\int_{0}^{t} v_{rot}(t') \\mathrm{d}t'$, in a Taylor series, we get :  \\begin{equation} \\Omega_{p} = \\frac{1}{R_{mean}} \\left(\\frac{(v_{rot})|_{0} t + \\frac{1}{2}  (\\frac{\\mathrm{d} v_{rot}}{\\mathrm{d}t'})|_{0} t^{2} + \\zeta}{t} - \\frac{d}{t} \\right), \\label{eqOmega_Ib} \\end{equation}  \\noindent where $\\zeta$ represents higher order terms in the expansion. According to theory \\citep[e.g.,][]{rob69,sly03}, the higher order terms may account for 20-30 km s$^{-1}$, a quantity that is of the same order as the mean rotation velocity error, after considering the uncertainty due to galaxy inclination (see Table~\\ref{tbl-param}). Hence, we can neglect the higher order terms in eq.~\\ref{eqOmega_Ib}, such that $v_{rot} \\sim$ constant, and obtain  \\begin{equation} \\label{eqOMEGA_II} \\Omega_{p} \\cong \\frac{1}{R_{mean}} \\left(v_{rot} - \\frac{d}{t} \\right). \\end{equation}  \\noindent $\\Omega_p$ is found by stretching the model (which gives $Q$ as a function of $t$) to fit the data (where $Q$ is a function of $d$). ",
        "conclusions": "If star formation is related to disk dynamics then, according to theory, the $R_{res} / R_{end}$ ratio must be close to one, for either the 4:1 resonance, the corotation radius, or the OLR. \\citet{con86}, based on orbital calculations, determined that the spiral pattern of strong spirals must extend to the 4:1 resonance. \\citet{pats91} defined strong spirals as those with large pitch angles (i.e., Sb and Sc galaxies, that are not tightly wound); their theory is considered ``non linear\". Conversely, the ``linear\" theory of spiral density waves concerns itself with tightly wound spirals and has concluded that the arms of some normal (nonbarred) spirals reach corotation \\citep{lin70} and are stationary, while those of others grow to the OLR \\citep{mar76,lin79,too81}. In the \\citet{con86} treatment, the ``linear\" theory is recovered when spiral arms are not strong (as in Sa galaxies).  In Fig.~\\ref{spiralEP}, we can see that most spirals in our sample extend to the OLR and one of them (NGC~578) reaches corotation. The mean of the ratio $R_{OLR}/R_{end}$, for all 13 points, is 0.95$\\pm$0.03. The reduced $\\chi^{2}/n$ of this result is 7.12;\\footnote{ For an expected value of $R_{OLR}/R_{end}=1$.} the probability of this result being due to chance for 13 degrees of freedom is less than 1 out of 10,000. If we remove the points corresponding to NGC~3938, NGC~7126, NGC~6951, and NGC~578, $R_{OLR}/R_{end} = 0.98\\pm0.04$, with a reduced $\\chi^{2}/n$ = 0.48, for $n=9$. This last result may indicate that some of our errors are overestimated. It is interesting to note that none of our objects is an Sa (weak spiral) galaxy. From this we may conclude that the ``linear\" result for the extent of spiral patterns applies to strong spirals too!  On the other hand, for most objects the spiral begins at the location of the ILR, as expected. For the analysis presented here, though, the location of this resonance should be taken with care, since we are employing flat rotation curves, even at small radii. When using real rotation curves, the positions of the ILRs for our sample may vary. The possibility also exists that there is no ILR in some objects.  In some regions, the downstream values of the $Q$ index are lower than the models. This is the case for NGC~4939~A, NGC~3162~B, NGC~1421~A, NGC~1421~C, NGC~7125~A, NGC~7125~C, NGC~578~A, and NGC~578~B. For the present analysis we have assumed that every star-forming region maintains its previous circular motion after the spiral shock. As already stated, in real galaxies the situation is different: variations in the velocity vector are present due to the shock itself and the resulting loss of angular momentum \\citep{yua81,fer08}. This effect could explain the ``downstream fall\" of the gradients. Metallicity is another factor that can play a significant part: a lower metallicity of the older stars may produce a steep drop of the gradients (see \\S~\\ref{rolmeta}).  We have shown that azimuthal color gradients are common in spiral arms of disk galaxies. Even inverse color gradients have been found in the arms of NGC~1421, NGC~4939, and NGC~7125. The observed picture is, of course, not as clean as originally envisioned, for several reasons. Dust and HII regions may mask the gradients (GG96); indeed, all the detected gradients have been satisfactorily fitted with models where $M_{upper}$ = 10 $M_\\odot$. Other star formation mechanisms, like self-propagating star formation, may take place simultaneously with density wave triggering in spiral arms and disks in general. Likewise, there may be substructure formation, as a result of non-linear and chaotic effects associated to the wave phenomenon itself \\citep{cha03,dob06,kim06,she06}. In the future, three aspects can be improved when carrying out this type of study. (1) A better determination of the inclination angle of the galaxy, and of the rotation curve at different radii; this will reduce the uncertainty in $\\Omega_p$. (2) The inclusion in the models of the variable density distributions of young and old stars, as well as of the stellar velocity changes near the spiral shocks. This will account for non-circular motions near the spiral shocks and for the higher order terms in equation \\ref{eqOMEGA_I}. Numerical simulations or semi-analytical treatments may be needed, particularly for the young stars; also, the contribution of red supergiants to the observed $K$-band surface brightness should be considered when modeling the distribution of old stars. (3) Spectroscopic studies could be very helpful as diagnostics of the population properties, particularly of the metallicity along the gradients.  According to the dynamic parameters derived from our analysis, spiral arms mostly extend to the OLR and sometimes to the corotation radius. Ten of thirteen objects (77\\%) support this conclusion. Similar results have been obtained by other authors as well \\citep[e.g.,][]{elm92,zha07}. In view of the consistency between the pattern speeds and resonance positions determined from the azimuthal stellar gradients, and the predictions of density wave theory, we conclude that disk dynamics do play an important role in large scale star formation in some spiral galaxies."
    },
    "0812/0812.3838_arXiv.txt": {
        "abstract": "{Young populations at Z$<$Z$_{\\odot}$ are being examined to understand the role of metallicity in the first phases of stellar evolution. For the analysis it is necessary to assign mass and age to Pre--Main Sequence (PMS) stars. While it is well known that the mass and age determination of PMS stars is strongly affected by the convection treatment, extending any calibration to metallicities different from solar one is very artificial, in the absence of any calibrators for the convective parameters. For solar abundance, Mixing Lenght Theory models have been calibrated by using the results of 2D radiative-hydrodynamical models (MLT-$\\alpha^{2D}$), that result to be very similar to those computed with non-grey ATLAS9 atmosphere boundary condition and full spectrum of turbolence (FST) convection model both in the atmosphere and in the interior (NEMO--FST models).} {While MLT-$\\alpha^{2D}$  models are not available for lower metallicities, we  extend to lower Z the NEMO--FST models, in the educated guess that in such a way we are simulating also at smaller Z the results of MLT-$\\alpha^{2D}$ models.}  % {We use standard stellar computation techniques, in which the atmospheric boundary conditions are derived making use of model atmosphere grids. This allows to take into account the non greyness of the atmosphere, but adds a new parameter to the stellar structure uncertainty, namely the efficiency of convection in the atmospheric structure, if convection is computed in the atmospheric grid by a model different from the model adopted for the interior integration.} {We present PMS models for low mass stars from 0.1 to 1.5 M$_{\\odot}$ for metallicities [Fe/H]= -0.5, -1.0 and -2.0. The calculations include the most recent interior physics and the latest generation of non-grey atmosphere models. At fixed luminosity more metal poor  isochrones are hotter than solar ones  by $\\Delta$ $\\log$ T$_{\\rm eff}$/$\\Delta$ $\\log$Z $\\sim$ 0.03-0.05 in the range in Z from 0.02 to 0.0002 and for ages from 10$^5$ to 10$^7$ yr.} {} ",
        "introduction": "The study of Pre-Main Sequence (PMS) stars is important to trace the modalities of star formation in space and time, to date young stellar systems by means of age tracers which do not suffer the  uncertainty in the physics of the upper main sequence stars, to derive the initial mass function of very low mass stars and brown dwarfs, and to understand the modalities  of the stellar rotational evolution and depletion of light elements. Theoretical tracks and isochrones provide an essential tool to understand and interpret the experimental data currently available for young objects. Since most of these sources are located in nearby Galactic star forming regions, this research has been generally limited so far to solar chemistry or close to it \\citep[e.g][]{siess2000}. Now that young populations at Z$<$Z$_{\\odot}$ are being examined  \\citep[e.g.][]{romaniello2006}, it becomes essential to understand the role of metallicity in the first phases of stellar evolution, thus demanding further investigations at lower metallicities.\\\\ On the other hand, the description of the PMS evolution of stars is one of the most delicate tasks in the context of stellar astrophysics: the location of the theoretical tracks on the Hertzprung-Russell (HR) plane depends critically on the ingredients used to calculate the models, the equation of state and  opacities \\citep[e.g.][]{dantona1993}, the boundary conditions used to match the integration of the interior with the atmospheric structure \\citep{baraffe2002}, and the treatment of convection \\citep{DM1994}. Detailed investigations by \\citet{montalban2004} showed that, at $T_{\\rm eff}$s where atmospheric convection is present, its modelling plays an important role in determining the radius of these structures, and thus the exact excursion of the theoretical track on the HR diagram; this is the most critical parameter in the computation down to T$_{\\rm eff}$$\\simeq$4000K. Below this temperature, both molecular opacities and convection treatment are the dominant uncertainties. Unfortunately, convection is a rather complex phenomenon, that is scarcely known from first principles, so that it is commonly treated by means of purely local approaches, the most popular of which is the Mixing Length Theory \\citep[MLT]{bohm1958}, where all the physical uncertainties are hidden below the unique parameter $\\alpha=l/H_p$, being $l$ the mixing scale and $H_p$ the pressure scale height; among other attempts to model convection locally, the Full Spectrum of Turbulence (FST) model by \\citet[CM]{CM1991} and \\citet[CGM]{CGM1996} has been largely used in recent years.\\\\ A potentially powerful tool to ``calibrate\" the convective model for the description of these evolutionary stages is the comparison between the theoretical loci in the HR plane and the position of PMS binaries, for which a rather precise estimation of the masses of the two components is possible \\citep[e.g][]{stassun2004,boden2005}. The results obtained so far are rather ambiguous, and evidentiate the difficulty to find out a unique description of convection, holding in all cases. \\citet{baraffe2002} found that several masses in PMS binaries demand an MLT parameter $\\alpha=1.9$, but some others, lying in the same region of the HR diagram, require a much lower efficiency, e. g. $\\alpha=1$. \\citet{stassun2004} also found that models with inefficient convection in the interior should be preferred. In a case, for the binary components of V1174Ori, they found a very good agreement with the models by \\citet{montalban2004}, employing the FST convection both in the interior and in the atmospheres computed by \\citet{NEMO2}. 2D and 3D radiative hydrodynamical simulations should provide more realistic results \\citep{tramp1999,ludwig2002}. Based on their 2D computations, \\citet{ludwig1999} provided a calibration of the MLT parameter $\\alpha$ for various effective temperatures and gravities: the meaning of this calibration is that the parameter describes the average efficiency of convection within the whole superadiabatic zone, in the region of T$_{\\rm eff}$  and gravity explored by the 2D models. This efficiency is  not constant, but varies with the position of the star in the HR diagram. \\citet{montalban2006} adopted such a calibration and computed the corresponding tracks (MLT-$\\alpha^{2D}$ tracks) for solar metallicity. Exam of binary location with respect to the MLT-$\\alpha^{2D}$ tracks  shows the same difficulties of previous track sets, although they result in excellent agreement with the components of the binary V773 Tau \\citep{boden2007}. The MLT-$\\alpha^{2D}$ tracks are not in agreement with the lithium depletion patterns in young clusters with the solar system abundance, but this additional feature can be attributed to the high metal abundance generally adopted for the solar model \\citep{montalban2006}. The same work shows that the FST non-grey tracks of solar composition, with boundary conditions based on model atmospheres in which convection adopts the same FST model \\citep{NEMO2}, present a striking similarity with the results of MLT-$\\alpha^{2D}$ models.\\\\ The scope of this paper is to extend the computations by \\citet{montalban2006} to lower metallicities, by naively assuming that FST models can approximate the results of  MLT-$\\alpha^{2D}$ simulations also at lower metallicity. As we wish to provide a more extended set of results, while the FST model atmospheres are available only at T$_{\\rm eff}$ $\\ge$ 4000K, we adopt a way to extend the computation to lower T$_{\\rm eff}$. We then  use MLT non grey models based on the atmospheric structures by \\citet[AH97]{AH1997} and choose the combination of atmospheric and interior convection efficiency that provides results similar to the FST tracks at T$_{\\rm eff}$ $\\ge$ 4000K. ",
        "conclusions": "A grid of stellar evolutionary tracks for low metallicity Pre--Main Sequence stars with masses between 0.1 and 1.5 M$_{\\odot}$ was presented. These models are based on up-to-date physics and updated non grey atmosphere models were used. A coherent treatment of convection in the interior and exterior region of the star was employed at  T$_{\\rm eff}$$\\ge$4000K. We extended our computations to models of smaller masses by using the AH97 grid of model atmospheres. The parameters $\\tau_{\\rm ph}$ and $\\alpha_{\\rm int}$ were  chosen to provide a smooth transition between the two model sets. This, of course, is only an educated guess to the problem of Pre--Main Sequence models at low metallicity. \\\\ The models (available in electronic form at the WEB location \\texttt{http://www.mporzio.astro.it/$\\%$7Etsa/}) can now be confronted to the complex realm of very young objects, providing important information on ages and star formation processes, and, on the other hand, providing some new constraints for PMS models."
    },
    "0812/0812.4303_arXiv.txt": {
        "abstract": "Thick layers of warm, low density ionized hydrogen (i.e., the warm ionized medium or WIM) in spiral galaxies provide direct evidence for an interaction between the disk and halo.  The wide-spread ionization implies that a significant fraction of the Lyman continuum photons from O~stars, produced primarily in isolated star forming regions near the midplane and often surrounded by opaque clouds of neutral hydrogen, is somehow able to propagate large distances through the disk and into the halo. Moreover, even though O~stars are the source of the ionization, the temperature and ionization state of the WIM differ significantly from what is observed in the classical O~star \\hii\\ regions. Therefore, the existence of the WIM and observations of its properties provide information about the structure of the interstellar medium and the transport of energy away from the midplane as well as place significant constraints on models. ",
        "introduction": "Forty-five years ago, Hoyle \\& Ellis (\\cite{HE63}) proposed the existence of an extensive layer of warm (10$^4$~K), low density (10$^{-1}$~cm$^{-3}$)  ionized hydrogen surrounding the plane of our Galaxy and having a power requirement comparable to the ionizing luminosity of the Galaxy's O and B stars. Their conclusion was based upon their discovery of a free-free absorption signature in the observations of the Galactic synchrotron background at frequencies below 10 MHz. However, the idea that a significant fraction of the Lyman continuum photons from the Galaxy's O~stars could travel hundreds of parsecs throughout the disk conflicted with the traditional picture in which the interstellar neutral hydrogen confined the ionizing radiation to small volumes (i.e., ``classical \\hii\\ regions'') near the hot stars.  Nevertheless, a few years later, the dispersion of radio signals from newly discovered pulsars (Hewish {\\em et al.\\/} \\cite{HBP68}) plus the detection of faint optical emission lines from the diffuse interstellar medium (Reynolds \\cite{Reynolds71}) firmly established warm ionized hydrogen as a major, wide-spread component of our Galaxy's interstellar medium.  But it was an additional two decades before deep \\ha\\ imaging with CCDs began to reveal similar warm plasmas permeating the disks and halos of other galaxies (Rand {\\em et al.\\/} \\cite{RKH90}; Dettmar \\cite{Dettmar90}), establishing the WIM as not just a peculiarity of the Milky Way, but a common property of galaxy disks. ",
        "conclusions": "The thick layers of warm, low density H$^+$ in spiral galaxies contain clues about the structure of the interstellar medium and the transport of energy away from regions of star formation at the midplane.  Some progress has been made in understanding this gas, but many fundamental questions remain:  How does the ionizing radiation propagate through the disk? Why is the H$^+$ layer so thick?  What is the distribution of the WIM within the interstellar medium? What is its relationship to the other phases of the medium? What is the source of the elevated temperatures in the WIM? Comparisons of the observed distribution, motions, and line ratios of the WIM with predictions of models have begun to address some of these questions (e.g., Hill {\\em et al.\\/} \\cite{Hill08}; Wood \\& Mathis \\cite{WM04}). Perhaps a next step is to compare the observations to more sophisticated multi-phase, dynamical models that included sources of heating and ionization with 3D radiation transfer.  The observations would constrain model parameters, and the models could predict properties of the gas to guide future observations.\\\\  \\thanks{RJR, LMH and AH acknowledge support from the National Science Foundation (NSF) through AST06-07512.  GJM is supported by the University of Sydney Postdoctoral Fellowship Program and the NSF grant AST04-01416.}"
    },
    "0812/0812.1183_arXiv.txt": {
        "abstract": "We present results of a semi-analytic model (SAM) that includes cold accretion and a porosity-based prescription for star formation. We can recover the puzzling observational results of low $V/\\sigma$ seen in various massive disk or disk-like galaxies, if we allow 18 $\\%$ of the accretion energy from cold flows to drive turbulence in gaseous disks at $z=2$. The increase of gas mass through cold flows is by itself not sufficient to increase the star formation rate sufficiently to recover the number density of $\\dot{M}_*>120$ M$_{\\odot}$ yr$^{-1}$ galaxies in our model.  In addition, it is necessary to increase the star formation efficiency. This can be achieved naturally in the porosity model, where star formation efficiency scales $\\propto \\sigma$, which scales as cloud velocity dispersion. As cold accretion is the main driver for gas velocity dispersion in our model, star formation efficiency parallels cold accretion rates, and allows  fast conversion into stars. At $z\\sim 2$, we find  a space density $~10^{-4}$ Mpc$^{-3}$ in star-forming galaxies with $\\dot{M}_*>120$ M$_{\\odot}$ yr$^{-1}$, in better agreement than earlier estimates form SAMs. However, the fundamental relation between $\\dot{M}_*$ and $M_*$ is still offset from the observed relation, indicating the need for possibly  more efficient star formation at high z perhaps associated with a role for AGN triggering. ",
        "introduction": "Multiple lines of evidence suggest that galaxy formation is taking place with a varying degree of efficiency over cosmic time. The overall star formation density in the universe declines at $ z\\le 2$ \\citep[e.g.][and references therein]{2004ApJ...615..209H}, a feature seen also in the normalization of the relation between stellar mass and star formation rate \\citep{2006ApJ...637..727C,2007ApJ...670..156D,2007ApJ...660L..43N}. Further support for the notion of rapid early formation of galaxies at $z \\geq 2$, especially massive ones, comes from high $\\alpha/$Fe ratios \\citep{2005ApJ...621..673T}, early metal enrichment \\citep{2008A&A...488..463M}, growth of super-massive black holes \\citep{2005A&A...441..417H} and the assembly rate of massive galaxies \\citep{2008ApJ...675..234P}. The redshift range $1.4 \\leq z \\leq 2.5$ has been targeted by various surveys combining broadband color selection techniques \\citep[e.g.][(D07)]{2004ApJ...617..746D,2004ApJ...604..534S,2007ApJ...670..156D} and recently by high-resolution H$\\alpha$ integral field spectroscopy \\citep{2006ApJ...645.1062F}. The latter reveals massive $ \\sim 10^{11}$M$_{\\odot}$, rotating disk galaxies with star formation rates of $\\sim 140$ M$_{\\odot}$ yr$^{-1}$ and velocity dispersions of $\\sim 60$ km s$^{-1}$   \\citep[e.g.][]{2006Natur.442..786G}. The origin of the high star formation rates and the large velocity dispersion in these galaxies is not presently understood.  On the theoretical side, detailed comparisons between observations at $z=2$ and N-body-SPH simulations or semi-analytical models (SAM) show  discrepancies. Although the overall slope of the $M_*-$SFR is well recovered in the models, the normalization is significantly smaller \\citep[D07;][see however Dekel et al. 2008]{2008MNRAS.385..147D}.  Theory and numerical simulations show that galaxies accrete cold gas in two modes. At early times and in small dark matter halos, most of the gas never shock heats to the  virial temperature and is accreted on a free fall time scale \\citep{2003MNRAS.345..349B,2005MNRAS.363....2K,2006MNRAS.368....2D,2008MNRAS.390.1326O,2009MNRAS.395..160K}. The general condition for a stable extended shock to be maintained is that the post-shock gas remains highly pressurized. In practice, this means that the loss of energy by radiative cooling needs to be at least compensated by the compression of post-shock infall material \\citep{2006MNRAS.368....2D}. This condition leads to a critical halo mass of M$_{shock} \\sim 10^{11.6}$ M$_{\\odot}$ \\citep{2008MNRAS.390.1326O} for shock heating of gas and hence a transition from fast and efficient \"cold accretion\" onto the galaxy to  a more  moderate \"hot accretion \" mode set by the radiative cooling time of the gas.  In an earlier study \\citet{2008ApJ...688..789G} showed that the number density and accretion rates of dark matter halos  suffice to provide enough high star forming galaxies, assuming that the newly accreted dark matter provides the universal fraction of baryons, and that these baryons are all transformed into stars. Alternatively, using a hydrodynamical simulation, \\citet{2009Natur.457..451D} showed that the gas accretion rates are comparable to the observed star formation rates thus requiring a very efficient transformation of accreted gas into stars.  In this Letter,  we aim at addressing the interlinked  questions  of high star formation rates and velocity dispersion observed in galaxies at $z=2$. Our work expands on the results previously presented by \\citet{2009Natur.457..451D} and \\citet{2008ApJ...688..789G}.  In particular, we model the effect of cold accretion on driving turbulence in the interstellar medium (ISM) and employ a porosity description for the ISM and star formation within it. The latter hypothesis is able to provide high star formation efficiencies in the context of a physically motivated model. ",
        "conclusions": "We have investigated a possible scenario for the formation of galaxies at high redshifts that is driven mainly by cold accretion flows. Massive galaxies at $z=2$ live in halos that are at the intersections of filaments that form the cosmic web. As shown in numerical simulations, these filaments provide a steady stream of dense cold gas that can reach the galaxy on a free-fall time scale and fuel ongoing star formation. Not only do these streams provide cold gas, but they can also drive turbulence in the gaseous disk by depositing some fraction of the accretion energy into the ISM. We find that between $10\\%-20\\%$ of this accretion energy is sufficient to increase the gas velocity dispersion to observed values. In particular we choose $\\eta=0.18$, which still leaves accretion driven turbulence dominant with respect to supernovae driven one.  A natural consequence of cold accretion driving turbulence is that it becomes less important in massive halos  $> 10^{12.5}$ $M_{\\odot}$ and at low redshifts, \\begin{figure} \\plotone{f3_color.eps}\\label{f3} \\figcaption{Number density of star-forming galaxies with SFR  $>60$ M$_{\\odot}$ yr$^{-1}$ (stars), $ > 120$ M$_{\\odot}$ yr$^{-1}$ (circles) and  $> 300$ M$_{\\odot}$ yr$^{-1}$ (triangles) . The square is the data of D07 for galaxies with  SFR $> 120$ M$_{\\odot}$ yr$^{-1}$.} \\end{figure} where the fraction of cold mode accretion drops and $V/\\sigma$ increases. Low redshift observations find velocity dispersions of 10 km s$^{-1}$, much lower than those at $z=2$. This can be understood as a natural consequence of the steady demise of cold accretion with redshift in halos that host disk galaxies today. The host halos of these galaxies live no longer at the intersection of filaments, but in the filaments itself. Those galaxies living in halos above the critical dark matter mass for stable shocks \\citep{2006MNRAS.368....2D} will no longer accrete from cold flows, but from hot halo gas, if at all, that radiatively cools and the main driver of turbulence will be supernovae. Such  a picture does fit in nicely with observation of low redshift disk galaxies, which have much larger $V/\\sigma$ than their high-$z$ counter parts \\citep{2007ApJ...660L..35K}.  Only increasing the amount of available fuel due to cold accretion is not enough to increase the number density of galaxies with large star formation rates. It is necessary to also increase the star formation efficiency at the same time. We \\begin{figure} \\plotone{f4_color.eps}\\label{f4} \\figcaption{Galaxy mass functions at $z=2$. Filled circles show results including cold accretion and the porosity based model for star formation.    } \\end{figure} tested an analytic model for a multiphase ISM that is based on a porosity prescription, that has been shown to capture some of the relevant physical mechanism related to star formation seen in high-resolution simulations \\citep{2005MNRAS.356..737S}. The star formation efficiency in this model is $\\alpha \\propto \\sigma$, resulting in an increasing efficiency that parallels the gas supply by cold accretion and allows to convert it into stars on a very short time-scale. With this model we find improved agreement with the observed number density of star-forming galaxies.   However our predicted star formation rates are lower than in the observed $z\\sim 2$ sample especially for the more massive galaxies. We have two comments on this point. Firstly,  the data are most likely biased to high star formation rates. Secondly, galaxies with high star-forming rates at high redshift commonly show  weak AGN activity (see for example  a recent study of a sample of $z\\sim 1.9 $ ULIRGs \\citep{2009huang} ). This most likely is the tip of the AGN iceberg  because of  possible obscured AGN. AGN outflows could trigger as well as play the more accepted  role of quenching star formation at high redshift. It remains to be investigated how well the agreement at different redshifts is, and we plan on addressing this in a future paper. \\\\  We would like to thank the referee Avishai Dekel for his very helpful comments that helped improve the manuscript, Andreas Burkert and the SINS team for useful discussions, and Emanuele Daddi as well as Giovanni Cresci for kindly providing their observational data."
    },
    "0812/0812.4559_arXiv.txt": {
        "abstract": "We consider the luminosity  and environmental dependence of structural parameters  of  lenticular galaxies  in  the  near-infrared $K$  band. Using  a  two-dimensional  galaxy  image decomposition  technique,  we extract  bulge and  disk  structural  parameters for  a  sample of  36 lenticular galaxies observed by us  in the $K$ band. By combining data from the literature  for field and cluster lenticulars  with our data, we study  correlations between parameters that  characterise the bulge and the  disk as  a function of  luminosity and environment.   We find that  scaling relations  such  as the  Kormendy relation,  photometric plane and other correlations  involving bulge and disk parameters show a luminosity dependence.  This dependence can be explained in terms of galaxy  formation models in  which faint  lenticulars ($M_T  > -24.5$) formed  via  secular  formation   processes  that  likely  formed  the pseudobulges  of late-type disk  galaxies, while  brighter lenticulars ($M_T <  -24.5$) formed through  a different formation  mechanism most likely involving  major mergers.  On probing variations  in lenticular properties as  a function of  environment, we find that  faint cluster lenticulars show  systematic differences  with respect to  faint field lenticulars.  These differences support  the idea  that the  bulge and disk  components fade  after the  galaxy falls  into a  cluster, while simultaneously undergoing  a transformation from  spiral to lenticular morphologies. ",
        "introduction": "Lenticular (S0) galaxies were  originally conceived as a morphological transition class between ellipticals  and early-type spirals by Hubble (1936). These galaxies have disks spanning a wide range in luminosity, contributing between five to fifty  percent of the total galaxy light. They are distinguished from spiral galaxies by the lack of conspicuous spiral arms.   For these reasons, lenticulars  can be thought  of as a population which  is intermediate  between ellipticals and  early type spirals.  Indeed, the placement of lenticulars by Hubble (1936) on the tuning fork  diagram clearly implies  such a relationship.  In several observable  properties such  as bulge-to-disk  luminosity  ratio, star formation rate  and colour,  lenticulars are, on  average, intermediate between ellipticals and early type spirals.  \\begin{table*} \\begin{center} \\begin{minipage}{166mm} \\caption{Best fit Bulge and Disk parameters for our sample.} \\label{table1} \\begin{tabular}{l l l c c c c c c c c c} \\hline Name & z & T  & \\multicolumn{4}{c}{Bulge parameters} & \\multicolumn{3}{c}{Disk parameters} & B/T & M$_T$ \\\\ \\hline & &  & $\\mubzero$ & $\\re$ & $n$ & $e_b$ & $\\mudzero$ & $r_d$ & $e_d$ &  & \\\\ & &  & (mag arcsec$^{-2}$) & (kpc) &  &  & (mag arcsec$^{-2}$) & (kpc) &  &  &  \\\\ \\hline UGC\\,00080 & 0.01019 &    -3 & 11.54 &  0.45 &  1.63 & 0.001 & 15.69 &  2.33 & 0.505 & 0.19 & -24.99  \\\\ UGC\\,00491 & 0.01664 &    -1 & 10.38 &  4.95 &  3.82 & 0.257 & 17.65 &  5.08 & 0.231 & 0.65 & -25.53  \\\\ UGC\\,00859 & 0.00711 &     0 & 11.34 &  1.39 &  3.36 & 0.141 & 16.40 &  1.22 & 0.403 & 0.43 & -23.25  \\\\ UGC\\,00926 & 0.01467 &  -2.5 & 09.74 &  9.81 &  4.39 & 0.348 & 14.93 &  0.80 & 0.353 & 0.93 & -26.17  \\\\ UGC\\,01250 & 0.01235 &    -2 & 09.51 &  1.92 &  3.79 & 0.397 & 15.92 &  2.66 & 0.696 & 0.32 & -25.24  \\\\ UGC\\,01823 & 0.01332 &    -3 & 11.41 &  1.96 &  2.28 & 0.416 & 16.81 &  6.56 & 0.398 & 0.34 & -26.34  \\\\ UGC\\,01964 & 0.01732 &    -2 & 10.06 &  4.03 &  3.96 & 0.103 & 15.60 &  1.86 & 0.585 & 0.58 & -25.33  \\\\ UGC\\,02039 & 0.01505 &    -2 & 10.18 &  1.92 &  3.37 & 0.007 & 17.19 &  5.52 & 0.644 & 0.29 & -25.52  \\\\ UGC\\,02187 & 0.01605 &    -2 & 13.11 &  0.45 &  0.76 & 0.276 & 15.13 &  1.99 & 0.554 & 0.14 & -25.17  \\\\ UGC\\,02322 & 0.01447 &    -1 & 10.51 &  4.18 &  3.86 & 0.008 & 16.84 &  2.08 & 0.015 & 0.75 & -24.89  \\\\ UGC\\,03178 & 0.01578 &    -2 & 09.84 &  2.55 &  3.89 & 0.014 & 14.76 &  1.17 & 0.479 & 0.47 & -24.90  \\\\ UGC\\,03452 & 0.01876 &    -2 & 10.53 &  3.31 &  3.47 & 0.242 & 15.91 &  2.89 & 0.516 & 0.46 & -25.70  \\\\ UGC\\,03536 & 0.01564 &    -2 & 13.18 &  0.55 &  0.99 & 0.012 & 14.80 &  2.03 & 0.576 & 0.10 & -25.48  \\\\ UGC\\,03567 & 0.02016 &    -2 & 09.96 &  4.87 &  4.15 & 0.062 & 15.99 &  1.06 & 0.503 & 0.87 & -25.04  \\\\ UGC\\,03642 & 0.01500 &    -2 & 11.94 &  1.03 &  2.02 & 0.010 & 17.74 &  8.45 & 0.206 & 0.17 & -25.72  \\\\ UGC\\,03683 & 0.01908 &    -2 & 12.13 &  1.56 &  1.83 & 0.283 & 17.01 &  5.69 & 0.171 & 0.37 & -25.91  \\\\ UGC\\,03699 & 0.01947 &    -2 & 10.35 &  3.13 &  3.32 & 0.336 & 16.54 &  1.65 & 0.489 & 0.87 & -25.36  \\\\ UGC\\,03792 & 0.01908 &     0 & 10.17 &  4.41 &  3.74 & 0.224 & 18.15 &  5.71 & 0.225 & 0.72 & -25.66  \\\\ UGC\\,03824 & 0.01787 &    -2 & 11.40 &  1.52 &  2.71 & 0.003 & 17.43 &  3.18 & 0.002 & 0.52 & -24.52  \\\\ UGC\\,04347 & 0.01494 &    -2 & 10.68 &  4.33 &  3.34 & 0.243 & 16.34 &  1.01 & 0.528 & 0.95 & -25.57  \\\\ UGC\\,04767 & 0.02413 &    -2 & 09.05 &  7.65 &  4.74 & 0.105 & 14.98 &  0.27 & 0.001 & 0.99 & -25.55  \\\\ UGC\\,04901 & 0.02810 &    -2 & 13.30 &  3.06 &  1.78 & 0.305 & 18.08 & 14.28 & 0.043 & 0.27 & -26.70  \\\\ UGC\\,05292 & 0.00511 &    -2 & 09.22 &  0.26 &  3.08 & 0.008 & 17.28 &  2.01 & 0.282 & 0.21 & -23.07  \\\\ UGC\\,06013 & 0.02190 &    -2 & 11.29 &  5.10 &  3.68 & 0.325 & 15.41 &  1.00 & 0.001 & 0.78 & -24.89  \\\\ UGC\\,06389 & 0.00663 &    -2 & 11.71 &  3.00 &  3.36 & 0.378 & 13.88 &  0.21 & 0.334 & 0.89 & -23.75  \\\\ UGC\\,06899 & 0.02249 &    -2 & 09.79 &  5.11 &  4.06 & 0.060 & 16.71 &  0.53 & 0.001 & 0.99 & -25.35  \\\\ UGC\\,07142 & 0.00328 &    -2 & 10.48 &  0.35 &  2.22 & 0.219 & 15.96 &  1.32 & 0.503 & 0.33 & -23.66  \\\\ UGC\\,07473 & 0.00414 &    -2 & 11.76 &  0.24 &  1.20 & 0.239 & 14.55 &  1.20 & 0.725 & 0.12 & -24.57  \\\\ UGC\\,07880 & 0.00388 &    -3 & 11.50 &  0.68 &  2.39 & 0.330 & 14.29 &  0.63 & 0.747 & 0.34 & -23.74  \\\\ UGC\\,08675 & 0.00361 &    -2 & 09.18 &  0.25 &  3.35 & 0.020 & 17.14 &  1.48 & 0.011 & 0.19 & -22.53  \\\\ UGC\\,09200 & 0.01081 &    -2 & 12.33 &  1.35 &  2.23 & 0.117 & 17.83 &  3.10 & 0.001 & 0.57 & -24.15  \\\\ UGC\\,09592 & 0.01791 &    -2 & 08.69 &  1.62 &  3.62 & 0.315 & 16.30 &  2.33 & 0.123 & 0.65 & -25.30  \\\\ UGC\\,11356 & 0.00820 &  -2.5 & 11.64 &  0.47 &  1.64 & 0.080 & 15.16 &  1.46 & 0.020 & 0.24 & -24.57  \\\\ UGC\\,11972 & 0.01463 &  -2.5 & 10.29 &  8.62 &  4.08 & 0.441 & 14.24 &  0.45 & 0.002 & 0.95 & -26.00  \\\\ UGC\\,12443 & 0.01463 &    -2 & 10.83 &  3.34 &  3.44 & 0.097 & 19.36 &  0.04 & 0.002 & 1.00 & -24.61  \\\\ UGC\\,12655 & 0.01727 &    -2 & 10.34 &  1.14 &  2.87 & 0.191 & 16.86 &  4.77 & 0.400 & 0.24 & -25.47  \\\\  \\hline \\end{tabular}  Notes.\\ Column (1) gives the UGC catalogue  number, column (2) and  column (3) give redshift and morphological  type respectively from NED, columns (4), (5), (6) and (7) give unconvolved bulge central surface brightness, bulge effective radius, S{\\'e}rsic index and bulge ellipticity respectively,  columns (8) (9) and (10)  give unconvolved disk  central   surface  brightness, disk  scale  length   and  disk ellipticity  respectively,   column  (11) gives   the  bulge-to-total luminosity ratio and  column (12) gives the absolute  magnitude in the $K$ band. \\end{minipage} \\end{center} \\end{table*}   Detailed study  of individual  lenticular galaxies indicates  that the situation is more complex in  reality.  It has been suggested (van den Bergh 1994) that there are different, but overlapping, sub-populations amongst  the  lenticulars.  Exploration  of  formation  scenarios  for lenticulars using  theory and numerical simulations  also suggest that lenticulars  may have  formed in  different  ways.  They  could be  of primordial  origin forming rather  rapidly at  early epochs,  or could have  been formed by  the slow  stripping of  gas from  spirals, which changes the  morphology (Abadi, Moore  \\& Bower 1999), or  through the mergers of unequal-mass spirals (Bekki 1998).  The two main components of lenticulars  -- the bulge and the  disk -- may have  their own {\\it different}     and     possibly     {\\it    independent}     formation history. Observationally, the disks seem to be younger than the bulges in both  spiral and lenticular  galaxies (Peletier \\&  Balcells 1996), with  the  age  difference   between  the  two  components  larger  in lenticulars (Bothun \\& Gregg 1990). Bars in lenticular galaxies are also of interest in many studies as they provide a clue about the evolutionary history of these galaxies. Recent studies reveal that bars in lenticular galaxies are  shorter, less massive and  have smaller bar torques (Buta \\etal 2006; Laurikainen \\etal 2005; Laurikainen \\etal 2006; Gadotti, D. A. \\etal 2007).  A detailed multiband study of the morphology  of  representative  samples of lenticulars  in  different environments,  and  comparison  of  their properties  with  those  of ellipticals,  and  with bulges  and  disks of  spirals will  be important  in  addressing   these  issues  observationally. However, comparing the predictions of models to observations is complicated for two  main reasons: (1)  The models  often use  simplifying assumptions that may not  hold for real galaxies and (2) Models  often do not make firm predictions about {\\it directly observable} quantities. Nevertheless, the {\\it statistical} properties of galaxy ensembles can be  compared  to model  predictions.  Such  a quantitative  comparison between  bulge and disk properties  of  galaxies is considerably simplified if  one assumes simple analytic profiles  for the  light distribution  of the bulge  and  the  disk.  In  practice,  the  S{\\'e}rsic  $r^{1/n}$  law (S{\\'e}rsic 1968)  adequately represents the  bulge surface brightness distribution while  an  exponential  best  represents  the  disk surface brightness distribution. The  total galaxy light is simply the sum of a S{\\'e}rsic bulge and an exponential disk.  In  this work,  we  obtain  and report  structural  parameters for  36 lenticular  galaxies observed  in the near-infrared  $K$ band.   Using results from the bulge-disk decomposition, we examine correlations among the bulge  and disk parameters  and discuss implications to  models of lenticular  formation. Our  goal is  mainly to  study  the constraints placed on  bulge formation in  lenticular galaxies by  these parameter correlations. Specifically, we want to investigate, in greater detail, the  discovery   of  Barway  \\etal  (2007)  who   found  two  distinct populations of  bulges in lenticulars  that were fainter  and brighter than a threshold luminosity.  This  paper   is  organised  as   follows:  the  sample,   details  of near-infrared  observations,  data  reduction  technique and  photometric calibration are   summarised  in   $\\S$   2.   The two-dimensional decomposition analysis is described in $\\S$ 3. In $\\S$ 4, we discuss the  luminosity and environment dependence of lenticular galaxy properties.  Throughout this paper, we use the standard concordance cosmology with $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$, and $h_{100} = 0.7$. ",
        "conclusions": "\\subsection{Luminosity dependence}  Barway \\etal \\ (2007) have presented evidence to support the view that the formation  history  of  lenticular  galaxies depends  upon  their luminosity.  According to  this view,  low-luminosity lenticular galaxies likely formed  by the  stripping of  gas from  the disk  of  late-type spiral galaxies, which in turn  formed their bulges through secular evolution processes. On the other  hand, more luminous lenticulars likely formed at early epochs through a  rapid collapse followed  by rapid star formation.  As we mentioned  earlier, the BAM06 galaxies complement  our sample in two ways: they extend  to fainter luminosities and provide lenticulars in  different  environments.    In  Figure~\\ref{f1}  we  show  the distribution of total  absolute magnitude ($M_T$) in the  $K$ band for the combined sample, which is seen to span a wide range in luminosity. We divide the combined sample into faint and bright groups, using $M_T = -24.5$  as a  boundary.  The bright sample has 37 lenticulars while 46 lenticulars belong to the faint sample according to this luminosity division. The  boundary at $M_T  = -24.5$  is somewhat arbitrary but our results do not critically depend on small shifts in the dividing luminosity.  For instance,  changing this value by half a magnitude on either side  maintains correlations significant at least at the 95\\% level.  Using the luminosity division at $M_T  = -24.5$, Barway \\etal \\ (2007) reported  markedly different correlations between bulge effective radius ($\\re$) and disk scale length ($\\rd$) for bright and faint lenticulars. A  positive correlation of bulge and  disk sizes is expected  if  the bulge  grows  over  time  through secular  evolution processes  (see  review  by  Kormendy  \\&  Kennicutt  2004).  No  such correlation  is  expected  if  the  bulge formed  via  merger  related processes.  Barway \\etal \\ (2007) found this positive correlation between bulge and disk sizes of low luminosity lenticulars; such a correlation was not observed for bright lenticulars. Moreover, they found that the correlation holds for faint galaxies, irrespective of whether they are situated in a field or cluster environment.  If  the differences between  low and  high luminosity  lenticulars are indeed fundamental,  then there should be  systematic differences seen between  the two  populations  in other  correlations among bulge and disk parameters. We study such correlations in the following sections.   \\begin{figure} \\centerline{\\psfig{figure=f4.eps,width=8.0cm,angle=0.}} \\caption{Photometric Plane for bright lenticulars (open circles) with restricted  $n$ value. Faint  lenticulars (filled  circles) are over plotted on the same photometric plane. These show a higher scatter than the bright ones.} \\label{f4} \\end{figure}   \\subsubsection{Kormendy relation and the effect of bars}  For samples of large elliptical galaxies, the central surface brightness $\\mubzero$ (or, equivalently, the mean surface brightness ($\\mmubre$) within $\\re$) is correlated with $\\log \\ \\re$.  This was first reported by Kormendy \\ (1977) and is known as the Kormendy relation.  In Figure~\\ref{f2}(a) we show the Kormendy relation for bright (as open circles) and faint lenticulars (as filled circles). The bright lenticulars show tight correlation with Pearson correlation coefficient $r = 0.95$ with significance greater than 99.99 \\%.  The Kormendy relation for faint lenticulars shows considerably greater scatter although the Pearson correlation coefficient is $0.75$ with significance greater than 99.99 \\%. Figure~\\ref{f2}(b) show the Kormendy relation for bright (as crosses) and faint (as dashes) lenticulars classified as barred lenticulars in RC3 catalogue to probe for any systematic effect caused due to presence of bars. It is clear from the figure that barred and non-barred show no systematic offset. Only three galaxies are classified as barred in our bright galaxy sample (indicated by crosses). Even after removing these three galaxies, the correlation is nearly unchanged with Pearson correlation coefficient $r = 0.96$ with significance greater than 99.99 \\%.  In the faint lenticular sample, there are as many as 16 galaxies classified as barred. Nevertheless, neglecting these 16 galaxies does not affect the correlation which has $r = 0.74$ with significance greater than 99.99 \\%. For all other correlations subsequently reported in this paper, we find that excluding barred galaxies does not significantly alter the correlation coefficient or its significance. Hence, in further discussion, we do not differentiate between barred and non-barred lenticulars, and include both types in the analysis.  The Kormendy relation is known  to be a projection  of the fundamental plane (Djorgovski \\& Davis 1987; Dressler \\etal 1987) of galaxies.  A tight  Kormendy relation indicates (near) virialized bulges like  those found  in elliptical galaxies.   In Figure 2a, ellipticals  from  the  Coma  cluster  are overplotted as  dots.   As expected,  these ellipticals  show  a tight  correlation with Pearson correlation  coefficient $r  =  0.81$ with  significance greater  than 99.99  \\%. The  slopes of  the best-fit  Kormendy relation  for bright lenticulars  and Coma  ellipticals are similar.   This similarity coupled with the systematic difference  in scatter  in the  Kormendy relation  between bright and faint lenticulars is an indicator  of the more virialized state of the bright lenticular bulges and their close relation to ellipticals.   \\subsubsection{Photometric Plane}  Analogous  to the  fundamental plane (FP),  ellipticals and  the  bulges of early type spiral  galaxies obey a single planar  relation of the form $\\log \\ n  = a \\ \\log \\ r_{\\rm e}+ b \\ \\mu_{\\rm b}(0)+ c $  which Khosroshahi \\etal  (2000a,b)  called the  photometric  plane  (PP). Assuming  that galaxy bulges behave  as spherical, isotropic, one-component systems, there is a  unique specific entropy that can  be associated with every galaxy  (Lima-Neto \\etal  1999).   This in  turn,  implies a  relation between  bulge   parameters  that   is  equivalent  to   the  observed photometric plane. An edge-on view of the best-fit PP is shown in the top  and bottom  panels of  Figure~\\ref{f3} for  bright  and faint lenticulars respectively. Table 2  lists the coefficients for best-fit PP relation.  It is  immediately evident from Figure~\\ref{f3} that the PP for faint lenticulars spans a limited range of $n$. We restrict our  bright  lenticular sample  to  the  range  spanned by  the  faint lenticulars  and then  obtain  the  best fit  PP  for this  restricted sample.   The rms  scatter  for the  (restricted)  bright galaxies  is considerably smaller than  that for the faint galaxies  (see Table 2). In  Figure~\\ref{f4}  we  show  the   PP  for  these  two  sets  of galaxies. Visually, the higher scatter of the faint galaxies about the PP is obvious. The PP is a natural state for galaxies that formed like ellipticals (at early epochs in  a burst of rapid star formation, with stellar orbits that are well  relaxed). The fact that bulges of bright lenticulars  have  a  tight  PP  is  an  indicator  of  homology  with ellipticals. Faint lenticulars seem to belong to a separate class.   \\begin{figure} \\centerline{\\psfig{figure=f5.eps,width=8.0cm,angle=0.}} \\caption{S{\\'e}rsic index $n$ plotted against bulge effective radius $r_e$ in kpc.  Symbols are as in Figure 2.} \\label{f5} \\end{figure}   \\subsubsection{Correlations involving the S{\\'e}rsic index ($n$)}  The  restricted  range  spanned  by  the  S{\\'e}rsic  index  in  faint lenticular  galaxies also  affects other  correlations  involving this index.  Figure~\\ref{f5} is a plot of the S{\\'e}rsic index ($n$) as a function  of effective radius  ($\\re$) which again  shows systematic differences  between  bright  and   faint  lenticulars.   Bright lenticulars   are  well  correlated  having   Pearson  correlation coefficient $r  = 0.79$  with significance greater  than 99.99  \\% but faint  lenticulars do not  show a  significant correlation.   A strong correlation  exists   between  $n$  and  the   bulge  central  surface brightness  ($\\mubzero$),  as  shown  in Figure~\\ref{f6}  for  the brighter lenticulars. The Pearson correlation coefficient in this case is 0.84 at a significance level better than 99.99\\%. Faint lenticulars also exhibit  good correlation having  Pearson correlation coefficient $r =  0.58$ at  significance 99.96 \\%,  but have a  markedly different slope.  Comparing Figure~\\ref{f5} and Figure~\\ref{f6} with Figure 3 and Figure  4 respectively of Khosroshahi \\etal (2000b), we observe a  close correspondence in the distribution  of the bulges of {\\it bright} lenticulars and early  type spirals. Bulges of early type spirals,  in  turn,  have  a  close  correspondence  with  ellipticals (Khosroshahi \\etal  2000b). On the other hand,  faint lenticulars seem to be a different population.  The  bulge-to-total  luminosity  ratio  ($B/T$)  increases  with S{\\'e}rsic  index $n$  as  shown in  Figure~\\ref{f7} for  brighter lenticulars  ($r   =  0.81$  with  significance   greater  than  99.99 \\%). Faint lenticulars show a weak correlation between $ \\log \\ (B/T)$ and $\\log \\ n$.  Again, the  trends for  bright lenticulars  are consistent with those  found by Khosroshahi \\etal (2000b) in early-type spirals.  Bright lenticulars plotted in Figure 3 of Barway \\etal (2007) include 5 galaxies that are obvious outliers in the anti-correlation of bulge and disk sizes. These bright lenticulars viz. UGC\\,80, UGC\\,2187, UGC\\,3536, UGC\\,7473 and UGC\\,11356  do follow the correlations involving the S{\\'e}rsic index. However, they are all characterised by a low S{\\'e}rsic index, low bulge central surface brightness,  low bulge effective radius and low bulge-to-total luminosity ratio. This places these lenticulars in the bottom left quadrant of Figures 5, 6 and 7. Their bulge/disk parameters indicate that these are disk-like systems, with relatively weak bulges, making them unlikely candidates of major merger induced formation. Nevertheless, they do follow the same correlations, involving the S{\\'e}rsic index, as the other bright galaxies. This discrepancy needs to be explored with a larger sample of bright, but disky lenticulars.    \\begin{figure} \\centerline{\\psfig{figure=f6.eps,width=8.0cm,angle=0.}} \\caption{S{\\'e}rsic  Index $n$  as  a function  of unconvolved  bulge central surface brightness. Symbols are as in Figure 2.} \\label{f6} \\end{figure}   \\begin{figure} \\centerline{\\psfig{figure=f7.eps,width=8.0cm,angle=0.}} \\caption{S{\\'e}rsic   Index   $n$   vs.   bulge-to-total   luminosity ratio. Symbols are as in Figure 2.} \\label{f7} \\end{figure}   \\begin{figure} \\centerline{\\psfig{figure=f8.eps,width=8.0cm,angle=0.}} \\caption{Unconvolved disk central  surface brightness $\\mu_d(0)$ as a function of  disk scale length. Early-type  spirals (as dots) are overplotted for comparison. Other symbols are as in Figure 2.} \\label{f8} \\end{figure}   \\begin{figure} \\centerline{\\psfig{figure=f9.eps,width=8.0cm,angle=0.}} \\caption{Dependence  of  the   mean  surface  brightness  within  bulge effective radius ($\\mmubre$) on  bulge effective radius $\\re$ as a  function of  environment.  Faint  field lenticulars are shown  as open triangles   while  faint   cluster  lenticulars are  plotted   as  filled squares.} \\label{f9} \\end{figure}   \\subsubsection{Disk correlations}  There exist  clear correlations between  the two main  disk parameters --central  surface  brightness  of  the disks  ($\\mudzero$)  and  disk scale-length  ($\\rd$)-- \\ for  both  bright and  faint lenticulars  (see Figure~\\ref{f8})  although they  occupy different  regions  of the plot.   Lenticulars with  large  disks have  a  lower central  surface brightness,  on average.   A  clear anti-correlation  is seen  between these two  disk parameters for  bright lenticulars but the  scatter is large.  The linear correlation  coefficient between  $\\mudzero$ \\ and log $\\rd$ is 0.43 with a  significance of 99.37 percent. Faint lenticulars show  less scatter  with  the  correlation coefficient  of  0.70 at  a significance  level  better  than  99.99  \\%.   Similar  statistically significant correlation was reported  by Khosroshahi \\etal (2000b) for early-type  spirals  and  M\\\"{o}llenhoff \\& Heidt  (2001)  for  late-type spirals for  these disk  parameters. In Figure~\\ref{f8},  we also plot as dots the central surface  brightness of the disks ($\\mudzero$) against the disk scale length ($\\rd$)  of galaxies from the $K$ band observations of Khosroshahi \\etal  (2000b) of early-type  spirals.   From the plot it is apparent that the  disk scale lengths of our sample and the Khosroshahi \\etal  (2000b) sample span the same  range, but $\\mudzero$ are  brighter,  on  the  average,  for  early-type  spirals  than  faint lenticulars.   \\begin{figure} \\centerline{\\psfig{figure=f10,,width=8.0cm,angle=0.}} \\caption{Unconvolved disk  central surface brightness  $\\mudzero$ is plotted   against  disk   scale  length   $\\rd$  as   a   function  of environment.  Symbols are as in Figure 9.} \\label{f10} \\end{figure}   \\subsection{Environmental dependence}  The dependence  of cluster environment  on formation and  evolution of galaxies  is well  known. The  environment  influence is  best visible  through the well-known   morphology-density  relationship,   according   to  which early-type  galaxies  are preferentially  found  in rich  environments (Dressler 1980),  with spirals are found predominately  in low density environments.    One   of   the    stronger   trends   seen   in   the morphology-density relation  at low redshift is  the dramatic increase in the number of lenticular galaxies in rich clusters (Dressler \\etal 1997), suggesting    that   the    dominant   process    in    defining   the morphology-density relation  is the transformation  of spiral galaxies into lenticular  galaxies within rich clusters  (Poggianti \\etal 1999; Kodama \\& Smail 2001). These spiral galaxies are continuously supplied by  accretion from  the surrounding  field  during the  course of  the assembly of  the cluster.  {\\it Hubble Space  Telescope (HST)} imaging of  10 clusters  at higher  redshift (z  $\\sim$\\,0.3-0.5)  by Dressler \\etal  1997  shows  similar  behaviour.  The  mechanisms  proposed  for morphological  transformation   of  spirals  to   lenticulars  include dynamical interactions  such as  galaxy harassment (Moore  \\etal 1996, 1998, 1999)  and gas dynamical processes e.g.   ram pressure stripping (Abadi \\etal 1999). Simulations have shown that galaxy harassment is a very efficient  mechanism for transforming early-type  disk galaxies to lenticulars (Moore 1999) with consequent reduction in disk sizes. Ram pressure  stripping truncates the  star formation by  removing the cold  neutral gas  reservoir causing the  disk component  of a spiral galaxy to fade within the cluster environment.  To  address   these possibilities   observationally,  it  is  necessary  to investigate how scaling relations defined by various galaxy properties vary  between field  and cluster  environments.  In  this  section, we examine     a     few     correlations     as    a     function     of environment.   Unfortunately,  we   are  unable   to   include  bright lenticulars in this comparison because  only 3 cluster galaxies out of the 38  in our sample,  are bright. For  faint galaxies, the  number of galaxies  in the  field and  in clusters  are almost  equal,  making a comparison possible.  The   Kormendy    relation   for   faint    lenticulars   plotted   in Figure~\\ref{f2}   shows  two   distinct   groups,  unlike   bright lenticulars. To investigate the origin  of this dichotomy in the faint lenticulars,  we plot  the  Kormendy relation  for faint  lenticulars, differentiating  them  by   environment,  in  Figure~\\ref{f9}.  We indicate field lenticulars with open triangles and cluster lenticulars with  filled  squares.   The  field  lenticulars seem  to  follow  the Kormendy relation reasonably well while the cluster lenticulars show a clear downward scatter with respect  to the relation. On average, this requires a 1.5 mag arcsec$^{-2}$ fading of the mean surface brightness within  $\\re$  of  faint  cluster  lenticulars, with  respect  to  the Kormendy  relation.  Boselli  \\&  Gavazzi (2006)  have suggested  that low-luminosity  lenticulars in  clusters might  be the  result  of ram pressure stripping of late-type galaxies, which causes fading and thus lowers surface brightness.  It is interesting to see if this fading is also seen for the disk component.  In Figure~\\ref{f10}, we plot the disk central surface brightness against the disk scale length. As above, we indicate faint field and cluster lenticulars with different symbols.  The faint field lenticulars exhibit a clear anti-correlation, whereas the cluster lenticulars occupy a limited region of the plot and show a downward scatter (indicating fading of the disk), very similar to the scatter seen in Figure~\\ref{f9}.  Our results obtained using only photometric data are consistent with the spectroscopic results of Barr \\etal (2007), which support the theory that lenticular galaxies are formed when gas in normal spirals is removed (from both bulge and disk), possibly when well formed spirals fall into the cluster.  It must be noted that the sample of Barr \\etal (2007) is a subset of the BAM06 sample, which is included in our analysis.  If lenticulars in clusters are indeed transformed spirals, it is likely that they  preserve other signatures  of their earlier  existence. For instance, if  they contained pseudobulges that  formed through secular evolution, their  bulge and disk  sizes should be  correlated (Courteau \\etal  1996). In  Figure~\\ref{f11},  we plot  the bulge  effective radius  against  the  disk  scale  length for faint lenticulars (a similar plot for bright lenticulars may be found in Barway \\etal (2007)). Faint field  and  cluster lenticulars  seems to both  follow the  same correlation, but inhabit different regions of  the plot. We also overplot data  for a sample of late-type  disk  galaxies  from  M\\\"{o}llenhoff \\&  Heidt  (2001). The correlation  for  faint lenticulars  and  late  type  spirals has similar slopes ($\\sim 1.40\\pm0.14$  ) but different intercept which is expected as late-type spirals have  larger disk scale length compared to faint lenticulars. This is  a natural expectation if late-type spirals are transforming  into faint lenticulars  due to interaction  with the cluster medium as well as with other galaxies in the cluster.    \\begin{figure} \\centerline{\\psfig{figure=f11,width=8.0cm,angle=0.}} \\caption{Dependence of the $\\re - \\rd$  relation on the  environment in $K$ band  for the faint lenticulars. Late-type  spirals (as asterisks) are overplotted for comparison. Other symbols are as in Figure 9. The dotted line is the best fit to the faint lenticulars (field as well as cluster) excluding the one obvious outlier while the dashed line is the best fit to late-type spirals. Note that these two lines have similar slopes but a different intercept.} \\label{f11} \\end{figure}"
    },
    "0812/0812.0466_arXiv.txt": {
        "abstract": "Rules for the transformation of time parameters in relativistic Langevin equations are derived and discussed. In particular, it is shown that, if a coordinate-time parameterized process approaches the relativistic J\\\"uttner-Maxwell distribution, the associated proper-time parameterized process converges to a modified  momentum distribution, differing by a factor proportional to the inverse energy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.2525_arXiv.txt": {
        "abstract": "Many, if not all, post AGB stellar systems swiftly transition from a spherical to a powerful aspherical pre-planetary nebula (pPNE) outflow phase before waning into a PNe. The pPNe outflows require engine rotational energy and a mechanism to extract this energy into collimated outflows. Just radiation and rotation are insufficient but a symbiosis between rotation, differential rotation  and large scale magnetic fields remains promising.  Present observational evidence for magnetic fields in evolved stars is suggestive of dynamically important magnetic fields, but both theory and observation are rife with  research opportunity. I discuss how magnetohydrodynamic outflows might arise in pPNe and PNe and distinguish different between approaches that address shaping vs. those that address both launch and shaping. Scenarios involving dynamos in single stars, binary driven dynamos, or accretion engines cannot be ruled out. One appealing paradigm involves accretion onto the primary post-AGB white dwarf core from a low mass companion whose decaying accretion supply rate owers first the pPNe and then the lower luminosity PNe. Determining   observational signatures of  different MHD engines is a work in progress. Accretion disk theory and large scale dynamos  pose many of their own fundamental challenges, some of  which I discuss in a broader context. ",
        "introduction": "Asymmetries were observed in planetary nebulae PNe before the Hubble Space Telescope but their ubiquity,  the rapidity with which they develop, and collimated outflow power have led to a cultural change in the field over the past decade (e.g. Balick \\& Frank 2002). Asymmetric  p/PNe have become the standard and understanding how these asymmetries arise is now fundamental rather than anecdotal. Perhaps most dramatic is  that collimated momenta of the pre-PNe (pPNe) (the reflection nebulae precursor to the ionized nebulae of PNe) exceed that which can be supplied by radiation pressure alone (\\cite{bujar01}). This is reminiscent of the young stellar object (YSOs) subject decades ago (\\cite{yso}). The recent cultural change toward the view that the end states of stars (including supernovae e.g. Wang \\& Wheeler 2008) are asymmetric, offers a plethora of research opportunity.  The beginning and ending of stars share the ingredients of  in-fall, collapse, turbulence, and angular momentum transport. The associated increased free energy in differential rotation can be tapped to amplify fields and produce outflows, so common underlying MHD principles  likely at work. Focusing on pPNe/PNe, I review the role of magnetic fields in these systems, and current evidence for their influence. I also address  broader  fundamental questions in magnetic field generation and angular momentum transport for which pPNe are yet another laboratory.  The role of large scale fields and the need to connect different  MHD engines with observations are central to the theme. ",
        "conclusions": ""
    },
    "0812/0812.1901_arXiv.txt": {
        "abstract": "The achromatic phase shifter (APS) is a component of the Bracewell nulling interferometer studied in preparation for future space missions (viz. \\emph{Darwin}/TPF-I) focusing on spectroscopic study of Earth-like exo-planets. Several possible designs of such an optical subsystem exist. Four approaches were selected for further study. Thales Alenia Space developed a dielectric prism APS. A focus crossing APS prototype was developed by the OCA, Nice, France. A field reversal APS prototype was prepared by the MPIA in Heidelberg, Germany. Centre Spatial de Li\\`ege develops a concept based on Fresnel's rhombs. This paper presents a progress report on the current work aiming at evaluating these prototypes on the {\\sc Synapse} test bench at the Institut d'Astrophysique Spatiale in Orsay, France. ",
        "introduction": "\\label{sec:intro}  %  Nulling interferometry, based on the concept suggested by Bracewell\\cite{Art_Bracewell_1978} in 1978, is one of the methods of future exoplanet research. It has been studied for the mission \\emph{Darwin}\\cite{2003toed.conf...41K} proposed to ESA as well as for the mission TPF-I (Terrestrial Planet Finder - Interferometry)\\cite{TPF_Coulter} proposed to NASA.  The basic performance parameter of a nulling interferometer is the ``nulling ratio'' (sometimes also referred to as ``stellar leakage'' although this term refers more properly to stray starlight due to the fact a star is not a point source) $nl(\\lambda,t)=I_{\\min}/I_{\\max}$ where $I_{\\min}$ and $I_{\\max}$ stand for the intensity of the on-axis dark fringe and of the off-axis bright fringe, respectively.  The goal of \\emph{Darwin}/TPF-I is to perform infra-red ($6-18\\,\\mu\\mathrm{m}$) imagery of extrasolar planetary systems and spectroscopic observations of exoplanetary atmospheres in view of evaluating the presence of biomarkers. In order to reach this goal, the required nulling performance has been estimated\\cite{Variability_Noise} as $\\left\\langle nl\\right\\rangle (\\lambda) =1.0\\,10^{-5}\\,\\left(\\lambda/7\\,\\mathrm{\\mu m}\\right)^{3.37},$ with a level of long-term stability expressed as $\\sigma_{\\left\\langle nl\\right\\rangle } (7\\,\\mathrm{\\mu m},10\\,\\mathrm{days})\\leq 3\\,10^{-9}.$   For a given geometric difference $\\delta$ between the two optical paths, the corresponding phase difference $\\Delta\\varphi$ can be expressed as $\\Delta\\varphi=2\\pi\\delta/\\lambda$ where $\\lambda$ is the wavelength. The phase shift therefore depends on the wavelength $\\Delta\\varphi=\\Delta\\varphi(\\lambda).$ In order to obtain a wavelength-independent phase shift of $\\pi$ at 0 OPD (optical path difference), an achromatic phase shifter (APS) has to be introduced into the setup. (An alternative approach, using an ``adaptive nuller'', has been studied at the JPL\\cite{Proc_Peters_2006}.)  Bracewell's nulling interferometer includes an APS as an optical means of eliminating the on-axis flux (stellar light), thus facilitating the detection of the off-axis flux (containing light from extrasolar planets). From the very inception, \\emph{Darwin}/TPF-I have opted for a Bracewell-based design, and therefore need an effective APS.   Many different concepts of achromatic phase shifters are available in the literature.\\cite{Report_Rabbia_2001,Proc_Rabbia_2000} Their merits need to be evaluated quantitatively. The \\emph{Darwin} collaboration studied ten APS concepts\\cite{ESA_APS}, promising $nl<10^{-6}$. The spectral range was that of $6 - 20\\,\\mu\\mathrm{m}$, with $6 - 18\\,\\mu\\mathrm{m}$ mandatory, the extension to $18 - 20\\,\\mu\\mathrm{m}$ priority number 2, and that to $4 - 6\\,\\mu\\mathrm{m}$ priority number 3. Four concepts were selected for further study:   \\paragraph{1. Dispersive prisms }  (Fig. \\ref{fig:APS_4} a). The principle of this method consists in introducing dielectric plates into each arm of the interferometer. Their number, composition and thickness are optimised, together with the OPD, so that the chromaticity of the resulting phase shift in a given spectral band is below a specified level. As we shall see in the description of {\\sc Synapse} (Sec. \\ref{sec:synapse}), a simple APS using two prisms of the same material in each arm of the interferometer is an integral part of the setup.  A model with three prisms of three different materials per interferometer arm is under development by Thales Alenia Space.  \\begin{figure} \\begin{center} \\subfigure[Dielectric prisms]{\\includegraphics[width=0.49\\textwidth]{Lames_dispersives.eps}} \\subfigure[Field reversal]{\\includegraphics[width=0.49\\textwidth]{Field_Reversal_Periscopes.eps}}  \\subfigure[Focus crossing]{\\includegraphics[width=0.49\\textwidth]{APS_FC_principle.eps}} \\subfigure[Fresnel's rhombs]{\\includegraphics[width=0.49\\textwidth]{fresnel_rhombs_new.eps}} \\end{center}   \\caption{\\label{fig:APS_4}The four APS concepts under study for \\emph{Darwin}.} \\end{figure}    \\paragraph{2. Field reversal at reflection}  (Fig. \\ref{fig:APS_4} b). From two reflections, arranged in such a way that s and p polarisation components are successively permuted, the flip of field vectors provides an achromatic phase shift, as well as a pupil rotation of $\\pi$. Suitably applied to a couple of interfering beams, a couple of periscope yields two lightwaves phase shifted by $\\pi$ (opposed field vectors). Then a modified constructive interferometer provides two nulled outputs by suitably mixing the beams. A prototype APS was developed at Max-Planck-Institut f\\\"ur Astronomie in Heidelberg in collaboration with Kayser-Threde GmbH in Munich and the IOF Fraunhofer Institute for Applied Optics in Jena.\\cite{Launhardt2008SPIE}  \\paragraph{3. Focus crossing}  (Fig. \\ref{fig:APS_4} c). The APS FC is based  on the so-called focus-crossing (FC) effect where light undergoes  a phase shift of $\\pi$ when crossing a focus. The phenomenon is independent of wavelength\\cite{Art_Gouy,Art_Boyd}. The principle of this approach is based on destructive interference. It is quite similar to the one of the Achromatic Interfero Coronagraph\\cite{Art_Gay_Rabbia,Art_Baudoz_2000a,Exp_Art_Baudoz_2000b, Proc_Rabbia_2006}.  APS FC comprises, one bearing two confocal half-parabolic off-axis mirrors, the other, two flat mirrors. The ``confocal'' module produces an achromatic  $\\pi$-phase shift, a pupil rotation by $\\pi$ and an extra optical path (twice the focal length of a paraboloid), while the role of the ``flat-flat'' module is to balance this pathlength and to reproduce the beam geometry of the confocal module. At beam combination, the resulting amplitude is nulled for an on-axis point-like source, provided the wavefronts are perfectly cophased. If not, nulling is incomplete and a residual flux remains.   \\paragraph{4. Fresnel's rhombs}  (Fig. \\ref{fig:APS_4} d). are based on the phase shift at total internal reflection. It occurs when the incidence at material/air interface exceeds a limit angle $\\vartheta_{\\lim}$. A linear polarisation with azimuth $\\pi/4$ at entry, becomes a circular polarisation at exit (phase shift by $\\pi/2$ between s and p components). Two cascaded rhombs then yield a $\\pi$ phase shift. A specific disposition of two coupled rhombs allows working with any azimuth of the polarisation at entry (dual polarisation is then allowed).  Centre Spatial de Li\\`ege has been working on a concept combined with sub-lambda (subwavelength, i.e., zero order) gratings on the active surfaces of the rhombs.\\cite{Mawet2006SPIE,Mawet2007OpticsExpress} ",
        "conclusions": "\\label{sec:disc_n_conc}   The primary objective of this paper was to present an update on our progress with the {\\sc Synapse} test bench.  Our work has so far demonstrated that the setup reaches a well reproducible monochromatic nulling level $nl=10^{-5}$ at $3.39\\,\\mu\\mathrm{m}$ which can be maintained for 10 minutes, thus reaching stability of $\\sigma_{\\left\\langle nl\\right\\rangle } (3.39\\,\\mathrm{\\mu m},100\\,\\mathrm{s})\\leq 4\\:10^{-7}$. The reproducible nulling level in the K band is currently of $nl=3\\:10^{-4}$.  We have verified that it can be maintained for at least several hours, reaching stability levels of $\\sigma_{\\left\\langle nl\\right\\rangle } (2.2\\,\\mathrm{\\mu m},2\\:\\mathrm{hrs})= 8\\:10^{-6}$.   Regarding evaluation of APS prototypes, at the moment, the bench allows for work with its ``intrinsic'' APS (a pair of dielectric prisms per branch) and with the APS FC.  The current primary objective of experiments with the APS FC is to provide us with more options for our diagnostic tests. The two APS concepts are in fact very different (one flips the field of view, while the other does not, one works only in a set spectral band outside of which nulling levels deteriorate fast, the other based on a physical principle which is independent of wavelength), and therefore using both gives us more information about the bench itself.   The APS FC was initially built to work in the spectral range of $6-18\\,\\mu\\mathrm{m}$ (proposed range for \\emph{Darwin}/TPF-I), and its expected nulling performance is  $nl\\leq 10^{-6}$, depending on the bench configuration. In the working conditions of {\\sc Synapse} the expected value is $10^{-5}$. Up to now this performance has been demonstrated with laser light at $3.39\\,\\mu\\mathrm{m}$. Rejection performance of the APS FC  in the K band ($2-2.5\\,\\mu\\mathrm{m}$) is to date limited by the performance of the {\\sc Synapse} bench itself.  Although evaluation of the performance of different APS's may be done using fringe dispersion or a monochromator (phase-shift measurement as a function of wavelength), clearly the most representative type of APS estimate calls for measurements with broad-band polychromatic light.  We are making progress towards eliminating all possible limiting factors to {\\sc Synapse} broad-band performance. The work is ongoing and we shall present an update at the SPIE Congress."
    },
    "0812/0812.4465_arXiv.txt": {
        "abstract": "As we resolve ever smaller structures in the solar atmosphere, it has become clear that magnetism is an important component of those small structures. Small-scale magnetism holds the key to many poorly understood facets of solar magnetism on all scales, such as the existence of a local dynamo, chromospheric heating, and flux emergence, to name a few. Here, we review our knowledge of small-scale photospheric fields, with particular emphasis on quiet-sun field, and discuss the implications of several results obtained recently using new instruments, as well as future prospects in this field of research. ",
        "introduction": "\\label{sec:intro} Magnetism on the sun occurs on all scales. It manifests itself at the largest scales as a mean-field component that covers an entire hemisphere, and on progressively smaller scales as active regions, sunspots, and pores. Magnetic field in the lower solar atmosphere has structure on the smallest observable scales, up to the diffraction limit of the best telescopes. Theoretical arguments and simulations indicate that there is structure well beyond what can be observed today or in the forseeable future.  A comprehensive ab-initio model of magnetic activity is currently impossible from a practical standpoint, and will remain so in the near future. The complex interaction of magnetic field, hydrodynamics, and radiative transfer requires sophisticated numerical analysis. A simulation would have to cover a substantial surface area over a good fraction of the convection zone in order to capture large-scale patterns such as supergranulation, yet also have sufficient resolution to capture interactions on scales of several kilometers or less. Such a simulation is prohibitively expensive in terms of computation time. In order to gain understanding of the physical processes involved in the creation, evolution, and eventual destruction of magnetic field, we must turn to observations to study the properties of magnetic structures.  Granular flows in the photosphere expunge field from cell interiors. Flux is swept into the intergranular lanes, where it clumps in small concentrations of mostly vertical field with strengths in excess of one kilogauss. Bright points and faculae, the most conspicuous features of magnetism in the lower solar atmosphere, correspond to these small concentrations of field. They are well-known in active regions, where they group together in plages. In the quiet sun, supergranular flows concentrate them in the magnetic network that incompletely outlines supergranular cells.  Internetwork quiet-sun magnetism has been somewhat ignored historically, largely due to a lack of observations with sufficient resolution and accuracy. However, it has attracted particular interest in the past years, and this subject is currently being studied vigorously. Isolated concentrations of strong field that produce bright points and faculae also exist in supergranular interiors. Outside the concentrations, much weaker field that is not predominantly oriented perpendicular to the surface is ubiquitously present. This more horizontal field typically does not produce bright points that are easily observed using proxy-magnetometry diagnostics such as imaging in the Fraunhofer G band. Instead, sensitive magnetometers are required to observe and study weak field. The development of new instrumentation and seeing-mitigating techniques (the SpectroPolarimeter instrument on the space-borne observatory Hinode is an excellent example of both), and advanced simulations facilitated by the steadily increasing processing power of computers have made it possible to study this subject in detail.  Here, we focus our attention on magnetic fine structure of the quiet solar photosphere. In particular, we will discuss internetwork field. Quiet-sun internetwork areas cover the majority of the solar surface. Four orders of magnitude more flux emerges in the internetwork than in active regions. Consequently, field in these areas is of importance in understanding certain aspects of solar magnetic activity, such as the existence and workings of a granular dynamo and the dynamical coupling of the photosphere to higher layers.  We aim to provide a comprehensive review of quiet-sun internetwork magnetic fine structure, starting with a general overview in Sect.~\\ref{sec:qsmf}. We then address several small-scale phenomena that have recently attracted particular attention as a result of new, high-resolution observations: properties of horizontal field in the photosphere (Sect.~\\ref{sec:thmf}), polar field (Sect.~\\ref{sec:pf}), and concentrations of strong vertical field (Sect.~\\ref{sec:bpme}). Section~\\ref{sec:urmf} concludes the chapter with a discussion on unresolved fields. ",
        "conclusions": "\\label{sec:conclusion} While we have attempted to give a comprehensive overview of small-scale magnetic field in the solar atmosphere, this review is by no means complete. In particular, we have not touched upon chromospheric fields, which besides being structured on small scales, also display dynamic behaviour on short timescales.  New instruments that are able to measure photospheric magnetic field in high-resolution and at sufficient cadence to study dynamics, either directly through (spectro)polarimetry, or indirectly through proxy-magnetometry, are now availble to observers. In particular, we have discussed several results from the Hinode mission. These results, while dealing with small-scale field, have great repercussions for important questions surrounding magnetism in the sun, and in particular for the existence and workings of both the local and global solar dynamos.  With new instruments and sophisticated modeling enabled by advances in computing, we have greatly improved our understanding of magnetic activity on all scales in the sun. Yet, we have also seen that the end is not yet in sight: field is likely structured on scales well beyond what can be observed or simulated today or in the forseeable future. Our understanding of the processes that give rise to small-scale magnetic field will continue to improve as more observations are analyzed, models become more sophisticated and lifelike, and new instruments are developed, such as the Sunrise baloon-borne observatory \\citep{Gandorfer2006} or the Advanced Technology Solar Telescope \\citep{Keil2000}."
    },
    "0812/0812.1847_arXiv.txt": {
        "abstract": "A full particle simulation study is carried out on the electron acceleration at a collisionless, relatively low Alfven Mach number ($M_A=5$), perpendicular shock. Recent self-consistent hybrid shock simulations have demonstrated that the shock front of perpendicular shocks has a dynamic rippled character along the shock surface of low-Mach-number perpendicular shocks. In this paper, the effect of the rippling of perpendicular shocks on the electron acceleration is examined by means of large-scale (ion-scale) two-dimensional full particle simulations. It has been shown that a large-amplitude electric field is excited at the shock front in association with the ion-scale rippling, and that reflected ions are accelerated upstream at a localized region where the shock-normal electric field of the rippled structure is polarized upstream. The current-driven instability caused by the highly-accelerated reflected ions has a high growth rate to large-amplitude electrostatic waves. Energetic electrons are then generated by the large-amplitude electrostatic waves via electron surfing acceleration at the leading edge of the shock transition region. The present result suggests that the electron surfing acceleration is also a common feature at low-Mach-number perpendicular collisionless shocks. ",
        "introduction": "Collisionless shocks are universal processes in space and are observed in laboratory, astrophysical, and space plasmas, including astrophysical jets, an interstellar medium, the heliosphere, and planetary magnetospheres. Particle acceleration is a common feature of at collisionless shocks, but is still a major unresolved issue. The most plausible mechanism of the acceleration is the diffusive shock acceleration (DSA), which explains broadband power-law spectrum with an index around 2 \\citep{drury1983,blandford1987}. Before the DSA phase in which electrons cross the shock front many times, they have to be pre-accelerated by unknown injection mechanism. Such ``injection problem'' is still unresolved.   One of the possible injection mechanisms for electrons is the electron surfing acceleration by electrostatic fields \\citep{Katsouleas_1983,Shimada_2000,Hoshino_2002}. \\citet{Shimada_2000} performed a one-dimensional (1D) electromagnetic full particle simulation and found the formation of large-amplitude electrostatic solitary structures during the cyclic reformation of a high-Mach-number perpendicular shock in a low-beta and weakly magnetized plasma. The coherent solitary structures trap electrons in their electrostatic potential well, resulting in significant surfing acceleration of electrons.   So far, the electron surfing acceleration has been argued for high-Mach-number shocks, and studied by many authors \\citep[e.g.,][]{McClements_2001,Hoshino_2002, Schmitz_2002a,Schmitz_2002b,Amano_2007}. A typical astrophysical example of such high-Mach-number shocks is the young supernova remnant (SNR), which is thought to be an accelerator of high-energy electrons, emitting bright synchrotron radiation in radio and X-ray bands at shocks with Mach number of ${\\cal M}_A\\gtrsim10^2$ \\citep[e.g.,][]{koyama1995,bamba2003,bamba2005}. However, in many astrophysical environments, the electron acceleration at lower-Mach-number (${\\cal M}_A\\lesssim10^2$) shocks is inferred from observations and/or expected theoretically. Possible examples are old SNRs or old SNRs interacting with giant molecular clouds \\citep{claussen1997,chevalier1999,bykov2000,yamazaki2006,brogan2006}, merging galaxy clusters \\citep{markevitch2002}, AGN outbursts at the galaxy clusters \\citep{fujita2007}, and so on. The electron acceleration is also observed at Earth's bow shock even with a low Mach number of ${\\cal M}_A\\sim6.4$ \\citep{oka2006}. Hence it is interesting to ask whether the electron surfing acceleration can occur at lower-Mach-number shocks or not.    It has been demonstrated by one-dimensional full particle simulations that the electron surfing acceleration takes place at a relatively-high-Mach-number shock of $M_A\\sim10$ but not at a very-low-Mach-number shock of $M_A\\sim3$ \\citep[e.g.,][]{Shimada_2000,Schmitz_2002a}. \\footnote{ Note that in the full particle simulations, a reduced ion-to-electron mass ratio of $m_i/m_e=20$--25, is usually adopted for computational efficiency. In general, different mass ratio sometimes changes the simulation results qualitatively, and the simulation results with a reduced mass ratio cannot be directly compared with astrophysical phenomena or results of hybrid simulations. % Therefore, to avoid a confusion, we use a notation $M_A$ for the Mach number of shocks with a reduced mass ratio, while we use another notation ${\\cal M}_A$ for the real mass ratio. } However, these previous works on the electron surfing acceleration are mainly based on 1D simulations, in which configuration of shock magnetic fields cannot be modified.   It is well known from resuts of two-dimensional hybrid-code simulations that there exist large-amplitude fluctuations in the magnetic field and density of the shock transition region of lower-Mach-number shocks (${\\cal M}_A=5\\sim10$) \\citep{Winske_1988}. These fluctuations at the shock surface have been analyzed in terms of turbulent ``ripples'' \\citep{Lowe_2003,Burgess_2006a}. \\citet{Burgess_2006b} has demonstrated, using a combination of self-consistent hybrid simulation and test particle calculation, that energetic electrons are produced by both magnetic mirroring by the rippled structure and by stochastic acceleration by magnetic fluctuations keeps particles within the shock transition region. However, electron-scale microscopic instabilities, i.e., current driven instabilities are neglected in the combination of hybrid simulation and test-particle calculation. Hence multi-dimensional full particle simulations are necessary for studying the effect of ion-scale rippling of a perpendicular shock on the electron surfing acceleration.   The purpose of this paper is to examine the effect of the dynamic rippled structuring of the shock front on electron acceleration processes by the first-principle full particle simulation. In order to take into account the rippling of a perpendicular shock, simulation domain is taken to be larger than the ion inertial length.  The paper is organized as follows. Section 2 describes the model and the parameters of the full particle simulation. Section 3 demonstrates the ion-scale structure of a perpendicular shock found in the full particle simulations and the resulting electron acceleration. Finally, section 4 gives summary of this paper. ",
        "conclusions": "We have studied electron acceleration at a low-Mach-number perpendicular collisionless shock by performing two-dimensional full particle simulations. In order to take into account the effect of the rippling of a perpendicular shock \\citep{Lowe_2003,Burgess_2006a,Burgess_2006b}, the simulation domain is taken to be larger than the ion inertial scale by using the shock-rest-frame model \\citep{Umeda_2008}. In the previous works, the electron acceleration by electron-scale microscopic instabilities at low-Mach-number perpendicular shocks has not focused on because the electron surfing acceleration by electrostatic waves \\citep[e.g.,][]{Shimada_2000,Hoshino_2002}, which is one of the possible injection mechanisms for electrons, is thought to be effective only in high-Mach-number perpendicular shocks. The present result suggests, by contrast, that the electron surfing acceleration is also a common feature at a perpendicular shock with $M_A = 5$.  The mechanism for generation of high-energy electrons is quite simple, and is summarized as follows: The perpendicular shock forms the rippled structures by ion temperature anisotropy; The rippled structures excite a strong electric field component in the shock-normal direction; Reflected ions in the shock transition region is strongly accelerated upstream by the electric field component of rippled structures; The strongly-accelerated reflected ions give a high growth rate of a current-driven instability to large-amplitude electrostatic waves in a localized region; Energetic electrons are generated by the electrostatic waves via surfing acceleration. Finally, high-energy non-thermal electrons are generated by both magnetic compression and scattering by ion-scale rippled structure in the shock transition region.  It has been demonstrated that multi-dimensionality sometimes weakens the electron surfing acceleration in the 2D simulations of uniform plasma models \\citep[e.g.,][]{Dieckmann_2006,Ohira_2007} and self-consistent shock models \\citep[e.g.,][]{Dieckmann_2008,Umeda_2008}. Very recently, however, \\citet{Amano_2008} have found the electron surfing acceleration in 2D simulation of a perpendicular shock with out-of-plane magnetic field. In their 2D simulation, the simulation domain is taken to be much larger than the ion inertial length along the shock surface, and kinetic effects of ions are fully included. The present result is another example showing the effect of ion kinetics to the electron surfing acceleration. It has been confirmed that the cross-scale coupling between an ion-scale mesoscopic instability and an electron-scale microscopic instability is important. Hence, a large-scale full particle simulation would be essential for studies of the electron acceleration at collisionless shocks.    Finally let us discuss on the reduced mass ratio. In the present simulation with $M_A=5$, $\\omega_{pe1}/\\omega_{ce1}=10$, and $\\beta = 0.125$, the Buneman-type mode becomes unstable at $\\omega_{UHR}\\simeq \\omega_{pe1}$ when we use a reduced mass ratio of $m_i/m_e=25$. When we use the real mass ratio of $m_i/m_e=1836$, the thermal velocity of upstream electrons becomes about 8.6 times as large as the case of $m_i/m_e=25$, the self-reformation process is suppressed and the Buneman-type mode is also stabilized \\citep{Scholer_2004}. To make the Buneman-type mode unstable the electron thermal velocity must be smaller than the relative velocity between incoming electrons and reflected ions. Thus the present simulation result can be applied to a much lower beta of $\\beta < 0.03$."
    },
    "0812/0812.2243_arXiv.txt": {
        "abstract": "The MiniBooNE Collaboration observes unexplained electron-like events in the reconstructed neutrino energy range from 200 to 475 MeV. With $6.46 \\times 10^{20}$ protons on target, 544 electron-like events are observed in this energy range, compared to an expectation of $415.2 \\pm 43.4$ events, corresponding to an excess of $128.8 \\pm 20.4 \\pm 38.3$ events. The shape of the excess in several kinematic variables is consistent with being due to either $\\nu_e$ and $\\bar \\nu_e$ charged-current scattering or to $\\nu_\\mu$ neutral-current scattering with a photon in the final state. No significant excess of events is observed in the reconstructed neutrino energy range from 475 to 1250 MeV, where 408 events are observed compared to an expectation of $385.9 \\pm 35.7$ events. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.2075_arXiv.txt": {
        "abstract": "Traditional ideas for testing unification involve searching for the decay of the proton and its branching modes. We point out that several astrophysical experiments are now reaching sensitivities that allow them to explore supersymmetric unified theories. In these theories the electroweak-mass dark matter particle can decay, just like the proton, through dimension six operators with lifetime $\\sim 10^{26} $ sec. Interestingly, this timescale is now being investigated in several experiments including ATIC, PAMELA, HESS, and Fermi. Positive evidence for such decays may be opening our first direct window to physics at the supersymmetric unification scale of $M_{GUT} \\sim10^{16}$ GeV, as well as the TeV scale.  Moreover, in the same supersymmetric unified theories, dimension five operators can lead a weak-scale superparticle to decay with a lifetime of $\\sim100$ sec. Such decays are recorded by a change in the primordial light element abundances and may well explain the present discord between the measured Li abundances and standard big bang nucleosynthesis, opening another window to unification. These theories make concrete predictions for the spectrum and signatures at the LHC as well as Fermi. ",
        "introduction": "One of the most interesting lessons of our times is the evidence for a new fundamental scale in nature, the grand unification (GUT) scale near $M_\\text{GUT} \\sim 10^{16} ~\\GeV$, at which, in the presence of superparticles, the gauge forces unify \\cite{Dimopoulos:1981yj, Dimopoulos:1981zb}.  This is a precise, quantitative (few percent) concordance between theory and experiment and one of the compelling indications for physics beyond the Standard Model.  Together with neutrino masses \\cite{GellMann:1976pg, Yanagida:1980xy}, it provides independent evidence for new physics near $10^{16} ~\\GeV$, significantly below the Planck mass of $M_\\text{pl} \\approx 10^{19} ~\\GeV$.  The LHC may considerably strengthen the evidence for grand unification if it discovers superparticles.  Furthermore, future proton decay experiments may provide direct evidence for physics at $\\Mgut$.  In this paper we consider frameworks in which GUT scale physics is probed by cosmological and astrophysical observations.  In grand unified theories the proton can decay because the global baryon-number symmetry of the low energy Standard Model is broken by GUT scale physics.  Indeed, only local symmetries can guarantee that a particle remains exactly stable since global symmetries are generically broken in fundamental theories.  Just as the proton is long-lived but may ultimately decay, other particles, for example the dark matter, may decay with long lifetimes.  If a TeV mass dark matter particle decays via GUT suppressed dimension 6 operators, its lifetime would be \\begin{equation} \\label{Eqn: dim 6 lifetime} \\tau \\sim 8 \\pi \\frac{\\Mgut^4}{m^5} = 3 \\times 10^{27} ~\\s \\left( \\frac{\\TeV}{m} \\right)^5 \\left( \\frac{\\Mgut}{2 \\times 10^{16} ~\\GeV} \\right)^4 \\end{equation} Similarly a long-lived particle decaying through dimension 5 GUT suppressed operators has a lifetime \\begin{equation} \\tau \\sim 8 \\pi \\frac{\\Mgut^2}{m^3} = 7 ~\\s \\left( \\frac{\\TeV}{m} \\right)^3 \\left( \\frac{\\Mgut}{2 \\times 10^{16} ~\\GeV} \\right)^2 \\end{equation} Both of these timescales have potentially observable consequences.  The dimension 6 decays cause a small fraction of the dark matter to decay today, producing potentially observable high energy cosmic rays.  The dimension 5 decays happen during Big Bang Nucleosynthesis (BBN) and can leave their imprint on the light element abundances. There is, of course, uncertainty in these predictions for the lifetimes because the physics at the GUT scale is not known.  If the dark matter decays through dimension 6 GUT suppressed operators with a lifetime as in Eqn. \\ref{Eqn: dim 6 lifetime} it can produce high energy photons, electrons and positrons, antiprotons, or neutrinos.  Interestingly, the lifetime of order $10^{27} ~\\s$ leads to fluxes in the range that is being explored by a variety of current experiments such as HESS, MAGIC, VERITAS, WHIPPLE, EGRET, WMAP, HEAT, PAMELA, ATIC, PPB-BETS, SuperK, AMANDA, Frejus, and upcoming experiments such as the Fermi (GLAST) gamma ray space telescope the Planck satellite, and IceCube, as shown in Table \\ref{Tab: astro limits}.  This is an intriguing coincidence, presented in section 2, that may allow these experiments to probe physics at the GUT scale, much as the decay of the proton and a study of its branching ratios would. Possible hints for excesses in some of these experiments may have already started us on such an  exciting path.   \\begin{table} \\begin{center} \\begin{math} \\begin{footnotesize} \\begin{array}{|c|c|c|c|c|c|} \\hline & \\text{Extragalactic } \\gamma \\text{-rays} & \\text{Galactic } \\gamma \\text{'s} & \\text{antiprotons} & \\text{positrons} & \\text{neutrinos} \\\\ \\text{Decay} & & & & & \\text{Super-K} \\\\ \\text{channel} & \\text{EGRET} & \\text{HESS} & \\text{PAMELA} & \\text{PAMELA} & \\text{AMANDA, Frejus}\\\\ \\hline q \\overline{q} & 4 \\times 10^{25} ~\\s & - & 10^{27} ~\\s & - & - \\\\ \\hline e^+ e^- & 8 \\times 10^{22} ~\\s & 2 \\times 10^{22} ~\\s ~ \\sqrt{\\frac{m_\\psi}{\\TeV}} \\left( \\text{K} \\right) & 10^{24} ~\\s & 2 \\times 10^{25} ~\\s \\left( \\frac{\\TeV}{m_\\psi} \\right) & 3 \\times 10^{21} ~\\s ~ \\left( \\frac{m_\\psi}{\\TeV} \\right)\\\\ \\hline \\mu^+ \\mu^- & 8 \\times 10^{22} ~\\s & 2 \\times 10^{22} ~\\s ~ \\sqrt{\\frac{m_\\psi}{\\TeV}} \\left( \\text{K} \\right) & 10^{24} ~\\s & 2 \\times 10^{25} ~\\s \\left( \\frac{\\TeV}{m_\\psi} \\right) & 3 \\times 10^{24} ~\\s ~ \\left( \\frac{m_\\psi}{\\TeV} \\right) \\\\ \\hline \\tau^+ \\tau^- & 10^{25} ~\\s & 10^{22} ~\\s ~ \\sqrt{\\frac{m_\\psi}{\\TeV}} \\left( \\text{K} \\right) & 10^{24} ~\\s & 10^{25} ~\\s \\left( \\frac{\\TeV}{m_\\psi} \\right) & 3 \\times 10^{24} ~\\s ~ \\left( \\frac{m_\\psi}{\\TeV} \\right) \\\\ \\hline W W & 3 \\times 10^{25} ~\\s & -  & 3 \\times 10^{26} ~\\s  & 4 \\times 10^{25} ~\\s & 8 \\times 10^{23} ~\\s ~ \\left( \\frac{m_\\psi}{\\TeV} \\right)  \\\\ \\hline & 9 \\times 10^{24} ~\\s \\left( m_\\psi = 100 ~ \\GeV \\right) & 2 \\times 10^{24} ~\\s ~ \\sqrt{\\frac{m_\\psi}{\\TeV}} \\left( \\text{K} \\right) & &  & \\\\ \\gamma \\gamma & 2 \\times 10^{22} ~\\s \\left( m_\\psi = 800 ~ \\GeV \\right) &  & 2 \\times 10^{25} ~\\s & 8 \\times 10^{23} ~\\s \\left( \\frac{\\TeV}{m_\\psi} \\right) & - \\\\ & 4 \\times 10^{23} ~\\s \\left( m_\\psi = 3200 ~ \\GeV \\right) & 5 \\times 10^{25} ~\\s ~ \\sqrt{\\frac{m_\\psi}{\\TeV}} \\left( \\text{NFW} \\right) & &  & \\\\ \\hline \\nu \\overline{\\nu} & 8 \\times 10^{22} ~\\s & - & 10^{24} ~\\s & 10^{23} ~\\s & 10^{25} ~\\s ~ \\left( \\frac{m_\\psi}{\\TeV} \\right) \\\\ \\hline \\end{array} \\end{footnotesize} \\end{math} \\caption[Astrophysical Limits on Decaying Dark Matter]{\\label{Tab: astro limits} The lower limit on the lifetime of a dark matter particle with mass in the range $10 ~\\GeV \\lesssim m_\\psi \\lesssim 10 ~\\TeV$, decaying to the products listed in the left column.  The experiment and the observed particle being used to set the limit are listed in the top row.  HESS limits only apply for $m_\\psi > 400 ~\\GeV$ and are shown for two choices of halo profiles: the Kravtsov (K) and the NFW.  PAMELA limits are most accurate in the range $100 ~\\GeV \\lesssim m_\\psi \\lesssim 1 ~\\TeV$.  All the limits are only approximate.  Generally conservative assumptions were made and there are many details and caveats as described in Section \\ref{Sec: Astro Limits}.} \\end{center} \\end{table}    GUT scale physics can also manifest itself in astrophysical observations by leaving its imprint on the abundances of light elements created during BBN.  For example neutrons from the decay of a heavy particle create hot tracks in the surrounding plasma in which additional nucleosynthesis occurs.  In particular, these energetic neutrons impinge on nuclei and energize them, causing a cascade of reactions.  This most strongly affects the abundances of the rare elements produced during BBN, especially $^6\\Li$ and $^7\\Li$ and possibly $^9\\Be$.  In fact, measurements of both isotopes of $\\Li$ suggest a discrepancy from the predictions of standard BBN (sBBN).  The observed $\\LiSeven$ abundance of $\\frac{^7\\Li}{\\text{H}} \\sim (1 - 2) \\times 10^{-10}$ is a factor of several below the sBBN prediction of $\\frac{^7\\Li}{\\text{H}} \\approx ( 5.2 \\pm 0.7 ) \\times 10^{-10}$ \\cite{Cyburt:2008kw}.  In contrast, observations indicate a primordial $\\LiSix$ abundance over an order of magnitude above the sBBN prediction.  The Lithium abundances are measured in a sample of low-metallicity stars.  The $\\Li$ isotopic ratio in all these stars is similar to that in the lowest metallicity star in the sample: $\\frac{^6\\Li}{^7\\Li} = 0.046 \\pm 0.022$ \\cite{Asplund:2005yt}.  This implies a primordial $\\LiSix$ abundance in the range $\\frac{^6\\Li}{\\Hy} \\approx (2 - 10) \\times 10^{-12}$, while sBBN predicts $\\frac{^6\\Li}{\\Hy} \\approx 10^{-14}$ \\cite{Jedamzik:2007cp, VangioniFlam:1998gq}.  The apparent presence of a Spite plateau in the abundances of both $\\LiSix$ and $\\LiSeven$ as a function of stellar temperature and metallicity is an indication that the measured abundances are indeed primordial.  Of course, though there is no proven astrophysical solution, either or both of these anomalies could be due to astrophysics and not new particle physics.  Nevertheless, the Li problems are suggestive of new physics because a new source of energy deposition during BBN naturally tends to destroy $\\LiSeven$ and produce $\\LiSix$, a nontrivial qualitative condition that many single astrophysical solutions do not satisfy.  Further, energy deposition, for example due either to decays or annihilations, during BBN most significantly affects the Li abundances, not the other light element abundances, making the Li isotopes the most sensitive probes of new physics during BBN.  Finally, a long-lived particle decaying with a $\\sim 1000$ s lifetime can naturally destroy the correct amount of $\\LiSeven$ and produce the correct amount of $\\LiSix$ without significantly altering the abundances of the other light elements $\\D$, $^3\\He$, and $^4\\He$.   In this paper we explore astrophysical signals of GUT scale physics in the framework of supersymmetric unified theories (often referred to as SUSY GUTs or supersymmetric standard models (SSM) \\cite{Dimopoulos:1981zb}), as manifested by particle decays via dimension 5 and 6 GUT suppressed operators.  In order to preserve the success of gauge coupling unification, we work with the minimal particle content of the MSSM with additional singlets or complete SU(5) multiplets.  The effects of GUT physics on a low energy experiment are generally suppressed. However, these effects can accumulate and lead to vastly different physics over long times depending on the details of the higher dimension operators generated by GUT physics. For example, particles that would have been stable in the absence of GUT scale physics can decay with very different lifetimes and decay modes depending on the particular GUT physics. Conversely, such decays are sensitive diagnostics of the physics at the GUT scale. The details of these decays depend both upon the physics at the GUT scale and the low energy MSSM spectrum of the theory. So, astrophysical observations of such decays, in conjunction with independent measurements of the low energy MSSM spectrum at the LHC, would open a window to the GUT scale - just as proton decay would.  The role of the water and photomultipliers that register a proton decay event  are now replaced by the universe and either the modified Li abundance, or the excess cosmic rays that may be detected in todays plethora of experiments.   To characterize the varieties of GUT physics that can give rise to decays of would-be-stable particles we enumerate the possible dimension 5 and 6 operators, each one of which defines a separate general class of theories that breaks a selection rule or conservation law that would have stabilized the particle in question. This approach has the advantage of being far more general than a concrete theory and encompasses all that can be known from the limited low-energy physics experiments available to us. In other words, all measurable consequences depend on the form of the operator and not on the detailed microphysics at the unification scale (e.g. the GUT mass particles) that give rise to it.  The presence of supersymmetry, together with gauge symmetries and Poincare invariance, and the simplicity of the near-MSSM particle content greatly reduces the number of possible higher dimension operators of dimension 5 and 6 and allows for the methodic enumeration of the operators and the decays they cause. This is what we do in Sections \\ref{Sec: Dim 6 decays}, \\ref{LithiumIntroduction}, and \\ref{LithiumPAMELA}.  We work in an SU(5) framework because any of the SU(5) invariant operators we consider can be embedded into invariant operators of any larger GUT gauge group so long as it contains SU(5).  Of course, such operators may in general contain several SU(5) operators so the detailed conclusions can be affected.  As an example, we study a model that is made particularly simple in an SO(10) GUT in Section \\ref{LithiumPAMELA}.  The dimension 6 operators of Section \\ref{Sec: Dim 6 decays} have many potentially observable astrophysical signals at experiments searching for gamma rays, positrons, antiprotons or neutrinos.  We separate them into R-conserving and R-breaking classes and in each case we also build UV models which give rise to these operators. The dimension 5 operators of Section \\ref{LithiumIntroduction} have potentially observable effects on BBN, may solve the Lithium problems, and give dramatic out-of-time decays in the LHC detectors. Section \\ref{LithiumPAMELA} presents frameworks and simple theories that have both dimension 5 and 6 operators destabilizing particles to lifetimes of both 100 sec and $10^{27}$ sec to explain Li and lead to astrophysical signatures today in PAMELA/ATIC and other observatories.  In Section \\ref{Sec: Astro signals} we similarly look at the consequences for Fermi/Glast and PAMELA/ATIC.  In Section \\ref{Sec: LHC signals} we outline LHC-observable consequences of some of these scenarios (operators) that could lead to their laboratory confirmation.  Finally, there are two more classes of theories that fit into the elegant framework of supersymmetric unification. The first of these are SUSY theories that solve the strong CP problem with an axion. The other is the Split SUSY framework. Both these frameworks provide particles that can have lifetimes long enough to account for the primordial lithium discrepancies without additional inputs, and are included in Section \\ref{LithiumIntroduction}. ",
        "conclusions": "In broad classes of theories, global symmetries that appear in nature, such as baryon number, are accidents of the low energy theory and are violated by GUT scale physics.  If the symmetry stabilizing a particle, such as the proton, is broken by short-distance physics then it will decay with a long lifetime.  In a sense, a particle is continually probing the physics that allows it to decay.  Thus, though the effects of GUT scale physics are suppressed at low energies, this suppression is compensated for by the long time scales involved.  Intriguingly, short-distance physics can be unveiled by experiments sensitive to long lifetimes.  Astrophysics and cosmology provide natural probes of long lifetimes.  The dark matter particle appears stable, but dimension 6 GUT suppressed operators can cause it to decay with a long lifetime $\\sim 10^{26} ~\\s$.  This can lead to observable signals in a variety of current experiments including PAMELA, ATIC, HESS, and Fermi.  Such observations can also probe the TeV scale physics associated with dark matter annihilations.  However, annihilations cannot proceed through GUT scale particles since such a cross section would be far too small.  Thus, the observation of annihilations can probe TeV physics but does not probe GUT physics directly.  This leads to interesting differences between the expected signals from dark matter decays and annihilations \\cite{MinimalDarkMatter, NimaDougTracyWeiner}.  Decaying dark matter can directly produce photons with an $\\OO(1)$ branching ratio, as seen in Section \\ref{Sec: Dim 6 decays}.  It can also produce light leptons without helicity or p-wave suppression and give signals in the range necessary to explain the PAMELA/ATIC excess without boost factors \\cite{Nardi:2008ix, MaximPospelov, YanagidaOne, YanagidaTwo, YanagidaThree}.  Additionally, in our framework the hadronic branching fractions can be naturally suppressed since sleptons tend to be lighter than squarks, perhaps explaining the lack of an antiproton signal at PAMELA. We have also found that the signals produced by decays may show many interesting features beyond just a bump in the spectrum of electrons and positrons. For example, the spectrum from decays may accommodate the secondary feature visible in the ATIC data at energies between 100 - 300 GeV.  The SUSY spectrum can allow for the direct production of a slepton, lepton pair from the decaying dark matter.  In this case, the subsequent decay of the slepton provides a secondary hot lepton that can give rise to the observed secondary ATIC feature.  It is also common in our framework to have two dark matter particles, scalar and fermion superpartners, that both decay at late times producing interesting features in the observed spectra of electrons or photons.  Such models will be tested at several upcoming experiments.  The spectrum of electrons and positrons will be measured to increasing accuracy by Fermi, HESS, and PAMELA, testing the observed excesses.  These measurements could potentially reveal not just the nature of dark matter, but also the mass spectrum of supersymmetry and GUT scale physics.  Excitingly, Fermi and ground-based telescopes such as HESS or MAGIC will be providing measurements of the gamma-ray spectrum in the near future.  These measurements have the potential to test many of our models, and in particular may be able to observe final state radiation from models that explain the electron/positron excess.  Further, they could distinguish decaying and annihilating dark matter scenarios by the different angular dependence of the gamma ray signals.  The supersymmetric GUT theories in which the dark matter decays through dimension 6 operators often have TeV mass particles that decay via dimension 5 operators with a lifetime of $\\sim$ 1000 sec.  This results in relic decays during BBN which could explain the observed primordial Lithium abundances.  This TeV mass particle can be produced at the LHC and if charged or colored will stop and decay inside the detector within $\\sim$ 1000 sec.  This will allow the measurement of its properties and directly make the connection with primordial nucleosynthesis."
    },
    "0812/0812.4925_arXiv.txt": {
        "abstract": "Radio synchrotron emission, its polarization and its Faraday rotation are powerful tools to study the strength and structure of interstellar magnetic fields. The total intensity traces the strength and distribution of total magnetic fields. Total fields in gas-rich spiral arms and bars of nearby galaxies have strengths of 20--30~$\\mu$Gauss, due to the amplification of turbulent fields, and are dynamically important. In the Milky Way, the total field strength is about $6~\\mu$G near the Sun and several $100~\\mu$G in filaments near the Galactic Center. -- The polarized intensity measures ordered fields with a preferred orientation, which can be regular or anisotropic fields. Ordered fields with spiral structure exist in grand-design, barred, flocculent and even in irregular galaxies. The strongest ordered fields are found in interarm regions, sometimes forming ``magnetic spiral arms'' between the optical arms. Halo fields are X-shaped, probably due to outflows. -- The Faraday rotation of the polarization vectors traces coherent regular fields which have a preferred direction. In some galaxies Faraday rotation reveals large-scale patterns which are signatures of dynamo fields. However, in most galaxies the field has a complicated structure and interacts with local gas flows. In the Milky Way, diffuse polarized radio emission and Faraday rotation of the polarized emission from pulsars and background sources show many small-scale and large-scale magnetic features, but the overall field structure in our Galaxy is still under debate. ",
        "introduction": "Interstellar magnetic fields were discovered already in 1932 by Karl Guthe Jansky who first detected diffuse low-frequency radio emission from the Milky Way, but the explanation as synchrotron emission was given only in 1950 by Karl Otto Kiepenheuer. The sensitivity of radio observations has improved by several orders of magnitude in the past decades, and synchrotron emission was detected from the interstellar medium (ISM) in almost all star-forming galaxies, in galaxy halos and the intracluster medium, proving that a large fraction of the Universe is permeated by magnetic fields. However, in spite of our increasing knowledge on interstellar magnetic fields, many important questions are unanswered, especially their first occurrence in young galaxies, their amplification when galaxies evolved, and their effect on galaxy dynamics.  As magnetic fields need illumination by cosmic-ray electrons to become observable by synchrotron emission, which are generated in star-forming regions or intracluster shocks, we do not know yet whether magnetic fields also exist in radio-quiet elliptical or dwarf galaxies or in the general intergalactic medium (IGM). Progress can be expected from using Faraday rotation which does not need cosmic rays, only magnetic fields and thin ionized gas. One of the research areas of the forthcoming radio telescopes (LOFAR, ASKAP, SKA) will be the search for Faraday rotation in these objects against polarized background sources (Gaensler, this volume). The SKA will also be needed to detect magnetic fields in young galaxies (\\cite{beck04,arshakian08}). ",
        "conclusions": ""
    },
    "0812/0812.3028_arXiv.txt": {
        "abstract": "Intermediate-mass black holes (IMBHs), with masses in the range $\\sim 10^{2-4}~M_\\odot$, will be unique sources of gravitational waves for LISA.  Here we discuss their context as well as specific characteristics of IMBH-IMBH and IMBH-supermassive black hole mergers and how these would allow sensitive tests of the predictions of general relativity in strong gravity. ",
        "introduction": "Conclusive evidence exists for black holes in the stellar-mass ($\\sim 3-30~M_\\odot$) and supermassive (SMBH; $\\sim 10^6-10^{10}~M_\\odot$) ranges.  This evidence comes from direct measurements of masses from orbital motion in, respectively, binaries and galactic nuclei.  In contrast, the range between these masses does not yet provide clear dynamical evidence for black holes of intermediate mass (IMBHs).  This is essentially because candidates are rare (hence the nearest possible IMBH in a binary is not easily observed) and have small radii of influence (hence SMBH-like observations of nearby stars are also challenging).  There are, however, numerous indirect suggestions of the $\\sim 10^{2-4}~M_\\odot$ mass range from the high fluxes of ultraluminous X-ray sources (ULXs; see \\cite{Fab89,CM99}), from their relatively low thermal temperatures \\cite{M03,M04a,M04b}, and from aspects of the dynamics of globular clusters \\cite{vdM02,GRH02,NGB08}.  See \\cite{MillerColbert04} for a recent summary of the data.  The lack of definitive dynamical evidence means that alternate scenarios have also been proposed for ULXs, including beamed emission \\cite{Rey97,Ki01} and super-Eddington emission \\cite{RB03,Beg06}.  Here we assume that IMBHs exist and explore various ways in which their interactions could lead to gravitational radiation detectable in the $\\sim 10^{-4}-10^{-2}$~Hz frequency range of the {\\it Laser Interferometer Space Antenna} (LISA). In \\S~2 we discuss the broader context of IMBHs and proposed formation mechanisms.  We also give basic formulae for the amplitude and frequency of gravitational waves from binaries, showing why IMBHs could be bridge sources between space-based and ground-based detectors.  In \\S~3.1 we evaluate the prospects for IMBH-IMBH mergers and what they might tell us.  In \\S~3.2 we discuss IMBH-SMBH mergers and their prospects as uniquely precise testbeds for strong gravity predictions of general relativity.  We present our conclusions in \\S~4. ",
        "conclusions": "The evidence for IMBHs is currently strong but circumstantial, pointing to the need for dynamical mass measurements of binary motion that will establish their existence definitively. Nonetheless, their likely formation mechanisms and dynamical interactions link them to many exciting topics in the current and early universe.  As gravitational wave sources they will be unique in several respects: as bridge objects between space-based and ground-based detectors, as comparable-mass binaries with palpable eccentricities, and as the events that potentially will yield the most precise tests of general relativity.  Many explorations need to be done, but current results are highly encouraging for their study.  \\ack  We thank Tal Alexander, Pau Amaro-Seoane, Mike Gill, Doug Hamilton, Clovis Hopman, Vanessa Lauburg, Fred Rasio, Derek Richardson, and Michele Trenti for stimulating conversations.  This work was supported in part by NASA ATFP grant NNX08AH29G. \\bigskip"
    },
    "0812/0812.3358_arXiv.txt": {
        "abstract": "We report on the results from {\\it Suzaku} broadband X-ray observations  of the galactic binary source {\\LS}. The {\\it Suzaku} data, which have continuous coverage of more than one orbital period, show strong modulation of the X-ray emission at the orbital period of this TeV gamma-ray emitting system. The X-ray emission shows a minimum at orbital phase $\\sim  0.1$, close to the so-called superior conjunction of the compact object, and a maximum at phase $\\sim 0.7$, very close to the inferior conjunction of the compact object. The X-ray spectral data up to 70~keV are described by a hard power-law with a  phase-dependent photon index which varies within $\\Gamma \\simeq 1.45$-- 1.61. The amplitude of the flux variation is a factor of 2.5, but is significantly less than that of the factor $\\sim$8 variation in the TeV flux. Otherwise the two light curves are similar, but not identical. Although  periodic  X-ray emission has been found from many galactic binary systems,  the {\\it Suzaku} result implies a phenomenon  different from the ``standard''  origin of X-rays related to the emission of the hot accretion plasma  formed around the compact companion object. The X-ray radiation of {\\LS} is likely to be linked to very-high-energy electrons which are also responsible for the TeV gamma-ray emission. While the gamma-rays are the result of  inverse Compton scattering by electrons on optical stellar photons, X-rays  are produced via synchrotron radiation. Yet, while the modulation of the TeV gamma-ray signal can be naturally explained by the photon-photon pair production and anisotropic inverse Compton scattering, the observed modulation of synchrotron X-rays requires  an additional process, the most natural one being adiabatic expansion in the radiation production region. ",
        "introduction": "{\\LS} is a high-mass X-ray binary \\citep{Motch97} with extended radio emission \\citep{Paredes00,Paredes02}. This system is formed by a main sequence O type star and a compact object of disputed nature that has been claimed to be both a black hole \\citep[e.g.,][]{casares05} and a neutron star/pulsar \\citep[e.g.,][]{martocchia05,Dubus06}. The compact object is moving around  the companion star in a moderately elliptic orbit (eccentricity $e=0.35$) with an orbital period of $P_{\\rm orb} = 3.9060$ days \\citep{casares05}.  As summarized in \\cite{vbosch07}, LS~5039 has been observed several times in the X-ray energy band for limited phases in the orbital period. Flux variations on time scale of days and sometimes on much shorter timescales have been reported. The spectrum was always well represented by a power-law model with a photon index ranging $\\Gamma=1.4$--1.6 up to $\\sim 10$~keV, with fluxes changing moderately around $\\sim 1\\times 10^{-11}~{\\rm erg}~{\\rm cm}^{-2}~{\\rm s}^{-1}$. Softer spectra and larger fluxes had been also inferred from {\\it RXTE} observations, although background contamination was probably behind these differences (see \\citealt{vbosch05}). Also, {\\it Chandra} data taken in 2002 and 2005 showed spectra significantly harder than 1.5 \\citep{vbosch07}, but such a hard spectrum was probably an artifact produced by photon pile-up. Recently, \\cite{hoffmann} reported the results of {\\it INTEGRAL} observations in hard X-rays. The source was detected at energies between 25 and 60~keV. The source was detected at energies between 25 and 60~keV. The flux was estimated to be  $(3.54 \\pm 2.30)  \\times 10^{-11}~{\\rm erg}~{\\rm cm}^{-2}~{\\rm s}^{-1}$ (90 \\% confidence level) around the inferior conjunction (INFC) of the compact object, and a flux upper limit of $1.45 \\times 10^{-11}~{\\rm erg}~{\\rm cm}^{-2}~{\\rm s}^{-1}$  (90 \\% confidence level)  was derived at the superior conjunction of the compact object (SUPC).  {\\LS} has also been detected in very-high-energy (VHE: $E \\geq 0.1$~TeV) gamma-rays  \\citep{HESS_5039}, exhibiting a periodic signal modulated with the orbital period \\citep{HESS_06}.  There are two other binary systems with robust detections in the VHE range: {\\PSR}  \\citep{HESS_1259} and  {\\LSI} \\citep{MAGIC_LSI,VERITAS_LSI}. Evidence for TeV emission has been found also in Cygnus~X-1 \\citep{MAGIC_CygX1}. {\\PSR} is a clear case of a high-mass binary system containing a non-accreting pulsar \\citep{Johnston92}, whereas Cygnus X-1 is a well known accreting black hole system \\citep{Bolton72}.  The nature of the compact object in {\\LS} is not yet established, and the origin of  the VHE emitting electrons is unclear.   They may be related to a pulsar wind or to a black hole with a (sub)relativistic jet.   In the standard  pulsar-wind scenario, the severe photon-photon absorption  makes the explanation of the detection  of the VHE radiation problematic, at least at the position corresponding  to the orbital phase $\\phi\\sim 0.0$ (see Figs. 16 and 4 from Sierpowska-Bartosik \\&  Torres 2008 and Dubus et al. 2008 and compare with Fig. 5 in Aharonian et al. 2006), in which the emitter is expected to be located between the compact object and the star.  However, particles may be accelerated in a relativistic outflow formed at the interaction of the pulsar and  the stellar winds \\citep{Bogovalov08}  and radiate far from the compact object, making the pulsar wind scenario a viable option.  In the microquasar scenario, the lack of accretion features in the X-ray spectrum may be a  problem unless the bulk of the accretion power is released in the form of kinetic energy of the outflow,  rather than thermal emission during accretion, as in the case  of SS~433 \\citep[see e.g.,][]{Marshall02}.   At this stage, we cannot give a preference to any of these scenarios, but new data, in particular those  obtained with the {\\it Suzaku} satellite, allow us to make an important step towards the understanding of  the nature of the non-thermal processes of acceleration and radiation in this mysterious object. ",
        "conclusions": "The X-ray emission observed with {\\it Suzaku} is characterized by (1) a hard power law with $\\Gamma \\simeq 1.5$ extending from soft X-rays to $\\sim 70$ keV, (2) clear orbital modulation in flux and photon index, (3) a moderate X-ray luminosity of $L_X \\sim 10^{33} (\\frac{D}{1 \\rm kpc})^2\\ \\rm erg\\ s^{-1}$, (4) a small and constant absorbing column density, and (5) a lack of detectable emission lines.  Although variable X-ray emission has been found from more than two hundred  galactic binary systems, the {\\it Suzaku} data hardly can be explained within the ``standard accretion''  scenario where  X-rays are produced by a hot thermal (comptonized)  accretion plasma around the compact object.  The lack of X-ray emission lines (at the level of sensitivity of {\\em Suzaku}) as well as  the hard $E^{-1.5}$ type energy spectrum of  the X-ray continuum, extending from soft X-rays up to 70~keV, favors  a non-thermal origin of the X-rays. This conclusion  is supported by the general similarities between the properties of the observed X-rays and TeV gamma-rays. Namely,  both radiation components  require a rather hard energy distribution of parent electrons with a power-law  index of $\\approx 2$.  This directly follows from the photon index of the synchrotron radiation  $\\Gamma \\approx 1.5$, and agrees quite well with the currently most favored interpretation of the TeV gamma-rays, in which they would be produced by IC scattering off  the anisotropic photon field of the massive companion star.  Assuming that the TeV gamma-ray production region is located at a distance from the companion star of $d \\sim 2\\times 10^{12}$~cm (i.e. the binary system size),  and taking into account that gamma-rays are produced in the deep Klein-Nishina (KN) regime with significantly suppressed cross-section,  for the well known  luminosity of the  optical star $L \\simeq  7 \\times 10^{38}\\,{\\rm erg}\\,{\\rm s}^{-1}$, one can estimate quite robustly the strength of the magnetic field in the emission region. The numerical  calculations show that the field  should be around a few Gauss (see e.g. Fig. 5). For such a magnetic field strength, the  energy intervals of  electrons responsible for the two emission components overlap substantially, as shown in Figure~\\ref{fig:diagram}. Therefore, we are most likely dealing with the same  population of parent electrons, which should be located  at large distances from the compact object, in the system periphery,  to prevent the severe absorption of the TeV radiation and the subsequent intense emission from the pair-created secondaries \\citep{Mitya_LS, vboshch08}.   \\begin{figure}[htbp] \\begin{center} \\epsscale{1.0} \\plotone{f5.eps} \\caption{ The radiative cooling times as a function of electron energy. The blue line corresponds to IC losses at the distance $d=2\\times10^{12}$~cm from the optical star (the star luminously and temperature were assumed to be $L_*=7 \\times 10^{38}$~erg/s and $T_*=3.8\\times10^{4}$~K, respectively). The red line corresponds to the synchrotron losses for B-field $B=3$~G. The filled regions reproduce the electron energy intervals relevant to the HESS (yellow),  {\\it Fermi} (light blue) and {\\it Suzaku}  (green) energy domains. Formally, two radiation channels, synchrotron radiation of very high energy electrons (light green) and IC scattering  of  low energy electrons in the Thompson regime (dark green) can produce X-ray photons in the {\\it Suzaku} energy domain.} \\label{fig:diagram} \\end{center} \\end{figure}   It should be noted that the observed X-ray emission is very difficult to explain as synchrotron emission produced by secondary (pair-produced) electron and positrons. Since the pair production cross-section has strong energy dependence with a distinct maximum, for the  target photons of typical energy of  $\\sim 10$~eV, the major fraction of the absorbed energy will be released in the form of $\\sim 100$~GeV electrons. Thus secondary pair synchrotron emission must show a spectral break in the {\\it Suzaku} energy band unless one assumes unreasonably high magnetic fields, $B \\geq 1$~kG, in the surroundings of the gamma-ray emission region \\citep{Mitya_LS,vboshch08}.  Figure~\\ref{fig:diagram} shows the synchrotron and IC cooling times of electrons, as a function of electron energy, calculated for the  stellar photon  density  at  $d = 2\\times 10^{12}\\ \\rm cm$ and for a magnetic field $B=3$~G. It can be seen that synchrotron losses dominate over IC losses at $E_{\\rm e} \\geq 1$~TeV.  Note that  the  TeV gamma-ray production takes place in the deep KN regime. This implies that the cooling time, $t_{\\rm cool} = \\gamma/\\dot{\\gamma}$, of electrons generating GeV gamma-rays via IC scattering (Thomson regime) is shorter than the cooling time of TeV electrons responsible for producing very high-energy gamma-rays (KN). The same applies for synchrotron cooling time of multi-TeV electrons  that produce low-energy (MeV) gamma-rays by synchrotron radiation.  One should therefore expect significantly higher  MeV (synchrotron) and GeV (IC) fluxes than at keV and TeV energies, provided that the acceleration spectrum of electrons extends from low energies to very high energies.  However,  in the case of  existence of low-energy and very-high-energy cutoffs in the acceleration spectrum, the gamma-ray fluxes $>10$~MeV and at GeV energies would be significantly suppressed.  To better understand the energy ranges of the electrons responsible for X-ray and gamma-ray production, we show in Figure~\\ref{fig:diagram} the energy zones of electrons relevant to the {\\it Suzaku}, {\\it Fermi}, and HESS radiation domains. Note that the reconstruction of the average energy of electrons responsible for the IC gamma-rays depends only on the well known  temperature of the companion star $T=3.8\\times 10^4$~K. The light green zone in Figure~\\ref{fig:diagram}  marked as ``{\\it Suzaku} synchrotron'', corresponds to electrons responsible  for the synchrotron photons produced in the energy interval $1~{\\rm keV}\\le \\epsilon_{\\rm syn}\\le40~{\\rm keV}$. For a reasonable range of magnetic field values, the energy interval of electrons relevant for {\\it Suzaku} data overlap on one hand with the HESS energy interval, and can overlap with the {\\it Fermi} one. This should allow us, in the case of detection of MeV/GeV gamma-rays by {\\it Fermi}, to considerably reduce the parameter space, in particular, to better localize the X- and gamma-ray production regions from electromagnetic cascade constraints, and derive the broadband energy spectrum of electrons and the strength of the magnetic field, both as a function of the orbital phase.  Formally, when X-rays and TeV gamma-rays are produced by the same population of  very-high-energy electrons,  one should expect a general correlation between the light curves obtained  by {\\it Suzaku} and HESS. In this regard, the  similarity between the {\\it Suzaku} and HESS light curves seems to be natural. However, such an interpretation is not straightforward in the sense that two major mechanisms that might cause modulation of the TeV gamma-ray signal are related to interactions of electrons and gamma-rays with the photons of the companion star, i.e. anisotropic IC scattering and photon-photon pair production \\citep{Mitya_LS, Dubus2008}, and thus cannot contribute to the X-ray modulation. The X-ray modulation requires periodic changes of the strength of the ambient magnetic field or the number of relativistic electrons. Note, however, that the change of magnetic  field would not have a strong impact as long as the radiation proceeds in the saturation regime and synchrotron losses dominate in the relevant energy interval. One would also expect modulation of the synchrotron X-ray flux if the energy losses of electrons are dominated by IC scattering, although in such a case we should observe significantly lower X-ray fluxes.  A more natural reason for the modulation of the synchrotron fluxes would come from dominantly adiabatic losses. The adiabatic cooling of electrons  in binary systems can be realized through complex (magneto)hydrodynamical  processes, e.g. due to interactions between a black hole jet or a pulsar wind with the dense stellar wind of a massive companion star (see e.g. \\citealt{Bogovalov08}, \\citealt{perucho08}). The orbital motion could naturally produce the modulation of  adiabatic cooling of electrons around the orbit (see e.g. \\citealt{khangulyan_dublin}). Note that because of the relatively small  variation of the X-ray flux  over the orbit, a factor of only two, the requirements for this scenario are quite modest. We note that dominant adiabatic losses have been invoked by \\cite{Khangulyan2007} to explain the variations of the X-ray and TeV gamma-ray  fluxes from the binary pulsar {\\PSR}.  The detected power-law spectrum of X-rays with photon index around $\\Gamma =1.5$ implies  that the established energy spectrum of electrons is also a power-law with index $\\alpha_{\\rm e}\\simeq 2$. This agrees well with the hypothesis of dominance of adiabatic losses, because the  adiabatic losses do not change the initial spectrum of electrons. Thus the required power-law index $\\alpha_{\\rm e}\\simeq 2$ implies  a reasonable acceleration spectrum $Q(E_{\\rm e})\\propto E_{\\rm e}^{-2}$. Otherwise, in an environment dominated by synchrotron losses, the acceleration spectrum should be very hard, with a power-law index $\\leq 1$, or should have an unreasonably large low-energy cutoff at $E \\geq 1$~TeV to explain the observed X-ray spectra.  Obviously, adiabatic losses modulate the IC gamma-ray flux in a similar manner. However, unlike X-rays, the TeV gamma-rays suffer significant distortion due to photon-photon absorption (see e.g. \\citealt{boettcher08}) and anisotropic IC scattering with its strong hardening of the gamma-ray spectrum \\citep{mitya_felix05}. All this leads to additional orbital modulation of the gamma-ray signal, and it is likely that these two additional processes are responsible for the strong change of gamma-ray flux, much more pronounced than that seen in X-rays (see Fig.~\\ref{fig:light-curve}). The {\\it Suzaku} data  presented in  this paper implies a key additional assumption, namely that the accelerated electrons must loose their energy adiabatically before they cool radiatively.  In order to demonstrate that the suggested scenario can satisfactorily explain the combined  {\\it Suzaku} X-ray and HESS  gamma-ray data,  we performed  calculations of  the broad-band  spectral energy distributions (SEDs) of the synchrotron and IC emission, assuming a simple model in which the same population of electrons is responsible for both X-rays and TeV gamma-rays. We also assumed that the emission region has homogeneous physical conditions. This is a reasonable assumption given that we deal with  very  short  cooling timescales  ($\\ll  100$~s), thus electrons cannot travel significant distances while emitting.  In the regime dominated by  adiabatic energy losses,  the  synchrotron X-ray flux is proportional to $t_{\\rm ad}$. The X-ray modulation seen by {\\it Suzaku} is then described by  the modulation of the adiabatic loss rate. In Fig.~\\ref{fig:phase}, we show $t_{\\rm ad}(\\phi)$ that is inferred from the X-ray data. The required adiabatic cooling timescales are $\\sim 1$~s. Any consistent calculation of the adiabatic cooling requires the solution of the corresponding hydrodynamical problem, and one needs to know in detail the nature of the source. At the present stage, we consider the simple example of adiabatic cooling in a relativistically expanding source. In such a case, the adiabatic loss rate can be written as: $\\dot{\\gamma}_{\\rm ad}(\\phi, \\gamma) = \\gamma / t_{\\rm ad}$ with  $ t_{\\rm ad} \\sim R/c  \\simeq 3 R_{11} $~sec,  where  $R_{11} \\equiv R/(10^{11}\\ \\rm cm)$ is the characteristic size of the source. The required variation of the adiabatic cooling is thus reduced to the modulation of the size of the radiation region ($R_{11} \\sim 0.3\\mbox{--}1$). The size in turn depends on the external pressure exerted by, e.g., the stellar wind from the massive star. The expected weaker external pressure around apastron implicitly assumed in our model would be broadly consistent with the radial dependence of the wind pressure.  \\begin{figure}[htbp] \\begin{center} \\epsscale{1.0} \\plotone{f6.eps} \\caption{Lightcurves.  (a) the phase dependence of adiabatic losses derived from the {\\it Suzaku} data; (b) solid line: calculated 1~TeV gamma-ray fluxes together with HESS data points, dotted line: the normalizations at 1~TeV of  power-law fits for the calculated spectra; (c)  calculated photon indices  around 1~TeV  together with power-law indices reported by HESS; (d) calculated 1--10 keV   X-ray fluxes  with the reported {\\it Suzaku} data; and (e) predicted 1-100 GeV integrated IC flux (note that the peak in the spectrum is at 10 GeV; see Fig. 7).  The calculations were performed for the following parameters: inclination of the orbit $i=30$~deg; elevation over the orbital plane $Z_0=1.8 \\times 10^{12}$~cm; magnetic field strength $B=2.5$~G. A power-law injection spectrum of electrons $Q(E) \\propto E_{\\rm e}^{-2}$ ($10{\\rm GeV}<E_{\\rm e}<100{\\rm TeV}$) was assumed. The total power injected in relativistic electrons was fixed at the level of $10^{37}$~erg/s. The cooled electron spectrum is formed due to radiative (IC and synchrotron) and non-radiative (adiabatic) losses, derived from the  {\\it Suzaku} data. \\label{fig:phase} } \\end{center} \\end{figure}   In Fig.~\\ref{fig:SED} we show the SEDs averaged over the INFC ($\\phi=0.45$--$0.9$) and SUPC ($\\phi \\leq 0.45$, $\\phi \\geq 0.9$) phase  intervals. The corresponding gamma-ray data have been previously reported by HESS \\citep{HESS_06}. Since both the absolute flux and the energy spectrum of TeV gamma-rays vary rapidly  with phase,  in order to compare the theoretical predictions with  observations, we should use smaller phase bins, ideally speaking with the time intervals  $\\Delta t \\leq 100$~s corresponding to the characteristic cooling timescales of electrons. Because of the lack of the relevant gamma-ray data available to us, we here use the X-ray and gamma-ray data integrated over $\\Delta \\phi=0.45$. While  this compromise does not allow us to perform  quantitative studies, it can be used  to make a qualitative comparison of the model calculations with observational data.  The theoretical calculations of the SEDs are in a reasonable agreement with the observed spectra, though they do not perfectly match the gamma-ray fluxes. One can improve the fits by introducing slight phase-dependent changes in the spectra of accelerated electrons, but it is beyond the scope of this paper given the caveat mentioned above. We should also note that the calculations of low energy gamma-rays (in the {\\it Fermi} domain) are performed assuming that the injection spectrum of electrons continues down to $10$~GeV. If this is not the case, the gamma-ray fluxes in the {\\it Fermi} energy domain may be significantly suppressed. On the other hand, the detection of gamma-rays by {\\it  Fermi} would allow us to recover the spectrum of electrons in a very broad energy interval, and thus distinguish between different acceleration models. Another important feature in this scenario is that a hard synchrotron spectrum extending up to a few MeV, is required by the robust detection of $\\sim10$~TeV gamma-rays from the system. A future detection of the emission at MeV energies may bring important information on the presence of highest energy particles in the system.  The  reproduction, at least qualitatively, of the observed spectral and temporal features of the nonthermal radiation with the simple toy model supports the production of X-ray and TeV gamma-rays by the same population of parent particles and allows us to derive several principal conclusions. The electron energy distribution  should be a power-law with an almost constant index of $\\alpha_{\\rm e} = 2 \\Gamma_{\\rm X} -1 \\simeq 2$ to explain the X-ray spectra. Note that in an \\textit{isotropic} photon gas when the Compton scattering takes place in the deep KN regime  such an electron distribution results in a quite steep TeV spectrum with photon index $\\Gamma_\\gamma \\simeq  1+ \\alpha_{\\rm e} \\simeq 3$. This does not agree with HESS observations. However, the anisotropic IC provides a remarkable hardening of the gamma-ray spectrum \\citep{mitya_felix05}, in particular, $\\Gamma_\\gamma \\sim 2$ would be expected for the INFC, and it has indeed been observed using HESS \\citep{HESS_06}. We also note that, to explain the VHE spectrum at SUPC, we have to assume that the emission region is located at a distance $\\approx 2\\times 10^{12}$~cm from the compact object, in the direction perpendicular to the orbital plane. In the standard pulsar scenario, the production region cannot be located far away from the compact object, and even invoking electromagnetic cascading (see e.g.  Fig. 16 in \\citealt{diego08}), one cannot reproduce the reported fluxes around  orbital phase 0.0 \\citep{Mitya_LS, vboshch08}. We have found that the magnetic field strength cannot deviate much from a few G. We can also derive a constraint on the size of the emission region,  imposing a maximum expansion speed of $\\sim$$c$ (the speed of light), and the Hillas criterion \\citep{Hillas}, in which the minimum size of a source capable of accelerating particles to a given energy $E_{\\rm e}$ is $R=R_{\\rm L}$ (where $R_{\\rm L}=E_{\\rm eV}/300 B_{\\rm G}$~cm),  the Larmor radius. This estimate yields a size of $10^{10}$--$10^{11}$~cm, which agrees quite well with the estimate based on the required timescale of adiabatic cooling. Note that  the  requirement of fast adiabatic losses imposes a strong constraint on the acceleration rate of electrons. Indeed, the acceleration timescale  can be expressed as: $t_{\\rm acc} = \\eta  R_{\\rm L}/{c} \\sim 0.1  \\eta  ({E_{\\rm e}}/{\\rm 1 TeV})  ({B}/{\\rm 1 \\ G})^{-1} \\ \\rm s$,  where $\\eta \\geq 1$ parametrizes the acceleration efficiency. In extreme accelerators with the maximum possible rate allowed by classical electrodynamics $\\eta =1$. The HESS spectrum provides evidence of electron acceleration well above 10~TeV. Therefore, $t_{\\rm acc} < t_{\\rm ad} \\sim 1$~s  is required  at $E_{\\rm e} = 10$~TeV, which translates into $\\eta <  3$  for  $B = 3$~G.  Thus we arrive at the conclusion that an extremely efficient acceleration with $\\eta < 3$  should operate in a compact region of $R \\sim 10^{11}\\ \\rm cm$.  Finally, we would like to emphasize that in the scenario described here different radiation energy intervals are characterized by fundamentally different light curves. While the synchrotron X-ray modulation is caused by  adiabatic losses, the light curve in gamma-rays depends critically, in addition, on effects related to interactions with the optical photons of the companion star. Two of these effects, photon-photon pair production and  anisotropic Compton scattering are equally important for the formation of the light curve of TeV gamma-rays. On the other hand, only  the effect of anisotropic Compton scattering has an impact on the formation of the light curve of GeV gamma-rays.  The difference of  light curves in the X-ray and GeV and TeV gamma-ray intervals in this scenario is shown  in Figure \\ref{fig:phase}.   \\begin{figure}[htbp] \\epsscale{1.0} \\begin{center} \\plotone{f7.eps} \\caption{ Model calculations of the non-thermal radiation spectra of  LS~5039, averaged over SUPC (low state) and INFC (high state) orbital phase intervals for $\\phi \\leq 0.45$, $\\phi \\geq 0.9$ (SUPC, Blue); and for $\\phi=0.45$--$0.9$ (INFC,Red). The flux points of $Suzaku$ X-ray spectra (1--40~keV band) are reconstructed with the best-fit function with correction for the interstellar absorption.  The HESS gamma-ray spectra are taken from Aharonian et.(2006). The calculations were performed for the same model parameters as in Fig. \\ref{fig:phase}. \\label{fig:SED}} \\end{center} \\end{figure}"
    },
    "0812/0812.0769.txt": {
        "abstract": "The weak gravitational lensing of distant galaxies by large-scale structure is expected to become a powerful probe of dark energy.  By measuring the ellipticities of large numbers of background galaxies, the subtle gravitational distortion called ``cosmic shear'' can be measured and used to constrain dark energy parameters.  The observed galaxy ellipticities, however, are induced not by shear but by {\\em reduced shear}, which also accounts for slight magnifications of the images.  This distinction is negligible for present weak lensing surveys, but it will become more important as we improve our ability to measure and understand small-angle cosmic shear modes.  I calculate the discrepancy between shear and reduced shear in the context of power spectra and cross spectra, finding the difference could be as high as 10\\% on the smallest accessible angular scales.  I estimate how this difference will bias dark energy parameters obtained from two weak lensing methods: weak lensing tomography and the shear ratio method known as offset-linear scaling.  For weak lensing tomography, ignoring the effects of reduced shear will cause future surveys considered by the Dark Energy Task Force to measure dark energy parameters that are biased by amounts comparable to their error bars.  I advocate that reduced shear be properly accounted for in such surveys, and I provide a semi-analytic formula for doing so.  Since reduced shear cross spectra do not follow an offset-linear scaling relation, the shear ratio method is similarly biased but with smaller significance. ",
        "introduction": "\\label{cha:intro}  Weak gravitational lensing of galaxies has rapidly become an essential cosmological tool.  Cosmic shear -- the large-scale pattern of gravitational distortion embedded in the observed shapes of galaxies -- is a particularly valuable effect that contains information about the growth of mass fluctuations as well as the Universe's expansion history and spatial curvature.  Cosmic shear was first detected by multiple independent groups at the turn of the century \\citep{bacon_refregier_etal_2000, kaiser_wilson_etal_2000,van-waerbeke_mellier_etal_2000, wittman_tyson_etal_2000}, and it could become one of the premier methods for studying dark energy \\citep{albrecht_bernstein_etal_2006}.  Subsequent surveys, such as the 100 Square Degree Weak Lensing Survey, have shown that cosmic shear measurements provide cosmological parameter constraints that are generally consistent with more mature observations such as CMB anisotropy, galaxy correlations, and galaxy cluster counting \\citep{benjamin_heymans_etal_2007}.  This survey was a compilation of four separate surveys which, together, provided 3.5 million galaxies with a median redshift of about $\\zmed\\sim 0.8$.  Ongoing surveys and those on the horizon are expected to provide wider, deeper multi-color images needed to extract information about dark energy.  For example, NASA's Joint Dark Energy Mission is a planned space-based telescope that would survey 4000 deg$^2$, measuring shapes and photometric redshifts for over a {\\em billion} galaxies with a median redshift of $\\zmed\\sim 1.5$.  As daunting as the observational systematics are for ambitious shear surveys, there are also important theoretical hurdles that must be overcome so that our ability to predict the shear signal will keep pace with our ability to measure it.  Various theoretical systematics (uncertainties or complications in the predictions) have been identified and in some cases addressed.  Determining how baryons and non-linear gravitational clustering affect the matter power spectrum on $\\sim 1 \\Mpch$ scales is a primary concern \\citep{huterer_takada_2005, rudd_zentner_etal_2008}.  Significant theoretical systematics also include the effects of non-gaussianity \\citep{cooray_hu_2001} and the spurious signals arising from the intrinsic alignment and clustering of source galaxies \\citep{schneider_van-waerbeke_etal_2002, hamana_colombi_etal_2002, heymans_heavens_2003, hirata_seljak_2004}.  Computational approximations are also being investigated more closely as we start to test regimes where simple first order shear calculations may no longer hold \\citep{vale_white_2003, cooray_hu_2002, shapiro_cooray_2006, dodelson_kolb_etal_2005, dodelson_zhang_2005}.  Ignoring the difference between shear and reduced shear is one example of a  computational approximation that will not be accurate enough for upcoming cosmic shear surveys \\citep{dodelson_shapiro_etal_2006}.  When a galaxy is weakly lensed, the change in ellipticity of its image is proportional to the {\\em reduced shear}\\,: \\begin{equation} \\label{eq:defineRS} g^a \\equiv \\frac{\\gamma^a}{1-\\kappa} \\end{equation} where $\\gamma$ is lensing shear and $\\kappa$ is lensing convergence, a measure of the matter fluctuations projected along the line of sight.  The superscripts denote a ``+'' or ``$\\times$'' shear polarization transverse to the line of sight.  When the distortion fields $\\kappa$ and $\\gamma$ are much less than unity, we can expand \\refeq{defineRS} to first order in the fields, making the approximation that $g^a\\approx\\gamma^a$, \\ie that galaxy ellipticities directly measure the shear field.  This approximation is widely used by theorists since it holds for the vast majority of galaxy images in a cosmic shear survey and since it makes calculating the cosmic shear power spectrum and bispectrum much more manageable.  Dodelson, Shapiro, and White (\\citeyear{dodelson_shapiro_etal_2006}) have found that this seemingly innocuous substitution can nevertheless change the cosmic shear spectrum by as much as 10\\% on arcminute angular scales, where lensing is enhanced by the non-linear matter power spectrum.  Thus, ignoring reduced shear in this regime could bias a survey's analysis, causing it to rule out a correct set of cosmological parameters.  This paper assesses the impact of the reduced shear approximation on dark energy parameter determination from cosmic shear.  Dodelson \\etal considered source galaxies at a single redshift, and hence only assessed the bias that reduced shear would have on \\LCDM models.  Here I generalize to multi-redshift weak lensing analyses that constrain dark energy models by using information about distances and/or the growth of massive structures.  The two methods I consider are weak lensing tomography and a shear ratio method.  Tomography is more well-known: using the CDM paradigm combined with some dark energy model, one calculates a prediction for cosmic shear observables for sources in multiple redshift bins, and then the predictions are compared with observation; model parameters can subsequently be fit to the data \\citep{hu_2002, albrecht_bernstein_etal_2006}.  Shear ratio methods forgo theoretical predictions of the shear signal and attempt to extract model-independent distance measurements by taking ratios of the shear signal at different redshifts.  I will focus on the ``offset-linear scaling'' shear ratio method of \\citet{zhang_hui_etal_2005}, which is an alternative to the ``cross-correlation cosmography'' method of \\citet{jain_taylor_2003}.  The structure of this paper is as follows.  In section 2, I will review some relevant weak lensing formalism and discuss how the next order correction to \\refeq{defineRS} alters cosmic shear power spectra.  In section 3, I estimate  the bias in dark energy parameters that various surveys would experience by using the reduced shear approximation with weak lensing tomography.  I review the offset-linear scaling method in section 4 and similarly calculate how it is biased by reduced shear. ",
        "conclusions": ""
    },
    "0812/0812.3444_arXiv.txt": {
        "abstract": "We report high angular resolution (3$\\arcsec$) Submillimeter Array (SMA) observations of the molecular cloud associated with the Ultra-Compact HII region G5.89-0.39. Imaged dust continuum emission at 870$\\mu$m reveals significant linear polarization. The position angles (PAs) of the polarization vary enormously but smoothly in a region of 2$\\times$10$^{4}$ AU. Based on the distribution of the PAs and the associated structures, the polarized emission can be separated roughly into two components. The component \"x\" is associated with a well defined dust ridge at 870 $\\mu$m, and is likely tracing a compressed B field. The component \"o\" is located at the periphery of the dust ridge and is probably from the original B field associated with a pre-existing extended structure. The global B field morphology in G5.89, as inferred from the PAs, is clearly disturbed by the expansion of the HII region and the molecular outflows.  Using the Chandrasekhar-Fermi method, we estimate from the smoothness of the field structures that the B field strength in the plane of sky can be no more than 2$-$3 mG. We then compare the energy densities in the radiation, the B field, and the mechanical motions as deduced from the C$^{17}$O 3-2 line emission.  We conclude that the B field structures are already overwhelmed and dominated by the radiation, outflows, and turbulence from the newly formed massive stars. ",
        "introduction": "One of the main puzzles in the study of star formation is the low star formation efficiency in molecular clouds. Since molecular clouds are known to be cold, the thermal pressure is small. Hence, if there are no other supporting forces against gravity, the free-fall time scale will be short and the star formation rate will be much higher than what is observed. Magnetic (B) fields have been suggested to play the primary role in providing a supporting force to slow down the collapsing process (see the reviews by Shu et al. (1999) and Mouschovias \\& Ciolek (1999)). In these models, the B field is strong enough and has an orderly structure in the molecular cloud. The B field lines, which are anchored to the ionized particles, will then be dragged in along the direction of accretion, only when the ambipolar diffusion process allows the neutral component to slip pass the ionized component. In the standard low-mass star formation model (Galli \\& Shu 1993; Fiedler \\& Mouschovias 1993), an hourglass-like B field morphology is expected with an accreting disk near the center of the pinched field. Alternatively, turbulence has also been suggested as a viable source of support against contraction (see reviews by Mac Low \\& Klessen (2004) and Elmegreen \\& Scalo (2004)). The relative importance of B field and turbulence continues to be a hot topic as the two methods of support will lead to different scenarios for the star formation process.  Compared with the low mass stars, the formation process of high mass stars is really poorly understood. High mass star forming regions, because of their rarity, are usually at larger distances and are always located in dense and massive regions, because they are typically formed in a group. Hence, both poor resolution and complexity have hampered past observational studies. Furthermore, the environments of high mass star forming regions are very different from the low mass case because of higher radiation intensity, higher temperature, and stronger gravitational fields. Will the B fields in massive star forming sites have a similar morphology to the low mass cases?  Polarized emission from dust grains can be used to study the B field in dense regions, because the dust grains are not spherical in shape. They are thought to be aligned with their minor axes parallel to the B field in most of the cases, even if the alignment is not magnetic (Lazarian 2007). Due to the differences in the emitted light perpendicular and parallel to the direction of alignment, the observed thermal dust emission will be polarized, the direction of polarization is then perpendicular to the B field. Although the alignment mechanism of the dust grains has been a difficult topic for decades (see review by Lazarian 2007), the radiation torques seem to be a promising mechanism to align the dust grains with the B field (e.g., Draine \\& Weingartner 1996; Lazarian \\& Hoang 2007). However, other processes such as mechanical alignments by outflows can also be important.  Polarized dust emission has been detected successfully at arcsecond scales. The best example might be the low mass star forming region NGC 1333 IRAS 4A (Girart, Rao \\& Marrone 2006), which reveals the classic predicted hourglass B field morphology. Results on the massive star formation regions, such as W51 e1/e2 cores (Lai et al. 2001), NGC2024 FIR5 (Lai et al. 2002), DR21 (OH) (Lai et al. 2003), G30.79 FIR 10 (Cortes et al. 2006) and G34.4+0.23 MM (Cortes et al. 2008), typically show an organized and smooth B field morphology.  However, this could be due to a lack of spatial resolution. Indeed, for the nearby high mass cases such as Orion KL (Rao et al. 1998) and NGC2071IR (Cortes, Crutcher, \\& Matthews 2006), abrupt changes of the polarization direction on small physical scales have been seen, which may suggest mechanical alignments by outflows as proposed by these authors. Whether high mass star forming regions will all show complicated B field structures on small scales remains to be examined.  In this study, we report on one of the first SMA measurements of dust polarization for a high mass star forming region, G5.89-0.39 (hereafter, G5.89). The linearly polarized thermal dust emission is used to map the B field at $\\sim$3$\\arcsec$ resolution, and the C$^{17}$O 3-2 line is used to study the structure and kinematics of the dense molecular cloud. The description of the source, the observations and the data analysis, the results, and the discussion are in Sec. 2, 3, 4 and 5, respectively. The conclusions and summary are in Sec. 6. ",
        "conclusions": "High angular resolution (3$\\arcsec$) studies at 870$\\mu$m have been made of the magnetic (B) field structures, the dust continuum structures, and the kinematics of the molecular cloud around the Ultra-Compact HII region G5.89-0.39. The goal is to analyze the role of the B field in the massive star forming process. Here is the summary of our results:  \\begin{enumerate} \\item The gas mass (M$_{gas}$) is estimated from the dust continuum and from the C$^{17}$O 3$-$2 emission line. The continuum emission at 870 $\\mu$m is detected with its total flux density of 12.6$\\pm$1.3 Jy. After removing the free-free emission from the detected continuum, the flux density of the 870 $\\mu$m dust continuum is 7.7 Jy, which corresponds to M$_{gas}$ $\\sim$300 M$_{\\sun}$. M$_{gas}$ derived from the detected C$^{17}$O 3-2 emission line is $\\sim$100 M$_{\\sun}$, which is 3 times smaller than the value derived from the dust continuum. The discrepancy of M$_{gas}$ derived from the dust continuum and the C$^{17}$O emission line is also seen in other UCHII regions, e.g. Hofner et al. (2000). The lower values measured from C$^{17}$O could be due to optical depth effects or abundance problems.  \\item The linearly polarized 870 $\\mu$m dust continuum emission is detected and resolved. The dust polarization is not uniformly distributed in the entire dust core. Most of the polarized emission is located around the HII ring, and there is no polarization detected in the southern half of the dust core except at the very southern edges. The position angles (PAs) of the polarization vary enormously but smoothly in a region of 2$\\times$10$^{4}$ AU (10$\\arcsec$), ranging from $-$60$\\degr$ to 61$\\degr$. Furthermore, the polarized emission is from organized patches, and the distribution of the PAs can be separated into two groups. We suggest that the polarization in G5.89 traces two different components. The polarization group \"x\", with its PAs ranging from $-$60$\\degr$ to $-$4$\\degr$, is located at the 870$\\mu$m \\textit{sharp dust ridge}. In contrast, the group \"o\", with its PAs ranging from 33$\\degr$ to 61$\\degr$, is at the periphery of the \\textit{sharp dust ridge}. The inferred B field direction from group \"x\" is parallel to the major axis of the 870$\\mu$m dust ridge. One possible interpretation of the polarization in group \"x\" is that it may represent swept-up B field lines, while the group \"o\" traces more extended structures. In the G5.89 region, both the large scale B field (group \"o\") and the compressed B field (group \"x\") are detected.  \\item By using the Chandrasekhar-Fermi method, the estimated strength of B$_{\\bot}$ from component \"o\" and from component \"x\" is in between 2 to 3 mG, which is comparable to the Zeeman splitting measurements of B$_{\\parallel}$ from the OH masers, ranging from $-$2 to 2 mG by Stark et al. (2007). The derived lower limit of B$_{\\bot}$ from the detected polarization without grouping and without modeling larger scale B field is $\\sim$1mG. Assuming that B$_{\\bot}$ and B$_{\\parallel}$ have the same strengths of 2 mG in the entire cloud, the derived $\\lambda$ increases from 0.1 to 2.5 toward the UCHII region, which is due to the high contrast of the column density across the cloud. Unless the actual B field strength differs by a factor of 25 across the region and compensates for the contrast in the column density, such a variation of $\\lambda$ in G5.89 is suggested. The corrected mass to flux ratio ($\\lambda_{corr}$) is closer to 1 near the HII region and is much smaller than 1 in the outer parts of the dust core. G5.89 is therefore most likely in a supercritical phase near the HII region.  \\item The kinematics of the molecular gas is analyzed using the C$^{17}$O 3-2 emission line. From the analysis of the channel maps, the position velocity plots and spectra, the molecular gas in the G5.89 region is expanding along with the HII region, and it is also possibly swept-up by the molecular outflows. Assuming the size along the line of sight is uniform in G5.89, P$_{kin}$ is in the range of $\\sim$ 1$\\times$10$^{-9}$ to 1.4$\\times$10$^{-6}$ dyne cm$^{-2}$. The calculated radiation pressure (P$_{rad}$) at a radius of 2$\\arcsec$ and the B field pressure (P$_B$) with a field strength of 3mG are 8.5$\\times$10$^{-7}$ and 3.6$\\times$10$^{-7}$ dyne cm$^{-2}$, respectively. Although the upper limit of P$_{kin}$ is 1.6 times larger than P$_{rad}$ at a radius of 2$\\arcsec$, any variation of the structure in the direction along the line of sight in G5.89 - which is most likely the case - will affect the estimated P$_{kin}$. Nevertheless, P$_{rad}$ is on the same order as P$_{kin}$ and P$_{B}$ at the radius of 2$\\arcsec$ around the O5 star. The scenario that the matter and B field in the 870$\\mu$m \\textit{sharp dust ridge} have been swept-up is supported in terms of the available pressure. \\end{enumerate}   G5.89 is in a more evolved stage as compared with the corresponding structures of other sources in the collapsing phase. The morphologies of the B field in the earlier stages of the evolution show systematic or smoothly varying structures, e.g. on the scale of 10$^{5}$ AU for W51 e1/e2 and G34.4+0.23 MM, and on the scale of 4$\\times$10$^{3}$ AU for NGC 2024 FIR 5. With the high resolution and high sensitivity SMA data, we find that the B field morphology in G5.89 is more complicated, being clearly disturbed by the expansion of the HII region and the molecular outflows. The large scale B field structure on the scale of 2$\\times$10$^{4}$ AU in G5.89 can still be traced with dust polarization. From the analysis of the C$^{17}$O 3-2 kinematics and the comparison of the available energy density (pressure), we propose that the B fields have been swept up and compressed. Hence, the role of the B field evolves with the formation of the massive star. The ensuing luminosity, pressure and outflows overwhelm the existing B field structure."
    },
    "0812/0812.1441_arXiv.txt": {
        "abstract": "{} {We study the emission of the hydrogen Lyman-$\\alpha$ line in the quiet Sun, its center-to-limb variation, and its radiance distribution. We also compare quasi-simultaneous \\lya and \\lyb line profiles.} {We used the high spectral and spatial resolution of the SUMER spectrometer and completed raster scans at various locations along the disk. For the first time, we used a method to reduce the incoming photon flux to a 20\\%-level by partly closing the aperture door. We also performed a quasi-simultaneous observation of both \\lya and \\lyb at the Sun center in sit-and-stare mode. We infer the flow characteristic in the \\lya map from variations in the calibrated $\\lambda$\\,1206~Si\\,{\\sc iii} line centroids.} {We present the average profile of Ly$-\\alpha$, its radiance distribution, its center-to-limb behaviour, and the signature of flows on the line profiles. Little center-to-limb variation and no limb brightening are observed in the profiles of the \\lya line. In contrast to all other lines of the Lyman series, which have a red-horn asymmetry, \\lya has a robust and - except for dark locations - dominating blue-horn asymmetry. There appears to be a brightness-to-asymmetry relationship. A similar and even clearer trend is observed in the downflow-to-asymmetry relationship. This important result is consistent with predictions from models that include flows. However, the absence of a clear center-to-limb variation in the profiles may be more indicative of an isotropic field than a mainly radial flow.} {It appears that the ubiquitous hydrogen behaves in a similar way to a filter that dampens all signatures of the line formation by processes in both the chromosphere and transition region.} ",
        "introduction": "\\lya is by far the strongest line in the VUV spectral range and must play a dominant role in many -- if not all -- processes  in the solar atmosphere. Irradiance at the line center excites atomic hydrogen in the cool material of our solar system by resonant scattering, playing an important role in planetary, cometary, and heliospheric physics. Despite its importance, high-resolution spectroscopic observations of this line are scarce and fundamental information is still missing or based on indirect methods. We present unprecedented observations of the solar emission in the \\lya line of neutral hydrogen with a detailed analysis of line profiles at various locations on the solar disk.  Three decades ago, the calibrated line profiles over the full disk of \\lya and \\lyb were obtained from {\\it OSO 8} data \\citep{Lemaire78}, while the {\\it Skylab} ATM \\citep{Reeves76} provided moderate resolution radiance maps at different locations on the disk. \\citet{Basri78} used \\lya rocket flight spectra obtained by the HRTS instrument to investigate the network contrast and center-to-limb variation in the line profile. One decade later, \\citet{Fontenla88} reported observations of the \\lya profile for various solar disk features obtained by the {\\it SMM}/UVSP instrument. However, all of these measurements were hampered by the strong, narrow geocoronal absorption at the line center.  With the advent of {\\it SoHO}, which offers uninterrupted and undisturbed observations of the Sun from its orbit around the Lagrangian Point L$_1$, new attempts were made to study this line with unprecedented spectral resolution. \\citet{Lemaire98} could deduce the average line profile of the solar irradiance by observing off disk the scattered light from the primary mirror of the SUMER spectrometer. They found a \\lya profile that had nearly no asymmetry. \\citet{Lemaire02} compared the profiles of \\lya and \\lyb and report calibrated irradiances with 10\\% uncertainty. A few years later, \\citet{Lemaire05} investigated the variation in \\lya during solar cycle 23. In contrast to Ly$-\\alpha$, which is generally too strong to be observed on disk by SUMER, \\lyb and the higher members of the Lyman series were repeatedly observed. \\citet{Warren98} completed a comprehensive analysis of these profiles. They found that (1) all lines from levels with main quantum number $n\\ge 2$ are self-reversed and have a shape that deviates substantially from Gaussian with a self-reversal that decreases with increasing $n$; (2) all lines are asymmetric, the red horn being higher than the blue horn; (3) interestingly, the red-horn asymmetry is more pronounced for the brighter network areas.  \\citet{Emerich05} established a new relationship between the line center H\\,{\\sc i}~\\lya irradiance and the total line irradiance derived from the scattered light from the primary mirror of the SUMER instrument. % This work superseded earlier work of \\citet{Vidal-Madjar80}, which used {\\it OSO 5} data and where the \\lya line had to be corrected for geocoronal absorption.  \\citet{Fontenla02} published a model to describe the \\lya line formation. They demonstrated that static models do not reproduce the observations and more reliable results can be obtained by assuming that flows and diffusion play a major role.  \\begin{figure*} \\centering \\sidecaption \\includegraphics[width=11.8cm]{f1_0868.eps} \\caption{{\\bf Left:} \\lya profiles of the 26 June data set for six different radiance bins obtained from all data points of raster scans at disk center, at a position close to the east limb, and at a position half-way to the limb. In logarithmic representation, the bins of log radiance / s cm$^2$ sr {\\AA} erg$^{-1}$  are equally spaced.\\newline {\\bf Right:} The Doppler flow in km/s$^{-1}$ for each pixel was established by the deviation of the $\\lambda$\\,1206~Si\\,{\\sc iii} line centroid from the rest wavelength and six equally spaced velocity bins were defined. Negative values correspond to upflows, and positive values to downflows. Here, we have excluded data points obtained off-disk. The different \\lya profiles for these six velocity bins are shown.\\newline The blue horn is always stronger than the red horn. The asymmetry increases significantly with the redshift. Near disk center, all line features (red horn, blue horn, central depresssion) are offset towards longer wavelenths with increasing downflows, as indicated by the red tracing lines. Close to the limb, the velocity bins are set up almost symmetrically relative to rest, and the variation in the radiance with velocity is significantly lower than on disk. This effect is expected for mostly radial plasma downflows in optically thin TR emission such as $\\lambda$\\,1206~Si\\,{\\sc iii}. } \\label{bins} \\end{figure*} ",
        "conclusions": "We have presented the first on-disk SUMER observations of \\lya with reduced incoming photon flux and some fundamental results derived thereof. Such observations are important in many respects as already demonstrated here. It is, however clear that this work can only be the starting point of greater efforts to understand the excitation and absorption processes responsible for the line formation, and eventually to constrain atmospheric models. From our work, we conclude that downflows must play a fundamental role in line profile formation. We plan further observations during the rise of solar cycle 24. These will reveal more details and indicate, whether there is an imprint of solar activity on the profile of this important line."
    },
    "0812/0812.3996_arXiv.txt": {
        "abstract": "Relativistic current carrying strings moving axisymmetrically on the background of a Kerr black hole are studied. The boundaries and possible types of motion of a string with a given energy and current are found. Regions of parameters for which the string falls into the black hole, or is trapped in a toroidal volume,  or can escape to infinity, are identified, and representative trajectories are examined by numerical integration, illustrating various interesting behaviors. In particular, we find that a string can start out at rest near the equatorial plane and, after bouncing around, be ejected out along the axis, some of its internal (elastic or rotational kinetic) energy having been transformed into translational kinetic energy. The resulting velocity can be an order unity fraction of the speed of light. This process results from the presence of an outer tension barrier and an inner angular momentum barrier that are deformed by the gravitational field. We speculatively discuss the possible astrophysical significance of this mechanism as a means of launching a collimated jet of MHD plasma flux tubes along the spin axis of a gravitating system fed by an accretion disk. ",
        "introduction": "In this paper we study the motion of an axisymmetric, current carrying relativistic string loop in the background of a Kerr black hole. This interesting dynamical system has been studied before in its own right~\\cite{Larsen:1993nt,Frolov:1999pj}), and our goal here is to better understand its behavior.  Our investigation is loosely motivated by the problem of the production of collimated astrophysical jets of matter. Systems that exhibit collimated jets range from accreting young stars, neutron stars and black holes to supermassive black holes in quasars and active galactic nuclei~\\cite{livio}. Magnetized plasmas interacting with an accretion disk are believed to play a central role in jet production, but despite much work and many prosposals for the mechanism, the process is still not yet well understood. Such plasmas are governed by magnetohydrodynamics (MHD), a complicated nonlinear field theory. Under some circumstances, plasmas exhibit associated string-like behavior, either via the dynamics of magnetic field lines embedded in the plasma~\\cite{Christensson:1999tp,semenov,Semenov:2004ib}, or in the dynamics of relatively thin isolated flux tubes of plasma which can be approximately described by an effective one-dimensional string~\\cite{spruit}. In either of these cases, the string would be described by an energy density and a tension, and might be carrying currents of mass and/or charge. The hope is that some essential aspects of the physics could be captured by string dynamics, which is tremendously simpler than MHD.   The dynamical formalism for describing relativistic current carrying strings is well-developed~\\cite{Hindmarsh:1994re,vilshell,Larsen:1993ha,Carter:1996rn}, and a generalization to a wider variety of equations of state is given in Ref.~\\cite{Carter:1997pb}. Here we restrict attention to a relatively simple class of strings, namely axisymmetric loops characterized by a constant tension and a conserved mass current on the string worldsheet, in order to begin developing some insight into the factors influencing the string dynamics. This could be generalized in future work to allow for different equations of state, coupling of a charge current to an external electromagnetic field~\\cite{Larsen:1992xz,Larsen:1992ya}, and deviations from axisymmetry~\\cite{Dyadechkin:2008zz, Dubath:2006vs, Snajdr:2002aa, Snajdr:2002rw}.  The string tension prevents a string loop from expanding beyond some radius, and the worldsheet current can produce an angular momentum barrier preventing the  loop from collapsing into the black hole. Thus, depending on its energy and current, a string may be trapped in a toroidal region surrounding the black hole, or its motion may be confined to a cylindrical shell that extends to infinity. In the latter case, one can envisage processes wherein internal energy of the string is converted into translational kinetic energy via the propagation on the black hole background. This is one of the phenomena we aim to understand. The other key question is what role may be played by the spin of the black hole, as a result of the associated dragging of inertial frames. ",
        "conclusions": "We have studied the motion of relativistic, current carrying strings moving axisymmetrically on the background of a Kerr black hole. First the equations of motion and conserved quantities in the conformal gauge were found. Next,  we determined how the energy of a string that is instantaneously at rest depends on the current and the position of the string. Contours of this quantity determine boundaries of the possible motion of a string with that energy and current. We compared this analysis with that for a Newtonian elastic ring moving around a point mass, and found that system to be qualitatively similar.  By considering a number of examples and various plots of the relevant functions, the possible types of motion were mapped out. Regions of parameters for which the string falls into the black hole, or is trapped in a toroidal volume,  or can escape to infinity, were identified. After this general analysis, we examined representative trajectories found by numerical integration, illustrating various interesting behaviors. In particular, we found that a string can start out at rest near the equatorial plane and, after bouncing around, be ejected out along the axis, some of its internal (elastic or rotational kinetic) energy having been transformed into translational kinetic energy. The resulting velocity can be an order unity fraction of the speed of light.  Finally, we addressed the question of possible astrophysical significance of this system as a simple model of MHD plasma flux tubes, which might conceivably play a role in jet formation and collimation."
    },
    "0812/0812.2959_arXiv.txt": {
        "abstract": "{ The morphology of a disk galaxy is closely linked to its kinematic state. This is because density wave features are likely made of spontaneously-formed modes which are allowed to arise in the galactic resonant cavity of a particular basic disk state. The pattern speed of a density wave is an important parameter that characterizes the wave and its associated resonances. Numerical simulations by various authors have enabled us to interpret some galaxies in terms of high or low pattern speeds. The potential-density phase-shift method for locating corotation radii is an effective new tool for utilizing galaxy morphology to determine the kinematic properties of galaxies. The dynamical mechanism underlying this association is also responsible for the secular evolution of galaxies. We describe recent results from the application of this new method to more than 150 galaxies in the Ohio State University Bright Galaxy Survey and other datasets. } ",
        "introduction": "The pattern speed \\Omegap\\ of a spiral or bar perturbation is an important physical parameter describing density wave features in galaxies.  Although often treated as a free parameter in passive numerical simulations, the pattern speed for a galaxy with a given basic state (described by the radial distribution of its star/gas surface densities, velocity dispersion, and rotation curve) is not arbitrary but is determined by the basic state as part of its spontaneously-formed density wave modal characteristics (Bertin et al. 1989), although interactions between galaxies can temporarily alter some of these characteristics.  The pattern speed of a density wave, together with the galaxy rotation curve, determines the locations of all wave-particle resonances. These resonances affect the shape of nearby periodic orbits, which can lead to the formation of rings (Buta and Combes 1996), and also limit the extent of the patterns themselves (e.g., Contopoulos 1980; Contopoulos \\& Grosbol 1986). Furthermore, pattern speed, which determines the location of the corotation resonance (CR), also influences the manner in which angular momentum is exchanged with the basic state (Zhang 1998) when the angular momentum flux is being transported outward by the density wave (Lynden-Bell and Kalnajs 1972).  This non-uniform angular momentum flux (or the non-zero radial gradient of the flux) leads to the secular morphological evolution of galaxies (Zhang 1996, 1998, 1999). The pattern speed may also influence the lifetime of a pattern (Merrifield et al. 2006).  To gauge connections between pattern speed and galaxy morphology, one can (1) directly measure pattern speeds or resonance locations and connect observed features to resonances; or (2) compare sequences of numerical simulations with actual galaxies, either through generic modeling or modeling of a specific galaxy. For example, Garcia-Burillo et al. (1993), in modeling M51, state that ``The gas response is very sensitive to the [pattern speed]; we obtain a spiral structure similar to what is observed only for a narrow range of \\Omegap.\" Included in (1) are applications of the Tremaine \\& Weinberg (1984) continuity equation method (see other papers in this volume), and also our potential-density phase-shift method (section 3), which taps into a more general connection between galaxy morphology and galaxy kinematic features. ",
        "conclusions": "Numerical simulations have been very useful for highlighting the impact of pattern speed on galaxy structure. From comparisons with ringed galaxies, there is a suggestion of the existence of pattern speed domains, where certain major resonances may or may not exist. The potential-density phase-shift method suggests that multiple pattern speeds are common in spiral and barred galaxies. Both bar-driven and non-bar-driven spirals are detected. The method also brings attention to the possibility of ``super-fast bars,\" where $r(CR) < r(bar)$."
    },
    "0812/0812.1920_arXiv.txt": {
        "abstract": "We present a detailed report on the experimental details of the Antarctic Impulsive Transient Antenna (ANITA) long duration balloon payload, including the design philosophy and realization, physics simulations, performance of the instrument during its first Antarctic flight completed in January of 2007, and expectations for the limiting neutrino detection sensitivity. Neutrino physics results will be reported separately. ",
        "introduction": "Cosmic neutrinos of energy in the Exavolt and higher (1 EeV = $10^{18}$~eV) energy range, though as yet undetected, are expected to arise through a host of energetic acceleration and interaction processes at source locations throughout the universe. However, in only one of these sources--the distributed interactions of the ultra-high energy cosmic ray flux-- does the combination of observational evidence and interaction physics lead to a strong requirement for resulting high energy neutrinos. Whatever the sources of the highest energy cosmic rays, their observed presence in the local universe, combined with the expectation that their sources occur widely throughout the universe at all epochs, leads to the conclusion that their interactions with the cosmic microwave background radiation (CMBR)-- the so-called GZK process (after Greisen, Zatsepin, and Kuzmin~\\cite{GZK})-- must yield an associated cosmogenic neutrino flux, as first noted by Berezinsky and Zatsepin~\\cite{BZ}. These neutrinos are often called the GZK neutrinos, as they arise from the same interactions of the ultra-high energy cosmic rays (UHECR) that cause the GZK cutoff, but they are perhaps more properly referred to as the BZ neutrinos.  In BZ neutrino production scenarios, current experimental UHECR measurements invariably point to the presence of an associated ultra-high energy neutrino flux. For UHECR above several times $10^{19}$~eV, intergalactic space is optically thick to UHECR propagation through the CMBR at a distance scale of several tens of Mpc. Each UHECR source at all epochs is thus subject to local conversion of its hadronic flux to secondary, lower energy particles over a distance scale of order 100~MPc in the current epoch. Neutrinos are the only secondary particle that may freely propagate to cosmic distances, and the resulting neutrino flux at earth is thus related to the integral over the highest-energy cosmic ray history of the universe, to the earliest epoch at which they occur.  Although local sources may also contribute to the EeV-ZeV neutrino flux at earth, the bulk of the flux is generally believed to arise from a much wider spectral convolution, and will thus be imprinted with the cosmological source distribution in addition to effects from local sources. This leads to strong motivations to detect the BZ neutrino flux: first, it is required by standard model physics, and thus its absence could signal new physics beyond the standard model. Second, it is the only way to directly observe the UHECR source behavior over cosmic distance scales. Finally, once established, the spectrum and absolute flux of such neutrinos may afford a calibrated ``test beam'' for both particle physics and astrophysics experimentation, providing center-of-momentum energies on target nucleons of 100-1000 TeV, an energy scale not likely to be reached by other methods in the near future.   The Antarctic Impulsive Transient Antenna was designed with the goal of measuring the BZ neutrino flux directly, or limiting it at a level which would provide compelling and useful constraints on the early UHECR source history. The BZ neutrino flux is potentially very low--of order 1 neutrino per square kilometer per week arriving over 2$\\pi$ steradians is a typical estimate. This flux presents an extreme challenge to detection, since the low neutrino interaction cross section also means that any target volume will have an inherently low efficiency for converting any given neutrino. ANITA's methodology centers on observing the largest possible volume of the most transparent possible target material: Antarctic ice, which has been demonstrated to provide extremely low-loss transmission of radio signals through its bulk over much of the continent. ANITA then exploits the Askaryan effect~\\cite{Ask62}, coherent, impulsive radio emission from the charge asymmetry in the electromagnetic component of a high energy particle cascade in a dielectric medium.  ANITA searches for cascades initiated by a primary neutrino interacting in the Antarctic ice sheet within its field of view from the Long-Duration Balloon (LDB) altitude of 35-37 km. The observed area of ice from these altitudes is of order 1.5~M km$^2$. Combining this with the electromagnetic field attenuation length of ice which is of order 1~km at ANITA's observation frequency range, ANITA is sensitive to a target volume of order 1-2~M km$^3$. The acceptance, however, is constrained by the fact that at any location within the target, the allowed solid angle of arrival for a neutrino to be detectable at the several-hundred-km average distance of the payload is a small fraction of a steradian. Folding in these constraints, the volumetric acceptance is still of order hundreds to thousands of km$^3$ steradians over the range of energy overlap--$10^{18.5-20}$~eV--with the BZ neutrino spectrum. This large acceptance, while tempered by the limited exposure in time provided by a balloon flight, still yields the largest sensitivity of any experiment to date for BZ neutrinos. In this report we document the ANITA instrument and our estimates of its sensitivity and performance for the first flight of the payload, completed in January of 2007. A separate report will detail the results on the neutrino flux. ",
        "conclusions": "We have presented the basis for a balloon-borne payload search for ultra-high energy neutrinos from an altitude of 37~km above the Antarctic Ice sheets. The ANITA payload's first flight resulted in a large database of impulsive events from locations all across Antarctica. When these events indicated a ground-based point source, it could be geolocated to a precision of order $0.3^{\\circ} \\times 1.0^{\\circ}$ in angle, which projects to error ellipses of several square km at typical distances of several hundred km. We have demonstrated that operation down to near-thermal-noise levels is achievable, and electromagnetic interference in Antarctical, while still problematic, is manageable and largely confined to a relatively small number of camps. Results from the analyses of all these data in the search for neutrino candidates will be presented separately.    \\appendix"
    },
    "0812/0812.0860_arXiv.txt": {
        "abstract": "{ We present a simple yet powerful method to test models of cosmic-ray (CR) origin using the distribution of CR arrival directions. The method is statistically unambiguous in the sense that it is binless and does not invoke scanning over unknown parameters, and general in the sense that it can be applied to any model that predicts a continuous distribution of CRs over the sky. We show that it provides a powerful discrimination between an isotropic distribution and predictions from the ``matter tracer'' model, a benchmark model that assumes small CR deflections and a continuous distribution of sources tracing the distribution of matter in the Universe. Our method is competitive or superior in statistical power to existing methods, and is especially sensitive in the case of relatively few high-energy events.  Applying the method to the present data we find that neither an isotropic distribution nor the \\modelsp can be excluded. Based on estimates of its statistical power, we expect that the proposed test will lead to meaningful constraints on models of CR origin with the data that will be accumulated within the next few years by the Pierre Auger Observatory and the Telescope Array. } ",
        "introduction": "\\label{sec:intro} The recent observation \\cite{Abbasi:2007sv,Abraham:2008ru} of a strong suppression in the energy spectrum of ultra-high energy cosmic rays (UHECRs) provides compelling evidence that the bulk of UHECRs come from relatively close (within $\\sim$$100$ Mpc) extragalactic sources. It seems inevitable that the distribution of cosmic-ray (CR) sources traces the distribution of matter on these length scales. One should therefore expect anisotropy in the UHECR flux due to the non-uniform distribution of matter in the nearby Universe. Although other experiments have not detected deviations from isotropy at the highest energies, the Pierre Auger Observatory (PAO) has recently reported anisotropy in the flux of UHECRs  above 57 EeV at a confidence level of more than 99\\% \\cite{Cronin:2007zz, Abraham:2007si}.   Exactly how CR anisotropy manifests itself depends crucially on two factors: the deflections of CRs in magnetic fields and the density of sources.  The deflections of UHECRs in Galactic and extragalactic magnetic fields may vary from a few to a few tens of degrees, the uncertainty being due to both unknown parameters of the magnetic fields and the uncertain chemical composition of CRs at the highest energies. The extragalactic magnetic fields are usually assumed to be small ($\\lesssim 10^{-9}$~G) in voids, and larger in clusters of galaxies and filaments. Simulations show that deflections in the extragalactic magnetic fields may be small for protons for most directions on the sky \\cite{Dolag:2004kp} (see, however, Ref. \\cite{Sigl:2004yk}). For iron they may reach several tens of degrees. The deflections in the Galactic magnetic field may also vary from a few to a few tens of degrees, depending on the chemical composition of UHECRs. However, even the largest deflections are unlikely to completely wash out the anisotropy created by the non-uniform source distribution.  Although the CR source density affects the anisotropy signature, deviations from isotropy are expected for any density: If the number of sources of observable CRs is large, anisotropies in the CR flux arise from the fact that the distribution of sources is inhomogeneous because it traces that of visible matter. If, on the other hand, there are few sources, they will produce anisotropy just because of that. Anisotropy is therefore expected for any source density.  Different assumptions on deflections and source densities lead to different predictions for the UHECR anisotropy.  The case of small deflections is of particular interest because in this case the uncertainties related to magnetic fields are not present. Under the assumption of small deflections, the existing data imply that the number of observable sources above $\\sim$$60$ EeV is large, namely several hundred or more \\cite{Abraham:2007si, Dubovsky:2000gv,Kachelriess:2004pc}. This situation may be approximated by a continuous distribution of sources following the matter distribution in the Universe. This model, which we refer to as the ``matter tracer'' model, predicts that local matter overdensities such as nearby galaxy superclusters should be visible in UHECRs. The matter tracer model may serve as a benchmark in anisotropy studies, as its main ingredients -- small deflections and a large number of observable sources -- represent a limiting case of many realistic situations. Apart from details on particle acceleration in CR sources, the model has no free parameters: for a given injection spectrum the distribution of UHECR over the sky is predicted uniquely. By comparing this prediction to observations, one may test to which extent the underlying assumptions --- small deflections in the magnetic fields, in the first place --- are correct.   The prediction of a continuous distribution of CRs over the sky calls for special statistical tests. In this paper we propose a new statistical test suitable in this situation. Several other methods have been discussed in the literature \\cite{Sigl:2004yk, Kashti:2008bw, Waxman:1996hp, Cuoco:2005yd, Kachelriess:2005uf, Kalashev:2007ph, Sommers:2000us, Singh:2003xr, Smialkowski:2002by, Evans:2001rv, Harari:2008zp, Takami:2008ri}. Our method is competitive or superior in statistical power to the existing methods (see Sect. \\ref{sec:powercomparison} for a comparison in power between different methods). It is especially sensitive in the case of a few high-energy events, which is a regime of physical interest.   \\FIGURE{  \\includegraphics[height=4.5cm]{skyplot_E40_15events}  \\caption{\\label{fig:skymap:example} Aitoff projection of the sky (galactic coordinates $l$ and $b$; see also Fig. \\ref{fig:skymap:galaxies}) showing the relative integral CR flux above $40$ EeV in the matter tracer model (grayscale) together with 15 fictitious CR events (`+' symbols). The gray bands are chosen such that each band contains 1/5 of the total flux, with darker bands indicating a larger flux. Inset: distribution of the 15 CR events over the five bands. } }   The idea of the method is illustrated in Fig.~\\ref{fig:skymap:example}, which shows the actual flux distribution above 40 EeV predicted in the \\modelsp together with 15 fictitious, uniformly distributed CR arrival directions. (We discuss how this figure is obtained in Sect.~\\ref{sec:CRdistmodel}). The contours of equal flux are drawn in such a way that they divide the sky into 5 areas, each containing an equal share of the total flux (1/5 in this case). This implies that events produced according to the \\modelsp should be distributed roughly equally over the five bands. On the contrary, if the sources are distributed in a different way, one generally expects deviations from equipartition. For instance, in the case of an isotropic CR flux most events will fall in the region with the smallest flux (the lightest band in Fig.~\\ref{fig:skymap:example}) just because this region covers the largest area. Comparing the distribution of observed events over the five bands to a model prediction then allows one to test the underlying model. This is illustrated in the inset of Fig.~\\ref{fig:skymap:example} with a histogram that shows the number of events in every band. For this specific example one would reject the matter tracer model as the distribution of events over the five bands deviates significantly from equipartition.  The method described above can be formulated in a binless way. This is an important advantage because binning introduces ambiguity which complicates the statistical analysis. Our numerical method starts by generating a (continuous) skymap of expected flux. We then extract flux values at the positions of data events, thus obtaining a set of flux values which we refer to as the ``data set''. Then a large number of simulated events is generated by Monte Carlo event generation from the model that is to be tested (e.g., an isotropic distribution or the \\model). The flux values at positions of simulated events form the (Monte Carlo) ``reference set''. Then a statistical test is performed to determine whether the data and reference sets are drawn from the same distribution. This final step can be done with the well-known Kolmogorov-Smirnov (KS) test. In this work we also propose a modified version of the KS test, which we show to be more sensitive than the standard KS test for the case at hand. Both the KS test and its modification do not require binning, rendering the whole procedure binless. Furthermore, the method does not invoke scanning over unknown parameters and hence avoids the use of penalty factors to estimate statistical significances.   Although the main  purpose of this paper is to introduce our statistical test and analyze its potential to infer properties of UHECR sources from future data, we have also applied it to the publicly available data from the Akeno Giant Air Shower Array (AGASA) and PAO to test the \\modelsp and an isotropic CR distribution. As neither of the two hypotheses can be ruled out, the present data give inconclusive results which do not significantly add to the ongoing debate on the sources of UHECRs. We find, however, that the proposed test offers a powerful discrimination between source models with more CR events, and we believe it is timely to present the method before such data become available.  The rest of this paper is organized as follows. In Sect.~\\ref{sec:modeling} we discuss our modeling of CR arrival directions in the \\model. This includes subsections on the modeling of the source distribution, CR propagation, and on the Monte Carlo event generation procedure. Sect.~\\ref{sec:tests} regards anisotropy tests. In this section we discuss the test statistics considered in this work, and recapitulate some essential concepts in statistical hypothesis testing. In Sect.~\\ref{sec:powercomparison} we compare the statistical power of the anisotropy test proposed in this work to other tests and we consider its dependence on model parameters. We apply our test to data from AGASA and PAO in Sect.~\\ref{sec:pval}. We discuss our work in Sect.~\\ref{sec:discussion} and conclude in Sect.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} In this paper we have presented a new and powerful method to test models of cosmic-ray origin that predict a continuous distribution of CRs over the sky. A prominent model of this kind is the ``\\model'', a simple yet general model that predicts that local matter overdensities are reflected in the distribution of CRs. Its main assumptions, a continuous distribution of sources tracing the distribution of matter in the Universe and small CR deflections, represent the limiting case of many detailed scenarios. We therefore consider this model as a useful benchmark model in CR anisotropy studies. The local matter density in the Universe, which is required as input for the \\model, is derived from a galaxy catalog. For this purpose we primarily use the catalog compiled by Kalashev et al. \\cite{Kalashev:2007ph}, a complete sample of galaxies up to 270 Mpc (the KKKST catalog; see Sect.~\\ref{sec:srcdist}). We have however recomputed some of our results using the PSCz catalog \\cite{Saunders:2000af}, which was used in various earlier studies (e.g., Ref.~\\cite{Kashti:2008bw}), for comparison.   We analyzed in detail (Sect.~\\ref{sec:powercomparison}) the statistical power of the two proposed tests (denoted as the $D$- and $Y$-tests) to reject an isotropic CR distribution when the \\modelsp is true (with matter densities reconstructed from the KKKST catalog). We find that the proposed tests are competitive or superior in power to other tests in the literature. They are especially sensitive in the physically interesting case of  relatively few events with very high energies, viz. $\\mathcal{O} (10)$ events above $E=100$ EeV (the test statistic proposed in Ref.~\\cite{Kashti:2008bw} provides a more powerful discrimination in the case of many events with lower threshold energy, $E \\gtrsim 40$ EeV). The proposed tests are completely binless, which is an important advantage because binning introduces an ambiguity that complicates the statistical analysis.   We have applied our tests to the AGASA data above 40 EeV \\cite{Takeda:1999sg, Hayashida:2000zr} and the PAO data above 57 EeV \\cite{Abraham:2007si}, and found that the present data are compatible with both the \\modelsp (with matter densities reconstructed from the KKKST catalog) and with an isotropic CR distribution (Sect.~\\ref{sec:pval}). The general trend is that the AGASA data show a greater compatibility with an isotropic distribution, whereas the PAO data favor the \\model. We expect that the proposed tests will provide meaningful constraints on models when more data become available. We estimate that $\\sim$$50$ events above 57 EeV are required to have a  50\\% probability of excluding an isotropic distribution with significance 0.01 if the \\modelsp is correct. The required data can be accumulated by PAO and TA within a few years.  Summarizing our method, both the $D$- and $Y$- tests are based on the (continuous) distribution of a parameter $\\chi$ --- which  is a measure of the amount of matter in the nearby Universe  --- over a set of observed CRs (Sect.~\\ref{sec:chi}). In the \\modelsp the probability that a CR comes from an overdense region is higher than  from an underdense region, so that the $\\chi$-parameter takes preferentially high values. For an isotropic distribution there is no such correlation and $\\chi$ is essentially a random variable. Correlating the distribution of $\\chi$ in a set of observed CR events with model distributions then gives a measure of the (in)compatibility between observations and model predictions. This last step may be done with the standard Kolmogorov-Smirnov test ($D$-test) or a modified version thereof proposed in this work ($Y$-test; Sect.~\\ref{sec:TS}). The latter is shown to be somewhat more powerful in discriminating between the \\modelsp and an isotropic CR distribution.  Reconstructing the matter density in the Universe from a galaxy catalog, as done in this study, necessarily introduces some ambiguity and uncertainties. For the KKKST catalog, which we have primarily used, the dominant source of errors is in the photometric redshift estimates for galaxies beyond 30 Mpc. The associated uncertainties are acceptable given the other uncertainties we are faced with, most notably in the energy resolution of present CR experiments and in modeling CR energy loss in the Universe. For the purpose of this study, the most important difference between the KKKST catalog and the PSCz catalog  is that the former exhibits stronger density contrasts, which improves the prospects of discriminating between the \\modelsp and an isotropic CR distribution. Whether or not actual CR sources exhibit such a strong density contrast remains an open question, to be answered by observations. We would like to point out that the increased contrasts cannot be an artefact of the random errors associated with photometric redshift estimates. Such errors will, on average, bring sources closer. Keeping the CR horizon fixed, this results in a more homogeneous flux distribution rather than a more contrasted one (as may be easily seen by noting the equivalence with increasing the horizon while keeping distances fixed).  Current and future data can test model predictions based on a given galaxy catalog. Using the present data from PAO, we argued in Sect.~\\ref{sec:discussion:X} that the CR source density seems to exhibit stronger contrasts than what is derived from the PSCz catalog, while the data are fully compatible with the KKKST catalog. This picture is however not supported by the AGASA data, which favor an isotropic model. In the future it may become possible to investigate the spatial distribution of CR sources in a greater amount of detail by comparing CR data with predictions from catalogs based on various galaxy types with different clustering properties. This would require more, and more precise, CR data as well as improved galaxy catalogs.   We comment on some possible improvements to and extensions of the proposed tests. A first possibility is to include the energies of individual events  rather than a single threshold energy in the analysis. A preliminary investigation has shown that this can improve the statistical power, but more work is needed to study this idea in detail. Second, we note that it is possible to choose the statistical parameter $\\chi$ in a different manner than eq. \\eqref{eq:def:chi}, e.g. to be equal to the detector exposure (so that $\\chi$ will take preferentially high values in an isotropic model). We have not investigated this possibility here. Finally the estimates for proton energy-loss lengths  may be updated using more sophisticated numerical propagation routines when future data call for a higher degree of precision."
    },
    "0812/0812.0926_arXiv.txt": {
        "abstract": "We calculate the major pair fraction and derive the major merger fraction and rate for 82 massive ($M_{*}>10^{11}M_{\\odot}$) galaxies at $1.7 < z < 3.0$ utilising deep HST NICMOS data taken in the GOODS North and South fields. For the first time, our NICMOS data provides imaging with sufficient angular resolution and depth to collate a sufficiently large sample of massive galaxies at z $>$ 1.5 to reliably measure their pair fraction history. We find strong evidence that the pair fraction of massive galaxies evolves with redshift. We calculate a pair fraction of $f_{m}$ = 0.29 +/- 0.06 for our whole sample at $1.7 < z < 3.0$. Specifically, we fit a power law function of the form $f_{m}=f_{0}(1+z)^{m}$ to a combined sample of low redshift data from Conselice et al. (2007) and recently acquired high redshift data from the GOODS NICMOS Survey. We find a best fit to the free parameters of $f_{0}$ = 0.008 +/- 0.003 and $m$ = 3.0 +/- 0.4. We go on to fit a theoretically motivated Press-Schechter curve to this data. This Press-Schechter fit, and the data, show no sign of levelling off or turning over, implying that the merger fraction of massive galaxies continues to rise with redshift out to z $\\sim$ 3. Since previous work has established that the merger fraction for lower mass galaxies turns over at z $\\sim$ 1.5 - 2.0, this is evidence that higher mass galaxies experience more mergers earlier than their lower mass counterparts, i.e. a galaxy assembly downsizing. Finally, we calculate a merger rate at z = 2.6 of $\\Re$ $<$ 5 $\\times$ 10$^{5}$ Gpc$^{-3}$ Gyr$^{-1}$, which experiences no significant change to $\\Re$ $<$ 1.2 $\\times$ 10$^{5}$ Gpc$^{-3}$ Gyr$^{-1}$ at z = 0.5. This corresponds to an average $M_{*}>10^{11}M_{\\odot}$ galaxy experiencing 1.7 +/- 0.5 mergers between z = 3 and z = 0. ",
        "introduction": "Hierarchical assembly is the established leading theory for the evolution of structure in the universe from the viewpoint of Cold Dark Matter (CDM) theory. This states that larger structures form from the merging of smaller structures. Thus CDM halos grow in size, with the largest structures forming latest in the history of the universe. Galaxies are often thought to form principally in line with halo mergers. If they do, it is likely that merging of smaller galaxies to create more massive ones is the major factor in the evolution of galaxies over cosmic time. Alternatives to a merger history of massive galaxies include rapid collapse mechanisms, whereby galaxies form over very short time-scales in the early universe, and then do not significantly merge during the rest of their lifetime.  The aim of this paper is to address the question: what mechanism drives galaxy evolution with regards to the most massive ($M_{*}>10^{11}M_{\\odot}$) galaxies in the universe? We can observationally test this by calculating the merger fraction of massive galaxies at different redshifts. Previous work by Conselice et al. (2007) explores this problem for similar mass galaxies out to z $\\sim$ 1.4, using morphological techniques. We extend this work using a close-pair method to z $\\sim$ 3, using data from the GOODS NICMOS Survey.  There are several different ways to locate merging galaxies which roughly relate to the stage of the merger. The most direct method is to look at morphological disturbances in galaxies, e.g. Conselice et al. (2003, 2006 and 2008). In this approach one selects galaxies with asymmetries, or distortions in their morphologies, above a certain threshold, and define these as merging systems. This can be carried out by eye, or via computational methods, e.g. CAS, Gini and $M_{20}$ parameters (see Conselice 2003, 2006, Conselice at al. 2008 and Lotz et al. 2008a). Alternatively, one could also look in principle for secondary features such as an abrupt increase in star formation rate. This is because the merging of gas rich galaxies will often give rise to an increase in star formation, and hence a distinctive spectral signature, see e.g. Lin et al. (2008) and Mathis et al. (2005).  Also one can look for galaxies in close proximity and count the number of galaxies in apparent pairs in a given sample. This pair counting method requires much less angular resolution than the morphological approaches, and can be extended to higher redshift more directly than the others. The close pair method does not precisely trace merging, however, but rather looks for likely potential future mergers. Thus, to go from a close pair fraction to a merger fraction requires that assumptions be made. Ultimately, the pair count method measures mergers before they happen, the morphological approach measures merging while it happens, and methods involving increased star formation rates trace merging during and after the event.  Merging is already known to be of paramount importance in driving the rate of star formation, the formation and evolution of supermassive black holes, and as a mechanism for increasing galaxy mass. Further, mergers can be used to trace the formation history of galaxies, probing where, when, and how galaxies form. Conselice et al. (2007) and Rawat et al. (2008) find that massive galaxies ($M_{*}>10^{11}M_{\\odot}$) out to z $\\sim$ 1.5 have an increasing merger fraction with redshift. It has been noted for lower mass systems that there is a levelling off and eventual turn around in the merger fraction at $z < 2$ (see Conselice et al. 2008). In order to probe whether the same levelling off and turn around in the merger fraction occurs for $M_{*}>10^{11}M_{\\odot}$ galaxies, it is necessary to explore the merger fraction of these massive galaxies at higher redshift.  Consequently, we have collated a sample of 82 galaxies with $M_{*}>10^{11}M_{\\odot}$, selected from the GOODS North and GOODS South fields. These are imaged as part of the GOODS NICMOS Survey (GNS). The GOODS field has the multiband coverage needed to estimate good photomentric redshifts, in addition the depth and angular resolution from NICMOS allows us to resolve significantly fainter objects around the primary targets to distances of a few kpc. Thus, for the first time, we can probe the pair fraction history to z $\\sim$ 3. We argue that the increase in merger fraction with redshift continues for massive galaxies in our sample out to z $\\sim$ 3. The merger fraction shows no sign of levelling off or turning around at $2 < z < 3$, unlike for lower mass systems observed by Conselice et al. (2008). Thus, we find that, complimentary to the star formation results of Bundy et al. (2006) and Pipino et al. (2008), the most massive galaxies appear to form by mergers at a higher rate in the early universe than less massive ones.  Throughout this paper we assume a $\\Lambda$CDM Cosmology with: H$_{0}$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7, and adopt AB magnitude units. ",
        "conclusions": "Conselice et al. (2008) find that galaxies with masses in the range $8<\\log(M_{*}/M_{\\odot})<9$ and $9<\\log(M_{*}/M_{\\odot})<10$ have a peak in their merger fraction at around z = 1.5, with galaxies with stellar masses $\\log(M_{*}/M_{\\odot})>10$ appearing to peak in their merger fraction later, at around z = 2. In this paper we have found that for galaxies with $\\log(M_{*}/M_{\\odot})>11$ the merger fraction continues to increase out to z $\\sim$ 3, indicating that the peak must occur at z $>$ 3. This provides further evidence that more massive galaxies undergo a greater number of major mergers earlier than less massive systems.  This method of calculating pair fractions is based on noting overdensities of galaxies around host galaxies, based on background counts. However, the clustering of galaxies increases with redshift (see Bell et al. 2006). So, this effect may be partly explained by the correlation function of galaxies increasing with redshift. Since the parameters of this function are poorly known observationally, no attempt has been made to correct for this issue. Consequently, our pair fraction results from both the GNS and POWIR surveys may be interpreted as the 2 point correlation at 30 kpc for these massive galaxies. This clearly evolves and it is likely true that the actual pair fraction evolves in line with this correlation. Contamination from line of sight projection at scales larger than 30 kpc will be minimised by the fact that we calculate our correction around the galaxies being measured. Further, we have recently acquired morphological CAS measurements on the GNS galaxies as well as the POWIR ones. This gives a consistent merger fraction ($f_{m} \\sim 0.3$) to the pair methods. We explore this in a forthcoming paper, and elude to the relation between morphologiocal and pair methods in Conselice, Yang \\& Bluck (2008).  Despite the large errors associated with calculating merger rates, we have demonstrated that there is little or no evolution in the rate of major mergers with redshift. The maximum evolution permitted by the errors would only amount to increasing the merger rate by a factor of a few. Of greater significance, however, is our calculation of the redshift dependence on the characteristic time between mergers, $\\Gamma(z)$. This quantity decreases significantly with redshift. As such, in the early universe (z $\\sim$ 3) there was a much shorter time between mergers than there is today."
    },
    "0812/0812.1569_arXiv.txt": {
        "abstract": "We analyze the available sample of double quasars, and investigate their physical properties. Our sample comprises 85 pairs, selected from the Sloan Digital Sky Survey (SDSS). We derive physical parameters for the engine and the host, and model the dynamical evolution of the pair. First, we compare different scaling relationships between massive black holes and their hosts (bulge mass, velocity dispersion, and their possible redshift dependences), and discuss their consistency. We then compute dynamical friction timescales for the double quasar systems to investigate their frequency and their agreement with the ``merger driven\" scenario for quasar triggering. In optical surveys, such as the SDSS, $N_{double, qso} / N_{qso} \\approx 0.1\\%$. Comparing typical merging timescales to expected quasar lifetimes, the fraction of double quasars should be roughly  a factor of 10 larger than observed. Additionally, we find that, depending on the correlations between black holes and their hosts, the occurrence of double quasars could be redshift-dependent. Comparison of our models to the SDSS quasar catalog suggests that double quasars should be more common at high redshift. We compare the typical separations at which double quasars are observed to the predictions of merger simulations. We find that the distribution of physical separations peaks at $\\sim 30$ kpc, with a tail at larger separations ($\\sim 100-200$ kpc). The peak of the distribution is roughly consistent with the first episode of quasar activity found in equal mass mergers simulations.  The tail of the quasar pairs distribution at large separations is instead inconsistent with any quasar activity predicted by published simulations. These large separation pairs are instead consistent with unequal mass mergers where gas is dynamically perturbed during the first pericentric passage, but the gas reaches the black hole only at the next apocenter, where the pair is observed. ",
        "introduction": "Interactions and mergers are an integral part of the process that builds up the galaxy population. Mergers are expected to be the dominant growth route of  galaxies in hierarchical scenarios linked to Cold Dark Matter cosmologies \\citep{White}. Recent high resolution observations have shown evidence of such phenomena, revealing that many of the most massive galaxies at $z \\sim 1-4$ are indeed interacting/merging galaxies \\citep[e.g.,]{Chapman}. Many of submillimeter detected galaxies also show signs of nuclear activity \\citep{Alexander2001,Alexander2003, Alexander2005}, and are found near radio galaxies, suggesting that quasar activity was common in dense environments at high redshift. Theoretical models indeed  link quasar activity to mergers  \\citep[e.g.,][]{Haehnelt, Cavaliere,Kauffmann2000,Wyithe2003,VHM, DiMatteo}. If most galaxies host a central massive  black hole (MBH), as implied by the local observations, and nuclear activity is  triggered in mergers, double/binary quasars should also be common in a hierarchically evolving Universe. The expected fraction of double quasars expected in a hierarchical scenario where mergers trigger nuclear activity was investigated by \\cite{VHM}. If accretion were assumed to be triggered on both MBHs involved in a major merger (that is, mergers between galaxies with similar masses, up to a factor of 10 difference), the fraction of binary quasars would be about 10\\%, in conflict with the existing data \\citep{VHM} . A rough agreement with the data could be reached assuming that only the MBH hosted in the larger halo accretes and radiates during a major merger. In this scenario at least two subsequent major mergers are required to give origin to a double quasar system, the satellite MBH being activated in the first merger event.  In these models, the mean timescale over which MBHs accrete and radiate is of order $\\sim 4\\times 10^7$ years, much shorter than the typical time between two major mergers \\citep{Somerville2001}.  Only 16 pairs are known in the Large Bright Quasar Survey  \\citep{Hewett} out of $10^4$ quasars, and among these 16, the confirmed physical associations are less than 10 \\citep{Kochanek}. Even the most recent optical samples imply the same fraction of binaries, 0.1--0.4\\%. Hennawi et al. (2006) analyze a sample of $\\sim 59.000$ quasars in the Sloan survey, identifying only 26 close ($<50$ kpc) systems.  On the other hand, about 30\\% of X-ray detected sub-mm galaxies at $z=2$ are in pairs (Alexander et al 2005).  There are only a few well studied cases of ongoing mergers at sub-galactic scales in which each companion is an AGN; e.g., NGC 6240 (Komossa et al. 2003) and Arp 229 \\citep{Ballo2004}. Both are advanced state merging systems, belonging to the class of ultraluminous infrared galaxies. Noticeably, none of the nuclei show signs of activity in the optical band, due to obscuration.  LBQS 0103-2753 is an example of an optically selected double quasar with separation below galactic scale, 3.5 kpc \\citep{Junkkarinen}.  Kochanek et al. (1999) analyzed the then available optical sample of double quasars in order to assess the correspondence between mergers and quasar activity. They concluded that the merger scenario \\citep[as proposed by][]{Bahcall} was in good agreement with the double quasar population. They supported the merger scenario by noticing that the maximum separations between pairs are characteristic of scales on which tidal perturbations become important ($\\sim 70$ kpc), and that the absence of small separation pairs is characteristic of dynamically evolving binaries. Mortlock et al. (1999) analyzes the same sample, providing a dynamical model that suggests that double quasars are associated with the very first stages of galaxy mergers. They compare the dynamical time-scale of the mergers, which results close to a Hubble time, to the expected quasar lifetime. Comparison of the time-scales implies that binary quasars are short-lived. They further suggest that the most likely explanation for the death of the binaries is that the in-flow of gas ceases, probably as the merger becomes more stable.  We build-up on the inspirational paper by Mortlock et al. (1999), and analyze the larger sample of double quasars available today. We focus on the currently favored merger scenario, comparing the double quasar population to models related to merger simulations, and deriving black hole and galaxy properties using the most recent results on quasar-host correlations. The redshift-dependencies of the observationally-determined relations linking MBHs to their hosts is an important factor in determining the merging timescale. We discuss the implications of the claimed observed redshift-dependent scaling relations in terms of quasar host and pair fraction evolution. ",
        "conclusions": "If most galaxies host a MBH \\citep[e.g.,][]{Richstone2004} and galaxy mergers trigger quasar activity (Haehnelt \\& Rees 1993; Wyithe \\& Loeb 2003;  Kauffmann \\& Haehnelt  2000; Cavaliere \\& Vittorini 2000; Di Matteo et al. 2005) then the observational sample of binary quasars can be used as a probe of quasar lifetimes. Let's start with Occam's razor and assume that every galaxy hosts a MBH and quasar activity is triggered by galaxy mergers. If  the lifetime of quasars equals the merger timescale, $t_{qso}=t_{DF}$, the probability of observing a double quasar is in principle 100\\%, if we do not consider additional factors, such as obscuration. If $t_{qso}\\ll t_{DF}$, then the probability  much smaller than 100\\%. Additionally, if there is a delay in the triggering of the two quasars, then one might have ceased its activity before the other started, for instance. Face value, our results require that the lifetime of quasars is extremely short, about an order of magnitude shorter than the Salpeter time, or the lifetime expected from numerical simulations.   In a merger driven scenario for quasar activity, with MBHs in most galaxies, where naively all sources are expected to be in binaries, the 0.1\\% binary fraction can be explained in different ways: \\begin{itemize} \\item Quasar pairs are obscured much more than expected from simulations, at Compton thick levels ($N_H>10^{24}{\\rm cm^{-2}}$), otherwise pairs would be easily detected in X-rays. \\item Activity is triggered on only one of the 2 black holes, or on the MBH(s) in the merger remnant. Even if there is a MBH binary in the remnant (Hopkins et al. 2006), we should expect the binary to be hard, or at least the MBHs to be bound by the time the galaxy merger remnant has relaxed enough to be considered a \"normal galaxy\". A pc-scale binary would not be resolved in the SLOAN or LBQS surveys. However, this hypothesis does not explain pairs occuring  at ten-kpc or hundred-kpc scales. \\item There is a time delay between the onset of activity in the 2 merging systems (e.g. the gas infall timescale depends on the dynamical time of each host, rather than the dynamical time of the merging system, see \\citet{Colpi2007}). If  the two quasars light up at different, random times, then  $N_{double, qso} / N_{qso} \\simeq (t_{qso}/t_{DF})^2\\simeq0.1$\\%, consistent with the observational result.  If the time delay is longer that the quasar timescale, then binarity is avoided completely. \\end{itemize}  The average mass ratio of the quasar pairs we analyzed is $\\simeq2$, while, since the mass function decreases steeply at large galaxy masses it is more probable that a very massive galaxy\\footnote{As all the MBHs in our sample have masses above $10^8\\msun$, they are hosted in large galaxies sitting at the steep end of the mass function.} interacts with a much smaller secondary. If the mass ratio is very large, quasar activity is not supposed to be triggered at the level of observability our sample requires. If the mass ratio is $\\simgt$ 4, quasar activity is expected \\citep{Colpi2007}, but the delay between the onset in the primary and in the secondary is too long for the pair to be detected as a ``double quasar\". Hence we prefer the third hypothesis above.   Clearly the occupation fraction (that is the fraction of galaxies hosting a MBH) is indeed an important factor. If not all galaxies host a MBH the fraction of binaries decreases. As long as we do not have a demography of inactive MBHs at $z>0$, this issue cannot be addressed properly.  This fair comparison has to await for the Laser Interferometric Space Antenna, probing the BH population via gravitational waves detections.  However, we can compare our hypothesis to theoretical expectations. In quasar evolution models that evolve the population {\\it ab initio} a minimum number density of MBHs is required to match the luminosity function of quasars. That is, one can not have less MBHs than quasars, at a given time. With this constraint in mind  models such as those described in Volonteri et al. 2008 (see also Marulli et al. 2007) require an occupation fraction higher than 90\\% for MBHs hosted in halos with mass above  $10^{12}\\msun$, the typical quasar host halo found in clustering analysis \\citep{MyersB,Coil}.   We compared the typical separations at which binary quasars are observed  (Figure~\\ref{fig:sep}) to the merger phase and separations that should correspond to the quasar phase in merger simulations. We find that the distribution of physical separations peaks at $\\sim 30$ kpc, with a tail at larger separations ($\\sim 100-200$ kpc). The peak of the distribution is roughly consistent with the first episode of quasar activity found in equal mass mergers simulations  \\citep[see, Figure 2 in][]{HopkinsC}.  The tail of the quasar pairs distribution at large separations does not correspond to any quasar phase predicted by published simulation results. However, the mass ratio of the black holes in the pairs does not match the 1:1 ratio that was typically selected in galaxy merger simulations. The average mass ratio is $\\langle 2.4\\pm1.3\\rangle$, a regrettably still unexplored range.  The pairs at large separation could be unequal mass mergers  seen at the first apocentric passage.  This scenario would follow naturally if gas is dynamically perturbed during the first pericentric passage, but the gas in the disk requires longer timescales to loose enough angular momentum to be driven at subparsec scales and fuel the MBH. This final stage likely occurs when the two galaxies are at the first apocenter.  Finally, we explored how the redhift dependence (or independence) of the correlations between MBHs and their hosts can be tested using large samples of quasar pairs. If MBHs follow a redshift independent correlation with either the velocity dispersion of the host (e.g., Tremaine et al. 2002, Ferrarese \\& Merritt 2000, case A), or with the mass of the bulge (e.g. Marconi \\& Hunt 2004, Haring \\& Rix 2004, case B), we expect merger timescales to decrease at high redshift, leading to a higher fraction of double quasars. If, on the other hand, the correlations between MBHs and their hosts evolve with time (cases C and D, e.g., McLure et al. 2006, Woo et al. 2006) the merger timescales vary extremely slowly with cosmic time, and an increase in the binary fraction is not expected."
    },
    "0812/0812.3931_arXiv.txt": {
        "abstract": "We have simulated a population of young spin-powered pulsars and computed the beaming pattern and light curves for the three main geometrical models: polar cap emission, two-pole caustic (``slot gap'') emission and outer magnetosphere emission. The light curve shapes depend sensitively on the magnetic inclination $\\alpha$ and viewing angle $\\zeta$.  We present the results as maps of observables such as peak multiplicity and $\\gamma$-ray peak separation in the $(\\alpha, \\zeta )$ plane. These diagrams can be used to locate allowed regions for radio-loud and radio-quiet pulsars and to convert observed fluxes to true all-sky emission. ",
        "introduction": "With the successful launch of the {\\it Fermi Gamma-ray Space Telescope}, formerly {\\it GLAST}, many pulsars are expected to be detected in the $\\gamma$-ray band. This is an important opportunity since $\\sim$GeV $\\gamma$-ray emission dominates the observed electromagnetic output of young pulsars. Also, some pulsars are expected to be detected in the $\\gamma$-ray band but undetectably faint in the radio. The natural interpretation is that the radio and $\\gamma$-ray antenna patterns differ. Indeed, the observed pulsars have very different radio and $\\gamma$-ray pulse profiles. Since the observed pulse profile is simply a cut through this antenna pattern for the Earth line-of-sight angle $\\zeta_E$, interpretation of the radio and $\\gamma$-ray profile requires a beaming model. When the spin geometry is known, observed pulses can be used to select the best model. Conversely, for a given model, the pulse pattern can restrict the spin orientation. Finally, beaming corrections can be used to relate the observed line-of-sight flux to the full-sky emission. Thus, aided by beaming models, $\\gamma$-ray pulsar data can give a new window into the energetics and evolution of neutron stars in the Galaxy.  There are three main geometrical models for the $\\gamma$-ray pulse beaming, each tied to locations in the magnetosphere where force-free perfect MHD (${\\vec E}\\cdot {\\vec B}=0$) conditions break down and the uncanceled rotation-induced EMF ${\\vec E}\\approx r {\\vec \\Omega} \\times {\\vec B}/c$ causes particle acceleration and radiation. These are the so-called magnetospheric gaps. The first $\\gamma$-ray (and radio) pulsar models traced emission to acceleration zones on the field lines at very low altitudes above each magnetic pole (although \\citealt{dh96} argued that such ``polar cap'' model emission could extend to  $\\sim 2 R_\\ast$). Another major class of models posits ``Holloway'' gaps above the null charge surface, extending toward the light cylinder (the ``outer gap'' model; see \\citealt{chr86}; \\citealt{r96}). More recently, it has been argued that the true null charge location depends on currents in the gap \\citep{h06}, allowing the start of the outer gap to migrate inward towards the stellar surface. Conversely, it has been argued that acceleration at the rims of the polar caps may extend to very high altitude \\citep{mh04}, effectively pushing the polar cap activity outward.  Together these effects suggest that emission can occur over a large fraction of the boundary of the open zone. One model that has been suggested for this geometry makes low-altitude emission visible from one hemisphere and higher-altitude emission (i.e., above the null charge surface) visible from the other. This is the ``two-pole caustic'' (TPC) model \\citep{dr03}, which is thus intermediate between the polar cap (PC) and outer gap (OG) pictures.  \\begin{figure*}[ht!] \\includegraphics[scale=0.40]{f1.eps} \\caption{\\label{Phaseplot} Sample antenna intensity patterns for $\\alpha=65\\arcdeg$ with the horizontal axis of pulsar phase ($\\phi$) and vertical axis of viewing angle ($\\zeta$). The TPC model for $w=0.05$ is shown in green, the OG model for $w=0.1$ is shown in red and the PC emission site is shown in blue. The cyan lines show the locus of the possible high altitude ($r=500$\\,km, here for $P=0.2$\\,s) radio emission, with the radiating front half shown solid and the back half dashed (see \\S 3.3).} \\end{figure*} ",
        "conclusions": "Our study allows a quick visual summary of model predictions, which may be compared with pulsar discoveries anticipated from the LAT. We have already seen that the classic polar cap (PC) picture is difficult to reconcile with the data. Both the TPC and OG pictures, however, seem viable. Certain observations would clearly discriminate. For example, detection of $\\gamma$-ray pulsars with $\\alpha+\\zeta < 90\\arcdeg$, but with nearly constant intensity through the period, would give a clear preference for the TPC model. In contrast, the prevalence of light curves with simple double pulses separated by $0.05 < \\Delta < 0.3$ is a hallmark of high-altitude OG emission. The distributions in peak widths and in the relative number of radio-loud and radio-quiet pulsars are also a strong discriminant; we will address these in a following paper concentrating on the population statistics of detectable objects (Watters et al. 2009).  Also, the predicted pulse profile complexities for the two models differ. OG models are complex for $\\alpha \\sim \\zeta \\sim 90\\arcdeg$; TPC models have complex profiles at intermediate $\\alpha$. In general, the radio-faint pulsars are double-peaked in both models, but the TPC picture has relatively more single- and zero-peaked solutions.  One area of agreement between the TPC and OG models is the preponderance of large $f_\\Omega$; this is a principal result of our study: inferred efficiencies should be increased by roughly an order of magnitude from those commonly assumed.  In practice this means that some young pulsars must have high ($>10\\%$) $\\gamma$-ray efficiencies. This argues that a study of the $\\gamma$-ray emission is an even better probe of magnetosphere structure and energetics than previously thought. Our $f_\\Omega$ corrections tighten the well-known trend for efficiency to scale as $\\eta \\propto {\\dot E}^{-1/2}$, although a few outliers remain. If, as more pulsars are discovered by the LAT, we find that this trend remains strong, we may even be able to use observed fluxes, after $f_\\Omega$ correction, for ``standard candle'' estimates of pulsar distances. Such arguments suggest relatively low distances for PSRs J2021+3651 and B1706$-$44.  Importantly, for models with $w({\\dot E})$ increasing with age, $f_\\Omega$ tapers off at small ${\\dot E}$. Thus the sky coverage decreases smoothly into a ``death zone'' where the gap shuts off. Near shut-off, the small solid angle on the sky ensures that old pulsars appear relatively bright and are visible to large $d$. This will be an important signature in the population sums (Watters et al. 2009). Examination of the distribution of pulse properties over the simulated pulsar sample is thus the key to testing the models and to producing the corrected luminosities that will make LAT data powerful in probing pulsar populations and spindown."
    },
    "0812/0812.1333_arXiv.txt": {
        "abstract": "We have used the visible integral-field replicable unit spectrograph prototype (VIRUS-P), a new integral field spectrograph, to study the spatially and spectrally resolved Lyman-$\\alpha$ emission line structure in the radio galaxy B2 0902+34 at $z=3.4$. We observe a halo of Lyman-$\\alpha$ emission with a velocity dispersion of $\\approx$250 km s$^{-1}$ extending to a radius of 50 kpc. A second feature is revealed in a spatially resolved region where the line profile shows blueshifted structure. This may be viewed as either HI absorption at $\\approx$-450 km s$^{-1}$ or secondary emission at $\\approx$-900 km s$^{-1}$ from the primary peak. B2 0902+34 is also the only high redshift radio galaxy with a detection of 21 cm absorption. Our new data, in combination with the 21 cm absorption, suggest two important and unexplained discrepancies. First, nowhere in the line profiles of the Lyman-$\\alpha$ halo is the 21 cm absorber population evident. Second, the 21 cm absorption redshift is higher than the Lyman-$\\alpha$ emission redshift. In an effort to explain these two traits, we have undertaken the first three dimensional Monte Carlo simulations of resonant scattering in radio galaxies. We have created a simple model with two photoionized cones embedded in a halo of neutral hydrogen. Lyman-$\\alpha$ photons propagate from these cones through the optically thick HI halo until reaching the virial radius. Though simple, the model produces the features in the Lyman-$\\alpha$ data and predicts the 21 cm properties. To reach agreement between this model and the data, global infall of the HI is strictly necessary. The amount of gas necessary to match the model and data is surprisingly high, $\\ge 10^{12}M_{\\odot}$, an order of magnitude larger than the stellar mass. The collapsing structure and large gas mass lead us to interpret B2 0902+34 as a protogiant elliptical galaxy. This interpretation is a falsifiable alternative to the presence of extended HI shells ejected through feedback events such as starburst superwinds. An understanding of these gas features and a classification of this system's evolutionary state give unique observational evidence to the formation events in massive galaxies. ",
        "introduction": "\\label{Intro_sec} \\par Progress in understanding high redshift radio galaxies (HzRGs, where we arbitrarily consider high to mean z$\\gtrapprox$2) will advance our picture of both massive galaxy formation and evolution. HzRG are thought to be the progenitors of massive local cluster giant elliptical galaxies. Their radio signal allows detection over a huge range of cosmic time. We will only discuss a few revelant trends of HzRGs as extensive summaries are given by \\citet{McCa93} and \\citet{Mile08}. Radio galaxies are found to be associated with some of the largest stellar mass galaxies at all redshifts \\citep{Lill84,Pent01}. The environments of massive radio galaxies have also been progressively traced to higher redshifts with radio galaxies typically residing in cluster environments at all redshifts \\citep{Long79,Pres88,Hill91,Cari97,Best00a}. HzRGs often show extended Lyman-$\\alpha$ emission up to and beyond their radio emission radii. The Lyman-$\\alpha$ emission probes warm, ionized gas. Against this emission is often seen line profile structure which serves as a probe of the neutral gas through scattering and absorption. The properties of Lyman-$\\alpha$ blobs (LABs) \\citep{Keel99,Stei00,Palu04} and the gas halos of quasars \\citep{Heck91,Bark03} are morphologically similar to those in HzRGs. \\par Cooling radiation may be an important source of extended Lyman-$\\alpha$ radiation, and an indication of ongoing galaxy formation as infalling gas continues to build the galaxy's mass \\citep{Haim00,Fard01}. The most direct signature of a high redshift cooling flow is extended Lyman-$\\alpha$ emission powered by the decrease in gravitational potential energy from infalling material. This is a challenging observation and interpretation as other sources of Lyman-$\\alpha$ generation are usually stronger than the cooling radiation, but \\citet{Nils06} and \\citet{Smit07} show promising cases. Low redshift cooling flows have a long history of investigation primarily through X-ray observations \\citep[e.g.][]{Fabi94}, but space based observations \\citep{Kaas01,Pete01,Tamu01} have given evidence against most cases of low redshift, strong cooling flows and diminished the importance of this mechanism for galaxy growth. Most massive galaxy growth is now thought to occur through major mergers. \\citet{Li07} show by a semi-analytic model, for example, the likely growth history of a galaxy like the SDSS quasar J1148+5251 at z=$6.42$ \\citep{Fan03} where gas rich major mergers dominate as the source of the mass growth and a starburst phase predates a quasar phase. Observationally checking and challenging theories for the relative contributions of mergers and steady gas accretion for different galaxy types and epochs is an important goal. \\par In terms of evolution, AGN feedback has been proposed \\citep{Vala99} as the mechanism to produce the tight black hole/stellar mass correlation \\citep{Mago98,Gebh00,Ferr00} and the difference between the theoretical \\citep{Kais86} and the observed \\citep{Mark98,Arna99} relation between cluster temperature and X-ray luminosity. Radio galaxies in particular and their powerful jets have been invoked as candidates for causing feedback \\citep{Inou01,Rawl04,Crot06} by depositing mechanical and thermal energy to their surroundings. A signature of outflowing gas at a velocity exceeding the halo escape speed would be an important observation of high redshift feedback. \\par B2 0902+34 \\citep{Lill88} is a powerful high redshift ($z\\approx3.39$) radio galaxy of moderate projected radio diameter (40 kpc, 5$\\arcsec$) with a huge, diffuse Lyman-$\\alpha$ halo ($\\approx100$kpc) and is one of the most thoroughly studied HzRGs. Three observations make B2 0902+34 a unique window into HzRG halos as a class. First, 21 cm absorption from HI gas against the radio source continuum has been observed at $z_{21 cm}=3.3962$ \\citep{Uson91,Brig93,Cody03,Chan04}, but any corresponding feature from this HI population has not shown in the Lyman-$\\alpha$ emission line structure. This is the only HzRG for which 21 cm HI absorption has been repeatably measured with seven other HzRGs giving null detections \\citep{Roet99,DeBr03}. Second, \\citet{Cari95} used radio spectral index and polarization mapping to constrain the approaching, northern radio jet inclination to the observer to between 30$\\arcdeg$ and 45$\\arcdeg$. This is a particularly low inclination angle for a radio galaxy, but it does offer very different lines of sight to the two radio jets and any other structures that may follow the radio jets. Third, \\citet{Vill07} have shown the emission line ratios in B2 0902+34 to be consistent with pure AGN photoionization through the Lyman-$\\alpha$, CIV$\\lambda$1549, and HeII$\\lambda$1640 lines. A subsample of their HzRGs show line ratios that requires stellar photoionization. These systems would have peculiar and complicated source distributions for Lyman-$\\alpha$ emission, but B2 0902+34 appears to display the relatively well understood geometry of Lyman-$\\alpha$ emission powered by an AGN. Until recently, longslit spectroscopy had detected only Lyman-$\\alpha$, CIV, and HeII emission lines \\citep{Mart95} and shown no sign of a bimodal Lyman-$\\alpha$ line profile which is often seen in other HzRGs \\citep{vanO97}. Deeper spectroscopy and a different slit position angle revealed a spatially resolved region with a bimodal Lyman-$\\alpha$ line profile \\citep{Reul07}, with most of the emission still from a simple and quiescent line profile. We will from here forward define a quiescent Lyman-$\\alpha$ line as one with $\\sigma<$1000 km s$^{-1}$ to indicate emission which does not display turbulent, radio jet dominated kinematics. \\citet{Reul07} have also reported a weak [O II]$\\lambda$3727 detection in the galaxy's center. Previous observations in deep Lyman-$\\alpha$ narrowband \\citep{Reul03} and K broadband \\citep{Eise92} imaging and spectroscopy of the [OIII] 5007\\AA\\ line \\citep{Eale93} have led some authors to tentatively interpret B2 0902+34 as a protogalaxy based on its flat spectral energy distribution. \\citet{Reul03} showed a diffuse Lyman-$\\alpha$ morphology with two emission peaks roughly perpendicular to the radio axis. Unlike the two other HzRGs in their sample, B2 0902+34 did not display filamentary structure in Lyman-$\\alpha$. \\par We present the first integral field spectroscopy (IFS) data of Lyman-$\\alpha$ emission in B2 0902+34. This is the fourth study of HzRGs halos using IFS after the works of \\citet{Adam97}, \\citet{Vill07b}, and \\citet{Nesv08}. \\citet{LeFe96} and \\citet{vanB05} have used IFS to search for Lyman-$\\alpha$ Emitters (LAEs) around HzRGs. With knowledge of the full spatial distribution of the bimodal Lyman-$\\alpha$ line profile, we create a simulation for resonant scattering of the Lyman-$\\alpha$ photons through a halo of HI. We find a model that reproduces most of the important spatial and spectral properties of our IFS observation and strictly requires a collapsing structure with large HI mass. The natural explanatory power of this picture for other observations, such as the 21 cm absorption data, leads us to claim that B2 0902+34 is a protogalaxy still experiencing initial collapse. \\par In \\S \\ref{Obs_sec} we describe our observations. In \\S \\ref{Mod_sec} we present our new resonant scattering model and analyze the impact of our model on 21 cm measurements. We show in \\S \\ref{in_or_out} that our data is best fit through resonant scattering against an infalling halo. In \\S \\ref{Conc_sec}, we propose further discriminating observations and review. Throughout, we use a standard cosmology with H$_0$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$ = 0.27, and $\\Omega_{\\Lambda}$ = 0.73. This gives look back times of 11.8 Gyr, angular scales of 7.5 kpc arcsec$^{-1}$, and a luminosity distance of 30 Gpc for B2 0902+34 at $z=3.4$. ",
        "conclusions": "\\label{Conc_sec} \\par We have made new spatially resolved IFS observations of Lyman-$\\alpha$ in B2 0902+34. We find a spatially resolved region of weaker and bluer emission superimposed on a symmetric halo of primary emission. We have shown that all properties of B2 0902+34 can best be fit with a resonant scattering infall model with a large, mostly neutral, hydrogen mass. The sign of the velocity field is relatively robust and model independent. Signal seen through higher optical depths becomes increasingly bluer indicative of infall and resonant scattering. We match our data with a model dependent HI mass of $\\ge10^{12}M_{\\odot}$ which is $\\ge$17$\\times$ larger than the stellar mass. For three reasons an outflowing HI scenario cannot fit our observations: 1) the strength of the observed 21cm absorption and the Lyman-$\\alpha$ profile are incompatible with a single HI population if optically thin and the velocity offsets between the 21cm signal and the higher optical depth Lyman-$\\alpha$ signal carry the wrong sign for outflow if optically thick, 2) the surface brightness profile, with the assumption that the gas follows the dark matter, would be steeper than observed if optically thin, and 3) the bimodal Lyman-$\\alpha$ line profile exists only in a small, non-central region of the galaxy which is incompatible with an outflowing shell geometry. The 21 cm HI absorption has been a powerful confirmation of the kinematics through its redshift offset from the Lyman-$\\alpha$ halo emission, and we believe observational progress in understanding this class of object will benefit from further deep 21 cm measurements on other HzRGs. This new scenario of a massive, infalling HI halo with resonant scattering places B2 0902+34 in a much earlier evolutionary state than the previously formulated starburst superwind picture, and we classify B2 0902+34 as a protogiant elliptical galaxy in an early collapse phase before significant feedback has disrupted star formation and AGN fueling. Observations of other HzRGs with IFS will be important in correctly classifying systems between formation and feedback modes and allow detailed observational tests of these two processes. We are observing a large sample of HzRGs with VIRUS-P in order to undertake a similar analysis on other systems. \\par Further types of data on this system will allow confirmation of our infalling, resonant scattering model. Polarization measurements of Lyman-$\\alpha$ could confirm the scattering process \\citep{Lee98}. The bright regions of HzRGs have been found to have modest levels of polarization ($\\approx$10\\%, \\citet{Vern01}) while scattering simulations in similar situations to our model \\citep{Dijk08} predict higher ($\\approx$40\\%) values. A measurement between HzRGs with and without emission line halos may demonstrate this split. The key, potentially falsifying observation for this resonant scattering, infalling model of B2 0902+34 will be radio interferometry spectral imaging to spatially resolve the distribution of HI gas through its 21 cm absorption. \\citet{Cari95} discussed the utility of such a measurement, and we have now given an observationally motivated model that predicts spatially extended HI."
    },
    "0812/0812.2510_arXiv.txt": {
        "abstract": "We identify SDSS\\,011009.09+132616.1, SDSS\\,030308.35+005444.1, SDSS\\,143547.87+373338.5 and SDSS\\,154846.00+405728.8 as four eclipsing white dwarf plus main sequence (WDMS) binaries from the Sloan Digital Sky Survey, and report on follow-up observations of these systems. SDSS\\,0110+1326, SDSS\\,1435+3733 and SDSS\\,1548+4057 contain DA white dwarfs, while SDSS\\,0303+0054 contains a cool DC white dwarf. Orbital periods and ephemerides have been established from multi-season photometry. SDSS\\,1435+3733, with $\\Porb=3$\\,h has the shortest orbital period of all known eclipsing WDMS binaries. As for the other systems, SDSS\\,0110+1326 has $\\Porb=8$\\,h, SDSS\\,0303+0054 has $\\Porb=3.2$\\,h and SDSS\\,1548+4057 has $\\Porb=4.4$\\,h. Time-resolved spectroscopic observations have been obtained and the H$\\alpha$ and \\Lines{Ca}{II}{8498.02,8542.09,8662.14} triplet emission lines, as well as the \\Lines{Na}{I}{8183.27,8194.81} absorption doublet were used to measure the radial velocities of the secondary stars in all four systems. A spectral decomposition/fitting technique was then employed to isolate the contribution of each of the components to the total spectrum, and to determine the white dwarf effective temperatures and surface gravities, as well as the spectral types of the companion stars.  We used a light curve modelling code for close binary systems to fit the eclipse profiles and the ellipsoidal modulation/reflection effect in the light curves, to further constrain the masses and radii of the components in all systems. All three DA white dwarfs have masses of $\\Mwd\\sim0.4-0.6\\,\\Msun$, in line with the expectations from close binary evolution. The DC white dwarf in SDSS\\,0303+0054 has a mass of $\\Mwd\\ga0.85\\,\\Msun$, making it unusually massive for a post-common envelope system. The companion stars in all four systems are M-dwarfs of spectral type M4 and later. Our new additions raise the number of known eclipsing WDMS binaries to fourteen, and we find that the average white dwarf mass in this sample is $<\\Mwd>=0.57\\pm0.16\\,\\Msun$, only slightly lower than the average mass of single white dwarfs. The majority of all eclipsing WDMS binaries contain low-mass ($<0.6\\,\\Msun$) secondary stars, and will eventually provide  valuable observational input for the calibration of the mass-radius relations of low-mass main sequence stars and of white dwarfs. ",
        "introduction": "White dwarfs and low-mass stars represent the most common types of stellar remnant and main sequence star, respectively, encountered in our Galaxy. Yet, despite being very common, very few white dwarfs and low mass stars have accurately determined radii and masses. Consequently, the finite temperature mass-radius relation of white dwarfs \\citep[e.g.][]{wood95-1,paneietal00-2} remains largely untested by observations \\citep{provencaletal98-1}. In the case of low mass stars, the empirical measurements consistently result in radii up to 15\\% larger and effective temperatures 400\\,K or more below the values predicted by theory \\citep[e.g.][]{ribas06-1, lopez-morales07-1}. This is most clearly demonstrated using low-mass eclipsing binary stars \\citep{bayless+orosz06-1}, but is also present in field stars \\citep{bergeretal06-1, moralesetal08-1} and the host stars of transiting extra-solar planets \\citep{torres07-1}.  Eclipsing binaries are the key to determine accurate stellar masses and radii \\citep[e.g.][]{andersen91-1, southworthetal07-4}. However, because of their intrinsic faintness, very few binaries containing white dwarfs and/or low mass stars are currently known.  Here, we report the first results of a programme aimed at the identification of eclipsing white dwarf plus main sequence (WDMS) binaries, which will provide accurate empirical masses and radii for both types of stars. Because short-period WDMS binaries underwent common envelope evolution, they are expected to contain a wide range of white dwarf masses, which will be important for populating the empirical white dwarf mass-radius relation. Eclipsing WDMS binaries will also be of key importance in filling in the mass-radius relation of low-mass stars at masses $\\la0.6\\,\\Msun$.  The structure of the paper is as follows. The target selection for this programme is described in Sect.\\,\\ref{s-targets}. In Sect.\\,\\ref{s-observations} we present our observations and data reduction in detail. We determine the orbital periods and ephemerides of the four eclipsing WDMS binaries in Sect.\\,\\ref{s-periods}, and measure the radial velocities of the secondary stars in Sect.\\,\\ref{s-k2}.  In Sect.\\,\\ref{s-spectralfit} we derive initial estimates of the stellar parameters from fitting the SDSS spectroscopy of our targets. Basic equations for the following analysis are outlined in Sect.\\,\\ref{s-equations}. In Sect.\\,\\ref{s-lcmodels} we describe our fits to the observed light curves, and present our results in Sect.\\,\\ref{s-results}. The past and future evolution of the four stars is explored in Sect.\\,\\ref{s-evolution}. Finally, we discuss and summarise our findings, including an outlook on future work in Sect.\\,\\ref{s-conclusions}. ",
        "conclusions": "\\label{s-conclusions}  \\begin{table*} \\centering \\setlength{\\tabcolsep}{0.7ex} \\begin{minipage}{185mm} \\caption{Physical parameters of the known eclipsing white dwarf main sequence binaries.  A dash sign means that no corresponding value was given by the respective authors. Errors on the parameters are also shown, if quoted by the authors. } \\label{allwdms} \\begin{tabular}{@{}lccccccccccl@{}} \\hline System & $P_{\\mathrm{orb}}$ [d] & \\Mwd [$\\Msun$] & \\Rwd [$\\Rsun$] & \\Twd [K] & log$g$ & \\Msec [$\\Msun$] & \\Rsec [$\\Rsun$] & $K_\\mathrm{WD}$ [\\kms] & $K_\\mathrm{sec}$ [\\kms] & Sp2 & Ref.\\\\ \\hline V471\\,Tau        & 0.521 & 0.79 & 0.01 & 32000 & - & 0.8 & 0.85 & - & 150.7$\\pm1.2$ & K2 & 1a \\\\ & 0.521 & 0.84$\\pm0.05$ & 0.0107$\\pm0.0007$ & 34500$\\pm1000$ & 8.31$\\pm0.06$ & 0.93$\\pm0.07$ & 0.96$\\pm0.04$ & 163.6$\\pm3.5$ & 148.46$\\pm0.56$ & K2 & 1b \\\\ RXJ2130.6+4710   & 0.521 & 0.554$\\pm0.017$ & 0.0137$\\pm0.0014$ & 18000$\\pm1000$ & 7.93\\tiny{$^{+0.07}_{-0.09}$} & 0.555$\\pm0.023$ & 0.534$\\pm0.017$ & 136.5$\\pm3.8$ & 136.4$\\pm0.8$ & M3.5-M4 & 2\\\\ DE\\,CVn          & 0.364 & 0.51\\tiny{$^{+0.06}_{-0.02}$} & 0.0136\\tiny{$^{+0.0008}_{-0.0002}$} & 8000$\\pm1000$ & 7.5 & 0.41$\\pm0.06$ & 0.37\\tiny{$^{+0.06}_{-0.007}$} & - & 166$\\pm4$ & M3V & 3 \\\\ GK\\,Vir          & 0.344 & 0.7$\\pm0.3$ & - & 50000 & - & - & - & - & - & M0-M6 & 4a \\\\ & 0.344 & 0.51$\\pm0.04$ & 0.016 & 48800$\\pm1200$ & 7.7$\\pm0.11$ & 0.1 & 0.15 & - & - & M3-M5 & 4b\\\\ SDSS\\,1212--0123 & 0.335 & 0.46-0.48 & 0.016-0.018 & $17700\\pm300$ & 7.5-7.7 & 0.26-0.29 & 0.28-0.31 & - & $181\\pm3$ & M4$\\pm1$ & 5 \\\\ SDSS\\,0110+1326  & 0.332 & 0.47$\\pm0.2$ & 0.0163-0.0175 & 25900$\\pm427$ & 7.65$\\pm0.05$ & 0.255-0.38 & 0.262-0.36 & - & 180-200 & M3-M5 & 6 \\\\ RR\\,Cae          & 0.303 & 0.365 & 0.0162 & 7000 & - & 0.089 & 0.134 & 47.9$^{*}$ & 195.8$^{*}$ & M6 & 7a \\\\ & 0.303 & 0.467 & 0.0152 & 7000 & - & 0.095 & 0.189 & - & 190$\\pm9$ & M6 & 7b \\\\ & 0.303 & 0.44$\\pm0.022$ & 0.015$\\pm0.0004$ & 7540$\\pm175$ & 7.61-7.78 & 0.183$\\pm0.013$ & 0.188-0.23 & 79.3$\\pm3$ & 190.2$\\pm3.5$ & M4 & 7c\\\\ SDSS\\,1548+4057  & 0.185 & 0.614-0.678 & 0.0107-0.0116 & 11700$\\pm820$ & 8.02$\\pm0.44$ & 0.146-0.201 & 0.166-0.196 & - & 274.7 & M5.5-M6.5 & 6 \\\\ EC13471-1258     & 0.150 & 0.77$\\pm0.04$ & - & 14085$\\pm100$ & 8.25$\\pm0.05$ & 0.58$\\pm0.05$ & 0.42$\\pm0.08$ & - & 241$\\pm8.1$ & M2-M4 & 8a \\\\ & 0.150 & 0.78$\\pm0.04$ & 0.011$\\pm0.01$ & 14220$\\pm300$ & 8.34$\\pm0.2$ & 0.43$\\pm0.04$ & 0.42$\\pm0.02$ & 138$\\pm10$ & 266$\\pm5$ & M3.5-M4 & 8b\\\\ CSS\\,080502$^\\dag$ & 0.149 & $0.35\\pm0.04$ & $0.02\\pm0.002$ & $17505\\pm516$ & $7.38\\pm0.12$ & 0.32 & 0.33 & - & - & M4 & 9,6\\\\ SDSS\\,0303+0054  & 0.134 & 0.878-0.946 & 0.0085-0.0093 & $<\\,$8000 & 8.4-8.6 & 0.224-0.282 & 0.246-0.27 & - & 339.7 & M4-M5 & 6 \\\\ NN\\,Ser          & 0.130 & 0.57$\\pm0.04$ & 0.017-0.021 & 55000$\\pm8000$ & 7-8 & 0.10-0.14 & 0.15-0.18 & - & 310$\\pm10$ & M4.7-M6.1 & 10a,b \\\\ & 0.130 & 0.54$\\pm0.05$ & 0.0189$\\pm0.001$ & 57000$\\pm3000$ & 7.6$\\pm0.1$ & 0.15$\\pm0.008$ & 0.174$\\pm0.009$ & 80.4$\\pm4.1$ & 289.3$\\pm12.9$ & M4.5-M5 & 10c \\\\ SDSS\\,1435+3733  & 0.125 & - & - & - & - & 0.15-0.35 & 0.17-0.32 & - & - & M4-M6 & 11 \\\\ & 0.125 & 0.48-0.53 & 0.0144-0.0153 & 12500$\\pm488$ & 7.62$\\pm0.12$ & 0.19-0.246 & 0.218-0.244 & - & 260.9 & M4-M5 & 6 \\\\ CSS080408$^\\ddag$ & - & $0.40\\pm0.05$ & $0.022\\pm0.003$ & $32595\\pm675$ & $7.36\\pm0.17$ & 0.26 & 0.27 & - & - & M5 & 9,6 \\\\ \\hline \\multicolumn{12}{p{\\textwidth}}{References: (1a)\\,\\citealt{young+nelson72-1}; (1b)\\,\\citealt{obrienetal01-1}; (2)\\,\\citealt{maxtedetal04-1}; (3)\\,\\citealt{vandenbesselaaretal07-1}; (4a)\\,\\citealt{greenetal78-1}; (4b)\\,\\citealt{fulbrightetal93-1}; (5)\\,\\citealt{nebotgomez-moranetal08-1}; (6)\\,this work; (7a)\\,\\citealt{bruch+diaz98-1}; (7b)\\,\\citealt{bruch99-1}; (7c)\\,\\citealt{maxtedetal07-1}; (8a)\\,\\citealt{kawkaetal02-1}; (8b)\\,\\citealt{odonoghueetal03-1}; (9)\\,\\citealt{drakeetal08-1}; (10a)\\,\\citealt{wood+marsh91-1}; (10b)\\,\\citealt{catalanetal94-1}; (10c)\\,\\citealt{haefneretal04-1}; (11)\\,\\citealt{steinfadtetal08-1}} \\\\ \\hline \\multicolumn{12}{p{\\textwidth}}{Notes: $^*$ Calculated from $K\\sin i$ provided by the authors; $^\\dag$ CSS080502:090812+060421\\,=\\,SDSS\\,J090812.04+060421.2, we determined the white dwarf and secondary star paremeters from  decomposing/fitting the SDSS spectrum (see Sect.\\,\\ref{s-spectralfit} and \\citealt{rebassa-mansergasetal07-1}), these parameters should be taken as guidance as no light curve modelling has been done so far; $^\\ddag$ CSS080408:142355+240925\\,=\\,SDSSJ142355.06+240924.3, also found by \\citet{rebassa-mansergasetal08-2}, white dwarf and secondary star parameters determined as for CSS080502.}\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}  Since \\citet{nelson+young70-1} reported the discovery of V471\\,Tau, as the first eclipsing white dwarf main sequence binary, only 7 additional similar systems have been discovered, the latest being SDSS\\,1435+3733 found by \\citet{steinfadtetal08-1}. Here, we presented follow-up observations of SDSS\\,1435+3733, which we independently identified as an eclipsing WDMS binary, and of three new discoveries: SDSS\\,0110+1326, SDSS\\,0303+0050, and SDSS\\,1548+4057. A fifth system has just been announced by our team \\citep{nebotgomez-moranetal08-1}, and two eclipsing WDMS binaries were identified by \\citet{drakeetal08-1} in the Catalina Real-Time Transient Survey. Table\\,\\ref{allwdms} lists the basic physical parameters of the fourteen eclipsing white dwarf main sequence binaries currently known.  Inspection of Table\\,\\ref{allwdms} shows that the orbital periods of the known eclipsing binaries range from three to 12 hours, with three of the systems presented here settling in at the short-period end.  The white dwarf masses cover a range $\\sim0.44-0.9\\,\\Msun$, with an average of $<\\Mwd>=0.57\\pm0.16\\,\\Msun$, which is only slightly lower than the average mass of single white dwarfs \\citep{liebertetal05-1, finleyetal97-1, koesteretal79-1}. This finding is somewhat surprising, as all these eclipsing WDMS binaries have undergone a common envelope evolution, which is expected to cut short the evolution of the primary star and thereby to produce a significant number of low-mass He-core white dwarfs.  Our estimates for the white dwarf masses in SDSS\\,0110+1326 and SDSS\\,1435+3733 suggest that they contain He-core white dwarfs. The white dwarf in SDSS\\,1548+4057 is consistent with the average mass of single white dwarfs. SDSS\\,0303+0054 contains a fairly massive white dwarf, potentially the most massive one in an eclipsing WDMS binary.  The secondary star masses are concentrated at very low masses, the majority of the fourteen systems in Table\\,\\ref{allwdms} have $\\Msec<0.6$\\,\\Msun, and $\\sim9$ systems (including the four stars discussed here) have $\\Msec\\la0.3\\,\\Msun$. This is a range where very few low-mass stars in eclipsing binaries are known \\citep{ribas06-1}, making the secondary stars in eclipsing WDMS binaries very valuable candidates for filling in the empirical mass-radius relation of low mass stars. This aspect will be discussed in a forthcoming paper.  The distribution of secondary star masses among the eclipsing PCEBs in Table~\\ref{allwdms} is subject to similar selection effects as the whole sample of known PCEBs. \\citet{schreiber+gaensicke03-1} found a strong bias towards late-type secondary stars in the pre-SDSS PCEB population, which arose as most of those PCEBs were identified in blue-colour surveys. The currently emerging population of SDSS PCEBs \\citep{silvestrietal06-1, rebassa-mansergasetal08-2} should be less biased as the $ugriz$\\,--\\,colour-space allows to identify WDMS binaries with a wide range of secondary spectral types and white dwarf temperatures.  Future work will need to be carried out on two fronts. On one hand, more accurate determination of the masses and radii will have to be made to unlock the potential that WDMS binaries hold for testing/constraining the mass-radius relation of both white dwarfs and low mass stars. The key to this improvement are high-quality light curves that fully resolve the white dwarf ingress/egress, and, if possible, the secondary eclipse of the M-dwarf by the white dwarf. In addition, accurate $K_\\mathrm{WD}$ velocities are needed, which will be most reliably obtained from ultraviolet intermediate resolution spectroscopy. On the other hand, it is clear that the potential of SDSS in leading to the identification of additional eclipsing WDMS binaries has not been exhausted. Additional systems will be identified by the forthcoming time-domain surveys, and hence growing the class of eclipsing WDMS binaries to several dozen seems entirely feasible."
    },
    "0812/0812.2022.txt": {
        "abstract": "% context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {We present the results from a comprehensive spectroscopic survey of the WINGS (WIde-field Nearby Galaxy-cluster Survey) clusters, a program called WINGS-SPE. The WINGS-SPE sample consists of 48 clusters, $22$ of which are in the southern sky and $26$ in the north. The main goals of this spectroscopic survey are: (1) to study the dynamics and kinematics of the WINGS clusters and their constituent galaxies, (2) to explore the link between the spectral properties and the morphological evolution in different density environments and across a wide range in cluster X-ray luminosities and optical properties.} % methods heading (mandatory) {Using multi-object fiber-fed spectrographs, we observed our sample of WINGS cluster galaxies at an intermediate resolution of 6-9 \\AA\\ and, using a cross-correlation technique, we measured redshifts with a mean accuracy of $\\sim45$\\,km\\,s$^{-1}$.} % results heading (mandatory) {We present redshift measurements for $6,137$ galaxies and their first analyses. Details of the spectroscopic observations are reported. The WINGS-SPE has $\\sim 30 \\%$ overlap with previously published data sets, allowing us to do both a complete comparison with the literature and to extend the catalogs.} % conclusions heading (optional), leave it empty if necessary {Using our redshifts, we calculate the velocity dispersion for all the clusters in the WINGS-SPE sample. We almost triplicate the number of member galaxies known in each cluster with respect to previous works. We also investigate the X-ray  luminosity vs. velocity dispersion relation for our WINGS-SPE clusters, and find it to be consistent with the form $L_x \\propto \\sigma_v^4$.} ",
        "introduction": "Rich clusters of galaxies have long been recognized as  valuable tools with which to study cosmology and galaxy formation. They are the most massive virialized objects in the universe and, as such, provide excellent laboratories for studying the influence of environment on galaxy formation and evolution, as manifested by the morphology-density relation (\\cite{Dressler80}, \\cite{Dressler97}, \\cite{Fasano00}, \\cite{Postman05}). The identification of the properties and content of clusters of galaxies is able to clarify their cosmic evolution since they can be detected at large distances. However clusters of galaxies are now recognized to not be simple relaxed structures; rather, they are evolving via merging processes in a hierarchical fashion from poor groups to rich clusters. In fact in the cluster outskirts, galaxies are still falling in for the first time, so it is possible also to explore environmental effects over a wide dynamic range in density.  Clusters contain large populations of galaxies at a common distance, which can be used to derive redshift-independent relative distance estimates (\\cite{Dressler87}). This allows the study of deviations from pure Hubble flow, i.e., the peculiar velocity field, and hence the dark matter distribution in the local universe (e.g., \\cite{LB88}; see also \\cite{dekel}; \\cite{strauss} for reviews).  Cluster velocity dispersions provide a measure of cluster mass (\\cite{fisher98}; \\cite{tran99}; \\cite{bor99a}; \\cite{lubin02}). The measurement of cluster velocity dispersions should be made using statistics insensitive to galaxy redshift outliers and the shape of the velocity distribution. However, the uncertainties in these studies are large, and there remains the possibility of systematic errors due to the heterogeneity of the spectroscopic data sets available for nearby clusters ($z<0.1$). In recent years,  a systematic investigation of a large sample of nearby clusters has become possible due to the advent of CCD mosaics and high multiplex multi-object fiber-fed spectrographs, allowing photometric and spectroscopic observations over large solid angles (\\cite{noaoI}, \\cite{fasano06}).  In the last few years, large redshift surveys such as the 2dF Galaxy Redshift Survey (hereafter 2dFGRS; \\cite{DePropris}) and especially the Sloan Digital Sky Survey (SDSS, \\cite{Goto}; \\cite{Bahcall}; \\cite{Miller}) have been the primary source for the compilation of data on nearby clusters of galaxies. However, both these surveys have their limitations. The 2dFGRS is only a spectroscopic survey while the SDSS, being a large area survey, is not deep enough to study the fainter end of the luminosity function. More recently, the NOAO Fundamental Plane Survey (\\cite{noaoI}) started to address these shortcomings, although its primary goal is not to study the evolution of galaxies, but rather to trace large scale velocity fields using the Fundamental Plane.  The need for accurate redshift measurements has also become evident from the great progress that has been made in recent years in the observations of the signatures of cluster merging processes. The presence of substructure, which is indicative of a cluster in an early phase of the process of dynamical relaxation or of secondary infall of clumps into already virialized clusters, occurs in a large percentage of clusters (\\cite{drshec}; \\cite{west}; \\cite{rhee}; \\cite{bird}; \\cite{escalera}; \\cite{west1}; \\cite{girardi97}; \\cite{solanes}; \\cite{biviano}; \\cite{burgett}; \\cite{flin}). Studying cluster substructure therefore allow us to investigate the process by which clusters form. In addition, it also enable us to better understand the mechanisms affecting galaxy evolution in clusters, which can be accelerated by the effects of a cluster-subcluster collision. An interesting result has been produced by our recent study (\\cite{ramella07})  in which we point out that the fraction of clusters with sub-clusters in the WINGS sample (73\\%) is higher than in most previous studies. In addition, in a following paper of the WINGS-SPE series (Cava et al. 2009, in prep.) a more detailed analysis of the sub-clustering in the WINGS sample will be included, exploiting and comparing different 2-D and 3-D approaches.  In this paper we present redshifts for galaxies in 48 clusters belonging to the Wide-field Imaging Nearby Galaxy-cluster Survey (WINGS) sample. In addition, we present cluster velocity dispersion measurements derived from our redshift data. These clusters comprise an almost complete X-ray selected sample of galaxy clusters at $z=0.04-0.07$ (see \\cite{fasano06} for details of the survey).  The main goal of the WINGS-SPE spectroscopic follow-up program is to supply a complete and uniform set of spectroscopic data such as redshift (present paper), line indices, and line widths useful to investigate the dynamics of the clusters and derive star formation histories, star formation rates (Fritz et al., \\cite{fritz}, and Fritz et al. in prep.), and other structural and physical properties for cluster galaxies. These new data will shed more light on the link between the evolution of star formation and galaxy morphology, as well as the dependence on the characteristics of the cluster and where the galaxies are located within the cluster. Furthermore, the availability of data on low redshift clusters can be used as a present-day reference for studies of distant clusters (see e.g. Poggianti et al. \\cite{poggianti06}).  This paper is organized as follows: In \\S 2 we give general information on the WINGS-SPE objectives and working strategy; in \\S 3 we present the spectroscopic observations and describe the reduction processes; in \\S 4 redshifts measurements and catalogs are presented, while in \\S 5 we check the data quality performing a comparison with data in the literature. Finally we draw our conclusions in \\S 6. ",
        "conclusions": "As part of the WINGS-SPE survey, we have carried out spectroscopic observations of galaxies in 48 clusters using the WHT/AF2 and AAT/2dF facilities. These observations have yielded redshifts for 6,137 galaxies, which have been used to derive precise redshifts and velocity dispersions for our clusters. By combining these data with those already available in the literature, we now have velocity dispersions for all the clusters in the WINGS sample.  In Table \\ref{table:catalogs} we present the complete and final set of redshift data now available for the WINGS galaxy clusters. A total of 3,647 galaxies turn out to be members of our clusters, thereby almost doubling the number of known members in this sample of nearby clusters. We have shown that our reduction and measurement procedures result in high quality redshift measurements. A comparison with data available in the literature to both check the accuracy and consistency of our measurements and to increase our overall redshift sample has been done. Using an iterative $3\\sigma$ clipping scheme, we have derived velocity dispersions for all the 48 WINGS-SPE clusters. The mean of the ratio of the number of galaxies we used to derive the cluster velocity dispersion and the number that were used for the literature values is $R=N_{wings}/N_{tab}=2.71$. This means that our velocity dispersion values are based on almost three times as many member galaxies than previous measurements. We found that the X-ray luminosity - velocity dispersion ($L_x$ - $\\sigma$) relation for our sample of 48 clusters has $L_x \\propto \\sigma^4$, although with a large scatter. Finally, we note that the WINGS clusters have a wide range of velocity dispersion values. The implied large range of masses therefore makes the WINGS cluster sample an unprecedented and unique dataset to study the processes affecting cluster galaxy evolution as a function of cluster mass. Future papers in this series will exploit this data set to perform a dynamical analysis of the WINGS sample of galaxy clusters and a substructure analysis."
    },
    "0812/0812.2984_arXiv.txt": {
        "abstract": "We investigate the cosmological evolution of large- and small-scale magnetic fields in galaxies at high redshifts. Results from simulations of hierarchical structure formation cosmology provide a tool to develop an evolutionary model of regular magnetic fields coupled to galaxy formation and evolution. Turbulence in protogalactic halos generated by thermal virialization can drive an efficient turbulent dynamo. The mean-field dynamo theory is used to derive the timescales of amplification and ordering of regular magnetic fields in disk and dwarf galaxies. For future observations with the SKA, we predict an anticorrelation at fixed redshift between galaxy size and the ratio between ordering scale and galaxy size. Undisturbed dwarf galaxies should host fully coherent fields at $z<1$, spiral galaxies at $z<0.5$. ",
        "introduction": "The observed polarized synchrotron emission and Faraday rotation showed the presence of regular large-scale magnetic fields with spiral patterns in the disks of nearby spiral galaxies (Beck 2005), which were successfully reproduced by mean-field dynamo theory (Beck et al 1996, Shukurov 2005). It is therefore natural to apply dynamo theory also in predicting the generation of magnetic fields in young galaxies at high redshifts.  We now have sufficient evidence that strong magnetic fields were present in the early Universe ($z<3$; Bernet et al. 2008, Seymour et al. 2008) and that synchrotron emission from distant galaxies should be detected with future radio telescopes such as the Square Kilometre Array (SKA). The SKA will allow us to observe an enormous number of distant galaxies at similar resolution to that achievable for nearby galaxies today (van der Hulst et al. 2004). The formation and evolution of regular large-scale magnetic fields is intimately related to the formation and evolution of disks in galaxies in terms of geometrical and physical parameters. A more robust understanding of the history of magnetism in young galaxies may help to solve fundamental cosmological questions about the formation and evolution of galaxies (Gaensler et al. 2004). ",
        "conclusions": ""
    },
    "0812/0812.4877_arXiv.txt": {
        "abstract": "We have mapped the central region of the HH 111 protostellar system in 1.33 mm continuum, \\cCO{} (J=2-1), \\bCO{} (J=2-1), and SO (\\SOta) emission at $\\sim$ \\arcs{3} resolution with the Submillimeter Array. There are two sources, VLA 1 (=IRAS 05491+0247) and VLA 2, with the VLA 1 source driving the HH 111 jet. Thermal emission is seen in 1.33 mm continuum tracing the dust in the envelope and the putative disks around the sources. A flattened, torus-like envelope is seen in \\cCO{} and \\bCO{} around the VLA 1 source surrounding the dust lane perpendicular to the jet axis, with an inner radius of $\\sim$ 400 AU (\\arcs{1}), an outer radius of $\\sim$ 3200 AU (\\arcs{8}), and a thickness of $\\sim$ 1000 AU (\\arcsa{2}{5}). It seems to be infalling toward the center with conservation of specific angular momentum rather than with a Keplerian rotation as assumed by \\citet{Yang1997}. An inner envelope is seen in SO, with a radius of $\\sim$ 500 AU (\\arcsa{1}{3}). The inner part of this inner envelope, which is spatially coincident with the dust lane, seems to have a differential rotation and thus may have formed a rotationally supported disk. The outer part of this inner envelope, however, may have a rotation velocity decreasing toward the center and thus represent a region where an infalling envelope is in transition to a rotationally supported disk. A brief comparison with a collapsing model suggests that the flattened, torus-like envelope seen in \\cCO{} and \\bCO{} could result from a collapse of a magnetized rotating toroid. ",
        "introduction": "Stars are formed inside molecular cloud cores by means of gravitational collapse. The details of the process, however, are complicated by the presence of magnetic fields and angular momentum. As a result, in addition to infall (or collapse), rotation and outflow are also seen toward star-forming regions. In theory, a rotationally supported disk is expected to form inside a collapsing core around a protostar, from which part of the material is accreted by the protostar and part is ejected away. Observationally, however, when and how a rotationally supported disk is actually formed are still unclear, because of the lack of detailed kinematic studies inside the collapsing core.  A rotationally supported disk has been seen with a radius of $\\sim$ 500 AU in the late (i.e., Class II or T Tauri) phase of star formation \\cite[see, e.g.,][]{Simon2000}. Formation of such a disk must have started early in the Class 0 phase. However, no such disk has been confirmed in the Class 0 phase, probably because it is still too small ($<$ 100 AU) for a detailed kinematic study with current instruments \\cite[see, e.g.,][]{Lee2006,Jorgensen2007}. In order to investigate how a rotationally supported disk is formed inside a collapsing core, we have mapped the protostellar system HH 111, which is in the Class I phase and is thus expected to have a bigger disk than that in the Class 0 phase. In this system, a Class I source IRAS 05491+0247 is found deeply embedded in a compact molecular cloud core in the L1617 cloud of Orion, driving the HH 111 jet \\citep{Reipurth1989}. The distance to the Orion Nebula Cluster has been recently refined to be 400 pc \\citep{Sandstrom2007,Menten2007}, the same that is generally adopted for the Orion~B cloud \\cite[see, e.g.,][]{Gibb2008} with which L1617 is associated. The HH 111 system has traditionally been assumed to be at a distance of 450~pc, but we here adopt this  new distance, at which the source has a luminosity of $\\sim$ 20 \\solarlum{} \\citep{Reipurth1989}. A rotating envelope has been seen around the source inside the core \\citep{Stapelfeldt1993} and may have infall motion toward the source \\citep{Yang1997}. The jet complex is highly collimated with a two-sided length of $\\sim$ 6.7 pc \\cite[][corrected for the new distance]{Reipurth1997}, and it produces a collimated molecular outflow around it \\citep{Nagar1997,Lee2000}. The source has also been detected with the VLA, associated with a tiny radio jet \\citep{Reipurth1999}, and a tiny dusty disk (with a radius of 30 AU) \\citep{Rodriguez2008} that can be the innermost part of a rotationally supported disk. Thus, this system is one of the best candidates to search for a rotationally supported disk and investigate how it is formed inside a collapsing core. Here, we present observations of the envelope in 1.33 mm continuum, \\bCO ($J=2-1$), \\cCO ($J=2-1$), and SO ($N_J=5_6-4_5$) emission obtained with the 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.} \\citep{Ho2004}. ",
        "conclusions": "We have mapped the central region of the protostellar system HH 111 in 1.33 mm continuum, \\cCO{} (J=2-1), \\bCO{} (J=2-1), and SO (\\SOta) emission at $\\sim$ \\arcs{3} resolution with the Submillimeter Array. There are two sources, VLA 1 (=IRAS 05491+0247) and VLA 2, with the VLA 1 source driving the HH 111 jet. Thermal emission is seen in 1.33 mm continuum tracing the dust in the envelope and the putative disks around the sources. A flattened, torus-like envelope is seen in \\cCO{} and \\bCO{} around the VLA 1 source surrounding the dust lane perpendicular to the jet axis, with an inner radius of $\\sim$ 400 AU (\\arcs{1}), an outer radius of $\\sim$ 3200 AU (\\arcs{8}), and a thickness of $\\sim$ 1000 AU (\\arcsa{2}{5}). It seems to be infalling toward the center with conservation of specific angular momentum rather than with a Keplerian rotation as assumed by \\citet{Yang1997}. An inner envelope is seen in SO, with a radius of $\\sim$ 500 AU (\\arcsa{1}{3}). The inner part of this inner envelope, which is spatially coincident with the dust lane, seems to have a differential rotation and thus may have formed a rotationally supported disk. The outer part of this inner envelope, however, may have a rotation velocity decreasing toward the center and thus represent a region where an infalling envelope is in transition to a rotationally supported disk. A brief comparison with a collapsing model suggests that the flattened, torus-like envelope seen in \\cCO{} and \\bCO{} could be a pseudodisk (a nonrotationally supported torus) resulting from a collapse of a magnetized rotating toroid."
    },
    "0812/0812.1349_arXiv.txt": {
        "abstract": "We investigate the spectrum of vector modes today which is generated at second order by density perturbations. The vector mode background that is generated by structure formation is small but in principle it contributes to the integrated Sachs-Wolfe effect, to redshift-space distortions and to weak lensing. We recover, clarify and extend previous results, and explain carefully why no vorticity is generated in the fluid at second order. The amplitude of the induced vector mode in the metric is around 1\\% that of the first-order scalars on small scales.  We also calculate the power spectrum and the energy density of the vector part of the shear at second order. ",
        "introduction": "Cosmological vector perturbations at linear order satisfy an evolution equation and a momentum constraint equation which can be found by calculating the $i-j$ and $0-i$ parts of the Einstein field equations respectively. In the case of a perfect fluid (or scalar field) there is no source in the vector evolution equation and its solutions decay as $1/a^2$. In the standard models of structure formation, inflation does not generate vector perturbations, and we can therefore ignore vector perturbations at first order. However, vector perturbations will be generated at nonlinear order by the growth of density perturbations.  Here we revisit the analysis of this cosmological vector background (see Refs.~\\cite{Matarrese:1997ay,Mollerach:1997up,Noh:2003yg, Mollerach:2003nq,Tomita:2005et,Mena:2007ve,Lu:2007cj} for previous work). As we show below, the amplitude of the induced vectors is surprisingly large on small scales, when compared to the amplitude of the first-order Newtonian potential. In order to investigate whether this is important dynamically we calculate the vector part of the shear and compare it to the usual first-order scalar shear.     We consider perturbations of a flat Robertson-Walker background up to second order: ${g}_{\\mu \\nu} = \\bar{g}_{\\mu \\nu} + \\delta^{(1)} g_{\\mu \\nu} + \\delta^{(2)} g_{\\mu \\nu}$. At first order, $\\delta^{(1)} g_{\\mu \\nu}$ only contains scalar perturbations: we neglect the tensor perturbations, and vector perturbations are not generated at first order in the standard model. For the second-order perturbations, $\\delta^{(2)} g_{\\mu \\nu}$, we project out the vector modes. In the Poisson gauge~\\cite{Matarrese:1997ay} \\begin{equation} ds^2 = -a^2 \\left[ 1 + 2 \\Phi_{(1)} \\right] d\\eta^2 - a^2 S_{i}^{(2)}dx^id\\eta + a^2  \\left [1 - 2 \\Phi_{(1)} \\right ] d\\vec{x}\\,^2, \\label{perturbed metric} \\end{equation} where $\\Phi_{(1)}$ is the first-order Newtonian potential (we assume zero scalar anisotropic stress at first order), and $S^{i}_{(2)}$ describes the gauge-invariant second-order vector modes, so that $\\partial_{i} S_{(2)}^{i} = 0$. We effectively ignore the second-order scalar mode $\\Phi_{(2)}$ since we are interested only in the second-order vector modes: second-order modes can be consistently split into scalar and vector, but the equations for the second-order vector modes will contain source terms that are quadratic in the first-order scalar modes. In what follows we will drop the order indices when there is no ambiguity. With this in mind, the fluid four-velocity is given by \\be u_\\mu=a \\left[-1-\\Phi+{1\\over 2}\\Phi^2-{1\\over2}v_{(1)}^jv_j^{(1)}\\,,~~ v^{(1)}_i+ {1\\over2} \\left\\{ v^{(2)}_i-S_i\\right\\} -2\\Phi v^{(1)}_i \\right], \\label{4vel} \\ee where $v^{(1)}_i=\\partial_iv_{(1)}$ and $\\partial_i v_{(2)}^i=0$.   The background dynamics are given by \\be \\mathcal{H}^2 = \\frac{8}{3} \\pi G a^2 \\rho + \\frac{1}{3} a^2 \\Lambda\\,,~~~  \\mathcal{H}' = -4\\pi G (1 + w)a^2\\rho + \\mathcal{H}^2\\,,\\label{bg} \\ee where $w=p/\\rho=\\,$const (later we specialize to $w=0$). The first-order perturbed field equations lead to~\\cite{Kodama} \\be v_{(1)} &=& - { \\Phi' + \\mathcal{H}  \\Phi \\over 4\\pi G \\rho a^2 (1 + w)}\\,, \\label{first order velocity}\\\\ \\delta \\rho &=& {{\\nabla^2} \\Phi-{3 \\mathcal{H}}( \\Phi'  +{\\cal H}\\Phi) \\over 4\\pi Ga^2} \\,, \\label{first order rho} \\ee where the Newtonian potential obeys the evolution equation \\begin{equation} \\Phi''+ 3 \\mathcal{H} (1 + w) \\Phi' +  \\left [ (1 + w)\\Lambda a^2 - w \\nabla^2 \\right] \\Phi= 0\\,. \\label{equation of motion for Bardeen potential} \\end{equation}  All linear scalar modes are determined by $\\Phi$, and terms quadratic in $\\Phi$ and its derivatives will source the second-order vector modes. Therefore we will require the power spectrum ${\\cal P}_\\Phi$. In the matter era ($w=0$), neglecting $\\Lambda$, we have $a=a_0(\\eta/\\eta_0)^2$, and the solution in Fourier space is $\\Phi_m({k},\\eta)=A_m({k})$, where we remove the decaying mode. Then the power spectrum today is computed using the growth suppression and transfer functions. (See Appendix~A for more details.)  Vector perturbations typically produce vorticity and a transverse shear in the fluid four-velocity, Eq.~\\eqref{4vel}. In order to compute them, we need the covariant definitions (see~\\cite{Tsagas:2007yx} for a recent review): \\be {\\omega}_{\\mu\\nu} &=& \\delta^{(1)} \\omega_{\\mu\\nu} + \\delta^{(2)} \\omega_{\\mu\\nu} = {h}^{\\alpha}_{[\\mu} {h}^{\\beta}_{\\nu]} {u}_{\\alpha;\\beta}\\,,\\label{vor}\\\\ {\\sigma}_{\\mu\\nu} &=& \\delta^{(1)} \\sigma_{\\mu\\nu} + \\delta^{(2)} \\sigma_{\\mu\\nu} = \\left\\{ {h}^{\\alpha}_{(\\mu} {h}_{\\nu)}^{\\beta} - \\frac{1}{3} {h}_{\\mu\\nu} {h}^{\\alpha \\beta} \\right\\}{u}_{\\alpha;\\beta}\\,, \\label{shear} \\ee where $h_{\\mu\\nu}=g_{\\mu\\nu}+u_\\mu u_\\nu$ is the projector into the instantaneous fluid rest space, and ${\\omega}_{\\mu\\nu}u^\\nu =0= {\\sigma}_{\\mu\\nu}u^\\nu $. The vorticity is a purely vector quantity. The vector part of the second order shear is defined via (see Appendix D for how to get $\\sigma_{j}$ from $\\delta^{(2)} \\sigma_{ij}$) \\be \\label{vshear} \\delta^{(2)} \\sigma_{ij}=a\\partial_{(i}\\sigma_{j)}\\,, ~~~ \\partial_i\\sigma^i=0\\,. \\ee At first order, there is only scalar shear: \\be \\delta^{(1)} \\omega_{ij} &=& 0 \\,, \\\\ \\delta^{(1)} \\sigma_{ij}& =& a\\left(\\partial_i \\partial_j   - \\frac{1}{3}{\\delta}_{ij} \\nabla^2\\right) v_{(1)} \\,. \\label{first order shear with velocity} \\ee ",
        "conclusions": "We have computed the power spectra for the metric vector perturbations and vector shear at second order, generated by first-order scalar perturbations in a $\\Lambda$CDM model. In addition, we used a covariant approach to show explicitly how vorticity is not generated in a perfect fluid at any perturbative order by first-order scalar perturbations, so that there is no vorticity in the matter. In order to generate vorticity, one requires either an imperfect fluid, or momentum exchange between the fluid and another fluid. Momentum exchange (via Compton and Coulomb interactions) between electrons, protons and photons in the radiation and recombination eras can generate vorticity and magnetic fields at second order~\\cite{Matarrese:2004kq,Gopal:2004ut,Takahashi:2005nd, Ichiki:2006cd,Siegel:2006px,Kobayashi:2007wd,Maeda:2008dv}.  In order to obtain the vector quantities $S^i$ and $v_{(2)}^i$, we used the vanishing of vorticity [Eq.~\\eqref{zerov}] and the $0i$ Einstein constraint [Eq.~\\eqref{second order velocity}]. Alternatively, one could also use the $ij$ Einstein equation [Eq.~\\eqref{vector evolution equation}] and the momentum conservation equation, which has the form \\be \\label{momentum conservation equation} \\left\\{v_{(2)i}-S_{i}\\right\\}'+\\mathcal{H}\\left(1-3c_s^2\\right) \\left\\{ v_{(2)i}-S_{i}\\right\\} = M_i\\,, \\ee where the source term $M^i$ is given in Appendix~B.  The cosmological background of vector modes is small, especially if measured in terms of the dimensionless shear density, as shown in Figs.~\\ref{power spectrum of vector modes} and \\ref{fig2}. However, given that the amplitude of the vector modes in the metric is as large as $\\sim\\,$1\\% of the metric first-order scalar modes, in principle these vector modes will have an effect on various cosmological observations. In particular:  \\begin{itemize}  \\item Redshift-space distortions~\\cite{Smith:2007sb}: the divergenceless velocity $v_{(2)}^i$ will make a contribution to radial peculiar velocities and thus to redshift-space distortions.  \\item Large-angle CMB temperature anisotropies~\\cite{Mollerach:1997up,Tomita:2005et}: the vector modes will contribute to the Doppler and integrated Sachs-Wolfe effects: \\be \\delta^{(2)}T = {1\\over2}\\left\\{v^{(2)}_i-S_i\\right\\}e^i\\Big|_E^O + {1\\over2}\\int_E^O d\\lambda\\,\\partial_i S_j \\,e^i e^j\\,. \\ee  \\item Weak lensing~\\cite{sef,Durrer:1994uu}: vector modes produce a deflection angle \\be \\vec\\alpha = \\int_E^O d\\lambda \\left(\\vec\\nabla \\times \\vec S \\right)\\times \\vec e\\,. \\ee  \\item CMB polarization~\\cite{Kamionkowski:1996zd,Mollerach:2003nq}: the vector modes will leave a characteristic imprint on CMB polarization.   \\end{itemize}  Further work is needed to compute the size of these vector corrections. They are likely to be significant mainly below $\\sim 10\\,$Mpc, but in this regime the scalar non-linear effects are important and likely to dominate.  The vector degree of freedom forms an integral part of the perturbative expansion when one goes beyond linear order. At the order we have considered, vectors must be present essentially through a constraint in the field equations arising at order $\\Phi^2$, even though vectors have no independent propagating degrees of freedom. This is distinct from the intrinsically propagating degree of freedom in the scalar-induced gravitational wave background~\\cite{ACW,BIST,MBMR}. We have shown that the vector mode background has maximum power around the equality scale, similar to the induced gravitational wave background, and the metric vector modes achieve their maximal fraction of the linear metric scalar modes on scales below the Silk scale. The size of the vector contribution to the full non-linear power spectrum relevant for structure formation remains to be calculated.     \\[ \\]{\\bf Acknowledgements:}\\\\ We thank Bruce Bassett, Marco Bruni, Ruth Durrer, Cyril Pitrou, Jean-Phillipe Uzan and David Wands for useful discussions and comments.  THCL, KA and CC are supported by the National Research Foundation (South Africa), and RM is supported by STFC (UK). THCL and CC acknowledge financial support from the University of Cape Town. KA is additionally supported by the Italian {\\it Ministero Degli Affari Esteri-DG per la Promozione e Cooperazione Culturale} under the joint Italy/ South Africa Science and Technology agreement. THCL, KA and CC thank the ICG Portsmouth for hospitality during various stages of the work; their visits were partly supported by a joint Royal Society (UK) - National Research Foundation (South Africa) grant UID 65329.   \\appendix"
    },
    "0812/0812.1663_arXiv.txt": {
        "abstract": "{ The Helios measurements of the angular momentum flux $L$ of the \\green{fast} solar wind \\green{lead to a tendency} for the fluxes associated with individual ion angular momenta of protons and alpha particles, $L_p$ and $L_\\alpha$, to be negative (i.e., in the sense of counter-rotation with the Sun). However, \\green{the opposite holds for the slow wind}, and the overall particle contribution $L_P = L_p+L_\\alpha$ tends to exceed the magnetic contribution $L_M$. These two aspects are at variance with previous models.} {We examine whether introducing realistic ion temperature anisotropies can resolve \\green{this} discrepancy.} {From a general set of multifluid transport equations with gyrotropic species pressure tensors, we derive the equations governing both the meridional and azimuthal dynamics of outflows from magnetized, rotating stars. The equations are not restricted to radial flows in the equatorial plane but valid for general axisymmetric winds that include two major ion species. The azimuthal dynamics \\green{are} examined in detail, using the empirical meridional flow profiles for the solar wind, constructed mainly according to measurements made in situ.} {The angular momentum flux $L$ is determined by the requirement that the solution to the total angular momentum conservation law is unique and smooth in the vicinity of the Alfv\\'en point, defined as where the combined Alfv\\'enic Mach number $M_T=1$. $M_T$ has to consider the contributions from both protons and alpha particles. Introducing realistic ion temperature anisotropies may introduce a change of up to $10\\%$ in $L$ and up to $\\sim 1.8$\\velunits\\ in azimuthal speeds of individual ions between 0.3 and 1~AU, compared with the isotropic case. The latter has \\green{strong} consequences on the relative importance of $L_P$ and $L_M$ in the angular momentum budget.} {However, introducing ion temperature anisotropies cannot resolve the discrepancy between in situ measurements and model computations. For the \\green{fast-wind} solutions, while in extreme cases $L_P$ may become negative,  \\green{$L_p$ never does}. On the other hand, for the slow solar wind solutions examined, $L_P$ never exceeds $L_M$\\green{,} even though $L_M$ may be \\green{less} than the individual ion contribution\\green{,} since \\green{$L_p$ and $L_\\alpha$ always have opposite signs for the slow and fast wind alike}. } ",
        "introduction": "The angular momentum loss of a rotating star due to its outflow influences the rotational evolution of the star \\green{considerably}, and is therefore of astrophysical significance in general~\\citep[see e.g.,][]{WD_67, BelcherMacGregor_76, MestelSpruit_87, Bouvier_etal_97}. However, direct tests of in situ measurements against theories such as \\green{those} presented by \\citet{WD_67} are \\green{only possible} for the present Sun. A substantial number of studies have been conducted and were compiled in the comprehensive paper by \\citet{Pizzo_etal_83}, who themselves paid special attention to the Helios measurements of specific angular momentum fluxes. The measurements, further analyzed by \\citet{MarschRichter_84a}, are unique in that they allow the individual ion contribution from protons $L_p$ and alpha particles $L_\\alpha$ to the solar angular-momentum loss rate {per steradian} $L$ \\green{to be examined}. For instance, despite the significant scatter, the data exhibit a distinct trend for ${L}_p$ to be positive (negative) for solar winds with proton speeds $v_p$ below (above) 400\\velunits. A similar trend for ${L}_\\alpha$ is also found on average. The magnetic contribution ${L}_M$, on the other hand, is remarkably constant. A mean value of ${L}_M = 1.6 \\times 10^{29}$\\specangunits\\ can be quoted for the solar winds of all flow speeds and throughout the region from 0.3 to 1~AU. For comparison, the mean values of angular momentum fluxes carried by ion flows in the slow solar wind are ${L}_p = 19.6$ and ${L}_\\alpha = 1.3$~$\\times 10^{29}$\\specangunits~\\green{\\citep[see Table~II of][]{Pizzo_etal_83}}. The overall particle contribution to $L$ is then ${L}_P = {L}_p + {L}_\\alpha= 20.9 \\times 10^{29}$\\specangunits, which tends to be larger than ${L}_M$. It is noteworthy that a more recent study by \\citet{Scherer_etal_01} showed how examining the \\green{long-term} variation of the non-radial components of the solar wind velocity and the corresponding angular momentum fluxes can help us \\green{understand the heliospheric magnetic field better}.   Alpha particles should be placed on the same footing as protons from the perspective of solar wind \\green{modeling}, given their non-negligible abundance and the fact that there tends to exist a substantial differential speed $v_{\\alpha p}\\equiv |\\vec{v}_{\\alpha p}|\\mbox{sign}(|\\vec{v}_\\alpha | -|\\vec{v}_p|)$. As shown by the Helios measurements, a $v_{\\alpha p}$ amounting to up to $20\\%-30\\%$ of the local proton speed may occur in both the fast and slow solar \\green{winds}~\\citep{Marsch_etal_82a, Marsch_etal_82b}, with the latter being exemplified by an event \\green{that} took place on day 117 of 1978, \\green{when} a positive $v_{\\alpha p}\\sim 100$\\velunits\\ was found at 0.3~AU~\\citep{Marsch_etal_81}. \\green{That} on the average $v_{\\alpha p}\\approx 0$ in the slow wind simply reflects that the events with positive and negative $v_{\\alpha p}$ occur with nearly equal frequency~\\citep{Marsch_etal_82a}. As for the alpha abundance relative to protons, a value of $4.6\\%$ ($0.4\\%-10\\%$) is \\green{well-established} for the fast (slow) solar wind~\\cite[e.g.,][]{McComas_etal_00}. Therefore alpha particles can play an important role as far as the energy and linear momentum balance of the solar wind are concerned. When it comes to the problem of angular momentum transport, it was shown that in interplanetary space not only the angular momentum flux carried by the alpha particles $L_\\alpha$ but also that convected by the protons $L_p$ are determined by the terms associated with $v_{\\alpha p}$~\\citep{LiLi06}. This essentially derives from the requirement that the proton-alpha velocity difference vector be aligned with the instantaneous magnetic field. As a consequence, these terms have no contribution to the overall angular momentum flux convected by the ion flow $L_P$, which turns out to be smaller than $L_M$ in all the models examined in the parameter study by \\citet{Li_etal_07}. This, together with the fact that $L_p$ is always positive (i.e., in the sense of corotation with the Sun), is at variance with the Helios measurements.  A possible means to reconcile the measurements and the model computation is to incorporate the species temperature anisotropies. This is because the total pressure tensor $\\tens{P}=\\sum_s \\tens{p}_s$ summed over all species $s$ participates in the problem of angular momentum transport via the component $P^\\Delta =P^\\parallel-P^\\perp$ where $\\parallel$ and $\\perp$ are relative to the magnetic field $\\vec{B}$~\\citep[see e.g.,][hereafter referred to as W70]{Weber_70}. While the overall loss rate per steradian $L$ may not be significantly altered, the azimuthal speed of the solar wind and therefore the particle part of $L$ may be when compared with the isotropic case. Note that in the treatment of W70 the solar wind was seen as a bulk flow and the ion species are not distinguished. On the other hand, the formulation by \\citet{LiLi06} did not take into account the pressure anisotropy, which is a salient feature of the velocity distribution functions for both protons and alpha particles as revealed by the Helios measurements~\\citep{Marsch_etal_82a,Marsch_etal_82b}. It therefore remains to be seen how introducing the pressure anisotropy influences individual ion azimuthal speeds. Moreover, the simple, prescribed functional form for $P^\\Delta$ assumed in W70 needs to be updated in light of the more recent particle measurements.  The aim of the present paper is to extend the W70 study in three ways. First, we shall follow a multicomponent approach and examine the angular momentum transport in a solar wind comprising protons, alpha particles and electrons where a substantial proton-alpha particle velocity difference exists. Second, although following W70 we use a prescribed form of $P^\\Delta$ for simplicity, this prescription is based on the Helios measurements, and also takes into account other in situ and remote sensing measurements. Third, unlike W70 where the model equations are restricted to the equatorial plane, the equation set we shall derive is appropriate for a rather general axisymmetrical, time-independent, multicomponent, thermally anisotropic flow emanating from a magnetized rotating star. We note that a similar set of equations, which was also restricted to radial flows, was derived by \\citet{I84} who worked in the corotating frame of reference and neglected the azimuthal dynamics altogether. The functional dependence on the radial distance and flow speed of the magnetic spiral angle was prescribed instead. His approach is certainly justifiable for the present Sun, but a \\green{self-consistent treatment of the azimuthal dynamics is required when flows from other stars are examined. This is because many stars either have a stronger magnetic field or rotate substantially faster than the Sun.}  The paper is organized as follows. We start with section~\\ref{sec_mathform} where a description is given for the general multifluid, gyrotropic transport equations, based on which the azimuthal dynamics of the multicomponent solar wind is examined. Then section~\\ref{sec_presflow} describes the adopted meridional magnetic field and flow profiles. The numerical solutions to the angular momentum conservation law are given in section~\\ref{sec_numres}. In section~\\ref{sec_discussion}, we shall discuss how examining the angular momentum transport in a multicomponent solar wind can also shed some light on the spectra of ion velocity fluctuations induced by Alfv\\'enic activities. Finally, section~\\ref{sec_summary} summarizes the results. The equations of and a discussion on the poloidal dynamics are presented in the appendix. ",
        "conclusions": "\\label{sec_discussion} As demonstrated by~\\citet{LiLi08}, the discussion on the angular momentum transport also allows us to say a few words on the frequency spectra $S_k (f)$ ($k=p, \\alpha$) of the ion velocity fluctuations during Alfv\\'enic activities in the fast solar wind in the super-Alfv\\'enic portion where $M_T^2 \\gg 1$. This is due to the well-known change of the properties of Alfv\\'enic fluctuations around some $f_c \\approx v_{cm, a}/(4\\pi r_a)$, where $v_{cm, a}$ is the speed of center of mass evaluated at the Alfv\\'en point $r_a$~\\citep[see e.g., ][]{HO80, LiLi08}. For typical fast wind parameters, $f_c \\approx 0.5-1 \\times 10^{-5}$\\frequnits. While the fluctuations with frequencies $f\\gtrsim f_c$ are genuinely wave-like and may be described by the WKB limit given the slow spatial variation of flow parameters in the region in question, those with $f\\lesssim f_c$ behave in a quasi-static manner and may be described by the solutions to the angular momentum conservation law which also governs the zero-frequency fluctuations. As shown by \\citet{LiLi08} who neglected the species temperature anisotropy, in the region $r\\gtrsim 0.2$~AU which will be explored by the Solar Orbiter and Solar Probe, the ratio of the alpha to proton velocity fluctuation amplitude $\\delta u_\\alpha/\\delta u_p$ can be an order-of-magnitude larger for $f<f_c$ than for $f>f_c$. Hence one may expect that, if the proton velocity fluctuation spectrum $S_p (f)$ is somehow smooth around $f_c$, then the alpha one $S_\\alpha (f)$ will show an apparent spectral break. Now let us revisit this problem in light of the discussion presented in this paper and see what changes the pressure anisotropies may introduce.  \\begin{figure} \\centering \\includegraphics[width=.45\\textwidth]{f06.eps} \\vskip 0.5cm \\caption{ Radial dependence of the ratio of the alpha to the proton velocity fluctuation amplitudes $\\delta u_\\alpha/\\delta u_p$ induced by Alfv\\'enic activities in super-Alfv\\'enic portions of the fast solar wind. The dashed curves correspond to the isotropic model, while the hatched areas give the possible range $\\delta u_\\alpha/\\delta u_p$ may occupy when the parameters $\\Gamma_{pE}$ and $\\Gamma_{\\alpha E}$ vary in the ranges given in text. Both the zero-frequency (upper portion) and WKB (lower) estimates are given. } \\label{fig_discuss_spectra} \\end{figure}  Restrict ourselves to either the high-latitude region or the region inside say $100$~$R_\\odot$ such that the magnetic field may be seen as radial. Furthermore, suppose that the waves are propagating parallel to the magnetic field in the empirical fast wind profiles detailed in section~\\ref{sec_presflow}. Figure~\\ref{fig_discuss_spectra} presents the radial dependence of $\\delta u_\\alpha/\\delta u_p$ in the region between $40$ and $100$~$R_\\odot$ for both the zero-frequency (upper part) and WKB (lower part) solutions. For comparison, the dashed curves represent the corresponding results in the isotropic model. To construct Fig.\\ref{fig_discuss_spectra}, all the possible values of $\\Gamma_{pE}$ and $\\Gamma_{\\alpha E}$ have been examined. As a result, at any radial location the ratio $\\delta u_\\alpha/\\delta u_p$ varies from model to model, and the range in which this ratio may occupy is given by the hatched area. The zero-frequency solutions are obtained by solving Eq.(\\ref{eq_tanp_standard}), while for hydromagnetic WKB Alfv\\'en waves it is well known that $\\delta u_\\alpha /\\delta u_p = |(v_{ph}-v_{\\alpha})/(v_{ph}-v_{p})|$, where $v_{ph}$ is the wave phase speed and given by \\citep[e.g.,][]{BarnesSuffolk_71, I84} \\begin{eqnarray*} v_{ph} = v_{cm} + \\sqrt{v_{A}^2\\left(1-\\frac{4\\pi P^\\Delta}{B^2}\\right)-\\hat{\\rho}_\\alpha \\hat{\\rho}_p v_{\\alpha p}^2} \\hspace*{0.5cm}, \\label{eq_phase_alfven} \\end{eqnarray*} in which $v_{cm} = \\hat{\\rho}_p v_{p} + \\hat{\\rho}_\\alpha v_{\\alpha}$ is the speed of center of mass, and $\\hat{\\rho}_k = \\rho_k/\\rho$ ($k=p, \\alpha$) defines the fractional ion mass density.  From Fig.\\ref{fig_discuss_spectra} one can see that the zero-frequency and WKB solutions are well separated from each other, in the isotropic and anisotropic cases alike. For the isotropic model, $\\delta u_\\alpha/\\delta u_p$ in the zero-frequency case increases monotonically from $3.79$ at 40~$R_\\odot$ to 4.68 at 100~$R_\\odot$. On the other hand, in the WKB case it decreases first from 0.22 at 40~$R_\\odot$ and attains its minimum of $0.083$ at $60.3$~$R_\\odot$ and then increases to 0.15 at 100~$R_\\odot$. The difference in $\\delta u_\\alpha/\\delta u_p$ between the zero-frequency and WKB solutions may be slightly smaller in the anisotropic than in the isotropic case for some combinations of $[\\Gamma_{p E}, \\Gamma_{\\alpha E}]$, but the difference is still quite significant. From this we can conclude that, with realistic ion temperature anisotropies included, the alpha velocity fluctuation spectrum $S_\\alpha (f)$ during Alfv\\'enic activities will also show an apparent break near $f_c$, if the proton one $S_p (f)$ is smooth there. This break is entirely a linear property, and has nothing to do with the nonlinearities that may also shape the fluctuation spectra.     This study has been motivated by the apparent lack of an analysis on the angular momentum transport in a multicomponent solar or stellar wind with differentially flowing ions and species temperature anisotropy. Moreover, there has been an outstanding discrepancy between available measurements and models concerning the relative importance of the particle $L_P$ and magnetic contribution $L_M$ to the solar angular momentum loss rate per steradian $L$. The Helios measurements indicate that for fast (slow) solar wind with $v_p\\gtrsim 600$ ($\\lesssim 400$)\\velunits, $L_P$ tends to be negative (positive), with the positive sign denoting the direction of corotation with the Sun. Furthermore, $L_P$ tends to be larger than $L_M$ in the slow wind. The behavior of $L_P$ derives from that of individual ion angular momentum fluxes, $L_p$ and $L_\\alpha$, thereby calling for a multifluid approach.  Starting with a general set of multifluid transport equations with gyrotropic species pressure tensors, we have derived the equations for both the angular momentum conservation (Eqs.(\\ref{eq_align_all}) and (\\ref{eq_tanp_standard}) in section~\\ref{sec_mathform}), and the energy and linear momentum balance (Eqs.(\\ref{eq_reduced_nk}) to (\\ref{eq_reduced_Tsper}) in the appendix). These equations are not restricted to radial outflows in the equatorial plane, instead they are valid for arbitrary axisymmetrical winds that include two major ion species, and therefore are expected to find applications in general outflows from late-type stars. To focus on the problem of angular momentum transport, we refrained from solving the full set of equations governing the meridional dynamics. Rather, we constructed, largely based on the available in situ measurements, the empirical profiles for the meridional magnetic field and flow parameters. Only the ion temperature anisotropies are considered, i.e., the electron temperature is seen as isotropic. For both the fast and slow solar wind profiles, we solved the angular momentum conservation law (Eqs.(\\ref{eq_align_all}) and (\\ref{eq_tanp_standard})) to examine how the azimuthal speeds of protons $v_{p\\phi}$ and alpha particles $v_{\\alpha\\phi}$, as well as the individual components in the solar angular momentum budget are influenced by the ion temperature anisotropies. To this end, solutions to the isotropic version are obtained for comparison.  Our main conclusions are: \\begin{enumerate} \\item From the derived equations governing the energy transport, a simple analysis given in the appendix yields that the adiabatic cooling may be considerably influenced with the introduction of the azimuthal components. Such an influence is understandably more prominent in the low-latitude regions. This means, when modeling the species temperature anisotropy, for a quantitative comparison of model computations to be made with the near-ecliptic measurements such as made by Helios, the spiral magnetic field has to be taken into account. \\item In agreement with the single-fluid case~\\citep{WD_70, Weber_70}, incorporating species temperature anisotropy leads to a situation where the total angular momentum loss rate per steradian $L$ is determined by the behavior of the solution to the angular momentum conservation law in the vicinity of the Alfv\\'en point where the combined Alfv\\'enic Mach number $M_T=1$. However, $M_T$ has to take into account the contribution from both ion species, as defined by Eq.(\\ref{eq_def_MT2_betaDelta}). \\item Relative to the isotropic case, the introduced species temperature anisotropy may enhance or decrease $L$ by up to $10$\\%, and introduce an absolute change of up to $\\sim 1.8$\\velunits\\ in individual ion azimuthal speeds in the region between $0.3$ and $1$~AU. While these changes seem modest, the corresponding changes in the angular momentum fluxes convected by protons $L_p$ or alpha particles $L_\\alpha$ may change substantially. In contrast, the flux associated with magnetic stresses $L_M$ hardly varies. \\item However, introducing ion temperature anisotropies cannot resolve the discrepancy between in situ measurements and models. For the fast wind solutions, while in extreme cases $L_P$ may become negative $L_p$ always stays positive. On the other hand, for the slow solar wind solutions examined, $L_P$ never exceeds $L_M$ even though $L_M$ may be smaller than the individual ion contribution. This is because, for both the slow and fast wind solutions, $L_p$ and $L_\\alpha$ always have opposite signs. \\item The discussion on the angular momentum transport has some bearing on the ion velocity fluctuation spectra $S_k (f)$ ($k=p, \\alpha$) during Alfv\\'enic activities in the super-Alfv\\'enic regions, which are likely to be explored by future missions such as Solar Orbiter and Solar Probe. In agreement with~\\citet{LiLi08} where species temperature anisotropies are neglected, an analysis based on the WKB and zero-frequency solutions yields that $S_\\alpha (f)$ will show an apparent break around some critical frequency $f_c$ if $S_p (f)$ is smooth there. This $f_c \\sim 0.5-1\\times 10^{-5}$\\frequnits\\ is the well-known frequency that separates the genuinely wave-like fluctuations from quasi-static ones. \\end{enumerate}"
    },
    "0812/0812.4270_arXiv.txt": {
        "abstract": "Coronal mass ejections (CMEs) are solar eruptions into interplanetary space of as much as a few billion tons of plasma, with embedded magnetic fields from the Sun's corona. These perturbations play a very important role in solar--terrestrial relations, in particular in the spaceweather. In this work we present some preliminary results of the software development at the Universidad Nacional Aut\\'onoma de M\\'exico to performe Remote MHD Numerical Simulations. This is done to study the evolution of the CMEs in the interplanetary medium through a Web--based interface and the results are store into a database. The new astrophysical computational tool is called the Mexican Virtual Solar Observatory (MVSO) and is aimed to create theoretical models that may be helpful in the interpretation of observational solar data. ",
        "introduction": "The Mexican Virtual Solar Observatory (MVSO) is a set of software tools that offer global solutions for Web development. The operating system is Linux, the Web server is Apache, SQL (Structure Query Language) and the relational database management system is MySQL, everything is programmed with PHP (Hypertext Pre--Processor). The computational backbone of the MVSO is structured into three stages. The first part is the related to the graphic user interface (GUI), the second part is associated to the remote numerical simulations (RNS) and the third part is the creation of the database and associated search tools. The implementation is explained by \\cite{Hernandez2008} ",
        "conclusions": ""
    },
    "0812/0812.3406_arXiv.txt": {
        "abstract": "Dark Matter candidates are natural in Technicolor theories. We introduce a general framework allowing to predict signals of Technicolor Dark Matter at colliders and set constraints from earth based experiments such as CDMS and XENON. We show that the associate production of the composite Higgs can lead to relevant signals at the Large Hadron Collider. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.1918_arXiv.txt": {
        "abstract": "Gaseous H$_2$O has been detected in several cold astrophysical environments, where the observed abundances cannot be explained by thermal desorption of H$_2$O ice or by H$_2$O gas phase formation. These observations hence suggest an efficient non-thermal ice desorption mechanism. Here, we present experimentally determined UV photodesorption yields of H$_2$O and D$_2$O ice and deduce their photodesorption mechanism. The ice photodesorption is studied under ultra high vacuum conditions and at astrochemically relevant temperatures (18--100~K) using a hydrogen discharge lamp (7-10.5 eV), which simulates the interstellar UV field. The ice desorption during irradiation is monitored using reflection absorption infrared spectroscopy of the ice and simultaneous mass spectrometry of the desorbed species. The photodesorption yield per incident photon, $Y_{\\rm pd}(T,x)$, is identical for H$_2$O and D$_2$O and its dependence on ice thickness and temperature is described empirically by $Y_{\\rm pd}(T,x)=Y_{\\rm pd}(T,x>8)(1-e^{-x/l(T)})$ where $x$ is the ice thickness in monolayers (ML) and $l(T)$ a temperature dependent ice diffusion parameter that varies between $\\sim$1.3~ML at 30~K and 3.0~ML at 100~K. For thick ices the yield is linearly dependent on temperature due to increased diffusion of ice species such that $Y_{\\rm pd}(T,x>8) = 10^{-3}\\left(1.3+0.032\\times T\\right)$ UV photon$^{-1}$, with a 60\\% uncertainty for the absolute yield. The increased diffusion also results in an increasing H$_2$O:OH desorption product ratio with temperature from 0.7:1.0 at 20~K  to 2.0:1.2 at 100~K. The yield does not depend on the substrate, the UV photon flux or the UV fluence. The yield is also independent on the initial ice structure since UV photons efficiently amorphize H$_2$O ice. The results are consistent with theoretical predictions of H$_2$O photodesorption at low temperatures and partly in agreement with a previous experimental study. Applying the experimentally determined yield to a Herbig Ae/Be star+disk model provides an estimate of the amount of gas phase H$_2$O that may be observed by e.g. \\textit{Herschel} in an example astrophysical environment. The model shows that UV photodesorption of ices increases the H$_2$O content by orders of magnitude in the disk surface region compared to models where non-thermal desorption is ignored. ",
        "introduction": "H$_2$O, in solid or gaseous form, is one of the most common species in molecular clouds, typically only second to H$_2$ and sometimes to CO. This makes H$_2$O, together with CO, the dominant reservoir of oxygen during the critical stages of star formation \\citep{vanDishoeck93}. H$_2$O is thus a key molecule in astrochemical models and its partitioning between the grain and gas phase therefore has a large impact on the possible chemical pathways, including the formation of complex organics \\citep{Charnley92, vanDishoeck06b}.  In cold, quiescent clouds H$_2$O forms through hydrogenation of O (O$_2$ or O$_3$) on cold (sub)micron-sized silicate grain surfaces forming icy layers \\citep{Tielens82,Leger85,Boogert04,Miyauchi08,Ioppolo08}. Other ices, like NH$_3$ and CH$_4$, probably form similarly, but observations show that H$_2$O is the main ice constituent in most lines of sight, with a typical abundance of $1\\times 10^{-4}$ with respect to the number density of hydrogen nuclei. Gas phase H$_2$O formation is only efficient above 300 K \\citep{Elitzur78,Elitzur78b,Charnley97}. At lower temperatures gas phase ion-molecule reactions maintain a low H$_2$O abundance around 10$^{-7}$ \\citep{Bergin95}. Any higher abundances require either thermal desorption of the H$_2$O ice above $\\sim$100 K or non-thermal desorption at lower temperatures \\citep{Bergin95, Fraser01}.  Gas phase H$_2$O is observed from the ground only with great difficulty. Still both isotopic and normal H$_2$O have been detected in astrophysical environments from ground based telescopes \\citep{Jacq88,Knacke91,Cernicharo90,Gensheimer96,vanderTak06}. The {\\it Infrared Space Observatory (ISO)} detected warm H$_2$O gas unambiguously toward several low- and high-mass young stellar objects \\citep{vanDishoeck96,Ceccarelli99, Nisini99,Boonman03}. {\\it ISO} was followed by two other space based telescopes, the {\\it Submillimeter Wave Astronomy Satellite (SWAS)} and {\\it Odin}. In difference to {\\it ISO}, {\\it SWAS} and {\\it Odin} are capable of detecting the fundamental ortho-H$_2$O 1$_{10}-1_{01}$ transition at 538.3 $\\mu$m and hence probe cold H$_2$O gas \\citep{Melnick00,Hjalmarson03}. Both telescopes have observed H$_2$O gas toward star forming regions and detected the expected enhancements near protostars and in outflows where thermal ice desorption or gas phase formation is possible \\citep{Hjalmarson03,Franklin08}. Critical for the present study, H$_2$O gas has also been detected toward photon dominated regions \\citep{Snell00,Wilson03} and is also more abundant toward diffuse than toward dense clouds. These two results point to an efficient ice photodesorption mechanism \\citep{Melnick05}. The importance of photodesorption at the edges of clouds has more recently  been modeled by Hollenbach et al. (ApJ, in press). They find that the H$_2$O gas abundance is enhanced by orders of magnitude at A$_{\\rm V}$=2--8 mag into the cloud when including photodesorption of H$_2$O ice in their model at a rate derived from the results by \\citet{Westley95}. Circumstellar disks is a second region where the impact of photodesorption is expected to be be large. \\citet{Willacy00} showed that a photodesorption yield of 10$^{-3}$ molecules per photon is enough to explain observed gas-phase CO abundances in the outer regions of flared disks. Employing a similar photodesorption yield for H$_2$O ice, the disk models of \\citet{Dominik05} and \\citet{Willacy07} both predict the existence of significant amounts of gas phase H$_2$O in a layer above the midplane region.  With the advent of the {\\it Herschel Space Observatory}, cold and warm H$_2$O gas observations on scales of protostellar envelopes and disks will for the first time become possible \\citep{vanKempen08}. In preparation for these and other observations, and to interpret data from {\\it Odin} and {\\it SWAS}, non-thermal processes need to be better understood. These non-thermal desorption processes include ion/electron sputtering, desorption due to the release of chemical energy and photodesorption. Of these, sputtering of ice by electrons and ions, has been investigated over a range of conditions during the past few decades \\citep[e.g.][]{Brown78, Fama08} and the dependencies of the yield on e.g. ice temperature, projectile type and energy are rather well understood. In contrast, only a handful of laboratory studies exists on the efficiency of ice photodesorption \\citep[e.g.][]{Westley95, Oberg07b, Oberg08b}.  \\citet{Westley95, Westley95b} determined the photodesorption rate of H$_2$O ice experimentally to be $3-8\\times10^{-3}$ molecules per UV photon for a 500 nm thick H$_2$O ice. In their experiment the photodesorption rate depends on UV fluence as well as temperature. The photon fluence dependence, together with the observed gas phase H$_2$ and O$_2$ during irradiation, was taken as evidence that H$_2$O photodesorption at low temperatures mainly occurs through desorption of photoproducts rather than of H$_2$O itself. Several questions remain regarding the applicability of their study to astrophysical regions due to the uncertainty of the proposed mechanism. In addition, their dependence on photon fluence is not reproduced in recent CO and CO$_2$ photodesorption experiments and cannot easily be explained theoretically \\citep{Andersson06, Oberg07b, Andersson08, Oberg08b}.  In a different experiment \\citet{Yabushita06} investigated H-atom photodesorption from H$_2$O ice during irradiation at 157 and 193 nm using time-of-flight mass spectrometry. They found that the temperature and hence the origin of the desorbed H atoms varies significantly between crystalline and amorphous ice at 100 K. This indicates that photodesorption depends on the ice morphology, which is in contrast to the findings of \\citet{Westley95}.  Desorption of recombined D$_2$ during irradiation of D$_2$O ice at 12 K has also been found by \\citep{Watanabe00}. Both experiments provide valuable input for models, but cannot directly be used to determine the total H$_2$O photodesorption rate.  UV irradiation of H$_2$O ice results in photochemistry products as well as photodesorption. A variety of photochemistry products are readily produced during irradiation of H$_2$O dominated ice mixtures \\citep{Dhendecourt82, Allamandola88}. Pure H$_2$O ice photolysis results in the production of H$_2$O$_2$, OH radicals and HO$_2$ at 10 K \\citep{Gerakines96,Westley95b}. After a fluence of $\\sim5\\times10^{18}$ UV photons cm$^{-2}$ \\citet{Gerakines96} found that the final band area of the formed H$_2$O$_2$  was only $\\sim$0.25\\% compared to the H$_2$O band area and the OH band area was even smaller.  No study exists for higher temperatures, but as photodissociation fragments become more mobile, e.g. O$_2$ formation would be expected \\citep{Westley95}.  Only a handful models of ice photodesorption exists in the literature. \\citet{Andersson06,Andersson08} have investigated H$_2$O photochemistry and photodesorption theoretically using classical dynamics calculations. In the simulations they followed H$_2$O dissociation fragments, after the absorption of a UV photon, in the top six monolayers of both crystalline and amorphous H$_2$O ice at 10 K. For each ice they found that desorption of H$_2$O has a low probability (less than 0.5\\% yield per absorbed UV photon) for both types of ice. The total H$_2$O photodesorption yield from the top six ice layers was calculated to be $\\sim4\\times10^{-4}$ molecules per incident UV photon.  In the present study we aim at determining experimentally the photodesorption yields of H$_2$O and D$_2$O and their dependencies on ice thickness, temperature, morphology, UV flux and fluence as well as irradiation time. We use these results as input for an astrophysical model of a typical circumstellar disk to estimate the impact of photodesorption and to predict the observable column densities of H$_2$O as relevant to e.g. upcoming {\\it Herschel} programs. ",
        "conclusions": "\\begin{enumerate} \\item The total D$_2$O  and H$_2$O photodesorption yields are indistinguishable within the experimental uncertainties and are empirically described by $Y_{\\rm pd}(T,x)=Y_{\\rm pd}(T,x>8)(1-e^{-x/l(T)})$ where $Y_{\\rm pd}(T,x)$ is the thickness and temperature dependent photodesorption yield, $x$ the ice thickness in monolayers and $l(T)$ an ice diffusion parameter that varies between 1.3 ML at 30~K and 3.0 ML at 100~K.  \\item For thick ices ($>$8 ML), the yield depends linearly on temperature such that $Y_{\\rm pd}(T,x>8) = 10^{-3}\\left(1.3+0.032\\times T\\right)$ photon$^{-1}$. The yields agrees, within the reported 60\\% uncertainty, with a previous experiment \\citep{Westley95}.  \\item The nature of the desorbed species is temperature dependent, with equal amounts of OH and H$_2$O detected at low temperatures. At higher temperatures the H$_2$O:OH fraction is $\\sim$ 2:1 and in addition about a fifth of the ice photodesorbs as heavier fragments like O$_2$. The fraction of the total photodesorption yield that results in H$_2$O molecules desorbing is described by $f_{\\rm H_2O} =  (0.42\\pm0.07)+(0.002\\pm0.001)\\times T$.  \\item We find no yield dependence on photon flux or fluence. The fluence independence is in contrast with a previous experiment \\citep{Westley95}.  \\item We also find no dependence on the ice structure i.e. whether the D$_2$O ice is amorphous or crystalline. This is consistent with spectroscopic evidence of fast destruction of crystalline ice into an amorphous state following UV irradiation.  \\item The photodesorption yield and dependencies found here are consistent with previous theoretical predictions of H$_2$O photodesorption, where the photodesorption is limited to the top few layers of the ice \\citep{Andersson06}. In addition we see that the photodesorption yield increases with ice temperature because of the increased mobility of the photolysis fragments, allowing desorption from deeper within the ice.  \\item Applying the experimental yield to a Herbig Ae/Be star+disk model we calculate that the predicted amount of cold ($<100$ K) gas phase H$_2$O, averaged over the entire disk, increases with orders of magnitude due to photodesorption. \\end{enumerate}"
    },
    "0812/0812.4226_arXiv.txt": {
        "abstract": "Measurements of clustering in large-scale imaging surveys that make use of photometric redshifts depend on the uncertainties in the redshift determination. We have used light-cone simulations to show how the deprojection method successfully recovers the real space correlation function when applied to mock photometric redshift surveys. We study how the errors in the redshift determination affect the quality of the recovered two-point correlation function. Considering the expected errors associated to the planned photometric redshift surveys, we conclude that this method provides information on the clustering of matter useful for the estimation of cosmological parameters that depend on the large scale distribution of galaxies. ",
        "introduction": "In recent years, photometric redshift surveys have been proposed as a way to extend large-scale structure studies towards higher redshifts than it is  possible using spectroscopic surveys. These surveys observe a region of the sky through a number of  filters, and use the photometry obtained to determine the redshifts, $z$, and spectral energy distributions (SED's) of galaxies. Using photometry instead of spectra allows them to get much deeper, but the uncertainty in the determination of redshifts is larger \\citep{bau62a, koo86a, con95a, bla05a, fer01a}.  Two projects of this kind are COMBO-17(Classifying Objects by Medium-Band Observations) and the ALHAMBRA (Advanced Large, Homogeneous Area Medium Band Redshift Astronomical) survey. COMBO-17 \\citep{wol03a} surveyed a total area of $\\sim 1 \\deg^2$ using a combination of 17 broad-band and medium-band filters. It provided photometric redshifts for $\\sim 25 000$ galaxies in $0.2 < z < 1.2$, with a typical error of $\\Delta z \\simeq 0.03$. The ALHAMBRA survey \\citep{mol06a, mol08a, fer08a}, currently ongoing, will observe a total area of $\\sim 4 \\deg^2$, in 16 $1\\degr \\times 0\\fdg25$ strips. It uses 20 medium-band, equal-width filters covering the optical range, plus the standard $J$, $H$, $K_s$ near-infrared filters. \\citeauthor{mol06a} expect to obtain photometric redshifts for $\\sim 3 \\times 10^5$ galaxies with $I_{AB} \\leq 24.7$ ($60$ per cent completeness level), $z_{med} = 0.74$, and $\\Delta z \\simeq 0.015 (1+z)$.  These surveys can therefore provide us with  large and deep samples of galaxies. In order to study large-scale structure using these samples, one has to deal with large redshift errors. These redshift errors produce uncertainties in the determination of distances to the galaxies, and hence in their three-dimensional positions \\citep{coe06a}. This uncertainty has to be added to the one produced by peculiar motions of galaxies. The latter is important for spectroscopic surveys, while the former dominates the uncertainties in photometric surveys.  In the present work, we focus on the two-point correlation function, $\\xi(r)$. We study how it is affected by redshift errors, and describe a method to recover its real-space value from photometric redshift survey data. The method we use is based in measuring the two-dimensional correlation function, $\\xi(\\sigma, \\pi)$ (where $\\pi$ is  the line-of-sight separation, and $\\sigma$ is the transverse separation), obtaining from it the projected correlation function, $\\Xi(\\sigma)$, and deprojecting it. This method (outlined in Section 3) was first proposed by \\citet{dav83a} as a way to avoid the uncertainties due to peculiar velocities in spectroscopic surveys, and has been used successfully in subsequent analyses \\citep*[e.g.][]{sau92a, haw03a, mad03a, zeh04a, zeh05a}.  \\citet{phl06a} studied the correlation function of galaxies in COMBO-17. They used $\\Xi(\\sigma)$ as a measure of real-space clustering, and compared it to the predictions of halo occupation models. However, they did not attempt to recover $\\xi(r)$ from their data.  We tested this method using data from the light-cone simulation of \\citet{hei05a}. From the simulation, we produced three mock photometric redshift catalogues, corresponding to different accuracies in the determination of redshifts. We then compared the correlation function $\\xi(r)$ obtained by our method in each case to the real-space one, computed from the original catalogue.   We describe the simulation and the way in which we created the mock catalogues in Section~\\ref{sec:data}. Section~\\ref{sec:meth} describes our method to calculate $\\xi(r)$ from the simulated data. In Section~\\ref{sec:res} we present our results, and we summarize our conclusions in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have shown the reliability of recovering the real-space two-point correlation function from photometric redshift surveys. We have used light-cone simulations to produce mock catalogues that have been distorted by randomizing along the line of sight the object positions following Gaussian distributions with different variances similar to the associated nominal errors of the photometric redshift surveys $\\Delta z/(1+z)$.  The method used to recover the real-space correlation function consists in obtaining the projected correlation function by integrating the two-dimensional correlation function along the line of sight. The projected correlation function is then deprojected assuming that redshift errors do not affect transverse distances.  The deprojection method applied on the distorted mock surveys provides quite satisfactory results for recovering the real-space correlation function. We have quantified the quality of the recovering process as a function of the errors in the photometric redshifts. Our method was able to recover the real-space correlation function within a $5$ per cent for $r < 10 \\hMpc$ from photometric catalogues with $\\Delta z \\leq 0.015 (1 + z)$. For larger redshift errors, the method is only valid (within a $7$ per cent) for smaller scales, $r < 2 \\hMpc$. Hence, our method allows the extraction of useful information on the clustering of galaxies through the correlation function. That information can be used for the estimation of cosmological parameters based on data from photometric redshift surveys.  We discuss now possible alternatives to the method exposed here. A variation of the method would be to use a smaller value of $\\pi_{\\rm max}$ in the integration of equation~(\\ref{eq:wp}), and multiply the result by a constant correction factor. This would reduce the extra noise introduced by the integration along a large range in the $\\pi$ direction. We found that, for $\\pi_{\\rm max} \\simeq \\Delta z$, a correction factor of $\\simeq 2$ works well in our simulation, generally reducing the error. However, the optimal value is slightly different for each mock photometric catalogue.  The main problem for the use of this method would be the accurate determination of the correction factor in each case, as any deviation from the optimal value would introduce a bias in the result. The Gaussian approximation used here is probably not so close to reality as to infer that constant from our simulations.  Another possible alternative to the method described in this work could make use of the new estimator $\\omega$ presented in \\citet*{pad07a}. They proposed $\\omega$ as an alternative to the projected correlation function to use with spectroscopic survey data. In principle, it can also be applied to photometric redshift survey data, in a way similar to the one presented here, but we did not investigate further this possibility."
    },
    "0812/0812.3589_arXiv.txt": {
        "abstract": "Shear flows have been prescribed in numerical models of coronal mass ejections and flares for decades as a way of energizing magnetic fields to erupt. While such shear flows have long been observed in the solar atmosphere, until recently, there was no compelling physical explanation for them. This paper will discuss the discovery that such shear flows are readily explained as a response to the Lorentz force that naturally occurs as bipolar magnetic fields emerge and expand in a gravitationally stratified atmosphere.  It will be shown that shearing motions transport axial flux, and magnetic energy from the submerged portion of the field to the expanding portion, strongly coupling the solar interior to the corona.  This physical process explains active region shear flows and why the magnetic field is found to be nearly parallel to photospheric polarity inversion lines where prominences form.  Finally, shear flows driven by the Lorentz force are shown to produce a loss of equilibrium and eruption in magnetic arcades and flux ropes offering a convincing explanation for CMEs and flares. ",
        "introduction": "Coronal mass ejections (CMEs) are energetic expulsions of plasma from the solar corona that are driven by the release of magnetic energy typically in the range of $10^{32-33}$ ergs. The majority of CMEs originate from the eruption of pre-existing large-scale helmet streamers \\citep[]{WM_hundhausen93}. Less common fast CMEs typically come from smaller, more concentrated locations of magnetic flux referred to as active regions.  In this case, the CMEs often occur shortly after the flux has emerged at the photosphere, but can also happen even as the active region is decaying. While CMEs occur in a wide range of circumstances, all CMEs originate above photospheric magnetic polarity inversion lines (neutral lines), which exhibit strong magnetic shear.  Shear implies that the magnetic field has a strong component parallel to the photospheric line that separates magnetic flux of opposite sign, and in this configuration, the field possesses significant free energy.  In contrast, a potential field runs perpendicular to the inversion line and has no free energy.  There is enormous evidence for the existence of highly sheared magnetic fields associated with CMEs and large flares. At the photosphere, magnetic shear is measured directly with vector magnetographs \\citep[e.g.][]{WM_hagyard84, WM_zirin93, WM_falconer2002, WM_yang2004, WM_liu2005}. Higher in the atmosphere, the magnetic field is difficult to measure directly, but its direction may be inferred from plasma structures formed within the field. Seen in chromospheric H$\\alpha$ absorption, filaments form only over photospheric inversion lines \\citep[]{WM_zirin83} along which the magnetic field is nearly parallel \\citep[]{WM_leroy89}. Fibrils and H$\\alpha$ loops that overlay photospheric bipolar active regions are also indicative of magnetic shear \\citep[]{WM_foukal71}. Comparisons between vector magnetograms and H$\\alpha$ images show that the direction of the sheared photospheric magnetic field coincides with the orientation of such fibril structures \\citep[]{WM_zhang95}.  Higher in the corona, evidence of magnetic shear is found in loops visible in the extreme ultraviolet \\citep[]{WM_liu2005} and X-ray sigmoids \\citep[]{WM_moore2001}.  These structures run nearly parallel to the photospheric magnetic inversion line prior to CMEs, and are followed by the reformation of closed bright loops that are much more potential in structure.  Finally, observations by the Transition Region and Coronal Explorer (TRACE) show that 86 percent of two-ribbon flares show a strong-to-weak shear change of the ribbon footpoints that indicates the eruption of a sheared core of flux \\citep{WM_su2007}.  Sheared magnetic fields are at the epicenter of solar eruptive behavior. Large flares are preferentially found to occur along the most highly sheared portions of magnetic inversion lines \\citep[]{WM_hagyard84, WM_zhang95}. More recent analysis by \\citet{WM_schrijver2005} found that shear flows associated with flux emergence drove enhanced flaring. Similarly, active region CME productivity is also strongly correlated with magnetic shear as shown by \\citet{WM_falconer2001, WM_falconer2002, WM_falconer2006}. It is not coincidental that large flares and CMEs are strongly associated with filaments, which are known to form only along sheared magnetic inversion lines \\citep[]{WM_zirin83}. The buildup of magnetic shear is essential for energetic eruptions, and for this reason, it is of fundamental importance to understanding solar activity.  Currently the majority of CME initiation models rely on the application of artificially imposed shear flows.  Examples of such models include \\citet{WM_antiochos99, WM_mikic94, WM_guo98, WM_amari2003}. Until recently, there was no theory to account for these large scale shear flows. In this paper, we discuss a series of simulations that illustrate a physical process by which these shear flows are self-organized in emerging magnetic fields.  We will give a close comparison of these simulations in the context of new observations that make a more complete and compelling picture of a fundamental cause of eruptive solar magnetic activity. ",
        "conclusions": "The build up of magnetic energy in active regions is essential to the onset of CMEs and flares.  The magnetic stress must pass from the convection zone into the corona in the form of non-potential fields, and effectively couple layers of the atmosphere. Observations show that such non-potential fields occurs along magnetic polarity inversion lines where the magnetic field is highly sheared.  The evidence that magnetic shear is essential to CMEs and flares is provided by the very strong correlation between photospheric shear flows, flux emergence and the onset of CMEs and large flares \\citep[]{WM_meunier2003, WM_yang2004, WM_schrijver2005}.  These shear flows are found to be strongest along the magnetic inversion line precisely where flares are found \\citep[]{WM_yang2004, WM_deng2006}. This complements earlier evidence that in the case of two-ribbon flares \\citep[]{WM_zirin84}, prior to the eruption, H$\\alpha$ arches over the inversion line are highly sheared, afterward, the arches are nearly perpendicular to the inversion line. \\citet{WM_su2006} found a similar pattern of magnetic shear loss in the apparent motion of footpoints in two-ribbon flares. There are even now observations of subphotospheric shear flows with a magnitude of 1-2 km/s 4-6 Mm below the photosphere that occur during flux emergence \\citep[]{WM_kosovichev2006}, which simulations show is fully consistent with the Lorentz force driving mechanism \\citep[]{WM_manchester2007}.  These ubiquitous shear flows and sheared magnetic fields so strongly associated with CMEs are readily explained by the Lorentz force that occurs when flux emerges and expands in a gravitationally stratified atmosphere.  This physical process explains and synthesizes many observations of active regions and gives them meaning in a larger context.  This shearing mechanism explains (1) the coincidence of the magnetic neutral line with the velocity neutral line, (2) the impulsive nature of shearing in newly emerged flux, (3) the magnitude of the shear velocity in different layers of the atmosphere, (4) the large scale pattern of magnetic shear in active regions, (5) the transport of magnetic flux, and energy from the convection zone into the corona, (6) eruptions such as CMEs and flares. With so much explained, it still remains a numerical challenge to model an active region with sufficient resolution to produce a large scale CME by this shearing mechanism.  The rope emergence simulation discussed here only produces a flux concentration that is one tenth the size of an active region, which at this scale simply can not produce an eruption the size of a CME.  However, this simulation illustrates the basic process by which CMEs and flares must occur, and current simulations already show a very favorable scaling of eruption size.  With increases in computer power, simulations of flux emergence should soon be producing CME size eruptions by shear flows driven by the Lorentz force."
    },
    "0812/0812.2949.txt": {
        "abstract": "We investigate the secular dynamics of a planetary system composed of the parent star and two massive planets in mutually inclined orbits.  The dynamics are investigated in wide ranges of semi-major axes ratios  (0.1--0.667), and planetary masses ratios (0.25--2) as well as in the whole permitted ranges of the energy and total angular momentum. The secular model is  constructed by semi-analytic averaging of the three-body system. We focus on  equilibria of the secular Hamiltonian (periodic solutions of the full system), and we analyze their stability. We attempt to classify  families of these solutions in terms of the angular momentum integral. We identified new equilibria, yet unknown in the literature. Our results are general and may be applied to a wide class of three-body systems, including configurations with a star and brown dwarfs and sub-stellar objects. We also describe some technical aspects of the semi-numerical averaging. The HD~12661 planetary system is investigated as an example configuration. ",
        "introduction": "% ----------------------------------------------------------------------------- Nowadays, about thirty extrasolar multi-planet systems have been detected\\footnote{See Jean Schneider's Extrasolar Planets Encyclopedia http://exoplanet.eu for frequent updates on the discoveries and orbital parameters}. Many of them seem either locked in or close to low-order mean motion resonances (MMRs). Moreover, there is a class of the so called hierarchical systems \\citep{Lee2003} which can be  characterized by relatively small ratio of semi-major axes.  Their planetary orbits are well separated and far from collision zones, hence the long-term, qualitative  dynamics of such systems may be  effectively investigated with secular theories. The Hamiltonian of a hierarchical system can be averaged out over mean longitudes which play the role of fast angles.  In the regime of small eccentricities and inclinations, this approach leads to the well known, classical Laplace-Lagrange (L-L) secular theory \\citep{Murray2000}. It relies on the expansions of the disturbing function in power series with respect to eccentricities and inclinations which are small parameters of the problem.  However, many multi-planet hierarchical systems do not satisfy the assumption of small eccentricities, and the L-L theory may fail.  Still, to deal with the observed diversity of orbital configurations, the secular theories relying on high-order expansions of the perturbations are used, e.g., the series in eccentricity \\citep[e.g.,][]{Murray2000,Rodriguez2005,Libert2006,Libert2007a,Libert2007b,Veras2007} or expansion to the third order in the ratio of semi-major axes $\\alpha$, known as the octuple theory \\citep{Ford2000,Lee2003}. This theory can be also generalized to higher orders \\citep[][and references therein]{Migaszewski2008a}. The analytical expansions are particularly suitable for studies of hierarchical systems.  Moreover, they are usually valid only in limited ranges of the orbital parameters, and special cases (like resonant configurations) must be treated individually. The alternative, recently developed quasi-analytical theory relies on averaging the perturbing Hamiltonian  numerically \\citep{Michtchenko2004,Michtchenko2006}. In this work, we are heavily  inspired by these papers and their idea of the semi-analytical technique. Because the method does not require any expansion of the perturbing Hamiltonian, basically, it has no limitations inherent in the analytical theory. For instance, with a help of this technique, \\cite{Michtchenko2004} found new, non-classic feature of the secular dynamics of  coplanar system of two planets (the so called non-linear secular resonance in the regime of large eccentricity). In the later work, \\cite{Michtchenko2006} consider more general three-dimensional (3-D) secular model of two-planet system, and present a systematic approach helpful to investigate the global dynamics of such configurations.  As an example to study, we choose the HD~12661 planetary system \\citep{Fischer2001,Fischer2003,Butler2006}.  The discovery paper \\citep{Fischer2001} announces two Jovian planets on well separated orbits with semi-major axes of $\\sim0.8$~au and $\\sim2.8$~au, respectively, and of moderate eccentricities. We analyzed the most recent, publicly available data from the catalogue of \\cite{Butler2006} and \\citep{Wright2008}, using the $N$-body model of the radial velocities (RV)  and the so called hybrid minimization  \\citep{Gozdziewski2006}. The results of our analysis  of the RV data published in \\citep{Butler2006} are illustrated in Fig.~\\ref{fig:fig1}. The best fit solution yielding $\\Chi\\sim 1.08$ and an rms $\\sim 7.5$~m/s is marked with a crossed circle in the dynamical map in terms of  the Spectral Number (SN) \\citep{Michtchenko2001}. The SN is the fast indicator making it possible to distinguish between chaotic and regular planetary configurations. The osculating elements of the best fit solution  at the epoch of the first observation are given in caption to Fig.~\\ref{fig:fig1}. In this figure, we mark the semi-major axis and  eccentricity of the outer planet  derived from an ensemble of fits within $1\\sigma$ of the best fit solution. Clearly, the available data already constrain  orbital elements of the outer planet very well.  The dynamical maps reveal orbits well separated from the low-order MMRs. Two most prominent MMRs within the vicinity of the best fit are 19:3 and 13:2~MMRs, respectively. Moreover, the best-fits within $1\\sigma$ confidence level span the region of small eccentricities in which the resonances are very narrow.  Hence, the HD~12661 system fits well assumptions of the secular theory.  This system has been studied already in a few papers: with the direct numerical integrations \\citep{Ji2003b}, with the  analytical octupole theory of hierarchical systems  \\citep{Lee2003}, with mapping of the phase space by  fast indicators \\citep{Gozdziewski2003a}, and with the classic analytical theory that relies on expansions of the perturbation in eccentricity \\citep{Rodriguez2005,Libert2006}. All the cited works assume that the HD~12661  system is coplanar and oriented edge-on.  However, we should keep in mind that a major limitation of the Doppler technique lies in the ambiguity of orbital inclinations, which cannot be well determined by far. The observational windows are still relatively narrow, and to remove the inclination degeneracy, several orbital periods of the outermost orbit are required. Moreover, the recent formation theories do not fully predict mutual inclinations in multi-planet systems. We cannot be certain yet whether the common assumption of coplanar orbits really holds true. Likely, many different forming scenarios are possible. For instance, the migration in low-order MMRs may end up with systems characterized by large mutual inclinations \\citep{Thommes2003}.  The dynamical relaxation of initially dense planetary systems of giant planets \\citep{Adams2003} may lead to scattering events  which produce wide distribution of the mutual inclinations. Indeed, recent simulations of  \\cite{Veras2008}  revealed that the outer planet of the HD~12661  system undergoes large oscillations for nearly all of the allowed two-planet orbital solutions. These authors conclude that it might be the effect of  a perturbation of planet~c, perhaps due to strong scattering of an additional planet that was subsequently accreted onto the star. Moreover, we stress that the inclination of the  HD~12661 system is still unknown, hence the understanding of basic features of its 3-D  dynamics  seems also important. This intriguing system is an excellent candidate for tests and numerical experiments regarding non-coplanar configurations. Moreover,  because we attempt to study the secular 3-D dynamics globally, our results are  general and valid for much wider class of three-body systems. The secular theory  which we consider here, covers planetary systems with different masses and semi-major axes ratios, and the full range of mutual inclinations.  The plan of this work is the following. In Sect.~\\ref{model}, we recall the general  mathematical model of the 3-D two-planet system. Subsection~\\ref{sec:average} is devoted to a short technical overview of the averaging approach and some computational details that may be useful in a practical implementation of the method. To make the paper self-contained, we also recall the notion of  representative planes, and energy levels calculated for fixed values of the total angular momentum integral (Sect.~\\ref{sec:equilibria}). To illustrate the precision of the semi-analytic approach, we compute the Poincar\\'e cross sections, and demonstrate chaotic behavior of the secular system \\citep{Michtchenko2006}. The main results are described in Sect.~\\ref{sec:families}  which is devoted to the analysis  of the existence and bifurcations of equilibria in the secular, spatial problem of two mutually interacting planets. In particular, we detect and investigate closely a few families of these equilibria in a  wide range of planetary mass ratio, $\\mu \\in\\{0.25,0.5,1,2\\}$, and the semi-major axes ratio, $\\alpha \\in \\{0.1,0.2,0.333,0.667\\}$. The results are general and valid as long as the partition of the Hamiltonian onto the Keplerian, integrable term, and the small perturbation is reasonable. In that part, our work extends the paper of \\cite{Libert2007b}. After introducing non-singular canonical variables, they investigate the existence, stability and bifurcations of stationary solutions emerging from the equilibrium at zero-eccentricities, the so called  {\\em Lidov--Kozai resonance} \\citep[][to mention a few papers in an endless list of references]{Lidov1961,Kozai1962,Lidov1976,Michel1996,Innanen1997,Kinoshita2007} which was found  and intensively investigated in the restricted  three body problem. The full three-body problem in the Hill case  (hierarchical configurations) was also intensively studied by many authors \\citep[e.g.,][]{Krasinsky1972,Krasinsky1974,Lidov1976,Ferrer1994,Miller2002, Fabrycky2007}. These works rely mostly on the second order expansion of the secular Hamiltonian in the semi-major axes ratio (the quadrupole approximation).  In the present work, we focus on two aspects of the problem: \\begin{itemize} \\item we consider the unrestricted problem in wide ranges of semi-major axes ratio $\\alpha$, up to 0.667, and mass ratio $\\mu$, \\item we study equilibria of the full secular Hamiltonian; the semi-analytic averaging helps us to compute the secular perturbations beyond convergence limits of the usual power series expansions. \\end{itemize} Thanks to the quasi-analytic averaging, we found new families  of equilibria of the secular 3D planetary problem which unlikely may be detected with the help of perturbation techniques. We also study the Lyapunov stability of these solutions in detail (or to an extent permitted by technical limits of the semi-analytic algorithm).  \\begin{figure*} \\centerline{\\includegraphics[width=5.6in]{fig1.png}} \\caption{ The dynamical map of the edge-on, coplanar HD~12661 system in the $(a_{\\idm{c}},e_{\\idm{c}})$-plane, for the best-fit solution to the RV data published in \\citep{Butler2006}. Large values of the Spectral Number (SN) marked in yellow indicate continuous spectrum  of the fundamental frequencies of the system and strongly chaotic motions, small SN (black) means discrete frequencies and ordered motions. Positions of low-order MMRs are labeled.  The best-fit solution, yielding $\\Chi \\sim 1.08$ and an rms $\\sim 7.46$~ms$^{-1}$, in terms of parameter tuples ($m$ [$\\mbox{m}_{\\idm{J}}$], $a$~[au], e, $\\omega$~[deg], ${\\cal M}$~[deg]) including planetary mass, semi-major axis, eccentricity, the argument of pericenter and the mean anomaly at the epoch of the first observation, $t_0$=JD2,450,831.608380, is the following (2.34, 0.831, 0.361, 296.24, 126.86) for planet~b, and (1.84, 2.888, 0.021, 52.66, 66.18) for planet~c, respectively. \\hide{Offsets are $0.073$~ms$^{-1}$ and  $-0.012$~ms$^{-1}$, for the Lick and Keck data, respectively.} The original errors are rescaled by adding stellar jitter of $3.5$~m/s in quadrature. The mass of the parent star is 1.11~$M_{\\sun}$. Solutions within $1\\sigma$ level of the best fit are marked with yellow circles, fits with marginally worse $\\Chi$ are marked in red. } \\label{fig:fig1} \\end{figure*}   % ----------------------------------------------------------------------------- ",
        "conclusions": "% ----------------------------------------------------------------------------- \\label{sec:concl} The semi-analytical averaging is a powerful technique helpful to reduce limitations of the analytical theories.  Although it relies on purely numerical algorithms, its solid theoretical background is vital for the interpretation and understanding of the results of numerical experiments. In this work, we demonstrate that it makes it possible to study the secular dynamics of two-planet system in wide ranges of  semi-major axes and masses ratio.  The appropriate scaling of the problem parameters  helps us to represent the phase space of the secular system globally. Assuming non-resonant configurations, and that orbits are distant enough from collision zones, the averaged system may be reduced to two degrees of freedom. Hence, to carry out the analysis,  we can apply geometric tools, like the representative planes of initial conditions, Poincar\\'e cross sections,  continuations of stationary solutions with respect to parameters, which are very helpful to  understand the structure of the phase space and, in turn, the long-term behavior of the planetary system.  Equilibria are the basic class of solutions which can be investigated with relatively simple tools. Our analysis reveals a number of families of stationary solutions in the 3D secular problem of two planets. To the best of our knowledge, some of them are yet  unknown in the literature and are related to unusual orbital configurations. For instance, we found the so called chained stationary configuration which are non-resonant,  can be found in the regime of {\\em small} mutual inclinations and large eccentricities, and  are located over geometric collision line of orbits. In spite of such extreme dynamical situation,  these secular solutions are Lyapunov stable and exist in wide ranges of semi-major and masses ratio. Simultaneously, such orbital configurations prohibit application of analytical methods  relying on power series expansion of the perturbations.  The semi-analytic averaging helps to generalize analytical results obtained for low-order expansions of the secular Hamiltonian.  We obtained some interesting results regarding the Lidov-Kozai equilibrium in the non-restricted problem. We found that this resonance may be associated with librations of  $\\Delta\\varpi$ around $\\pi$ in the neighboring trajectories (not only with librations of the inner pericenter around $\\pm\\pi/2$). These librations are possible for relatively large ratios of  semi-major axes and planetary masses. We found that the L-K resonance may also  appear in the regime of large eccentricity of the inner orbit, and then it would be associated with librations of $\\omega_1$ around $\\pm\\pi/2$ with simultaneous librations of $\\Delta\\varpi$ around $0$. The parametric evolution of the L-K resonance is related to the stability of the zero-eccentricity equilibria. There is a link between bifurcations of this family (and changes of its stability) with an appearance of new families of stationary solutions in other parts of the phase space (hence, with changes of its global topology).  Our work illustrates qualitatively different view of the 3D dynamics as compared to the coplanar configurations. It is already known \\citep{Michtchenko2004} that the  non-resonant, coplanar systems of two point-mass planets fulfilling  the averaging theorem and interacting through Newtonian forces are integrable. The secular dynamics of such systems are basically trivial and may be reduced to one degree of freedom. Under the same assumptions, the spatial configurations may exhibit strong chaos and extremely complex secular phenomena.  In the approximation of Newtonian, point-mass interactions, the dynamics depend on masses and semi-major axes ratios, hence our results are valid both for planetary systems with small planets, as well as for systems comprising of brown dwarfs or even sub-stellar companions. However, the dynamics of real systems may strongly depend on the magnitude of the mutual interactions. Moreover,  many stable equilibria are found for large values of $\\nAMD$. According to the notion of $\\AMD$ \\citep{Laskar2000}, in such cases the real configurations are unlikely long-term stable even close to secularly stable equilibria.  In that sense, the results of stability  analysis may be too optimistic. We skip a study of such effects in this work, because it would make the paper necessarily very lengthly. We should keep in mind that introduction of relativistic, tidal, and stellar quadrupole-moment perturbations,  may affect the secular dynamics dramatically \\citep{Migaszewski2008c,Migaszewski2008d}.  The quasi-global technique applied in this paper has been proved an efficient and effective tool for the analysis of the secular 3-D model. Nevertheless, we learned from the work that the quasi-analytical approach is not a perfect tool. Due to limitations of the numerical algorithms, the continuation of  families of equilibria and analysis of their stability is particularly difficult when we reach limits of permitted motion or $\\Hsec$ is a weakly varying function. We can also overlook some solutions. Unfortunately, that leave us sometimes with open questions. %______________________________________________________________________________"
    },
    "0812/0812.2770_arXiv.txt": {
        "abstract": "The properties of compact stars made of massive bosons with a repulsive selfinteraction mediated by vector mesons are studied within the mean-field approximation and general relativity. We demonstrate that there exists a scaling property for the mass-radius curve for arbitrary boson masses and interaction strengths which results in an universal mass-radius relation. The radius remains nearly constant for a wide range of compact star masses.  The maximum stable mass and radius of boson stars are determined by the interaction strength and scale with the Landau mass and radius. Both, the maximum mass and the corresponding radius increase linearly with the interaction strength so that they can be radically different compared to the other families of boson stars where interactions are ignored. ",
        "introduction": "White dwarfs, neutron and quark stars, collectively dubbed compact stars, are the final result of stellar evolution.  White dwarfs are stabilized by the Fermi degeneracy pressure. There exists an upper limit for the mass of a white dwarf, the Chandrasekhar mass, which is about 1.4 times the mass of the sun \\cite{1931ApJ....74...81C}.  Beyond this limit the white dwarf is unstable against gravitational collapse. Neutron stars are stable mainly due to the repulsive nature of the interactions between nucleons. Therefore, the precise value of the maximum mass for a neutron star is less certain, but is presumably close to the predictions based on the Landau consideration \\cite{Landau32}. As shown in Ref.~\\cite{2006PhRvD..74f3003N}, Landau's argument can be extended to a general compact star made of fermions, a fermion star, with arbitrary fermion mass and interaction strength.  In the following we are studying compact stars made of bosons. Unlike fermion stars, boson stars have no observational evidence, yet. Besides this the existence of any stable scalar particle has never been experimentally verified. Wheeler \\cite{1955PhRv...97..511W} introduced a gravitational electromagnetic entity called a geon. The gravitational attraction of its own field energy confines the geon in a certain region. Later Kaup \\cite{1968PhRv..172.1331K} has solved the Klein-Gordon Einstein equations for scalar fields and found a new class of solutions for gravitating objects. These boson stars are stable with respect to spherically-symmetric gravitational collapse. Ruffini and Bonazzola \\cite{1969PhRv..187.1767R} demonstrated that boson stars describe a family of self-gravitating scalar field configurations within general relativity.  In ref.~\\cite{1984ZPhC...26..241T} Takasugi and Yoshimura calculated boson stars within an approach similar to the one conventionally adopted for neutron stars by solving the Tolman-Oppenheimer-Volkoff (TOV) equation \\cite{Tolman34,Tolman39,OV39} with a separate equation of state describing the properties of matter.  Boson stars with selfinteractions have been also considered, in particular as candidates for dark matter \\cite{Colpi:1986ye}. For reviews on boson stars we refer to \\cite{Lee:1991ax,1992PhR...220..163J,2003CQGra..20R.301S}.  In the following, we consider boson stars as localized, gravitationally bound objects made of self-interacting bosons at zero temperature. The interactions between the bosons is described by vector meson exchange in the relativistic mean-field approximation. The resulting equation of state is used as input to solve the TOV equation for boson stars, similar to the approach of ref.~\\cite{1984ZPhC...26..241T} but with an equation of state based on a field-theoretical approach.  We demonstrate that there are scaling relations for the mass-radius curve. In particular, we show that the maximum mass is controlled by the interaction strength and the Landau mass, not by the boson mass. We compare our results to previous works and to the case of fermion stars with self-interactions. ",
        "conclusions": "We have constructed an equation of state for a system of massive bosons interacting by the exchange of vector mesons. By solving the TOV equations for such boson stars, we have demonstrated that there exists a universal mass-radius curve independent of the boson mass and the interaction strength.  The maximum mass and the corresponding radius of boson stars are scaled with the Landau mass and Landau radius times the interaction strength. For masses much smaller than the maximum mass, the radius stays constant and is only slightly larger than the one for the maximum mass configuration.  The maximum mass and the corresponding radius can be computed with the simple formulae $M_{\\max }=0.164 y\\cdot M_{P}^{3}/m_{b}^{2}$ and $R_{\\min}=0.763 y\\cdot M_{P}/m_{b}^{2}$ for any given boson mass $m_b$ and interaction strength $y$. The possible masses and radii for boson stars can therefore cover a wide range and can be similar to the ones found for astrophysical compact objects, be it neutron stars or black hole candidates. For example, for a boson with QCD-type interaction strength and a boson mass of 100~GeV the maximum mass is $M_{\\max }\\sim 0.3 M_{\\odot}$ with a radius of about 2~km. A boson with a typical axion-like mass of $10^{-5}$~eV and an interaction scale of $10^{12}$~GeV will give a maximum mass of the boson star of $30M_\\odot$ with a radius of 200~km. The compactness of boson stars for the maximum mass configuration is about $R/(2GM)\\approx 2.3$ which is close to the value found for fermion stars $R/(2GM)\\approx 2.4$ in Ref.~\\cite{2006PhRvD..74f3003N}. It is interesting that these values are below the radius of the innermost stable circular orbit of nonrotating black holes $R/(2GM)=3$.  Finally, we mention that the full problem addressed here involves solving the Einstein equations with the coupled system of Klein-Gordon and Proca equations which we leave to address as an interesting extension for future work."
    },
    "0812/0812.0069_arXiv.txt": {
        "abstract": "The Christodoulou memory is a nonlinear contribution to the gravitational-wave field that is sourced by the gravitational-wave stress-energy tensor.  For quasicircular, inspiralling binaries, the Christodoulou memory produces a growing, nonoscillatory change in the gravitational-wave ``plus'' polarization, resulting in the permanent displacement of a pair of freely-falling test masses after the wave has passed. In addition to its nonoscillatory behavior, the Christodoulou memory is interesting because even though it originates from 2.5 post-Newtonian (PN) order multipole interactions, it affects the waveform at leading (Newtonian/quadrupole) order. The memory is also potentially detectable in binary black-hole mergers. While the oscillatory pieces of the gravitational-wave polarizations for quasicircular, inspiralling compact binaries have been computed to 3PN order, the memory contribution to the polarizations has only been calculated to leading order (the next-to-leading order 0.5PN term has previously been shown to vanish).  Here the calculation of the memory for quasicircular, inspiralling binaries is extended to 3PN order. While the angular dependence of the memory remains qualitatively unchanged, the PN correction terms tend to reduce the memory's magnitude. Explicit expressions are given for the memory contributions to the plus polarization and the spin-weighted spherical-harmonic modes of the metric and curvature perturbations. Combined with the recent results of Blanchet et al.~[Class.~Quantum Grav.~{\\bf 25}, 165003 (2008)], this completes the waveform polarizations to 3PN order. This paper also discusses: (i) the difficulties in extracting the memory from numerical relativity simulations, (ii) other nonoscillatory effects that enter the waveform polarizations at high PN orders, and (iii) issues concerning the observability of the memory in gravitational-wave detectors. ",
        "introduction": "Introduction} The primary purpose of this paper is to compute the post-Newtonian (PN) corrections to the nonlinear memory piece of the gravitational-wave (GW) polarizations for quasicircular\\footnote{\\emph{quasicircular} means that the binary orbit is circular up 2.5PN order radiation-reaction effects that cause the binary to slowly inspiral.}, inspiralling binaries. We begin by reviewing the gravitational-wave memory (Sec.~\\ref{sec:memintro}) and then motivate its further study (Sec.~\\ref{sec:motivation}). A summary of results is given in Sec.~\\ref{sec:summary}. \\subsection{\\label{sec:memintro}What is gravitational-wave memory?} Gravitational-wave memory refers to the permanent displacement of an ``ideal'' GW detector after the GW has passed.  An ideal detector is one which is only sensitive to gravitational forces\\textemdash e.g., a ring of freely-falling test-masses\\textemdash and is isolated from local tidal interactions. After the passage of a GW \\emph{without} memory, the detector returns to its initial state of internal displacement (its state long before the passage of the wave). After the passage of a GW \\emph{with} memory, the initial and final displacement states differ.  There are two types of GW memory: linear and nonlinear. The linear memory has been known since the 1970s \\cite{zeldovich-polnarev,braginskii-grishchuk,braginskii-thorne} and arises from gravitational sources that produce a net change in the time derivatives of one or more of their \\emph{source-multipole moments}. For example, at leading order in a PN expansion the linear memory causes a net change in the GW field given by \\be \\label{eq:hTTmem} \\Delta h_{jk}^{\\rm TT} = \\frac{2}{R} \\Delta(\\ddot{\\mathcal I}_{jk}^{\\rm TT}) , \\ee where ${\\mathcal I}_{ij}$ is the source mass-quadrupole moment, $R$ is the distance to the source, TT means to take the transverse-traceless projection [Eq.~\\eqref{eq:TT}], and the $\\Delta$ refers to the difference between late and early times. Geometric units ($G=c=1$) are used here and throughout.  A simple example of a source with linear memory is a binary on a hyperbolic orbit (gravitational two-body scattering) \\cite{turner-unbound,turner-will,kovacs-thorne-IV}. To see this consider the TT piece of the second-time derivative of the mass-quadrupole moment at Newtonian order and in the center-of-mass frame: \\be \\label{eq:ddIij} \\ddot{\\mathcal I}_{jk}^{\\rm TT} = 2 \\mu \\left[ \\dot{x}_j \\dot{x}_k - \\frac{M}{r^3} x_j x_k \\right]^{\\rm TT} , \\ee where $\\mu=m_1 m_2/M$ is the reduced mass, $M=m_1+m_2$ is the total mass, $x_j$ is the relative orbital separation vector with magnitude $r$, $\\dot{x}_j$ is the relative orbital velocity, and we have replaced second-time derivatives with the equation of motion $\\ddot{x}_j = -M x_j/r^3$. If we take $t=0$ to be the time of closest approach, then at very early ($t\\rightarrow -\\infty$) and very late ($t\\rightarrow +\\infty$) times the relative velocities will be finite ($v_{\\infty} = \\sqrt{2E/\\mu}$ where $E$ is the orbital energy), while the second term in Eq.~\\eqref{eq:ddIij} falls off like $1/r$ and is negligible. The memory for a hyperbolic binary thus arises from the difference in the direction of the velocity vectors at late and early times: \\be \\Delta h_{jk}^{\\rm TT} = \\frac{4 \\mu}{R} \\Delta [\\dot{x}_j \\dot{x}_k]^{\\rm TT} . \\ee  Other systems with linear memory are those whose components change from being bound to unbound (or vice versa). These include binaries whose components are captured, disrupted, or undergo mass loss. Gravitational-waves with linear memory have been studied in the context of supernova explosions and their resulting neutron star kicks \\cite{burrows-hayesPRL96,buonanno-stochasticBGmemory,muller-GWcorecollapse,kotake-sato-SNmemory,davies-king-SNkick-memory,cuesta-pulsarkickmem,kotake-etal-sasiApJL09}, asymmetric mass loss due to neutrino emission \\cite{turner-neutrinomemory,epstein-neutrinomemory,loveridge-pulsarkickmemory,burrows-hayesPRL96,kotake-etal-sasiApJL09} (see Ref.~\\cite{ott-corecollapsereview} for a recent review), or gamma-ray-burst jets \\cite{sago-GRBmemory,hiramatsu-kotake-GRBmemory,segalis-ori-GRBmemory}. Linear memory can also arise from GW recoil in binary black-hole mergers \\cite{favata-lisa7confproc}. Sources whose components remain bound generally do not display linear memory (but see Sec.~\\ref{sec:linmem} for a caveat).  A general formula for the linear memory produced by a system of $N$ bodies with changing masses $M_A$ or velocities ${\\bm v}_A$ is given by Thorne \\cite{kipmemory} (also Ref.~\\cite{braginskii-thorne}): \\be \\label{eq:hijlinmem} \\Delta h_{jk}^{\\rm TT} = \\Delta \\sum_{A=1}^N \\frac{4 M_A}{R\\sqrt{1-v_A^2}} \\left[ \\frac{v_A^j v_A^k}{1-{\\bm v}_A \\cdot {\\bm N}} \\right]^{\\rm TT} \\,, \\ee where the masses are unbound in their initial or final states (or both), the $\\Delta$ means to take the difference between the final and initial values of the summation, and ${\\bm N}$ is a unit vector that points from the source to the observer. For example, the masses $M_A$ and velocities ${\\bm v}_A$ might refer to two or more particles on a scattering orbit (as discussed above), or they might refer to the various pieces of a star that becomes unbound. This formula is essentially the standard Li\\'{e}nard-Wiechert solution of the space-space part of the linearized Einstein equations (with a source term given by the stress-energy tensor of $N$ noninteracting point particles).  The nonlinear memory was discovered independently by Payne \\cite{payne-zfl},  Blanchet and Damour \\cite{blanchet-thesis,blanchet-damour-hereditary}, and Christodoulou \\cite{christodoulou-mem}. It is often referred to as the ``Christodoulou memory.''  The nonlinear memory\\footnote{Throughout this paper the terms ``nonlinear memory'' and ``Christodoulou memory'' are used interchangeably. When not otherwise specified, ``memory'' refers to the ``Christodoulou memory''.} arises from a change in the \\emph{radiative-multipole moments} that is sourced by the energy flux of the radiated gravitational waves. One can heuristically understand the origin of the nonlinear memory as follows: Consider the Einstein field equations in harmonic gauge \\cite{kiprmp}: \\bs \\label{eq:EFEeqs} \\be \\label{eq:EFE} \\Box \\bar{h}^{\\alpha \\beta} = -16 \\pi (-g) ( T^{\\alpha \\beta} + t_{\\rm LL}^{\\alpha \\beta} ) -  {\\bar{h}^{\\alpha \\mu}}_{\\;\\;\\; , \\nu} {\\bar{h}^{\\beta \\nu}}_{\\;\\;\\; , \\mu} + \\bar{h}^{\\mu \\nu} {\\bar{h}^{\\alpha \\beta}}_{\\;\\;\\; , \\mu \\nu} , \\ee \\be \\label{eq:harmonicgauge} {{\\bar h}^{\\alpha \\beta}}_{\\;\\;\\; , \\beta}= 0, \\ee \\es where \\be \\label{hbardef} \\bar{h}^{\\alpha \\beta} \\equiv \\eta^{\\alpha \\beta} - \\sqrt{-g} g^{\\alpha \\beta} \\ee is the gravitational-field tensor\\footnote{\\label{ftnt:sign-h}Different notations for $\\bar{h}^{\\alpha \\beta}$ are used in the literature. For example, Blanchet \\cite{blanchetLRR} and Will and Wiseman (WW) \\cite{will-wiseman-2pn} both use the symbol $h^{\\alpha \\beta}$, but their conventions differ by a sign: $\\bar{h}^{\\alpha \\beta} = h^{\\alpha \\beta}_{\\rm WW} = - h^{\\alpha \\beta}_{\\rm B}$. However, all agree on the sign convention for the waveform $h_{jk}^{\\rm TT}$.}, $g$ is the determinant of the metric $g_{\\alpha \\beta}$, $T^{\\alpha \\beta}$ is the matter stress-energy tensor, $t_{\\rm LL}^{\\alpha \\beta}$ is the Landau-Lifshitz pseudotensor, $\\Box \\equiv -\\partial^2_t + \\nabla^2$ is the flat-space wave operator, a comma denotes a partial derivative (${}_{,\\mu}\\equiv \\partial_{\\mu}$), and $\\nabla^2$ is the flat-space Laplacian. While there are many nonlinear terms that enter the right-hand side of Eq.~\\eqref{eq:EFE}, there is a particular piece of $t_{\\rm LL}^{jk}$ [see the last term of Eq.~(2.7) in Ref.~\\cite{will-wiseman-2pn}] that is equal to the GW stress-energy tensor \\cite{isaacson,mtw}: \\be T^{\\rm gw}_{jk} = \\frac{1}{32\\pi} \\left\\langle h^{\\rm TT}_{ab, j} h^{\\rm TT}_{ab, k} \\right\\rangle \\approx T_{00}^{\\rm gw} n_j n_k = \\frac{1}{R^2} \\frac{dE^{\\rm gw}}{dt d\\Omega} n_j n_k , \\ee where $\\frac{dE^{\\rm gw}}{dt d\\Omega}$ is the GW energy flux, $n_j$ is a unit radial vector, the angle-brackets mean to average over several wavelengths,  $h_{ab}^{\\rm TT} = \\bar{h}_{ab}^{\\rm TT}$, and we have used the plane-wave approximation $h^{\\rm TT}_{ab} \\approx F_{ab}(t-R)/R +O(R^{-2})$.   When applying the standard Green's function to the right-hand side of Eq.~\\eqref{eq:EFE}, this piece---proportional to the energy flux of the emitted GWs---yields the following correction term to the GW field \\cite{wiseman-will-memory}: \\be \\label{eq:hTTnonlinmem} \\delta h^{\\rm TT}_{jk} = \\frac{4}{R} \\int_{-\\infty}^{T_R} dt'\\, \\left[ \\int \\frac{dE^{\\rm gw}}{dt' d\\Omega'} \\frac{n'_j n'_k}{(1-{\\bm n}' \\cdot {\\bm N})} d\\Omega' \\right]^{\\rm TT}, \\ee where $T_R$ is the retarded time. This equation shows that part of the distant GW field is sourced by the loss of GW energy.  Note that while the magnitude of the nonlinear memory approximately scales with the total radiated GW energy $\\Delta E^{\\rm gw}$, the angular-dependent unit vectors in the integrand of Eq.~\\eqref{eq:hTTnonlinmem} imply that the nonlinear memory is \\emph{not} directly proportional to $\\Delta E^{\\rm gw}$. The memory should not be mistaken as a change in the monopolar piece of the $1/R$ expansion of the metric. Rather it is a change in the quadrupolar (and higher-order) pieces of the $1/R$ part of the TT projection of the asymptotic spatial metric (see Sec.~\\ref{sec:hered} below).  Since the nonlinear memory occurs in any system that radiates GWs, systems that are usually considered to have vanishing linear memory (such as bound binaries) have a nonvanishing nonlinear memory. Thorne \\cite{kipmemory} has shown that the nonlinear memory [Eq.~\\eqref{eq:hTTnonlinmem}] can be described by his formula for the linear memory [Eq.~\\eqref{eq:hijlinmem}] if the unbound objects in the system are taken to be the individual gravitons with energies $E_A=M_A/(1-v_A^2)^{1/2}$ and velocities $v^j_A = c \\, n_A'^j$. The correspondence between these two equations also holds for null sources that contribute only to the linear memory. For example, the linear memory from a massless neutrino (or any other null particle) can be described either as a discrete sum over the memory from each individual particle [Eq.~\\eqref{eq:hijlinmem}], or by replacing the GW energy flux in Eq.~\\eqref{eq:hTTnonlinmem} with the energy flux of neutrinos \\cite{epstein-neutrinomemory}.  The leading-order PN expansion of Eq.~\\eqref{eq:hTTnonlinmem} is proportional to [see Eqs.~\\eqref{eq:hijTT}, \\eqref{eq:Uij}, and \\eqref{eq:canonical}] \\be \\label{eq:hTTmempropto} \\delta h^{\\rm TT}_{jk} \\propto \\frac{1}{R} \\left[ \\int_{-\\infty}^{T_R} {\\mathcal I}_{aj}^{(3)}(\\tau) {\\mathcal I}_{ka}^{(3)}(\\tau) d\\tau  \\right]^{\\rm TT}, \\ee where $ {\\mathcal I}_{ij}^{(3)}$ is the third-time derivative of the source-quadrupole moment. If we specialize to quasicircular, inspiralling binaries with total mass $M=m_1+m_2$, reduced mass ratio $\\eta=m_1 m_2/M^2$, and angular orbital frequency $\\omega(t)$, the integrand of Eq.~\\eqref{eq:hTTmempropto} contains oscillatory terms proportional to $\\eta^2 (M\\omega)^{10/3} e^{\\pm 4i\\omega t}$ and nonoscillatory terms proportional to  $\\eta^2 (M\\omega)^{10/3}$. When performing the time integration, the oscillatory terms are effectively multiplied by the \\emph{orbital timescale} $T_{\\omega}\\propto 1/\\omega$: \\bs \\label{eq:memintscaling} \\begin{align} \\label{eq:memintscaling1} \\delta h_{jk}^{\\textrm{TT, osc}} &\\propto \\frac{1}{R} \\int_{-\\infty}^{T_R} dt' \\, \\eta^2 (M\\omega)^{10/3} e^{\\pm 4i\\omega t'} \\nonumber \\\\ &\\propto \\frac{\\eta^2}{R} (M\\omega)^{10/3} e^{\\pm 4i\\omega t} T_{\\omega} \\propto \\frac{\\eta^2 M}{R} (M\\omega)^{7/3} e^{\\pm 4i\\omega t}, \\end{align} while the nonoscillatory terms are effectively multiplied by the \\emph{radiation-reaction timescale} $T_{\\rm rr} \\propto (M/\\eta) (M\\omega)^{-8/3}$: \\begin{align} \\label{eq:memintscaling2} \\delta h_{jk}^{\\textrm{TT, non-osc}} &\\propto \\frac{1}{R} \\int_{-\\infty}^{T_R} dt' \\, \\eta^2 (M\\omega)^{10/3} \\propto \\frac{\\eta^2}{R} (M\\omega)^{10/3} T_{\\rm rr} \\nonumber \\\\ &\\propto \\frac{\\eta M}{R} (M\\omega)^{2/3}. \\end{align} \\es This shows that while the oscillatory pieces of $\\delta h^{\\rm TT}_{jk}$ are a 2.5PN correction to the waveform amplitudes, the nonoscillatory piece enters at the same order as the familiar quadrupole-order piece of the waveform:\\footnote{This also follows from the approximate scaling of the nonlinear memory with the total radiated energy, $\\delta h_{jk}^{\\rm TT} \\sim \\Delta E^{\\rm gw}/R$. Since the energy radiated during the inspiral is equal to the change in the orbital energy, $\\Delta E^{\\rm gw} \\approx (\\eta M/2)(M/r)$, the nonlinear memory has the same scaling as the quadrupolar waveform $h_{jk}^{\\rm TT,(0)} \\propto (\\eta M/R)(M/r)$, where $r$ is the orbital separation.} \\be \\label{eq:quadapprox} h_{jk}^{\\rm TT,(0)} = \\frac{2}{R} \\ddot{\\mathcal I}_{jk}^{\\rm TT} \\propto \\frac{\\eta M}{R} (M \\omega)^{2/3} . \\ee \\subsection{\\label{sec:motivation}Motivation} The Christodoulou memory is a unique and interesting manifestation of the nonlinearity of general relativity. Analytic computations of gravitational radiation involve many types of nonlinearities, the origins of which are often obscured by the complex, iterative algorithms involved in solving the Einstein field equations. The Christodoulou memory, however, has a clear physical interpretation: It arises from the loss of GW energy from the system and the effect of this loss (through the GW stress-energy tensor) on the system's radiative mass-multipole moments. For comparison, GW tails are another interesting nonlinear, general-relativistic effect (see Sec.~3.4 of Ref.~\\cite{blanchet-marck-lasota} for references). They arise from the last term in Eq.~\\eqref{eq:EFE}, which modifies the flat-spacetime wave operator on the left-hand side and causes backscattering of the gravitational radiation as it propagates through the curved spacetime around the binary \\cite{wiseman-tails}.  Tail effects enter the waveform at 1.5PN order. The Christodoulou memory, on the other hand, arises from nonlinear interactions at 2.5PN order, but affects the gravitational waveform at leading (0PN) order. Both tails and the Christodoulou memory are \\emph{hereditary}---their contribution to the GW field at any given retarded time depends on the entire past history of the source [see e.g., Eq.~\\eqref{eq:Uij}]. However, the Christodoulou memory is especially sensitive to the motion of the source in the distant past (see Sec.~\\ref{sec:NRmemory}), while tails are primarily sensitive to the recent past (see also Sec.~4 of Ref.~\\cite{arun25PNamp}). Tails and most other nonlinear effects that enter the waveform cause oscillatory corrections to both the $+$ and $\\times$ polarizations. In contrast, the Christodoulou memory causes a nonoscillatory shift in the amplitude of the $+$ polarization only\\footnote{This statement is only true for standard choices of the polarization triad (see Sec.~\\ref{sec:polmodes}). Since a rotation of this triad by an angle $\\Psi$ about the propagation direction transforms the GW polarizations via $h_+ - i h_{\\times} \\rightarrow (h_+ - i h_{\\times}) e^{2i \\Psi}$, a purely $+$ polarized wave can become mixed-polarized.}. This amplitude shift starts small at early times (when the binary is widely separated) and slowly grows during the inspiral. As the binary components merge the memory rapidly grows and then saturates to a final value during the ringdown phase. The details of how the memory reaches its saturation value are explored in Refs.~\\cite{favata-memory-saturation,favata-lisa7confproc}. Unlike tails, the nonoscillatory pieces of the memory do not affect the orbital phase of a quasicircular binary up to 3.5PN order (the highest order to which the phase has been computed), but they could modify the phase at higher PN orders\\footnote{For example (and using notation defined later), the GW luminosity ${\\mathcal L}=-\\dot{E}\\propto \\sum_{lm} |\\dot{h}_{lm}|^2$ will have a term proportional to $|\\dot{h}_{20}|^2$ which will affect the luminosity and hence the phase at relative 5PN order (see Sec.~\\ref{sec:selecrules} below for scalings). However, one should note that the \\emph{oscillatory} terms that arise from the memory integral [cf.~Eq.~\\eqref{eq:memintscaling1} and associated discussion] affect the phasing beginning at relative 2.5PN order.}.  Over the past three decades the PN corrections to the oscillatory pieces of the GW polarizations have been computed \\cite{wagoner-will,blanchet-iyer-will-wiseman-CQG-2PNwaveform,arun25PNamp,kidder-blanchet-iyer-25PNwaveform}, most recently at the 3PN order \\cite{blanchet3pnwaveform}. These computations were motivated by the development---and finally the operation---of a global network of ground-based GW interferometric detectors \\cite{ligoweb,geoweb,virgoweb,tamaweb}.  These detectors are primarily sensitive to the oscillatory components of the GW signal, as most of the signal power lies in the lowest-order oscillatory modes of the radiation (for nearly circular orbits). More recently, successes in numerical relativity (NR) \\cite{pretorius-PRL2005,pretorius-CQG2006,pretorius-BBHreview,baker-etal-PRL2006,baker-etal-PRD2006,campanelli-etal-PRL2006,campanelli-etal-PRD2006,herrmann-etal-CQG2007,sperhake-PRD2007,scheel-etal-PRD2006,koppitz-etal-PRL2007,pollney-etal-spinorbitrecoil,scheel-merger,samurai09,ninja09} have necessitated the need for accurate PN waveforms to compare with the results of binary black-hole (BH) merger simulations \\cite{buonanno-cook-pretorius,boyle-etal-PRD2007,mroue-kidder-saul-PRD2008,boyle-etal-Efluxcomparison,baker-etal-PRL2007-NRPN,baker-etal-PRD2007-NRPN,hinder-etal-eccentricPN-NR,campanelli-etal-spinning-NR-PN-compare,hannam-etal-PN-NR-meet,berti-etal-multipolarnonspinning,berti-etal-mulitpolarspinning,hannam-gopa-NRPN-spin,damour-nagar-jena,damour-nagar-AEI,damour-nagar-caltechcornell,damour-nagar-PRD09,gopa-jena-eccentricPNNR,pan-buonanno-baker-etal-NRPN,buonanno-pan-baker-etal-nonspinningEOB,buonanno-caltechEOB09}.  Computations of the nonlinear memory's contribution to the waveform have not progressed as far. Wiseman and Will \\cite{wiseman-will-memory} first calculated the nonlinear memory's leading-order effect on the waveform polarizations for a quasicircular, inspiralling binary. Their result was confirmed\\footnote{The factor of $2$ discrepancy in the memory waveform computed in Ref.~\\cite{wiseman-will-memory} is either a typo or due to a different choice of normalization for the polarization tensors in Eq.~\\eqref{eq:epluscross} below. Note that the relative amplitudes shown in Fig.~1 of Ref.~\\cite{wiseman-will-memory} have the correct values at large separation.} in Refs.~\\cite{kennefick-memory,arun25PNamp,blanchet3pnwaveform}. The 0.5PN corrections to the memory were calculated in Ref.~\\cite{blanchet3pnwaveform} and found to vanish.  The primary purpose of this paper is to compute the corrections to the leading-order formula for the nonlinear memory to 3PN order.  Why are these corrections needed? First, while GW interferometers are mostly sensitive to the oscillatory parts of the GW, the memory piece of the signal could be detectable for certain sources. During a signal's observation time, the nonlinear memory causes a growing change in the signal's amplitude. This change contributes power at low frequencies. For many sources this power is swamped by the detector's low-frequency noise. But for sources with large signal-to-noise ratios---especially supermassive black-hole binaries in the low-frequency LISA \\cite{lisaweb} band---the memory can be detectable. For example, the memory from the merger of two $10^6 M_{\\odot}$ black holes should be detectable by LISA out to a redshift of $z\\approx 2$ (see Fig.~3 of Ref.~\\cite{favata-memory-saturation}).  The detectability of the nonlinear memory was previously considered by Thorne \\cite{kipmemory} and Kennefick \\cite{kennefick-memory} and is discussed further in Sec.~\\ref{sec:memdetec} below and in Refs.~\\cite{favata-memory-saturation,favata-lisa7confproc}. Central to the issue of detecting the memory is knowing its magnitude, and computing higher-PN corrections to the memory will help to determine this more accurately.  Second, the comparison of PN waveforms with numerical relativity simulations could also benefit from more accurate expressions for the memory. While numerical relativity accounts for all of the nonlinear effects of general relativity, it is difficult to compute the memory accurately in these simulations: The modes of the waveform that have memory are very small and depend sensitively on the initial separation of the binary (see Sec.~\\ref{sec:NRmemory} for details). The PN corrections to the memory computed here could supply initial conditions for the waveform modes in the numerical simulations. As has been done with the oscillatory pieces of the waveform \\cite{buonanno-cook-pretorius,boyle-etal-PRD2007,mroue-kidder-saul-PRD2008,boyle-etal-Efluxcomparison,baker-etal-PRL2007-NRPN,baker-etal-PRD2007-NRPN,hinder-etal-eccentricPN-NR,campanelli-etal-spinning-NR-PN-compare,hannam-etal-PN-NR-meet,berti-etal-multipolarnonspinning,berti-etal-mulitpolarspinning,hannam-gopa-NRPN-spin,damour-nagar-jena,damour-nagar-AEI,damour-nagar-caltechcornell}), it will also be insightful to compare the memory calculated in future simulations with the PN expressions computed here.  Lastly, it is aesthetically pleasing to have all of the pieces---both oscillatory and nonoscillatory---of the waveform amplitudes computed consistently to the same post-Newtonian order. Combined with the results of Ref.~\\cite{blanchet3pnwaveform}, this work completes the waveform polarization amplitudes to 3PN order. \\subsection{\\label{sec:summary}Summary} The bulk of this paper is devoted to calculating the nonlinear memory's contribution to the $+$ waveform polarization and its PN corrections. Readers uninterested in the details should skip to the results presented in Sec.~\\ref{sec:results}. Readers uninterested in lengthy formulas can skip directly to Figs.~\\ref{fig:H-theta} and \\ref{fig:x-hlm}, the discussion (Sec.~\\ref{sec:discussion}), and conclusions (Sec.~\\ref{sec:conclusion}). An outline of the paper and a summary of its main results are presented here:  The necessary PN wave-generation formalism is briefly reviewed in Sec.~\\ref{sec:pnformalism}. Section \\ref{sec:polmodes} gives expressions for the waveform polarizations $h_{+,\\times}$ in terms of the radiative-multipole moments. Expressions for the spin-weighted spherical-harmonic mode decomposition of the waveform in terms of the ``scalar'' radiative-multipole moments are also introduced. Post-Newtonian calculations of the radiative-multipole moments are usually performed entirely in terms of symmetric-trace-free (STF) tensors, and are then (if needed) decomposed into spherical-harmonic modes at the end of the calculation. However, the computation of the memory is significantly easier if one works almost entirely from the start in terms of ``scalar'' quantities that are the coefficients of tensor quantities decomposed on the basis of scalar or spin-weighted spherical harmonics. Sections \\ref{sec:MPM} and \\ref{sec:canonicalmoments} review the multipolar-post-Minkowskian (MPM) formalism which relates the radiative-multipole moments ${\\mathcal U}_L$ and ${\\mathcal V}_L$ to the canonical moments ${\\mathcal M}_L$ and ${\\mathcal S}_L$ and the source moments ${\\mathcal I}_L$ and ${\\mathcal J}_L$. Section \\ref{sec:hered} focuses on the hereditary contributions to the radiative multipoles: tails and memory. The results of Blanchet and Damour \\cite{blanchet-damour-hereditary}, which serve as the starting point of the memory calculation, are reviewed. The main result of this section is Eq.~\\eqref{eq:Ulmmem}, which expresses the memory piece of the radiative mass-multipole moments in terms of time and angular integrals over the GW energy flux from the source.  The main part of the memory calculation---explicitly evaluating the nonlinear memory contributions to the radiative-mass multipoles for quasicircular binaries---is detailed in Sec.~\\ref{sec:evalUmem}. The computation of the angular integrals is discussed in Sec.~\\ref{sec:angularint}. Section \\ref{sec:selecrules} discusses which $(l,m)$ modes need to be calculated to determine the memory contribution to the waveform to 3PN order. In Sec.~\\ref{sec:timeintegrals} we discuss how the time integral over the past history of the source is computed. In particular we need to use a model for the adiabatic inspiral that is accurate to at least 3PN order. This model is contained in Eq.~\\eqref{eq:xdot35pn}, which gives the time evolution of the PN parameter $x\\equiv (M\\omega)^{2/3}$ to 3.5PN order. This formula (which easily follows from the results of Refs.~\\cite{blanchet35PNphase,blanchet35PNphaseerratum}) expresses the well-known result for the 3.5PN frequency evolution in a form that is useful for computations involving the PN parameter $x$.  The main results of this work are listed in Sec.~\\ref{sec:results}: the 3PN memory contributions to the $+$ polarization $h_{+}$ of the gravitational waveform and its spin-weighted spherical-harmonic modes $h_{lm}$. Those expressions can be directly combined with the oscillatory pieces of the waveform given to 3PN order in Ref.~\\cite{blanchet3pnwaveform}. Figure \\ref{fig:H-theta} illustrates the memory's angular dependence and the relative sizes of the various PN corrections. The PN corrections have little effect on the memory's angular dependence, but tend to decrease the overall magnitude of the memory.  Aside from the Christodoulou memory, there are additional nonoscillatory (DC or ``direct current'') time-varying terms present in the waveform. Section \\ref{sec:crosmem} discusses a new type of \\emph{nonlinear, nonhereditary} DC term discovered by Arun et al.~\\cite{arun25PNamp}. Unlike the Christodoulou memory, this nonlinear, nonhereditary DC term affects the $\\times$ polarization and first enters the waveform at 2.5PN order. Section \\ref{sec:linmem} discusses an additional class of \\emph{linear, nonhereditary} DC terms that enter the waveform at 5PN and higher orders. These terms arise from the effects of radiation-reaction or nongravitational forces on the source-multipole moments.  Section \\ref{sec:NRmemory} discusses the challenges in extracting the nonlinear memory from numerical simulations of binary black holes. It will be difficult for current simulations to accurately extract the memory waveform. One of the reasons for this is illustrated in Fig.~\\ref{fig:x-hlm}. Numerical relativity simulations can most accurately compute the $(l,m)=(2,2)$ mode of the waveform; but the nonlinear memory is present only in the $m=0$ modes (for binaries orbiting in the $x$-$y$ plane). For simulations that can directly compute the metric-perturbation modes $h_{lm}$, the largest memory mode $h_{20}$ is an order-of-magnitude smaller than the $h_{22}$ mode. For the majority of simulations that compute the spin-weighted spherical-harmonic modes of the $\\Psi_4$ Weyl scalar, $\\psi_{lm}=\\ddot{h}_{lm}$, the situation is significantly worse: The largest memory mode $\\psi_{20}$ is nearly \\emph{three orders of magnitude} smaller than the dominant $\\psi_{22}$ mode in the late inspiral. There are also difficulties in determining the two integration constants needed when computing the $h_{lm}$ modes from the $\\psi_{lm}$ modes. The memory's strong dependence on the past history of the source also introduces significant errors unless the simulations start with very large binary separations. The results of this paper could help to alleviate some of these difficulties by providing initial conditions for the $m=0$ modes.  Section \\ref{sec:memdetec} discusses the detection of the memory. While the memory from a signal in the distant past is unobservable, it is possible to detect the buildup of the memory from a passing GW. Previous work by Thorne \\cite{kipmemory} and Kennefick \\cite{kennefick-memory} has estimated the nonlinear memory's signal-to-noise ratio in laser interferometers. However, the signal-to-noise ratio is sensitive to the details of how the memory rises to its saturation value, which previous works have not properly modeled.  Conclusions and suggestions for further work are presented in Sec.~\\ref{sec:conclusion}. Some results relegated to the appendices are potentially useful in other applications: Appendix \\ref{app:angularint} gives an explicit formula for the angular integral of the product of three spin-weighted spherical harmonics.  Appendix \\ref{app:psilm} gives an explicit prescription for calculating the Weyl scalar modes $\\psi_{lm}$ from the known expressions for $h_{lm}$ presented here and in Ref.~\\cite{blanchet3pnwaveform}. \\subsection{\\label{sec:notation}Notation} This paper borrows notations and conventions from Kidder \\cite{kidder08}, Thorne \\cite{kiprmp}, and Blanchet, et al.~\\cite{blanchet3pnwaveform}. We generally set $G=c=1$, except in certain equations where we wish to make explicit the post-Newtonian or post-Minkowskian order of the various terms. We use the notation $O(n)$ to denote post-Newtonian (PN) correction terms of order $O(c^{-n})$. Spacetime indices are denoted with Greek letters, spatial indices with Latin letters.  In Sec.~\\ref{sec:MPM} we use Blanchet's \\cite{blanchetLRR} notation for the gravitational field, $h^{\\alpha \\beta} = - \\bar{h}^{\\alpha \\beta}$ (see footnote \\ref{ftnt:sign-h}). Multi-index notation is denoted with a capital subscript: $A_L = A_{i_1 i_2 \\ldots i_l}$. A multi-index $L$ on a vector denotes a product of $l$ vectors: $x_L = x_{i_1} x_{i_2} \\ldots x_{i_l}$. Repeated spatial indices and multi-indices are summed regardless of their relative positions. Symmetric-trace-free (STF) spatial tensors are denoted with capital script letters  (as in ${\\mathcal U}_L$). The corresponding ``scalar'' versions of these moments [their coefficients on the basis of the STF spherical-harmonic tensors ${\\mathcal Y}_L^{lm}$ (defined below)] are denoted with nonscript capital letters (as in $U_{lm}$). Spherical-harmonic indices $(l,m)$ are raised or lowered arbitrarily (i.e.,~$U_{lm} = U^{lm})$.  Symmetrization and STF projection are denoted by enclosing the relevant indices by $()$ or $<>$, respectively. Indices that are left out of the symmetrization or STF projection are displayed with an underbar: $A_{<i}B_{\\underline{a}j>}$. Time derivatives are denoted by an overdot or by superscript parenthesis: ${\\mathcal M}_L^{(p)}= d^p{\\mathcal M}_L/{dT^p}$. To avoid confusion with the azimuthal harmonic index $m$, we denote the sum of the binary's component masses as $M\\equiv m_1+m_2$. This is distinct from the total binary (ADM) mass or mass-monopole moment, which we denote ${\\mathcal M}$. Euler's constant is denoted by $\\gamma_{E} = 0.577\\,21\\ldots$. The symbol $\\phi$ denotes an azimuthal angle while $\\varphi$ denotes an orbital phase. ",
        "conclusions": "Conclusions} The nonlinear (Christodoulou) memory is an interesting and potentially observable effect. Previous work has computed only the leading-order memory contributions to the gravitational-wave polarizations \\cite{wiseman-will-memory,kennefick-memory,arun25PNamp,blanchet3pnwaveform}. Here the post-Newtonian corrections to the memory for quasicircular, inspiralling binaries have been computed. The main results are the equations listed in Sec.~\\ref{sec:results} which gives the memory's contribution to the waveform's spin-weighted spherical-harmonic modes and the $+$ polarization. These results are illustrated graphically in Fig.~\\ref{fig:H-theta} and show that the post-Newtonian corrections to the memory are important in the late inspiral. The calculations presented here complete the memory to 3PN order. Other nonoscillatory effects on the waveform were also discussed. In particular, a nonlinear, nonhereditary DC (nonoscillatory) term is present in the $\\times$ polarization at 2.5PN order \\cite{arun25PNamp}, while linear DC terms can enter the waveform at 5PN and higher orders.  While most studies have focused on the memory in quasicircular orbits, the memory should also be calculated for other types of orbits. For hyperbolic orbits with small deflection angles, the nonlinear memory has been calculated in Ref.~\\cite{wiseman-will-memory}. In future work the memory for inspiralling eccentric binaries \\cite{favata-eccentricmemory} will also be computed. This is especially important because circularized binaries have growing eccentricity in the past, and the nonlinear memory is sensitive to the binary's past history.  One of the main conclusions of this paper is that it will be difficult (but not impossible) to directly extract the memory from numerical relativity waveforms. Nonetheless, numerical relativity will be critical in providing a full understanding of the memory. These simulations will eventually better determine the size of the memory in the late inspiral, and they will also be able to compute the evolution and saturation value of the memory during the merger and ringdown phases of coalescence \\cite{favata-memory-saturation}. The formulas presented here could be compared with the inspiral phase of those future simulations. They could also be used to set the memory's initial value at the start of a simulation.  Detecting the memory in binary black-hole coalescences could serve as a test of general relativity. However, the nonoscillatory nature of the memory makes its detection difficult. Detections will only be likely for sources whose primary waves will have high signal-to-noise ratios, such as supermassive black-hole binaries in the LISA band. This paper argues that previous signal-to-noise ratio estimates of the memory may contain large errors. This issue is investigated further in Ref.~\\cite{favata-memory-saturation}.  Aside from the formal mathematical descriptions in Refs.~\\cite{christodoulou-mem,frauendiener-memnote}, the memory has only been explored in detail with post-Newtonian theory. It would be interesting to also explore the memory using second-order black-hole perturbation theory. Recent work along these lines has demonstrated a memory effect in the second-order quasinormal modes of the Schwarzschild spacetime \\cite{nakano-ioka-2ndorderQNM}. It would also be interesting to see how the nonlinear memory manifests itself in alternative theories of gravity."
    },
    "0812/0812.2034.txt": {
        "abstract": "We  use photometric data from the  Spitzer mission to explore the mid- and far-infrared  properties of 10 red QSOs ($J-K>2$, $R-K>5$\\,mag)  selected by  combining the  2MASS in  the near-infrared with  the SDSS  at optical  wavelengths. Optical  and/or near-infrared spectra are available for  8/10 sources.  Modeling the Spectral Energy Distribution (SED)  from UV to  far-infrared shows that  moderate dust reddening   ($A_V = 1.3 - 3.2$) can explain  the red  optical  and near-IR  colours of  the sources in the sample. There is  also evidence that red QSOs have $\\rm 60/12\\mu  m$  luminosity  ratio  higher  than PG  QSOs  (97  per  cent significance).   This  can  be  interpreted as  a  higher  level  of star-formation  in these  systems (measured  by the  $\\rm 60  \\,\\mu m$ luminosity) for a  given AGN power (approximated by  the $\\rm 12 \\,\\mu m$  luminosity).  This  is consistent  with a  picture where  red QSOs represent an early phase of AGN evolution, when the supermassive black hole is  enshrouded in dust and  gas clouds, which  will eventually be blown out (possibly by AGN driven outflows) and the system will appear as typical  optically luminous QSO.  There is  also tentative evidence significant  at the  96\\% level  that red  2MASS QSOs  are  more often associated  with radio  emission  than optically  selected SDSS  QSOs. This  may  indicate  outflows,  also  consistent with  the  young  AGN interpretation.   We  also estimate  the  space  density  of red  QSOs relative  to optically  selected SDSS  QSOs, taking  into  account the effect of dust extinction and the intrinsic luminosity of the sources. We estimate  that the fraction of  red QSOs in  the overall population increases from  3\\% at  $M_K=-27.5$\\,mag to 12\\%  at $M_K=-29.5$\\,mag. This suggests that  reddened QSOs become more important  at the bright end of the Luminosity Function.  If red QSOs are transition objects on the way to becoming typical optically luminous QSOs, the low fractions above  suggest  that these  systems  spent  less  than 12\\%  of  their lifetime at the ``reddened'' stage. ",
        "introduction": "In  the  last  few  years  an increasing body  of evidence  points to an  intimate link  between the formation of galaxies  and the growth of the  supermassive black holes (SBH) at  their centres.   For example, the  accretion history  of the Universe (Barger et al.  2005; Hasinger et al.  2005) looks strikingly similar   to  the   evolution  of   the  cosmic   star-formation  rate \\citep[e.g.][]{Hopkins_A2004,   Hopkins_A2006}.    Additionally,  most nearby  spheroids contain  a  SBH \\citep[SBH;  e.g.][]{Magorrian1998}, with a  mass that scales with  the stellar velocity  dispersion of the host   galaxy    bulge   \\citep[e.g][]{Ferrarese2000,   Gebhardt2000}, suggesting similar formation epochs and/or mechanisms.  Recent cosmological simulations account  for the observations above by assuming  that  both  the  galaxy  and  the  SBH  form  during  galaxy interactions  and mergers  \\citep[e.g][]{Kauffmann2000, Kauffmann2002, DiMatteo2003, Menci2004, Cattaneo2006}.   Such violent events are very efficient in driving gas to  the nuclear regions, where it forms stars or it  is consumed  by the  SBH.  Because the  nuclear regions  of the interacting system  are dusty, both  processes above most  likely take place within a cocoon of dust and gas clouds. Energetic considerations underline the  significance of the energy  released by the  AGN in the process    above,    which    can    regulate    the    gas    inflows \\citep[e.g][]{Silk1998,  Fabian1999, King2003}.  Simulations  of major mergers indeed show that the feedback from the AGN can be particularly violent close  to the final  merging, leading to  the blow out  of the dust and gas clouds and the termination of both the star-formation and eventually the AGN activity \\citep[e.g.][]{DiMatteo2005, Springel2005, Hopkins2006}.  The key  prediction of the  picture above is  that there should  be an association  between SBH accretion,  starbursts and  mergers. Although there   are  examples   of  such   systems  \\citep[e.g.][]{Genzel1998, CidFernandes2001, Franceschini2003, Georgakakis2004, Alexander2005}, they appear to be the  exception rather than the rule.  X-ray surveys, arguably the  most efficient method  for finding AGN, have  shown that the  majority of  X-ray sources  are  hosted by  red evolved  galaxies \\citep{Nandra2007},  with prominent  bulges  and little  morphological evidence  for  ongoing  interactions  \\citep{Grogin2003,  Pierce2007}. This  may  suggest  that  X-ray  wavelengths, because  of  their  high sensitivity to AGN, select a large number of systems after the peak of their activity (i.e.   final merging) and during the  decline stage of the SBH accretion.   In order to understand the  interplay between SBH and  galaxy formation  one should  study the  properties of  AGN hosts (e.g. star-formation  rate, optical morphology)  close to the  peak of the activity, when  the system appears as luminous  QSO.  Such studies however, are  hampered by  the enormous brightness  of the  QSO nuclei relative to the  host galaxy.  Nevertheless progress has  been made in the last  few years.   Mid-Infrared spectroscopy with  Spitzer reveals PAH  features, i.e.   evidence for  star-formation, in  about  30\\% of optically  selected luminous QSOs  \\citep{Schweitzer2006, Netzer2007}. HST  high  resolution  imaging  and  ground  based  observations  with adaptive optics find  that a fraction of the  QSO population is hosted by either  elliptical galaxies with fine structure  indicative of past interactions  \\citep[80\\%;][]{Bennert2008}, or  galaxies  with obvious signs of  disturbance \\citep[30\\%;][]{Guyon2006}.  The  evidence above suggests  that  at  least  some  QSOs  are  formed  in  major  mergers accompanied by starburst events.  An important development  in studies of the interplay  between AGN and galaxies  has been  the  identification of  QSOs  with optical  and/or near-infrared  (near-IR) colours  much  redder than  those of  typical UV/optically  selected QSOs.   \\cite{Cutri2001}  and \\cite{Wilkes2002} for example,  have found a  previously undetected population  of broad line AGN  among sources  with $J-K>2$\\,mag in  the Two Micron  All Sky Survey  (2MASS).   \\cite{White2003}   studied  an  $I$-band  magnitude limited ($I<20.5$\\,mag)  sample of radio  sources and showed  that the QSOs in that sample have  redder $B-R$ colours, on average, than those selected  at bluer  optical filters.   \\cite{Glikman2004, Glikman2007} demonstrated that selecting radio  (1.4\\,GHz) sources in the region of the colour  space $J-K>1.7$  and $R-K>4.0$\\,mag yields  a 50  per cent success rate  of discovering QSOs  that are substantially  redder than those  found in optical  surveys. Similar  results were  obtained more recently by \\cite{Urrutia2008a} using a slightly modified colour wedge for  selecting  radio  sources  ($J-K>1.3$ and  $r-K>5.0$\\,mag).   The optical/near-IR continua  of the red QSO population  identified in the studies above  are consistent with moderate amounts  of dust reddening \\citep[$A_V\\approx1.5-3$][]{White2003, Glikman2007, Urrutia2008a}.  An intrinsically  flat UV/optical spectrum  \\citep[e.g.][] {Richards2003} or   synchrotron    emission   with   a    high   turnover   frequency \\citep{Whiting2001} can also explain  the red colours without invoking dust.  It is unlikely  however, that these alternative scenarios apply to the entire  population of red QSOs.  It  is therefore proposed that this new  population of dust  reddened QSOs represent  systems shortly before or during the blow out stage of AGN evolution, when the central source is  still cocooned in dust  and gas clouds.   If true, reddened QSOs  are  intermediate  between  dusty  starbursts  and  UV/optically luminous QSOs. Clues on the evolutionary stage of these systems can be obtained  by studying  the properties  of their  host  galaxies.  Deep optical imaging for example, has  shown that the morphology of the red QSO hosts are in most cases disturbed, suggesting a higher fraction of ongoing  interactions compared to  typical UV/optically  luminous QSOs \\citep{Hutchings2003, Hutchings2006, Urrutia2008}, consistent with the young AGN picture.  Another property  of red QSOs that  has not yet been  addressed and is particularly  relevant to  their evolutionary  stage is  the  level of star-formation rate in the host galaxy.  In this paper we address this issue using Spitzer to constrain and to model the mid- and far-IR SEDs of a  sample of red QSOs identified  in the Two Micron  All Sky Survey (2MASS).  Under reasonable  assumptions, the shape of the  IR SED is a powerful diagnostic  of the physical processes  (star-formation vs AGN activity) responsible  for heating  the dust in  extragalactic sources \\citep[e.g.][]{rowan2008}.  The  fraction  of  red QSOs  in  the  overall  population is  also  an important parameter.  In the  scenario where these systems represent a stage  of the  AGN phenomenon,  their fraction  provides clues  on the timescale of this stage.  Currently, there is large uncertainty in the fraction  of red QSOs  missing from  optically selected  samples, with values ranging from 15  to over 50 per cent \\citep[e.g.][]{Wilkes2002, Richards2003, White2003,  Glikman2004, Glikman2007}. The uncertainty is associated with the  selection of the appropriate comparison sample of  optically selected QSOs  and with  difficulties in  accounting for dust extinction  in the selection of  reddened QSOs. In  this paper we address these  issues by  using the 1/Vmax  formalism to  estimate the maximum volume  that a source  is detectable, given the  survey limits and  taking  into  account  the  effect of  dust  extinction  and  the intrinsic luminosity  of the source.  For comparison we use  the Sloan Digital   Sky  Survey   (SDSS)  optically   selected  QSO   sample  of \\cite{Schneider2005} , which is  large and with well defined selection criteria.  Throughout  the paper  we adopt  $\\rm H_{0} =  70 \\,  km \\, s^{-1} \\, Mpc^{-1}$, $\\rm  \\Omega_{M} = 0.3$ and $\\rm \\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "We combine optical and near-IR photometry from the ground with Spitzer mid- and  far-IR broadband data to explore  the SEDs of a  sample of 10 QSOs with  red optical and  near-IR colours selected by  combining the 2MASS with the SDSS. The main conclusions are  \\begin{itemize}  \\item The UV/optical part of the  SED can be fit assuming an intrinsic spectrum similar  to that of typical optically  selected QSOs reddened by moderate amounts of dust.  \\item  There is  evidence that  red 2MASS  QSOs have  higher  level of star-formation  (approximated by  the  $\\rm 60\\mu  m$ luminosity)  for their AGN power  (measured by the $\\rm 12\\mu  m$ luminosity), compared to optically selected PG QSOs.  \\item We  find that  2MASS QSOs are  more often associated  with radio emission,  compared to  optically  selected SDSS  QSOs.  This  excess, although interesting, is only significant at the 96\\% level  \\item The  fraction of  red QSOs in  the overall  population increases from    $3\\pm2$\\%    at    $M_K=-27.5$\\,mag   to    $12\\pm8$\\%    at $M_K=-29.5$\\,mag,   suggesting  that   these  systems   become  more important at the bright luminosities.  \\item We argue that these findings are consistent with a picture where reddened QSOs are young AGN shortly before they shed their cocoon of dust  and  gas  clouds  and  become  typical  UV/optically  luminous QSOs.  Comparison with  Seyfert-1 ULIRGs,  another class  of sources proposed  to be  young QSOs,  suggests  that red  2MASS QSOs  likely represent an earlier evolutionary stage.  \\end{itemize}"
    },
    "0812/0812.0145.txt": {
        "abstract": "In 2006 June, the obscured low luminosity active galactic nucleus in the nearby Seyfert 1.9 galaxy NGC~4258 was observed with Suzaku for $\\sim$ 100 ks. Utilizing the XIS and the HXD, the nucleus emission was detected over a $\\sim$ 2 to $\\sim$ 40 keV range, with an unabsorbed 2--10 keV luminosity of $\\sim 8 \\times 10^{40}$~erg~s$^{-1}$, and varied by a factor of $\\sim$ 2 during the observation. Its 2--40 keV spectrum is reproduced by a single power law with photon index of $\\Gamma$~$\\sim$ 2.0, absorbed by an equivalent hydrogen column of $\\sim$ 1.0 $\\times$ $10^{23}$ cm$^2$. The spectrum within $4'$ of the nucleus required also a softer thin-thermal emission, as well as an intermediate hardness component attributable to integrated point sources. A weak neutral Fe-K$\\alpha$ florescence line was detected at an equivalent width of $\\sim 40$ eV. The cold reflection component was not required by the data, with the reflector solid angle $\\Omega$ seen from the nucleus constrained as $\\Omega / 2 \\pi$ $\\lesssim 0.3$ assuming a general case of 60$^\\circ$ inclination. The results suggest that the cold reflecting material around the nucleus is localized along our line of sight, rather than forming a thick torus. ",
        "introduction": "\\label{sec:intro} %================== Although active galactic nuclei (AGNs) exhibit a wide variety of spectral properties, a comprehensive classification of them, called ``Unified Scheme'' (e.g., \\cite{Antonucci1993}), has been developed. In addition to the black hole mass and the mass accretion rate which are obvious parameters, this scheme employs two more fundamental parameters: radio loudness and our viewing direction to the accretion plane. A ``type~1'' AGN refers to an object with a nearly pole-on viewing angle, while a ``type 2'' AGN with a roughly edge-on aspect. The spectrum differs between the two classes depending on whether our line of sight is blocked by some part of the accreting material or not. The blocking material is often assumed to have a shape of an inflated torus around the nucleus, so called ``molecular torus'' or simply ``torus\". This scheme can explain much of the observational differences among various kinds of AGNs. Nevertheless, the validity of the Unified Scheme and the nature of the putative torus are still open.   X-ray observations of type~2 AGNs provide a powerful tool in these attempts, because the strength of photoelectric absorption affecting the nuclear hard X-rays provides the most direct information on the column density of the obscuring material. Indeed, X-ray observations with Ginga (e.g., Koyama et al. 1989; Awaki et al. 1991), and ASCA ( Ueno et al. 1996) of typical type~2 AGNs demonstrated that their nuclei are heavily obscured with a column density of $10^{23}$ to $10^{24}$ cm$^{-2}$. Subsequently, BeppoSAX investigators introduced a new concept called ``Compton thick AGNs\" (\\cite{Maiolino1998}, \\cite{Matt2000}), as the most extreme class of type 2 AGNs, wherein the absorber is opaque not only to photoelectric absorption, but also to Compton scattering. Because of the strong suppression of the direct component, these objects are expected to provide further clues to the circumnuclear matter distribution, via detection of various secondary components, including in particular the Compton reflection hump, and fluorescent Fe-K lines.   Using Suzaku, we observed the Compton-thick AGN, NGC~4945, which is one of the brightest AGNs above 20 keV. Analyzing the data, Itoh et al. (2008) found the reflection features to be unusually weak, as judged from both the Fe-K edge feature in the XIS spectra and the Compton hump in the data from the HXD. Reinforced by the detection of clear hard X-ray variations, Itoh et al. (2008) thus constrained the reflection fraction \\ensuremath{f_\\mathrm{refl}} $\\equiv$ $\\Omega$/2$\\pi$ to be less than a few percent, where $\\Omega$ is the solid angle of the reflector as seen from the nucleus. This makes a significant contrast to many other Compton-thick Seyfert 2 AGNs, where we usually find \\ensuremath{f_\\mathrm{refl}} $\\sim$ 1-2 (e.g. De Rosa et al. 2008). Therefore, the reflector in NGC~4945 is suggested to have a geometrically-thin, disk-like structure rather than that of a thick torus, and the reflector, absorber, and the water maser source are thought to be provided by the same structure, namely a flat accretion disk (Madejski et al. 2000, 2006). Although this means a clear deviation from the Unified Scheme, it is not yet clear whether such a property is specific to NGC~4945, or more or less common to a certain class of AGNs. One of effective ways to examine this issue is to study AGNs with low mass-accretion rates, namely low luminosity AGNs (LLAGNs), because NGC~4945 has a rather low intrinsic (absorption corrected) luminosity as $\\sim 1 \\times 10^{43}$erg s$^{-1}$ in 2--10 keV.  NGC~4258 (M~106), one of the typical LLAGNs, is a highly inclined (72\\degree, \\cite{{Tully1988}}) SABbc spiral galaxy located at a nearby distance of 7.2~Mpc (\\cite{{Herrn1999}}), where $1''$ corresponds to 35~pc. The presence of an obscured AGN was suggested by the strong polarization of the relatively broad optical emission lines (\\cite{Wilkes1995}), and H$\\alpha$ and X-ray emitting ``anomalous arms\" (Cecil et al. 1995). The observation with ASCA up to 10 keV (Makishima et al. 1994) clearly demonstrated the presence of an LLAGN, absorbed by an equivalent hydrogen column of $N_\\mathrm{H} \\sim 1.5 \\times 10^{23}$~cm$^{-2}$, with an absorption-corrected 2--10~keV luminosity of $\\sim 4 \\times 10^{40}$~erg~s$^{-1}$. The obscured nucleus was later reconfirmed by a Chandra image (\\cite{Wilson2001}; \\cite{Yang2007}). The epoch-making water maser observation (\\cite{Miyoshi1995}; \\cite{Greenhill1995a}; \\cite{Greenhill1995b}) established the AGN mass as $3.6 \\times 10^7$~\\ensuremath{M_\\odot}, and its disk inclination as $\\sim 83\\degree$. The X-ray flux is known to be variable on time scales of hours to years (\\cite{Reynolds2000}; \\cite{Terashima 2002}; \\cite{Fruscione2005}; \\cite{Fiore2001}). According to observations with XMM-Newton (Pietsch and Read 2002) and Chandra (Yang et al. 2007), contributions by other softer X-ray components associated with NGC~4258 can be spectrally separated from the harder nuclear emission. Among them, the soft emission extending out from the nucleus, along ``anomalous arms'' is of particular interest, because of its possible association with the AGN jets (e.g., \\cite{Wilson2001}), but it is beyond the scope of the present paper.  The most outstanding property of the NGC~4258 nucleus is its extremely low luminosity, not only in the absolute sense but also relative to the Eddington value. Indeed, its absorption-corrected 2--10 keV luminosity is only $\\sim 8 \\times 10^{-6}$ of the Eddington value, which is much lower than those of Seyfert galaxies (typically thought to be of the order of $\\sim$ 1\\%). Therefore, its accretion flow could be significantly different from those of Seyferts; e.g., the concept of Advection Dominated Accretion Flow (ADAF; Ichimaru 1977; Narayan and Yi 1994, 1995) may be more applicable to these LLAGNs than to Seyfert galaxies. Nevertheless, broad-band X-ray properties of NGC~4258 are not much different from those of radio-quiet Seyferts, as reported by Fiore et al. (2001) based on a BeppoSAX detection up to $\\sim 70$ keV with a loosely constrained photon index of $\\sim 2.1$. Then, we may need to look for more subtle differences, e.g. in the reflection hump and Fe-K lines. Obviously, Suzaku is most suited for this purpose. In the present paper, we report on a tight upper limit on the reflection component in this LLAGN, and spectral changes on a time scale of $\\sim$ 50 ks, both derived from a Suzaku observation made in 2006 June. ",
        "conclusions": "\\label{section:discussion} %======= 4 =========  %------------ 4.1 --------------------------- \\subsection{Summary of the results} %------------ 4.1 ---------------------------  In 2006 June, we observed NGC~4258  for  $\\sim 100$ ks with  Suzaku. In addition to the soft thermal plasma emission, and a medium-hardness component presumably due to point sources, the absorbed hard emission from its LLAGN was detected with the XIS and HXD-PIN over a $\\sim 2$ keV to $\\sim 40$ keV band altogether. The acquired wide-band spectra of the nucleus have been described successfully in terms of a power-law with $\\Gamma = 1.88$, absorbed by $N_{\\rm H} = 1.1 \\times 10^{23}$ cm$^{-2}$. The absorption-corrected 2--10 keV luminosity of the nucleus, estimated as $\\sim 8 \\times 10^{40}$ erg s$^{-1}$, is typical of this LLAGN (e.g., \\cite{Max1994}, \\cite{Reynolds2000}, \\cite{Pietsch2002}, \\cite{Yang2007}). A narrow Fe-K line at 6.4 keV was detected at  a relatively low EW of $\\sim 50$ eV. The reflection component was insignificant, with an upper limit of  $f_{\\rm refl}< 0.3$ if the reflector is assumed to have $i=60^\\circ$, although it is relaxed to  $f_{\\rm refl}< 2.0$ if  the water-maser configuration of $i=83^\\circ$ is adopted \\citep{Miyoshi1995}.  On a time scale of  $\\sim 100 $ ks, we detected a clear hard X-ray variation by 20\\%, which can be explained by a change in the nuclear component. As the hard X-ray flux increased, the power-law component steepened by $\\Delta \\Gamma \\sim 0.3$ with  marginal evidence for an increase in $N_{\\rm H}$. The varying component was described by a relatively steep power-law with $\\Gamma \\sim 2.8$  %------------ 4.2 ---------------------------------------- \\subsection{Implications of the time-averaged spectrum} %------------ 4.2 -----------------------------------------  Being a prototypical LLAGN, the most outstanding feature of the NGC~4258 nucleus is of no doubt its very low X-ray luminosity, both in the absolute and relative sense. Hence, one of the most intriguing questions is whether it exhibits any X-ray signatures (other than its low luminosity) that are considered specific to LLAGNs. When trying to answer this question, the wide-band spectral coverage of Suzaku, with its good energy resolution,  is considered essential. As part of such attempts, we may  quote X-ray photon indices of the intrinsic nuclear component of  several representative Seyferts, determined  with Suzaku through careful separation of the reflection component. The results include  $\\Gamma \\sim$ 1.6 in NGC~4945 (\\cite{Itoh2008}), $\\Gamma \\sim 1.7$  in NGC~4388 (\\cite{Shirai2008}), $\\Gamma \\sim 1.8$  in MCG 5--23-16 (\\cite{Reeves2007}), and $\\Gamma \\sim$ 1.9  in NGC 3516 (\\cite{Markowitz2008}). Thus, our results on  NGC~4258,  $\\Gamma \\sim 1.9$, does not distinguish this LLAGN from the other Seyferts.  Then, what do we find about circum-nuclear matter distribution in NGC~4258? Generally, this can be probed using reprocessed signals, including in particular the  fluorescence iron lines and the reflection component. The measured EW ($\\sim 50$ eV) of the former component, which agrees with previous measurements from this LLAGN \\citep{Max1994,Fiore2001}, is significantly smaller than those generally found among Seyfert galaxies of both type 1 (typically $\\sim 150$ eV) and type 2 (up to $\\sim 1$ keV; \\citet{Ueda2007}). The latter component, i.e. reflection,  consists of three spectral features; a deep Fe-K edge at $\\sim 7$ keV, a Compton hump above 10 keV, and a  hard continuum approximated by $\\Gamma \\sim -0.3$ (rising with the energy; figure~\\ref{widebandspec}e) below the Fe-K edge. While the last feature is confused in our spectrum with the point-source contribution, the iron edge is clearly very weak (figures~5 and 6). Furthermore, the HXD-PIN data, combined with the XIS flux points, rule out the presence of a strong Compton hump. Assuming a general configuration of $i=60^\\circ$, we have thus  obtained a tight constraint as $f_\\mathrm{refl}< 0.3$, which falls much below those typically found with Seyferts ($f_\\mathrm{refl} \\sim 1-2$; \\cite{DeRosa2008}).  The above results on the Fe-K line and reflection both means that the nucleus of NGC~4258 is devoid of thick materials with  large solid angles. As a result, an inflated torus (for which $i=60^\\circ$ is regarded  as a reasonable approximation) is unlikely to be present around it. Nevertheless,  our direct view to the nucleus is obscured by a thick material with $N_{\\rm H} \\sim 1\\times 10^{23}$ cm$^{-2}$. Then, the obscuring material must be localized to a limited solid angle including our line of sight, rather than forming a thick torus. In fact, our data are consistent with a slab-like matter distribution with  $\\ensuremath{f_\\mathrm{refl}} \\sim 1$, as long as it is viewed from an edge-on direction (section~\\ref{subsec:above1.7}). One likely scenario is that the  nearly edge-on and warped accretion disk, which has been revealed by the water-maser results, causes a grazing obscuration of the nucleus.  After NGC~4945 \\citep{Itoh2008}, the NGC~4258 nucleus thus becomes a second example in which the circum-nulear matter  is inferred to have a thin disk geometry rather than a thick toroidal shape. While such objects  are obviously considered relatively rare in a statistical sense, we must investigate the reason why these objects lack a thick torus that is usually considered to be a common attribute of an AGN. The 2--10 keV luminosity of $8 \\times 10^{40} $ erg s$^{-1}$, measured from NGC~4258, translates to a bolometric luminosity of $L_{\\mathrm{bol}} \\sim 6 \\times 10^{41}$, if we use the argument by \\citet{Marconi2004}, assuming that the multi-wavelength spectra of LLAGNs are not much different from those of typical AGNs. The estimated $L_{\\mathrm{bol}}$ is still only $\\sim 1 \\times10^{-4}$ of the Eddington limit for an object of $3.6 \\times 10^7~M_\\odot$. In contrast, the Eddington-normalized bolometric luminosity of NGC~4945 amounts to 0.3--1.0 (\\cite{Itoh2008}). Therefore, a low {\\it normalized} luminosity is unlikely to provide a common account for the lack of a torus.  When we  consider the absolute luminosity (or mass accretion rate) instead of the normalized one, not only NGC~4258 but also NGC~4945 has a considerably lower value ($7 \\times 10^{42}$ erg s$^{-1}$ in 2--10 keV after correcting for the absorption and Compton scattering; \\cite{Itoh2008}) than typical AGNs (e.g., $10^{43-44}$ erg s$^{-1}$; \\cite{Tueller2008}): the high normalized luminosity of NGC~4945 is simply a results of its low black-hole mass  ($1.4 \\times 10^6~M_\\odot$; \\cite{Greenhill1997}). Thus, a common feature of the two objects is very low accretion rates. We hence speculate that the presence/absence of a thick torus is related to the mass accretion rate rather than to the normalized luminosity. Considering causality, one possible scenario  is that a high accretion rate onto a giant black hole at the galaxy nucleus is realized when the matter distribution has (for some unspecified reasons) an inflated toroidal geometry, whereas the rate remains low when the matter has a thin disk configuration.  An accretion flow under an extremely low rate, like in the present case, may be described by the ADAF model \\citep{Narayan1995}, which assumes the optically-thick accretion disk to be truncated at a large radius. \\citet{Lasota1996} applied the ADAF prediction to the X-ray spectrum of NGC~4258 measured by \\citet{Max1994}, and argued that the innermost radius of the  accretion disk, $r_{\\rm{in}}$, should be larger than $\\sim100~ r_{\\rm{g}}$ in this LLAGN, where $r_{\\rm{g}}$ is the Schwarzschild radius ($= GM/c^2$) Since our spectrum is basically similar to that measured by \\citet{Max1994}, the result by \\citet{Lasota1996}  may also be applicable to our case. Such a condition is consistent with our data, because they also allow a solution with  $i=82^\\circ$ and $\\ensuremath{f_\\mathrm{refl}} \\sim 0$. Incidentally, \\citet{Max2008} measured $\\ensuremath{f_\\mathrm{refl}} \\sim 0.4$ from Cygnus X-1, together with $r_{\\rm{in}} \\sim 15~r_{\\rm{g}}$. Therefore, the concept of ADAF may be more applicable to NGC~4258 than to Cygnus X-1.  %------------ 4.2 ------------------ \\subsection{The time variations} %------------ 4.2 ------------------  The time variation we detected, on a time scale of  $\\sim 100$ ks, is one of the most rapid changes ever observed from NGC~4258. We discovered  that  the power-law spectrum softens by $\\Delta \\Gamma \\sim 0.3$ when the nucleus brightens up by 20\\%. This result is apparently reminiscent of the positive correlation between $\\Gamma$ and flux, measured repeatedly from other AGNs; e.g., from NGC~4151 with EXOSAT and Ginga \\citep{Warwick1989,Yaqoob1992}, and from NGC~4051 with ASCA (\\cite{Guainazzi1996}). Such  positive $\\Gamma$ vs. flux correlations allow two alternative interpretations. One is to assume that this is an intrinsic property of the nuclear emission, possibly related to cooling of  hot electron ``coronae\" that produce  hard X-rays via thermal Comptonization of some soft seed photons (e.g., \\cite{ST1979}). The other is to consider the relation as an artifact, produced by a superposition of a  power-law component which varies without changing its slope, and a constant (and harder) reflection component.  As to black hole binaries in so-called  Low-Hard state, including Cygnus X-1 in particular, Suzaku observations \\citep{Max2008} revealed that the former is actually occurring at least on time scales longer than 1 s. (The latter mechanism does not work on this time scale, since the reflection catches up with the intrinsic nucleus variation). In the case of AGNs, in contrast, the latter is likely to be dominant, because their ``difference\" spectra can be generally expressed by a single power-law, of which the photon index is close to what is found with the time-averaged spectrum after separating the reflection component. Such examples include NGC~4945 (\\cite{Itoh2008}) and MCG 5--23-16 (\\cite{Reeves2007}). Returning to NGC~4258, our results prefer the former interpretation, because the reflection component is basically insignificant in this object. Then, the mechanism responsible for the variation of this LLAGN is suggested to be closer to that of black-hole binaries in the Low-Hard state than to those of Seyfert galaxies.  An intriguing result is the weak variation in the 1--2 keV band, apparently anti-correlated with that in the absorbed power-law. One possible interpretation is to invoke a variation of a small number of luminous point sources, as suggested by the analysis conducted in section \\ref{subsec:hl_ave}. However, the implied luminosity change amounts to $\\sim 5 \\times 10^{38}$ erg s$^{-1}$, which would be too large to be explained by such a mechanism. Furthermore, we would have to invoke a chance coincidence to explain the apparent anti-correlation between the 1--2 keV and hard-band variations. Therefore, it is more natural to consider that the soft-band signals are contributed not only by point sources, but also by the nucleus at some level, and the variation therein is attributable to the nucleus contribution. Since the absorbed power-law component falls far below ($\\sim 1/30$) the overall signal intensity at 1.5 keV (figure~\\ref{ape_eeuf}), this interpretation requires that a small  (e.g., $\\sim 1\\%$) fraction of the intrinsic emission reaches us without being absorbed. Below, we construct a possible explanation based on this idea.  Let us  recall  that  the hard X-rays from Seyfert-like objtects are likely to be emitted from a hot corona with a relatively large scale height, via Compton scattering by thermal electrons \\citep{Max2008}. If such a corona is viewed from sideways, its line-integrated Compton emissivity is expected to decrease rapidly as a height away from the accretion plane (like in the limb regions of solar coronae). Then, the highest part of the corona, which carries only a small fraction of the overall emissivity (but geometrically rather extended), could be rising above the grazing absorber and hence directly visible. If  the height of this corona is  time variable, the hard X-ray intensity will increase when the corona becomes less tall, because the Compton optical depth of the corona will increase due to compression. At the same time, the directly-visible  fraction of the coronal top region will decrease, and will reduce the unabsorbed X-ray flux. Furthermore, the absorber may have a gradient in the column density, in such a way that it decreases away from the accretion plane. Then, a more compressed corona will sample on average lower heights of the absorber, leading to a slight enhancement  in  $N_{\\rm H}$ as suggested by the data (table \\ref{hl_tbl}).  As a consistency check of the above scenario, a typical scale of such a corona is estimated as several tens $r_{\\rm g}$ if an analogy to Cygnus X-1 is adopted \\citep{Max2008}. In the present case of NGC~4258, such a corona can vary just on $\\sim 50$ ks, assuming a Keplerian time scale at such radii. Furthermore, the Comptonizing corona of Cygnus X-1 was revealed to be highly inhomogeneous \\citep{Max2008}, and  its hard X-ray intensity was suggested to increase when the corona becomes ``less porous\". This is consistent with the height variation we invoked above, because an inhomogeneous corona will naturally become less porous when it becomes vertically shorter. For a final speculation, a corona may become vertically shorter and less porous (thus leading to an increased hard X-ray intensity), as its internal magnetic pressure is released by magnetic reconnection.  The authors would like to express their thanks to the Suzaku team members. Our work was supported by Grant-in-Aid for JSPS Fellow."
    },
    "0812/0812.1015_arXiv.txt": {
        "abstract": "Probably not, because there are lots of manifestly unanthropic ways of producing entropy.  We demonstrate that the Causal Entropic Principle (CEP), as a replacement for the anthropic principle to explain the properties of the observed universe, suffers from many of the same problems of adopting myopic assumptions in order to predict that various fundamental parameters take approximately the observed values. In particular, we demonstrate that four mechanisms -- black hole production, black hole  decay,  phase transitions, and dark matter annihilations or decays -- will manifestly change the conclusions of the CEP to predict that we should live in a universe quite different than the one in which we find ourselves. ",
        "introduction": "Recently, it was proposed \\cite{Bousso:2007kq} that the correct anthropic measure for the probability of observing a given universe is tied to the amount of entropy production in that universe. According to this proposal, known as the Causal Entropic Principle (CEP), life depends on the production of entropy and we are more likely to live in a universe in which the entropy production is as large as possible. The entropy should be calculated within a causally connected 4-volume known as the causal diamond \\cite{Bousso:2006ev}. The authors of \\cite{Bousso:2007kq} determined that the most important sites of  entropy production in our universe are the radiating dust grains heated by stars. Life should therefore cluster around stars -- a postdiction consistent with the lone observed nexus of living organisms -- and the parameters of our universe should maximize the number density of stars and associated dust clouds. Keeping all other parameters fixed, the value of the cosmological constant that maximizes the dust entropy within the causal diamond is in reasonable agreement with the observed low but non-zero value.  However, it is not clear that interstellar dust is indeed required for life, e.g., if we strip off all the dust from the solar system then the entropy production would be reduced but life on earth would be virtually unaffected. Life does seem to need a steady gentle (over several billion years) source of free energy, which intuitively is associated with semi-adiabatic processes that produce little entropy. Violent energy release and concomitant entropy production is unlikely to be conducive to the formation of life which depends on complicated fragile molecules.  Furthermore, if we do assume ourselves to be typical representatives of observers, it seems that baryons are essential to the formation of life. The CEP does not accommodate this fact. The baryon to entropy ratio (B/S) can be very small, (e.g., in our universe B/S$\\sim 10^{-9}$) or even vanishingly small if the Sakarov conditions  (baryon number violation, T violation and CP violation) are not satisfied. Such a baryon poor universe is certainly not a hospitable environment for life. Considering the ~$10^{500}$ possible vacua in the string landscape, there are surely many examples of universes with multiple first order phase transitions and lots of entropy production, but without enough baryons for nucleosynthesis to proceed and no life. Such large entropy producing universes would be a detriment to the formation of life. By assuming ourselves to be typical representatives of observers we essentially limit the above discussion to carbon-based forms of life similar to those we find on earth. We are adopting here a conservative approach: relaxing the assumption of carbon-based life opens up an additional volume of parameter space where life can develop and the entropy measure is simply unknown.   Three sources of entropy which were either not accounted for or were discounted in \\cite{Bousso:2007kq} are the entropy associated with black holes, the entropy produced at reheating after a phase transition, and the entropy produced by out of equilibrium decays or annihilations of dark matter particles, presumably because none seems to be directly related to life. One could argue that as long as the entropy of the black hole is behind the horizon it is irrelevant to life and to the CEP. However, if there is a way to ``mine\" this entropy, or if it is radiated out of the black hole, then in principle this is no different than other sources of entropy. As we explicitly show below, this source of entropy is by far the most significant source, and changes the CEP predicted value of the cosmological constant. The entropy of phase transitions, reheating, decays and annihilations is what is referred to in \\cite{Bousso:2007kq} as ``matter sector entropy\", which is not different in nature than say the heated dust entropy. We present a calculation showing that if we optimize the entropy production by reheating then inflation will not result in the required number of e-folds of growth of the universe. We also show that dark matter annihilation can lead to far greater entropy production than currently results from stellar burning.   The paper is organized as follows, the entropy of black holes and its accessibility are considered in section \\ref{bh}. In section \\ref{inf} we maximize entropy produced after a phase transition, and show it is inconsistent with the observed universe. Out of equilibrium processes are discussed in section \\ref{ooe}, and we conclude in section \\ref{disc}. ",
        "conclusions": "\\label{disc} In our universe, life does not seem to be associated with either the times or locations of the greatest entropy or the greatest entropy production -- formation of, or accretion onto, massive black holes at the centers of galaxies, post-inflationary (p)reheating, or synchrotron radiation from  charged particles produced in late-time dark matter annihilation and decay.  Furthermore, we can look at our own universe and ask how we could increase the rate of entropy production. A higher  dark matter density and higher amplitude of primordial fluctuations would have resulted in the formation of denser  more massive galaxies and clusters. Energy loss by dark matter annihilation (or decay) products in those structures would have resulted in greater entropy production. Denser galaxies would also have a higher rate of accretion onto their central black hole, resulting in both larger black holes, and more entropy released into the environment during the accretion. These denser galaxies would have occupied richer, denser clusters  resulting in more galaxy mergers, and thus higher final black hole masses.    More massive black holes implies a corresponding increase in the black hole entropy -- entropy which is ultimately released through Hawking radiation of ordinary photons (and other standard model particles) into a single causal diamond. These denser galaxies and clusters could also have accommodated a higher value of the cosmological constant than we currently see without being disrupted.  Meanwhile, if we keep the dark matter density and the amplitude of primordial fluctuations fixed at their values in our universe,  a lower value of the cosmological constant would have allowed the formation  of larger gravitationally bound systems -- super clusters, super-duper-clusters and so-on.  These more massive bound systems would eventually  have had more massive black holes at their centers, eventually resulting in more entropy release within a single causal diamond than in our universe.  Thus, life does not seem to have demanded that the universe optimize the production of entropy.  Most of the entropy seems not to be particularly closely associated with the assumed concentration of life around  stars. A modification of the CEP to count only standard model entropy could no doubt be formulated, but would not evade the problem that dark matter annihilations into standard model particles results in enormous entropy production via synchrotron radiation. Moreover,  as with much anthropic reasoning it would seem to be another example of designing the questions to give the already known answers."
    },
    "0812/0812.1223_arXiv.txt": {
        "abstract": "We present a study of the dynamics and magnetic activity of M dwarfs using the largest spectroscopic sample of low-mass stars ever assembled. The age at which strong surface magnetic activity (as traced by H$\\alpha$) ceases in M dwarfs has been inferred to have a strong dependence on mass (spectral type, surface temperature) and explains previous results showing a large increase in the fraction of active stars at later spectral types. Using spectral observations of more than 40000 M dwarfs from the Sloan Digital Sky Survey, we show that the fraction of active stars decreases as a function of vertical distance from the Galactic plane (a statistical proxy for age), and that the magnitude of this decrease changes significantly for different M spectral types. Adopting a simple dynamical model for thin disk vertical heating, we assign an age for the activity decline at each spectral type, and thus determine the activity lifetimes for M dwarfs. In addition, we derive a statistical age-activity relation for each spectral type using the dynamical model, the vertical distance from the Plane and the H$\\alpha$ emission line luminosity of each star (the latter of which also decreases with vertical height above the Galactic plane). ",
        "introduction": "M dwarfs are the smallest and least luminous stars on the main sequence, yet they are the most numerous of all stellar constituents and have lifetimes longer than the current age of the Universe. Determining the age of an individual main sequence, field M dwarf is quite challenging.  Stellar clusters would seem to provide the ideal environments for calibrating the ages of M dwarfs (because the age of the cluster is determined). However, the large distances to clusters older than a few Gyr (and resulting faintness of cluster M dwarfs), preclude any detailed observations of the intermediate or oldest cluster resident M dwarfs. Age estimates for field M dwarfs must therefore rely on a number of techniques that piggyback on other age dating methods.  Although several methods exist for estimating the ages of M dwarfs (e.g. contribution by S. Catalan), we focus this contribution on what the magnetic activity of M dwarfs can tell us about their age.  M dwarfs are host to intense magnetics dynamos that give rise to chromospheric and coronal heating, producing emission from the X-ray to the radio. This magnetic activity has long been thought to be tied to the age of the host star.  Almost 40 years ago, Wilson \\& Woolley (1970) found a link between magnetic activity (as traced by the Ca II H \\& K emission lines) and the orbital elements (namely the inclination and eccentricity) of more than 300 late-type dwarfs; stars in near-circular orbits with small inclinations had the strongest activity.  They concluded that as stars age, their orbits get more inclined and more eccentric (due to dynamical interactions), and thus, magnetic activity must also decline with age.  Subsequent studies over the following decades have found similar connections between age and activity in low-mass stars (Wielen 1977; Giampapa \\& Liebert 1986; Soderblom et al. 1991; Hawley et al. 1996, Hawley et al. 1999, Hawley et al. 2000).  Although activity has been observed for decades, the exact mechanism that gives rise to the chromospheric heating is still not well-understood.  In the Sun, magnetic field generation and subsequent heating is closely tied to the Sun's rotation.  From helioseismology, we know that there is a rotational boundary between the inner solid-body rotating radiative zone and the outer differentially rotating convective zone (Parker 1993; Ossendrijver 2003; Thompson et al. 2003).  This boundary, dubbed the tachocline, creates a rotational shear that likely allows magnetic fields to be generated, preserved and eventually rise to the surface where they emerge as magnetic loops.  These loops bring heat to the Sun's chromosphere and corona, driving both large stellar flares, as well as lower-level quiescent (or persistent) magnetic activity.  The faster a star (with a tachocline) rotates, the stronger its magnetic heating and surface activity. Therefore, the angular momentum evolution of solar-type stars should play an important role in determining the magnitude of the observed activity.  Angular momentum loss from magnetized winds has been shown to slow rotation in solar-type stars; as a result magnetic activity decreases. Skumanich (1972) found that both activity and rotation decrease over time as a power law ($t^{-0.5}$).  Subsequent studies confirmed the Skumanich results and empirically demonstrated a strong link between age, rotation and activity in solar-type stars (Barry 1988; Soderblom et al. 1991; Pizzolato et al. 2003; Mamajek \\& Hillenbrand 2008; see contribution by E. Mamajek).  There is strong evidence that the rotation-activity (and presumably age) relation extends from stars more massive than the Sun to smaller dwarfs (Pizzolato et al. 2003; Mohanty \\& Basri 2003; Kiraga \\& Stepien 2007). Therefore, rotation derived ages (using gyrochronology) may provide a useful independent estimate of age for M dwarfs when large enough calibration samples can be acquired (see contributions by S. Meibom, J. Irwin, and S. Barnes).  However, at a spectral type of $\\sim$M3 (0.35 M$_{\\odot}$; Reid \\& Hawley 2005; Chabrier \\& Baraffe 1997), stars become fully convective and the tachocline presumably disappears.  Despite this changes, magnetic activity persists in late-type M dwarfs; the fraction of active M dwarfs peaks around a spectral type of M7 before decreasing into the brown dwarf regime (Hawley et al. 1996; Gizis et al. 2000; West et al. 2004).  It is unclear if the rotation-age-activity relation extends to the fully convective M dwarfs.  A few empirical studies have uncovered evidence that activity and rotation might be linked in late-type M dwarfs (Delfosse et al. 1998; Mohanty \\& Basri 2003; Reiners \\& Basri 2007).  In addition, recent simulations of fully convective stars find that rotation may play a role in magnetic dynamo generation (Dobler et al. 2006; Browning 2008).  However, a recent study found detectable rotation in a few inactive M7 dwarfs, indicating that the situation may in fact be more complicated (West \\& Basri 2009).  Irrespective of rotation, many studies have found evidence that the age-activity relation extends into the M dwarf regime.  Eggen (199) observed a Skumanich-type decrease in activity as a function of age. Larger samples of M dwarfs have added further evidence that magnetic activity decreases with age (Fleming et al. 1995; Gizis et al. 2002). There is also evidence that M dwarfs may have finite active lifetimes. Stauffer et al. (1994) suggested that activity may not be present in the most massive M dwarfs in the Pleiades.  Hawley et al. (2000) confirmed the Stauffer et al. (1994) claim by observing a a sample of clusters that spanned several Gyr in age.  They were able to calibrate the lifetimes for early-type M dwarfs by observing the color at which activity (traced by H$\\alpha$ emission) was no longer present. Because of the small sample size, the derived Hawley et al. (2000) activity ages are lower limits at a given color or spectral type. As mentioned above, the Hawley et al. (2000) study was limited to younger, nearby clusters due to the intrinsic faintness of M dwarfs; thus, it could only probe ages of a few Gyr and spectral types as late as $\\sim$M3. Larger samples of M dwarfs are required to statistically derive age-activity relations that extend both in the ages they probe and the range of spectral types they cover.  In this contribution, we review recent studies that have used large spectroscopic samples to investigate the relationship between age and activity in M dwarfs.  In addition, we derive a statistical H$\\alpha$ age-activity relation for M2-M7 dwarfs. ",
        "conclusions": "Large astronomical surveys have produced M dwarf samples of unprecedented size.  We have shown that using the statistical foothold of the largest low-mass spectroscopic sample ever assembled, a strong tie between magnetic activity (as traced by H$\\alpha$) and age has been established and that M dwarfs have finite active lifetimes. These lifetimes are a strong function of spectral type, dramatically increasing as the M dwarf interiors become fully convective.  We derived age-activity relations for M2-M7 dwarfs based on the average activity as a function of height above (or below) the Galactic plane. Future studies will extend these relations to other spectral types. In addition, these relations can be confirmed and calibrated using both white dwarf cooling ages in wide binary systems that have both a white dwarf and an M dwarf (see contributions by S. Catalan, M. Salaris, and J. Kalirai; Silvestri et al. 2006), and deep cluster observations of M dwarfs (when such observations become possible).     \\nocite{*}"
    },
    "0812/0812.4892_arXiv.txt": {
        "abstract": "A new approach to extraction of quantum vacuum energy, in the context of the accelerated expansion, is proposed, and it is shown that experimentally realistic orders of values can be derived. The idea has been implemented in the framework of the Friedmann--Lema\\^{\\i}tre--Robertson--Walker geometry in the language of the effective action in the relativistic formalism of Schwinger's proper time and Seeley--DeWitt's heat kernel expansion. ",
        "introduction": "The following three, well-known problems of modern physics and cosmology, accelerated expansion of the Universe (Riess {\\em et al.\\/} \\cite{Riess1998}; Perlmutter {\\em et al.\\/} \\cite{Perlm1999}), very small but non-vanishing cosmological constant or dark energy (Weinberg \\cite{Wein1989}), (Carroll \\cite{Carr2001}; Padmanabhan (\\cite{Padm2003}, \\cite{Padm2006})), and theoretically extraordinarily huge quantum vacuum energy density (Zel`dovich \\cite{Zeld1967}), (Volovik (\\cite{Vol2005}, \\cite{Vol2006})) can be treated as mutually related or as independent problems. An old approach to the issue of the cosmological constant $\\Lambda$ utilizes quantum vacuum energy as a solution of this issue, but unfortunately it does not work properly. Namely, it appears that directly calculated, Casimir-like value of quantum vacuum energy is more than one hundred orders greater than expected. Such a huge value of quantum vacuum energy is a serious theoretical problem in itself. Lowering the UV cutoff scale down from the planckian to the supersymmetric one is a symbolic improvement (roughly, it cuts the order by two (Weinberg \\cite{Wein1989})). A more radical reduction of the cutoff could cure the situation but it would create new problems. Sometimes, it is claimed that vacuum energy for one or another reason does not influence gravitational field.   In this paper, following ideas presented in (Broda {\\em et al.\\/} \\cite{Broda2008}), we show in what sense  quantum vacuum energy influences gravitational field, and in what sense it does not. Actually, we propose a reasonable derivation of a contribution of quantum vacuum energy which influences gravitational field, and which assumes an experimentally reasonable value. ",
        "conclusions": "In the framework of standard quantum field theory, without any additional more or less exotic assumptions we are able to derive an experimentally reasonable result \\eqref{eq:Final Density 3}. This numeric value corresponds to only a single mode. Therefore in the real world it should be multiplied by a small natural number."
    },
    "0812/0812.0363_arXiv.txt": {
        "abstract": "{Recent progresses and constraints brought by helio\\- and astero\\-seismology call for a better description of stellar interiors and an accurate description of rotation-driven mechanisms in stars.}% {We present a detailed analysis of the main physical processes responsible for the transport of angular momentum and chemical species in the radiative regions of rotating stars. We focus on cases where meridional circulation and shear-induced turbulence only are included in the simulations (i.e., no internal gravity waves nor magnetic fields). We put special emphasis on analyzing the angular momentum transport loop and on identifying the contribution of each physical process involved.}% {We develop a variety of diagnostic tools designed to help disentangle the role of the various transport mechanisms. Our analysis is based on a 2-D representation of the secular hydrodynamics, which is treated using expansions in spherical harmonics. By taking appropriate horizontal averages, the problem reduces to one dimension, making it implementable in a 1D stellar evolution code while preserving the advective character of angular momentum transport. We present a full reconstruction of the meridional circulation and of the associated fluctuations of temperature and mean molecular weight along with diagnosis for the transport of angular momentum, heat and chemicals. In the present paper these tools are used to validate the analysis of two main sequence stellar models of 1.5 and 20~\\Ms{} for which the hydrodynamics has been previously extensively studied in the literature.}% {We obtain a clear visualization and a precise estimation of the different terms entering the angular momentum and heat transport equations in radiative zones of rotating stars. This enables us to corroborate the main results obtained over the past decade by Zahn, Maeder, and collaborators concerning the secular hydrodynamics of such objects. We focus on the meridional circulation driven by angular momentum losses and structural readjustements. We confirm quantitatively for the first time through detailed computations and separation of the various components that the advection of entropy by this circulation is very well balanced by the barotropic effects and the thermal relaxation during most of the main sequence evolution. This enables us to derive simplifications for the thermal relaxation on this phase. The meridional currents in turn advect heat and generate temperature fluctuations that induce differential rotation through thermal wind thus closing the transport loop. We plan to make use of our refined diagnosis tools in forthcoming studies of secular (magneto-)hydrodynamics of stars at various evolutionary stages.}% {} ",
        "introduction": " ",
        "conclusions": "The paper presents a complete set of diagnostic tools that provide a comprehensive and coherent understanding of the secular hydrodynamical transport processes operating in the radiative zones of rotating stars. The framework is that of ``type I rotational mixing'', where angular momentum and nuclides are transported by the same mechanisms, namely large-scale meridional circulation and shear-induced turbulence.  In order to validate our new approach, the first analysis presented here is performed on two well-studied cases of a low-mass (1.5~\\Ms{}) \\citep[see][]{PA03} and a massive (20~\\Ms{}) star \\citep{MM00}. It allows us to disentangle what are the main processes at the origin of meridional circulation and driving angular momentum, heat and chemicals transports in the stellar interior and to confirm results previously established in the literature. In particular, the angular momentum loss, either by radiation-driven wind (the massive star) or by magnetic braking (the low-mass star) combined with structural readjustments are the sources responsible for the generation of meridional circulation (Fig.~\\ref{fig:loopth}). The direction of the flow is such that it transports angular momentum outward to the surface, except in the deep interior of the 1.5~\\Ms{} star where it can reverse, producing a steeper angular velocity gradient.  This circulation advects heat, and thus generates latitude dependent temperature fluctuations; these tend to be damped out by radiative diffusion, until they establish a subtle balance between advection and diffusion, that allows them to induce differential rotation through the baroclinic torque (Eq. \\ref{ThermalWind}). This differential rotation is shear-unstable, and generates turbulence that participates to the transport of angular momentum, thus closing the loop (Fig.~\\ref{fig:loopth}). This picture thus differs drastically from the classical Eddington-Sweet circulation, where the circulation is deemed to originate from the thermal imbalance due to the centrifugal force. If turbulent transport and extraction of angular momentum were absent, the circulation would die altogether, as was pointed out by \\citet{busse82}.  Thanks to the diagnostic tools we have developed, the present investigation helped to clarify a number of other points. We confirm that it is mainly the meridional circulation that transports angular momentum in the low-mass star \\citep[see also][]{PA03}; it is thus crucial to describe that transport as an advective process, as in the approach developed by Zahn, Maeder, and collaborators. In the massive star, where the redistribution of angular momentum due to the star's expansion plays a much more important role than its extraction by the radiative wind, advection and diffusion make comparable contributions, though in different regions: the advective transport dominates in the inner layers of the star, whereas the diffusive transport is the prime actor in the upper part and, moreover, works there in the opposite direction.  \\begin{figure}[tbp] \\begin{center} \\includegraphics[width=0.45\\textwidth]{fig18} \\caption{Rotational mixing of type I in the radiative zone of a rotating star, where the transport of angular momentum is achieved by meridional circulation and turbulent diffusion.} \\label{fig:loopth} \\end{center} \\end{figure}  For the first time, we confirm through detailed computation and separation of the different components, that the advection of entropy by meridional circulation is almost exactly balanced by the thermal relaxation (barotropic and thermic terms) during most of the main sequence. This was generally assumed prior to the \\citet{MZ98} work integrating the time-dependency of the temperature fluctuations, but not actually verified through the computation. Grounding our detailed analysis on advanced graphical tools, we are also able to propose a simplified expression for the thermal relaxation on the main sequence.  On the other hand, the transport of chemicals is shown to be dominated by the shear-induced turbulence while the star evolves on the main sequence. This is particularly clear for a massive star such as the 20~\\Ms{} star.  In the near future, we plan to apply the same tools to Type II rotational mixing, where we shall include the transport of angular momentum by internal gravity waves and magnetic stresses and to advanced phases of stellar evolution."
    },
    "0812/0812.0649_arXiv.txt": {
        "abstract": "This paper describes the Seventh Data Release of the Sloan Digital Sky Survey (SDSS), marking the completion of the original goals of the SDSS and the end of the phase known as SDSS-II.  It includes 11663 deg$^2$ of imaging data, with most of the $\\sim 2000$ deg$^2$ increment over the previous data release lying in regions of low Galactic latitude. The catalog contains five-band photometry for 357 million distinct objects.  The survey also includes repeat photometry on a 120$^\\circ$ long, 2.5$^\\circ$ wide stripe along the Celestial Equator in the Southern Galactic Cap, with some regions covered by as many as 90 individual imaging runs.  We include a coaddition of the best of these data, going roughly two magnitudes fainter than the main survey over 250 deg$^2$.  The survey has completed spectroscopy over 9380 deg$^2$; the spectroscopy is now complete over a large contiguous area of the Northern Galactic Cap, closing the gap that was present in previous data releases. There are over 1.6 million spectra in total, including 930,000 galaxies, 120,000 quasars, and 460,000 stars.  The data release includes improved stellar photometry at low Galactic latitude. The astrometry has all been recalibrated with the second version of the USNO CCD Astrograph Catalog (UCAC-2), reducing the rms statistical errors at the bright end to 45 milli-arcseconds per coordinate.  We further quantify a systematic error in bright galaxy photometry due to poor sky determination; this problem is less severe than previously reported for the majority of galaxies.  Finally, we describe a series of improvements to the spectroscopic reductions, including better flat-fielding and improved wavelength calibration at the blue end, better processing of objects with extremely strong narrow emission lines, and an improved determination of stellar metallicities. ",
        "introduction": "\\label{sec:introduction} The Sloan Digital Sky Survey (SDSS; York \\etal\\ 2000) saw first light a decade ago, with the goals of obtaining CCD imaging in five broad bands over 10,000 deg$^2$ of high-latitude sky, and spectroscopy of a million galaxies and one hundred thousand quasars over this same region.  With this, its seventh public data release, these goals have been realized.  The survey facilities have also been used to carry out a comprehensive imaging and spectroscopic survey to explore the structure, composition, and kinematics of the Milky Way Galaxy (Sloan Extension for Galactic Understanding and Exploration; SEGUE; Yanny \\etal\\ 2009), and a repeat imaging survey that has discovered almost 500 spectroscopically confirmed Type Ia supernovae with superb light curves (Frieman \\etal\\ 2008; Holtzman \\etal\\ 2008).  The SDSS uses a dedicated wide-field 2.5m telescope (Gunn \\etal\\ 2006) located at Apache Point Observatory (APO) near Sacramento Peak in Southern New Mexico.  The telescope uses two instruments.  The first is a wide-field imager (Gunn \\etal\\ 1998) with 24 $2048 \\times 2048$ CCDs on the focal plane with $0.396^{\\prime\\prime}$ pixels, that covers the sky in drift scan mode in five filters in the order $riuzg$ (Fukugita \\etal\\ 1996).  The imaging is done with the telescope tracking great circles at the sidereal rate; the effective exposure time per filter is 54.1 seconds, and 18.75 deg$^2$ are imaged per hour in each of the five filters.  The images are mostly taken under good seeing conditions (the median is about $1.4''$ in $r$) on moonless photometric nights (Hogg \\etal\\ 2001); the exceptions are a series of repeat scans of the Celestial Equator in the Fall for a supernova search (Frieman \\etal\\ 2008), as is described in more detail in \\S~\\ref{sec:stripe82}.  The 95\\% completeness limits of the images are $u,g,r,i,z = 22.0, 22.2, 22.2, 21.3, 20.5$, respectively (Abazajian \\etal\\ 2004), although these values depend as expected on seeing and sky brightness.  The images are processed through a series of pipelines that determine an astrometric calibration (Pier \\etal\\ 2003) and detect and measure the brightnesses, positions and shapes of objects (Lupton \\etal\\ 2001; Stoughton \\etal\\ 2002).  The astrometry is good to 45 milli-arcseconds (mas) rms per coordinate at the bright end, as described in more detail in \\S~\\ref{sec:astrometry}.  The photometry is calibrated to an AB system (Oke \\& Gunn 1983), and the zeropoints of the system are known to 1--2\\% (Abazajian \\etal\\ 2004).  The photometric calibration is done in two ways, by tying to photometric standard stars (Smith \\etal\\ 2002) measured by a separate 0.5-m telescope on the site (Tucker \\etal\\ 2006; Ivezi\\'c \\etal\\ 2004), and by using the overlap between adjacent imaging runs to tie the photometry of all the imaging observations together, in a process called ubercalibration (Padmanabhan \\etal\\ 2008).  Results of both processes are made available; with this data release, the ubercalibration results, which are uncertain at the $\\sim 1\\%$ level in $griz$ and 2\\% in $u$, are now the default photometry made available in the data release described in this paper.  The photometric catalogs of detected objects are used to identify objects for spectroscopy with the second of the instruments on the telescope: a 640-fiber-fed pair of multi-object double spectrographs, giving coverage from 3800\\AA\\ to 9200\\AA\\ at a resolution of $\\lambda/\\Delta \\lambda \\simeq 2000$.  The objects chosen for spectroscopic follow-up are selected based on photometry corrected for Galactic extinction following Schlegel, Finkbeiner, \\& Davis (1998; hereafter SFD), and include: \\begin{itemize} \\item A sample of galaxies complete to a Petrosian (1976) magnitude limit of $r=17.77$ (Strauss \\etal\\ 2002); \\item Two deeper samples of luminous red ellipticals selected in color-magnitude space to $r = 19.2$ and $r=19.5$, respectively, which produce an approximately volume-limited sample to $z = 0.38$, and a flux-limited sample extending to $z=0.55$, respectively (Eisenstein \\etal\\ 2001); \\item Flux limited samples of quasar candidates, selected by their non-stellar colors or FIRST (Becker \\etal\\ 1995) radio emission to $i=19.1$ in regions of color space characteristic of $z < 3$ quasars, and to $i=20.2$ for quasars with $3 < z < 5.5$ (Richards \\etal\\ 2002); \\item A variety of ancillary samples, including optical counterparts to ROSAT-detected X-ray sources (Anderson \\etal\\ 2007); \\item Stars for spectrophotometric calibration and telluric absorption correction, as well as regions of blank sky for accurate sky subtraction; \\item A variety of categories of stellar targets with a series of color and magnitude cuts for measurements of radial velocity, metallicity, surface temperature, and Galactic structure as part of SEGUE (Yanny \\etal\\ 2009). \\end{itemize} These targets are arranged on {\\em tiles} of radius $1.49^\\circ$, with centers chosen to maximize the number of targeted objects (Blanton \\etal\\ 2003).  Each tile contains 640 objects, and forms the template for an aluminum spectroscopic plate, in which holes are drilled to hold optical fibers that feed the spectrographs. Spectroscopic exposures are 15 minutes long, and three or more are taken for a given plate to reach pre-defined requirements of signal-to-noise ratio (S/N), namely $(S/N)^2 > 15$ per 1.5\\AA\\ pixel for stellar objects of fiber magnitude $g = 20.2, r=20.25$ and $i=19.9$.  For the SEGUE faint plates, the exposures are considerably deeper, and typically consist of eight 15-minute exposures, giving $(S/N)^2 \\sim 100$ at the same depth (Yanny \\etal\\ 2009).  Spectra are extracted and calibrated in wavelength and flux.  The typical S/N of a galaxy near the main sample flux limit is 10 per pixel.  The broad-band spectrophotometric calibration is accurate to 4\\% rms for point sources (Adelman-McCarthy \\etal\\ 2008), and the wavelength calibration is good to 2 \\kms.  The spectra are classified and redshifts determined using a pair of pipelines (Stoughton \\etal\\ 2002; Subbarao \\etal\\ 2002), which give consistent results 98\\% of the time; the discrepant objects tend to be of very low S/N, or very unusual objects, such as extreme BALs, superposed sources, and so on. The vast majority of the spectra of galaxies and quasars yield reliable redshifts; the failure rate is of order 1\\% for galaxies, and slightly larger for quasars.  The stellar targets are further processed by a separate pipeline (Lee \\etal\\ 2008a,b; Allende Prieto \\etal\\ 2008a) which determines surface temperatures, metallicities, and gravities.  The resulting catalogs are stored and distributed via a database accessible on the web (the Catalog Archive Server, CAS\\footnote{\\tt http://cas.sdss.org/astro}; Thakar \\etal\\ 2008), and the images and flat files are available in bulk through the Data Archive Server (DAS)\\footnote{\\tt http://das.sdss.org}.  The SDSS saw first light in May 1998, and started routine operations in April 2000.  It was originally funded for five years of operations, but had not completed its core goals of imaging and spectroscopy of a large contiguous area of the Northern Galactic Cap by 2005.  The survey was extended for an additional three years, with the additional goals of the SEGUE and the supernova surveys mentioned above.  The extended program is known as SDSS-II, and the component of SDSS-II that represents the completion of SDSS-I is known as the Legacy Survey.  SDSS-II observations were completed in July 2008.  The SDSS data have been made public in a series of yearly data releases (Stoughton \\etal\\ 2002; Abazajian \\etal\\ 2003, 2004, 2005, 2006; Adelman-McCarthy \\etal\\ 2007, 2008; hereafter the EDR, DR1, DR2, DR3, DR4, DR5, and DR6 papers, respectively).  The most recent of these papers described the Sixth Data Release (DR6), which included data taken through July 2006.  The present paper describes the Seventh Data Release (DR7), including data taken through the end of SDSS-II in 2008 July, and thus represents two additional years of data.  The data releases are cumulative; DR7 includes all data included in the previous releases as well.  In \\S~\\ref{sec:sky_coverage}, we describe the footprint of this survey; most importantly, we have completed our goals of: \\begin{itemize} \\item contiguous imaging and spectroscopy over 7500 deg$^2$ of the Northern Galactic Cap (the Legacy survey); \\item imaging and spectroscopy of stellar sources over an additional 3500 deg$^2$ at  lower Galactic latitudes to study the structure of the Milky Way, and \\item repeat imaging of $>250$ deg$^2$ on the Celestial Equator in the Fall months to discover Type Ia supernovae with $0.1 < z < 0.4$. \\end{itemize}  In \\S~\\ref{sec:runsDB}, we describe the repeat scans on the Celestial Equator, including a co-addition of the images to reach about two magnitudes deeper than the main survey.  In \\S~\\ref{sec:imaging}, we present improvements in the processing of the imaging data, including improved stellar photometry at low Galactic latitudes, an astrometric recalibration, and improvements in our photometric redshift algorithms for galaxies.  The DR6 paper described a problem with the photometry of bright galaxies; we explore this further in \\S~\\ref{sec:skysub}. In \\S~\\ref{sec:spectroscopy}, we discuss improvements in the spectroscopic processing of the data.  The DR6 paper described improvements in the wavelength and spectrophotometric calibration; we have implemented further refinements which are important in the determination of accurate stellar parameters from the spectra.  We conclude in \\S~\\ref{sec:conclusions} with a discussion of the future of the SDSS project. ",
        "conclusions": ""
    },
    "0812/0812.2150_arXiv.txt": {
        "abstract": "In the present paper we construct maps of polarized synchrotron radio emission of a whole galaxy, based on local models of the cosmic ray (CR) driven dynamo. We perform numerical simulations of the dynamo in local Cartesian domains, with shear-periodic boundary conditions, placed at the different galactocentric radii. Those local solutions are concatenated together to construct the synchrotron images of the whole galaxy. The main aim of the paper is to compare the model results with the observed radio continuum emission from nearly edge-on spiral galaxy.  On the basis of the modeled evolution of the magnetic field structure, the polarization maps can be calculated at different time-steps and at any orientation of the modeled galaxy. For the first time a self-consistent cosmic-ray electron distribution is used to integrate synchrotron emissivity along the line of sight. Finally, our maps are convolved with the given radiotelescope beam. We show that it is possible to reconstruct the  extended magnetic halo structures of the edge-on galaxies (so called X-shaped structures). ",
        "introduction": "\\label{sec:intro}  Recent deep observations of polarized radio-continuum emission at centimeter wavelengths of edge-on spiral galaxies \\cite[]{tul00,soida05,heesen05} revealed characteristic  X-shaped magnetic field structure. Many galaxies showed magnetic field oriented in the halo at a large angle to the galaxy disk. Earlier observations, mostly limited by sensitivity to narrow area along the major axis, demonstrated mainly plane-parallel component \\citep[e.g.][]{dumke95}. The polarization vectors of  NGC~5775 close to the galactic plane are parallel to the disk, but in the halo they extend up to 10~kpc above the disk forming so called $X$-shaped structure in all four quadrants. Such picture most probably reflects  the quadrupole configuration of magnetic field in NGC~5775 \\cite{soida08}. The observed degree of polarization is very high (locally reaching even up to 50\\%). It resembles high regularity of magnetic field there. Such structures seem to be common among edge-on spiral galaxies (NGC~5775, NGC~4666, NGC~4217 NGC~3628, NGC~253, NGC~891, \\citep[see e.g.][]{dahlem97,sukumar91,soida05}. In addition, \\cite{tul00} detected that NGC~5775 shows differential rotation in the direction perpendicular to the disk. So far, there is no physical explanation for those extended structures.  An attempt to explain the structure of polarized vectors by the dipolar dynamo wave in the NGC~5775 was made by \\cite{sokol02}. He used a solution of the dipolar classical turbulent dynamo equation and as a result he got structures not confirmed by the observations so far \\cite{soida08}.  The principle of the action of the CR-driven dynamo \\citep{parker92,hanasz00,hanasz04} is based on the cosmic ray (CR) energy supplied continuously by supernova (SN) remnants. Due to the anisotropic diffusion of cosmic rays along the horizontal magnetic field lines, cosmic rays tend to accumulate within the disc volume. However, the configuration stratified by the vertical gravity is unstable against the Parker instability. Buoyancy effects induce the vertical and horizontal motions of the fluid and the formation of undulatory patterns -- magnetic loops in the frozen-in, predominantly horizontal magnetic fields. The presence of rotation in galactic disks implies a coherent twisting of the loops by means of the Coriolis force leading to the generation of the small-scale radial magnetic field components. The next phase is merging the small-scale loops by the magnetic reconnection process which forms the large scale radial magnetic fields. Finally, the differential rotation stretches the radial magnetic field amplifying the large-scale azimuthal magnetic field. The coupling of amplification processes of the radial and azimuthal magnetic field components results in exponential growth of the large scale magnetic field with timescales of galactic rotation period (140~Myr) as shown by \\cite{hanasz06}.  In this paper we present the results of galactic disk modeling which is constructed from a set of local volumes placed at different distances from the galactic center. We show the results of integration of synchrotron emissivities (Stokes parameters I, Q and U) along the line of sight (LOS), and compare them to the observed radio images of the galaxy NGC~5775. Our results are compared to polarized radio continuum observations made at 6\\,cm \\citep{tul00}. ",
        "conclusions": "\\label{sec:conclusions}  In the present paper we compare the maps of polarized synchrotron radio emission of galaxy NGC~5775 with those constructed from the numerical model of the cosmic ray driven dynamo. The main conclusions can be summarized as follow: \\begin{itemize} \\item All models of the local cosmic-ray driven dynamo computed at the galactic radii between 4\\,kpc and 10\\,kpc indicate the fast growth of the magnetic flux and the total magnetic energy. \\item The synthetic radio maps of polarized emission computed on the basis of our local models exhibit vertical magnetic field structures similar to those observed in numerous edge-on galaxies. \\item The polarization vectors in the disk plane (face-on) form a spiral pattern with the pitch angles about 5$\\degr$. The pitch angles from the model are slightly smaller than we normally observe in galaxies. However, the pitch angle of the mean magnetic field depends to some extent on the actual parameters like e.g. the magnitude of the CR diffusion coefficients. \\item As it was expected for observation at relatively high frequencies, the Faraday effects considered in our model have no significant importance on the final results. \\item The major simplification of our model is caused by assumed axial symmetry, which is necessary for construction of the global model from local-box simulations. \\end{itemize}"
    },
    "0812/0812.0964_arXiv.txt": {
        "abstract": "Recent PAMELA data show that positron fraction has an excess above several GeV while anti-proton one is not. Moreover ATIC data indicates that electron/positron flux have a bump from 300 GeV to 800 GeV. Both annihilating dark matter (DM) with large boost factor and decaying DM with the life around $ 10^{26} s$ can account for the PAMELA and ATIC observations if their main final products are charged leptons ($e$, $\\mu$ and $\\tau$). In this work, we calculated the neutrino flux arising from $\\mu$ and $\\tau$ which originate from annihilating/decaying DM, and estimated the final muon rate in the neutrino telescopes, namely Antares and IceCube. Given the excellent angular resolution, Antares and IceCube are promising to discover the neutrino signals from Galactic center and/or large DM subhalo in annihilating DM scenario, but very challenging in decaying DM scenario. ",
        "introduction": "The main components of our Universe are `dark'. Many astrophysical observations have confirmed the existence of dark matter (DM) which contributes roughly $23\\%$ to the energy density of the Universe. If the DM are weakly interacting massive particles (WIMPs), they would annihilate into gamma rays, neutrinos and anti-matter particles, which could be detected experimentally. Generally speaking, the DM should be stable compared with the life of the Universe. However it is possible that the DM is decaying at an extremely low rate. Similarly, the DM decay products can be observable.  Recently PAMELA released their first cosmic rays results on the positron and anti-proton ratios \\cite{Adriani:2008zr}. Usually the anti-matter particles are expected to be produced when the cosmic rays propagate in the Galaxy and interact with the interstellar medium. The PAMELA data show a clear upturn (excess) on the positron ratio above $\\sim 10 $ GeV up to $\\sim 100 $ GeV. Besides PAMELA observation, ATIC collaboration also reported an electron/positron excess around $300 \\sim 800$ GeV \\cite{Chang:2008zz}. Contrary to the excess of position ratio, PAMEMA anti-proton ratio is consistent with the predictions. Both PAMELA and ATIC observations indicate some unknown high energy primary positron sources in the Galaxy. Studies showed that the pulsars or other nearby astrophysical sources may account for the PAMELA results \\cite{pulsar} .  Other possibility to induce the positron excess can be the annihilating \\cite{annidm,Cirelli:2008jk,ArkaniHamed:2008qn} or decaying \\cite{decaydm} DM in the Galaxy. In the annihilating DM scenario, if the cross section $ \\left\\langle {\\sigma v} \\right\\rangle$ is taken as the typical annihilation cross section $\\sim 3\\times10^{-26}cm^3s^{-1}$ for thermal relic, a large extra boost factor (BF) is needed to produce sufficient positron flux to fit the PAMELA results. Many ideas are proposed to induce the large BF. For example, the new `long' range force would enhance the annihilation cross section for today low velocity DM, which is dubbed as `Sommerfeld effect' \\cite{Cirelli:2008jk,ArkaniHamed:2008qn,sommerfeld,Lattanzi:2008qa}. For the decaying DM scenario, by adjusting the lifetime of the DM, no BF is needed. Contrary to positron excess, no obvious anti-proton excess has been observed. This suggests that the main final products from the annihilation or decay of DM are charge leptons ($e$, $\\mu$ and $\\tau$) for $m_{DM}< 1TeV$. Otherwise one needs much heavy DM, say $O(10TeV)$ \\cite{Cirelli:2008jk}, or some special cosmic-ray propagation parameters.  Provided that charged leptons are generated by DM annihilating and/or decaying DM, it is quite interesting to investigate which charged leptons are response for the PAMELA and ATIC observations. It is natural to expect that the high energy charged leptons are correlated to other observable signals, such as gamma rays, synchrotron emissions or neutrinos. Detecting neutrino signals from DM is very attractive and useful to pin down the nature of DM. In this paper, we will calculate the neutrino flux from the $\\mu$ and $\\tau$ which are producing by DM. Unlike other DM related products, neutrinos have less trajectory defection and less but finite energy loss \\cite{itoh} during propagation, due to their weakly interaction nature. As a result, the neutrinos can keep the important information of the DM. For the same reason, the neutrinos are difficult to be detected if the flux is low (for some recent works, see \\cite{Beacom:2006tt,Barger:2007xf,PalomaresRuiz:2007ry,Liu:2008kz,Yin:2008mv,Covi:2008jy} ). Usually the neutrinos detected by neutrino telescopes, such as Super-Kamiokande \\cite{Desai:2004pq,Ashie:2005ik}, AMANDA \\cite{Ackermann:2004aga} etc. are expected to be the atmospheric neutrinos, which are produced by interaction between the cosmic rays and atmosphere. The atmospheric neutrinos are the main backgrounds for detecting the extra neutrinos from DM. Currently, no high energy astrophysical neutrinos have ever been detected \\cite{Kistler:2006hp}.  If the charged leptons are indeed produced by DM, it is a good news for neutrino detection. Firstly, the neutrino flux can be much larger than that of usual expected. The PAMELA and ATIC data need higher primary positron production rate than usual production, thus it is natural to expect that large neutrino flux can be produced from $\\mu$ and $\\tau$ decays. Secondly, the heavy DM (e.g. $\\sim TeV$) is favored by ATIC result, thus the neutrinos at the similar energy scale are expected. As the main backgrounds, the atmospheric neutrino flux decreases logarithmically with the increment of energy, thus higher energy neutrinos correspond to less backgrounds. On-going and future kilometer size neutrino telescopes such as Antares \\cite{Aslanides:1999vq}, IceCube \\cite{Ahrens:2003ix} and KM3NeT \\cite{Carr:2007zc} can better explore these signals. Thirdly, in the annihilating DM scenario, the flux of neutrinos depends on the square of the DM number density. If one searches the Galactic center (GC) or large DM subhalos which contain the denser DM than other areas, large neutrino flux can be produced. Moreover if the Sommerfeld effect of DM annihilation is true, the DM annihilation in the subhalo can acquire even larger BF due to the lower velocity dispersion than that in the DM halo \\cite{ArkaniHamed:2008qn,Lattanzi:2008qa}. As a side remark, this feature provides a potential method to distinguish two DM scenarios, namely GC or DM subhalos can be higher flux neutrino sources in the annihilating DM scenario than those in the decaying DM one.  This paper is organized as following. In section II, we calculated the neutrino flux from the GC and DM subhalo, in the model-independent manner, in the annihilating and decaying DM scenarios. In section III, we estimated the muon rate at Antares and IceCube for the two DM scenarios. We pointed out that Antares and IceCube might detect neutrino signals in annihilating DM scenario and distinguish different DM scenarios. Section IV contains our discussions and conclusions. ",
        "conclusions": "Based on the PAMELA and ATIC observations, there may exist extra sources of charged leptons, namely electron, muon and tau. In this paper, we calculated the neutrino flux from muon and tau which arise from the annihilating/decaying DM. The final muon rate at neutrino telescopes Antares and IceCube are also estimated. Our results show that Antares is promising in discovering the neutrino signal from GC in annihilating DM scenario, but challenging in decaying DM scenario. Moreover, massive DM subhalo is also the promising high energy neutrino source in the annihilating DM scenario. For ten years operation, IceCube can reach five sigma significant neutrino signal from subhalos for heavy DM. Note that in the subhalo the extra Sommerfeld enhancement from lower velocity dispersion can further improve the neutrino signal.  In this paper we focus on the muon neutrino detection. Actually the tau neutrino is deserved further investigation because the atmospheric neutrino contains less tau neutrino than muon neutrino. If one only counts the tau numbers in the neutrino detector, it can further suppress the atmospheric background \\cite{Covi:2008jy} and this can be done in future kilometer neutrino detector like KM3NeT. Finally, it is worthy to mention that pulsars are difficult to be detected by upcoming neutrino detector \\cite{Bhadra:2008cg}, thus one may distinguish it from the DM scenarios."
    },
    "0812/0812.2016_arXiv.txt": {
        "abstract": "{% Observationally, both the 3.4$\\mum$ aliphatic hydrocarbon C--H stretching absorption feature and the 9.7$\\mum$ amorphous silicate Si--O stretching absorption feature show considerable variations from the local diffuse interstellar medium (ISM) to Galactic center (GC): both the ratio of the visual extinction ($A_V$) to the 9.7$\\mum$ Si--O optical depth ($\\tausil$) and the ratio of $A_V$ to the 3.4$\\mum$ C--H optical depth ($\\tauahc$) of the solar neighborhood local diffuse ISM are about twice as much as that of the GC. In this work, we try to explain these variations in terms of a porous dust model consisting of a mixture of amorphous silicate, carbonaceous organic refractory dust (as well as water ice for the GC dust). } ",
        "introduction": "The interstellar extinction law is one of the primary sources of information about the interstellar grain population, and one often obtains direct information on the composition of interstellar dust from spectral features in extinction (Draine 2003). These spectral features also provide strong constraints on interstellar grain models. With the advent of ground-based and space borne infrared (IR) telescope facilities, the IR extinction continuum and absorption features have been receiving increasing attention and play an essential role in recovering the intrinsic energy distribution of celestial objects and inferring the characteristics of interstellar dust.  In the interstellar extinction curve, the 2175$\\Angstrom$ bump is outstanding in the ultraviolet (UV), while in the IR, there are a number of prominent absorption features as well: (1) the ubiquitous 9.7$\\mum$ and 18$\\mum$ features respectively due to the Si--O stretching and O--Si--O bending modes of amorphous silicates; (2) the 3.4$\\mum$ feature due to the C--H stretching mode of aliphatic hydrocarbon dust, as ubiquitously present in the ISM of the Milky Way and external galaxies as the 9.7$\\mum$ and 18$\\mum$ silicate bands, except this feature is not seen in dense molecular clouds (see Pendleton 2004 for a review); (3) the 3.3$\\mum$ and 6.2$\\mum$ weak features seen in both local sources and GC sources (Schutte et al.\\ 1998; Chiar et al.\\ 2000), respectively due to the C--H stretching and C--C stretching modes of polycyclic aromatic hydrocarbon (PAH) molecules; and (4) in dense clouds the 3.1$\\mum$ feature due to the O--H stretching mode of water ice as well as a number of weaker features at 4.68$\\mum$ (CO), 7.68$\\mum$ (CH$_4$), 4.28$\\mum$, 15.2$\\mum$ (CO$_2$), 3.54$\\mum$, 9.75$\\mum$ (CH$_3$OH).  \\begin{table*}[t] % \\renewcommand{\\arraystretch}{1.2} \\vspace{0cm} \\caption{Observational values of $A_V/\\tauahc$ and $A_V/\\tausil$} \\vspace{-0.1cm} \\begin{center} \\begin{tabular}{cccc}\\hline\\hline & $A_V/\\tauahc$ & $A_V/\\tausil$ &  References\\\\ \\hline & 240$\\pm$40 & &  Sandford et al.\\ 1991 \\\\ & 250$\\pm$40 & &  Pendleton et al.\\ 1994 \\\\ & 333 & 16.7   &  Adamson et al.\\ 1990 \\\\ \\raisebox{1.5ex}[0cm][0cm]{\\textbf{Local diffuse ISM}} &  & 18.5  & Roche \\& Aitken 1984 \\\\ &  & 18.5$\\pm$2 & Draine 2003 \\\\ &  & 19.2$\\pm$0.6& Bowey et al.\\ 2004 \\\\ mean for local ISM &  274  & 18.2 & \\\\ \\hline  &   150$\\pm$20 &      &  Pendleton et al.\\ 1994 \\\\ & 143 &  8.31    & McFadzean et al.\\ 1989 \\\\ \\raisebox{1.5ex}[0cm][0cm] {\\textbf{Galactic Center}} &  & 8$\\pm$3 & Becklin et al.\\ 1978 \\\\ &  &  9 & Roche \\& Aitken 1985 \\\\ mean for GC  & 146 & 8.4 &  \\\\ \\hline \\end{tabular} \\label{tab1} \\end{center} \\end{table*}   The 9.7$\\mum$ silicate extinction profile varies among different sightlines; in particular, its optical depth $\\tausil$ (relative to the visual extinction $A_V$) shows considerable variations from the local diffuse ISM (LDISM) to the Galactic center (GC): $A_V/\\tausil \\approx 18.2$ for LDISM differs from that of the GC ($A_V/\\tausil \\approx 8.4$) by a factor of $\\simali$2.2 (see Table 1; also see Draine 2003).\\footnote{% $^1$Cohen et al.\\ (1989) argued that the observed 9.7$\\mum$ dip on which $\\tausil$ was measured may be partly contributed by PAHs (i.e. the red tails of the PAH 7.7$\\mum$ and 8.6$\\mum$ bands and the blue tail of the 11.3$\\mum$ band could form an ``artificial'' 10$\\mum$ dip). But this would result in a smaller $\\tausil$ for the GC and a larger $\\tausil$ for the LDISM (since the PAH emission is more likely to present in the LDISM while toward the GC PAHs are seen in absorption), quite on the opposite. } Roche \\& Aitken (1985) argued that $A_V/\\tausil$ varies because there are fewer carbon stars in the central regions of the Galaxy and the production of carbon-rich dust may be substantially reduced compared with the outer Galactic disk. However, as shown in Table~1, the 3.4$\\mum$ C--H feature of aliphatic hydrocarbon dust also exhibits a similar behavior: $A_V/\\tauahc \\approx 274$ for LDISM is higher than that of the GC ($A_V/\\tauahc \\approx 146$) by a factor of $\\simali$1.9 (see Table 1). If the argument of Roche \\& Aitken (1985) was valid, one would expect a much smaller $A_V/\\tauahc$ ratio in the LDISM than that of the GC. Sandford et al.\\ (1995) tried to quantitatively explain this phenomena by assuming that the abundance of the C--H carrier (relative to other dust components) gradually increases from the local ISM toward the GC.\\footnote{% $^2$They also pointed out that the C--H and Si--O carriers may be coupled, perhaps in the form of silicate-core organic-mantle grains. This idea is challenged by the nondetection of the 3.4$\\mum$ feature polarization along sightlines where the 9.7$\\mum$ feature polarization is detected (Adamson et al.\\ 1999, Chiar et al.\\ 2006; also see Li \\& Greenberg 2002). It is possible that the 3.4$\\mum$ feature may be not produced by a carrier residing in a mantle on a silicate core but by very small (unaligned) grains (Chiar et al.\\ 2006). } However, this requires that amorphous silicate dust and aliphatic hydrocarbon dust should not be solely responsible for the visual extinction. If one has to invoke an additional dust component (most likely a population of carbon dust which does not show the characteristic 3.4$\\mum$ feature, say, graphite) making an appreciable contribution to $A_V$, one would encounter a severe carbon budget problem (see Snow \\& Witt 1996).  Along the lines of sight toward the GC, there are dense molecular clouds.\\footnote{% $^3$The sightline toward the Galactic center source $\\sgrA$ suffers about $\\simali$30$\\magni$ of visual extinction (e.g. see McFadzean et al.\\ 1989), to which molecular clouds may contribute as much as $\\simali$10$\\magni$ (Whittet et al.\\ 1997). } In cold, dense molecular clouds, interstellar dust is expected to grow through coagulation (as well as accreting an ice mantle) and the dust is likely to be porous (Jura 1980). In this work, we demonstrate that the observed variations of $A_V/\\tausil$ and $A_V/\\tauahc$ from the LDISM to the GC could be explained in terms of composite porous dust. ",
        "conclusions": "We calculate $A_V/\\tauahc$ and $A_V/\\tausil$ for various dust models with a range of porosities and dust sizes. The results are tabulated in Table~2. For a given dust size $a=0.1\\mum$, porous dust results in smaller $A_V/\\tauahc$ and $A_V/\\tausil$ values than compact dust (see \\S2). This is encouraging that porous dust which is likely present in the dense molecular clouds along the sightlines toward the GC indeed is on the right track to account for the observed $A_V/\\tauahc$ and $A_V/\\tausil$ variations from the local diffuse ISM to the GC. More specifically, from Table~2 we see that the dust with $P\\sim 0.8-0.9$ and $a\\sim 0.5-1\\mum$ can approximately explain the observed variations of $A_V/\\tauahc$ and $A_V/\\tausil$ (by a factor of $\\simali$2) from the local diffuse ISM to the GC. Both high porosities ($P\\sim 0.8-09$) and large sizes ($a\\ge 0.5\\mum$) are required for the GC dust to account for the lower $A_V/\\tauahc$ and $A_V/\\tausil$ ratios.  In Figure~4, we show $A_V/\\tauahc$ and $A_V/\\tausil$ as a function of dust size for $P=0.8,\\,0.9,0.95$. It is clearly seen that for a given porosity, the variation of $A_V/\\tauahc$ with dust size exhibits a tendency similar to that of $A_V/\\tausil$. This suggests that with a dust size distribution taken into account, we would still maintain the variation tendency. For small, highly porous grains ($a<0.05\\mum$), $A_V/\\tauahc$ and $A_V/\\tausil$ are nearly independent of size since they are more or less in the Rayleigh regime in the optical-IR and therefore $A_V/V_{\\rm dust}$, $\\tauahc/V_{\\rm dust}$ and $\\tausil/V_{\\rm dust}$ are independent of grain size (where $V_{\\rm dust}$ is the dust volume). The visual extinction $A_V$ reaches its maximum at $a\\sim \\lambda/\\left[2\\pi\\,a\\,(n-1)\\right]$ (where $n$ is the real part of the refractive index of the porous dust at wavelength $\\lambda$; see Li 2008). At even large sizes, while $A_V$ reaches a constant (i.e. in the geometric optics regime) $\\tauahc$ and $\\tausil$ increase with $a$ (till they reach their respective peak values). This explains why $A_V/\\tauahc$ and $A_V/\\tausil$ decrease with $a$ after they reach their peak values.  For $A_V/\\tauahc$, our model with $m_{\\rm carb}/m_{\\rm sil} = 0.7$ (and $P=0.8$, $a=0.5-1.5\\mum$) is consistent with the observed factor-of-two variations in the local ISM and toward the GC (see Tables~1,2). However, the model values of $A_V/\\tausil$ for both the diffuse ISM dust ($A_V/\\tausil\\approx 38.2$) and the GC dust ($A_V/\\tausil\\approx 16.3$) are higher by a factor of $\\simali$1.5--2 than that observed in the local diffuse ISM ($A_V/\\tausil\\approx 18.2$) and the GC ($A_V/\\tausil\\approx 8.4$). This discrepancy may result from the underestimation of the silicate mass fraction.  \\begin{figure}[t] \\vspace{-0.8cm} \\centerline{\\hspace{0.2cm}\\includegraphics[scale=.38,clip]{fig4.eps}} \\vspace{-1.0cm} \\caption{% Variations of $A_V/\\tauahc$ and $A_V/\\tausil$ with dust size. \\label{fig4} } \\vspace{-0.2cm} \\end{figure} With an increased silicate mass fraction, say, $m_{\\rm carb}/m_{\\rm sil} = 0.5$ which is consistent with the {\\it in situ} measurements of comet Halley (Jessberger \\& Kissel 1991) and widely adopted in cometary dust modeling (Greenberg 1998; Greenberg \\& Li 1999; Kimura et al.\\ 2006, 2008a; Kolokolova et al.\\ 2004, Kolokolova \\& Kimura 2008; Mann et al.\\ 2006), we obtain $A_V/\\tauahc\\approx252$ and $A_V/\\tausil\\approx27.1$ for the local ISM (assuming compact dust), and $A_V/\\tauahc\\approx154$ and $A_V/\\tausil\\approx11.3$ for the GC (assuming porous dust). These values are closer to that observed. It is expected that with a smaller $m_{\\rm carb}/m_{\\rm sil}$ (i.e. a larger silicate mass fraction), one would obtain a smaller $A_V/\\tausil$ while $A_V/\\tauahc$ does not change much. Thus the observed variations of $A_V/\\tauahc$ and $A_V/\\tausil$ from the local ISM to the GC could be explained. It is worth noting that, based on a detailed analysis of the GC 5--8$\\mum$ absorption spectra obtained from the Kuiper Airborne Observatory, Tielens et al.\\ (1996) argued that silicate dust may contribute as much as 60\\% of the interstellar dust volume. This would translate to $m_{\\rm carb}/m_{\\rm sil} \\approx 0.34$ if we assume that the remaining 40\\% of the interstellar dust volume is all from the 3.4$\\mum$ C--H feature carrier (which is indeed a very generous assumption).  Admittedly, the proposed explanation is oversimplifed. In the future we will consider more realistic models in which more dust species (e.g. hydrogenated amorphous carbon with a range of C/H ratios), the distribution of dust along the line of sight toward the GC (e.g. see Sandford et al.\\ 1995), a distribution of dust sizes, and the possible porous nature of the diffuse ISM dust (e.g. see Mathis \\& Whiffen 1989) will be considered."
    },
    "0812/0812.0013_arXiv.txt": {
        "abstract": "Modifications of general relativity provide an alternative explanation to dark energy for the observed acceleration of the universe. We calculate quasilinear effects in the growth of structure in $f(R)$ models of gravity using perturbation theory. We find significant deviations in the bispectrum that depend on cosmic time, length scale and triangle shape. However the deviations in the reduced bispectrum $Q$ for $f(R)$ models are at the percent level, much smaller than the deviations in the bispectrum itself.  This implies that three-point correlations can be predicted to a good approximation simply by using the modified linear growth factor in the standard gravity formalism. Our results suggest  that gravitational clustering in the weakly nonlinear regime is not fundamentally altered, at least for a class of gravity theories that are well described in the Newtonian regime by the parameters $G_{\\rm eff}$ and $\\Phi/\\Psi$. This approximate universality was also seen in the N-body simulation measurements of the power spectrum by Stabenau \\& Jain (2006), and in other recent studies based on simulations. Thus predictions for such modified gravity models in the regime relevant to large-scale structure observations may be less daunting than expected on first principles. We discuss the many caveats that apply to such predictions. ",
        "introduction": "The energy contents of the universe pose an interesting puzzle, in that general relativity (GR) plus the Standard Model of particle physics can only account for about $4\\%$ of the energy density inferred from observations.  By introducing dark matter and dark energy, which account for the remaining $96\\%$ of the total energy budget of the universe,  cosmologists have been able to account for a wide range of  observations, from the overall expansion of the universe to the large scale structure of the early and late universe~\\cite{Reviews}.  The dark matter/dark energy scenario assumes the validity of GR at galactic and cosmological scales and introduces exotic components of matter and energy to account for observations. Since GR has not been tested independently on these scales, a natural alternative is that GR itself needs to be modified on large scales. This possibility, that modifications in GR on galactic and cosmological scales can replace dark matter and/or dark energy, has become an area of active research in recent years.  Attempts have been made to modify GR with a focus on galactic~\\cite{MOND} or cosmological scales~\\cite{DGP,fR,Sahni2005}.   Modified Newtonian Dynamics (MOND) and its relativistic version (Tensor-Vector-Scalar, TeVeS) ~\\cite{MOND} attempt to explain observed galaxy rotation curves without dark matter (but have problems on larger scales). The DGP model~\\cite{DGP}, in which gravity lives in a 5D brane world,  naturally leads to late time acceleration of the universe.  Adding a correction term $f(R)$ to the Einstein-Hilbert action \\cite{fR} also allows late time acceleration of the universe   to be realized.  In this paper we will focus on modified gravity (MG) theories that are designed as an alternative to dark energy to produce the present day acceleration of the universe. In these models, such as DGP and $f(R)$ models, gravity at late cosmic times and on large-scales departs from the predictions of GR.  By design, successful MG models are difficult to distinguishable from viable DE models against observations of the expansion history  of the universe. However, in general they predict a different growth of perturbations which can be tested using observations of large-scale structure (LSS) \\cite{Yukawa,Stabenau2006,Skordis06,Dodelson06,DGPLSS,consistencycheck,Koyama06,fRLSS,Zhang06,Bean06,MMG,Uzan06,Caldwell07,Amendola07, HuSaw3}.  In this paper we consider the quasilinear regime of clustering in which perturbation theory calculations are valid. We explore what $f(R)$ modifications to gravity predict about the behavior of the three-point correlation function. In \\S II we outline the particular type of $f(R)$ gravity model we will be using for our calculations. In \\S III we introduce the fundamentals of perturbation theory and in particular how it applies to modified gravity. In \\S IV we focus on second order corrections and in particular the Bispectrum. In \\S V we present our results and compare with other studies of the nonlinear regime. In \\S VI we discuss the implications for observations. ",
        "conclusions": "In this paper we studied the growth of structure in an $f(R)$ modified gravity model \\cite{HuSaw}. In the quasilinear regime, we used perturbation theory to calculate the three-point correlation function, the bispectrum, of matter. Our results are applicable up to scales at which the derivation of the quasi-linear perturbation theory is still valid.  We found that in the bispectrum the dominant behavior is due to the difference in the linear growth factors between the modified gravity model and regular gravity. The bispectrum itself shows significant departures in scale and time compared to the predictions of GR. However the reduced bispectrum, which is independent of the linear growth factor in perturbation theory for GR, remains within a few percent of the regular gravity prediction. It does show interesting signatures of modified gravity at the percent level.  Our results are consistent with studies of the nonlinear regime via N-body simulations, which have found that on a wide range of scales the nonlinear power spectrum can be predicted using the (modified gravity) linear growth factor in the standard formulae developed for Newtonian gravity. Our results imply that three-point correlations follow this trend at the few percent level. (The regime around and inside halos probed by \\cite{LimaHu} to test the chameleon behavior is not included in our perturbative study.) It would be interesting to compare perturbative and simulation results for the bispectrum for the models considered here and other modified gravity models.  Upcoming surveys in the next five years will not attain the percent level accuracy at which the reduced bispectrum shows distinct signatures of $f(R)$ gravity. In this time-frame, the bispectrum will be useful as a consistency check on potential deviations from GR found in the power spectrum. Such a check is useful since measurement errors and scale dependent biases of tracers can mimic some of the deviations in the power spectrum. Next generation surveys, to be carried out in the coming decade, will provide sufficient accuracy to test the distinct signatures seen in our results for the reduced bispectrum. With these surveys, the bispectrum can provide truly new signatures of gravity.  \\bigskip  {\\it Acknowledgments:} We are grateful to Jacek Guzik, Wayne Hu, Mike Jarvis, Justin Khoury, Matt Martino, Fabian Schmidt, Roman Scoccimarro, Ravi Sheth, Fritz Stabenau, Masahiro Takada and Pengjie Zhang.  BJ is supported in part by NSF grant AST-0607667."
    },
    "0812/0812.0007_arXiv.txt": {
        "abstract": "We use high resolution cosmological hydrodynamical simulations to demonstrate that cold flow gas accretion, particularly along filaments, modifies the standard picture of gas accretion and cooling onto galaxy disks.  In the standard picture, all gas is initially heated to the virial temperature of the galaxy as it enters the virial radius.  Low mass galaxies are instead dominated by accretion of gas that stays well below the virial temperature, and even when a hot halo is able to develop in more massive galaxies there exist dense filaments that penetrate inside of the virial radius and deliver cold gas to the central galaxy.  For galaxies up to $\\sim$L$^*$, this cold accretion gas is responsible for the star formation in the disk at all times to the present.  Even for galaxies at higher masses, cold flows dominate the growth of the disk at early times.  Within this modified picture, galaxies are able to accrete a large mass of cold gas, with lower initial gas temperatures leading to shorter cooling times to reach the disk.  Although star formation in the disk is mitigated by supernovae feedback, the short cooling times allow for the growth of stellar disks at higher redshifts than predicted by the standard model. ",
        "introduction": "\\label{intro}  The classic model for galaxy formation within the cold dark matter (CDM) paradigm assumes that gas within galaxy halos initially reaches the temperature of the virialized halo.  During the collapse of the halo, the process of violent relaxation of the dark matter shocks the gas component to the virial temperature.  Any subsequently accreted gas is also shocked to the virial temperature as it enters the virial radius \\citep{rees77,silk77,wr78}. This assumption has been routinely adopted by many analytical works \\citep[e.g.,][]{wf91,kauffmann93,Cole94,AR98,sp99,sb07,ks07,silk07}. Given a density profile, a cooling radius can be calculated, inside of which gas can radiate away its energy while conserving its specific angular momentum to form a centrifugally supported disk that grows from the inside out \\citep[e.g.,][]{FE80,dalcanton97,mmw,mao98,vdb00,galics, somerville08}. In general, these semi-analytic models (SAMs) have been successful in reproducing many of the present day properties of galaxies.  Some studies have called this virial temperature assumption into question, however.  Recent theoretical work has found that, in fact, the majority of baryons may never reach the virial temperature as they accrete onto galaxies \\citep{BD03,DB06,K05,ocvirk08,K08,D08}.  \\citet[][hereafter K05 and K08, respectively]{K05,K08} examined how gas gets into galaxies using a N-Body + Smoothed Particle Hydrodynamics (SPH) simulation of a cosmological volume.  They found that galaxies with halo masses below a few 10$^{11}$ \\Msun ~are dominated by ``cold'' gas accretion.  This cold gas is never shock heated to the virial temperature of the galaxy, and is thought to be accreted predominantly along filaments in the cosmic web \\citep[e.g.,][]{massey07}. Further support for this new paradigm has come from \\citet[][hereafter OPT08] {ocvirk08} using a similarly large cosmological volume simulated with an adaptive mesh refinement (AMR) code.  For galaxy halos to shock heat infalling gas, the cooling rate of gas behind the shock must be slower than the compression rate of the infalling gas \\citep{BD03, DB06}.  This requirement leads to two regimes in which gas is able to avoid shock heating, and instead cools onto the galaxy on a free fall timescale instead.  First, galaxies below a certain critical mass are not massive enough to support a stable shock because their dynamical time is shorter than the cooling time of the halo gas.  OPT08 find that the halo mass where cold and shocked accretion contribute equally is in good agreement with the theoretical analysis of \\citet[][hereafter DB06] {DB06}, with a critical mass of 10$^{11.6}$ \\Msun.  This critical mass is roughly independent of redshift.  Second, even when a shock is present, cold gas accretion can occur along filaments that penetrate deep inside the hot halo.  The ability of filaments to feed galaxy growth has been noted for some time \\citep{katz93, KW93,katz94,bj94,bond96,shen06, harford08}. K05 find that their cold flow gas accretion is largely along filaments. These filaments are efficient at funneling gas to the galaxy, and can bring material to the galaxy from a much larger radius than predicted in the standard spherical accretion model.  The mass contributed by filaments is a strong function of redshift (DB06, OPT08).  Note that in the first case above, when a galaxy is in the low mass regime incapable of supporting a stable shock, cold gas accretion can occur both in a quasi-spherical fashion and along filaments, while in the second, high mass regime in which a hot halo exists, cold gas accretion occurs primarily along preferred directions in filaments.   \\begin{deluxetable*}{lccccccccr} \\tablecaption{Simulated Galaxy Properties \\label{simsum} } \\tablewidth{0pt} \\tablehead{\\colhead{simulation} & \\colhead{M$_{vir}$} & \\colhead {R$_{vir}$} & \\colhead{V$_{circ}$} &  \\colhead{T$_{vir}$} & \\colhead{$\\lambda$} & \\colhead{z$_{LMM}$} &  \\colhead{$\\epsilon$} & \\colhead{N within R$_{vir}$} \\\\ & \\colhead{\\Msun} & \\colhead {kpc} & \\colhead{km/s} & \\colhead{K} &  &  & \\colhead{kpc} & \\colhead{dm+star+gas} } \\tabletypesize{\\footnotesize} \\startdata H579\t& 3.4$\\times$10$^{10}$ & 85  & 40  & 5.1$\\times$10$^4$ & 0.03 & $>$5 & 0.17 & $\\sim$9.5$\\times$10$^5$ \\\\ DWF1\t& 1.4$\\times$10$^{11}$ & 135 & 70  & 1.3$\\times$10$^5$ & 0.01 & 1.5 & 0.15 & $\\sim$5.3$\\times$10$^6$ \\\\ MW1\t& 1.1$\\times$10$^{12}$ & 270 & 130 & 5.1$\\times$10$^5$ & 0.05 & 2.5 & 0.3 & $\\sim$4.8$\\times$10$^6$ \\\\ GAL1\t& 3.3$\\times$10$^{12}$ & 385 & 190 & 1.1$\\times$10$^6$ & 0.02 & 2.8 & 0.3 & $\\sim$3.7$\\times$10$^6$ \\\\ H277\t& 7.1$\\times$10$^{11}$ & 230 & 115 & 3.8$\\times$10$^5$ & 0.03 & 2.5 & 0.35 & $\\sim$2.3$\\times$10$^6$ \\\\ MW1.lo  & 1.1$\\times$10$^{12}$ & 265 & 130 & 5.0$\\times$10$^5$ & 0.07 & 2.5 & 0.6 & $\\sim$5.8$\\times$10$^5$ \\\\ MW1.ad  & 1.1$\\times$10$^{12}$ & 270 & 130 & 5.1$\\times$10$^5$ & 0.07 & 2.5 & 0.6 & $\\sim$3.0$\\times$10$^5$ \\\\ \\enddata \\end{deluxetable*}  Galaxies above the critical mass for shocking appear to be dominated by hot gas accretion as in the standard model, but they may have been fed by cold flow filaments even after a shock develops, particularly at high $z$, and they are also built in part by the mergers of smaller galaxies that were dominated by cold flows.  Thus, cold flow gas accretion is an important component of galaxy formation that has remained largely ignored until recently.  Previous work on cold flow gas accretion in simulations have addressed the issue using large cosmological volumes, furnishing a statistical description of how galaxies get their gas.  This provides a useful starting point for modifying semi-analytic models of galaxy formation.  However, with large volumes and many simulated galaxies, numerical resolution for any individual galaxy must be sacrificed.  Hence, these simulations were unable to resolve the internal structure of galaxies and investigate the impact of cold flows on the formation of galactic {\\it disks}.  OPT08 follow gas as it enters the virial radius of galaxies, and examined the temperature history of that gas as it flows toward the central galaxy, down to 0.1 R$_{vir}$.  However, limited resolution prevented them, or K05 or K08, from being able to resolve disks in their galaxies.  Thus, the impact of cold flows to disk assembly, and in particular the growth of the stellar disk component, has yet to be addressed.  Rather than sample a  cosmological volume at low resolution, in this paper we investigate gas accretion in several very high resolution simulations, in order to study the impact of cold gas accretion on the assembly of stellar disks.  Each of these galaxies has at least a million particles within the virial radius by redshift zero.  This high resolution provides two advantages over previous simulation work on this subject. First, it allows us to resolve the internal structure of our galaxies and their disks.  Second, we overcome artificial broadening of shocks in SPH so that we are able to identify entropy increases associated with accretion shocks, allowing us to explicitly search for shocks rather than simply using a temperature threshold as in previous works (K05, K08, OPT08).  We are thus capable of examining the temperature and entropy history of accreted gas that eventually settles to the disk and forms stars.  These simulated galaxies are run within a $\\Lambda$CDM context and span two orders of magnitude in galaxy halo mass, allowing us to study the role of cold flows in galaxies from 0.01 L$^*$ to L$^*$.  This paper is outlined as follows: In Section \\ref{sims} we briefly discuss the simulations used in this study.  In Section \\ref{gasaccr} we describe our method to identify shocked and unshocked gas accretion to galaxies, and find that the fraction of accreted gas that is shocked is a strong function of galaxy mass, in agreement with previous works.  In Section \\ref{diskgrowth} we extend the results of Section \\ref{gasaccr} to investigate the role that unshocked, cold gas accretion has in building the stellar disk of galaxies, as a function of galaxy mass.  We briefly compare our results on disk growth to analytic models of disk galaxy formation.  In Section \\ref{issues} we discuss the various effects of our numerical scheme on the results, and conclude that our results are robust (those readers interested strictly in the science results may wish to skip this section). Finally, in Section \\ref{conclusions} we discuss the implications of our results for existing analytic models of disk galaxy formation, and conclude. ",
        "conclusions": "\\label{conclusions}  We have tracked the history of gas accreted to four high resolution disk galaxies that span two orders of magnitude in galaxy mass. For all of these galaxies, early disk growth can be predominantly attributed to cold gas, and up to Milky Way masses the SF of the stellar disk is dominated by unshocked, cold flow gas accretion at essentially all times.  This dominance of cold gas in the formation of disk stars can exist even when the the majority of gas accreting to the virial radius is being shocked and reaching high temperatures.  At galaxy masses above L$^*$, an early phase of disk growth prior to $z$ = 1 can be attributed to the presence of cold gas accretion.  We have examined gas temperature histories in relation to the virial temperature of the halo that they are accreted to.  While the shocked component is capable of reaching temperatures near the virial temperature of the halo, the unshocked component of gas accretion remains below 3/8 of the virial temperature at all times.  The important role of unshocked, cold flow gas in the growth of galaxies is consistent with previous studies that adopt an explicit temperature cutoff, determined by the cooling curve, to identify cold flow gas. We confirm that the ability of smoothly accreted gas to shock is a strong function of mass.  However, our focus here has been to study the subsequent role of this cold gas accretion in the growth of stellar disks.  Because the unshocked gas does not reach the same high temperatures as shocked gas, it is capable of cooling to the disk faster and forming stars.  It is important to note that this early cooling of unshocked gas to the disk is not a manifestation of the historic ``overcooling'' problem.  This can be seen by examining the results from our stellar mass - metallicity relationship, which matches the observed relation at both low and high redshifts \\citep{Brooks07,maiolino08}.  As discussed by both \\citet{Brooks07} and \\citet{maiolino08}, this match is achieved by overcoming the overcooling problem with both high resolution and a physically motivated SN feedback model.  Previous works consistently produced metallicities too high when compared to observations at high $z$, due to overcooling and a rapid consumption of gas and subsequent metal enrichment.  Instead, our SN feedback scheme regulates the SFRs of our galaxies, preventing them from too quickly consuming their gas. The early SF seen in our more massive galaxies is due instead to the accretion of cold gas that can cool rapidly to the disk to form stars.  Despite running our simulations in a fully cosmological context that includes galaxy mergers, the dominant gas supply is from smoothly accreted gas that has never belonged to another galaxy halo.  For a Milky Way mass galaxy with a quiescent merging history like the galaxies shown here, only $\\sim$25\\% of the accreted gas is acquired from other galaxy halos, resulting in a similar fraction of the disk stars forming from this gas.  It is our goal to assess the impact that cold flow gas accretion may have in altering current models of disk galaxy formation.  A direct comparison with semi-analytic models is difficult to do and beyond the scope of this paper.  Current SAMs include the possibility of rapid gas cooling for low mass galaxies. It has been suggested that this regime is similar to the cold flow gas accretion seen in simulations \\citep{croton06b}. \\citet{cattaneo07} performed a detailed comparison of the GalICS SAM \\citep{galics} to the SPH results of K05, and found that the cold gas accretion rates in the two schemes were in agreement. This implies that the amount of cold gas currently available in SAMs may agree well with simulations, though this needs to be confirmed with further comparisons.  However, the results in this paper may emphasize a separate problem within the cold flow regime.  Even in the case of rapid cooling in SAMs, the temperature of the gas is initially set to the virial temperature of the galaxy halo, and cooling times are based on this assumption. Prior to the development of a hot halo, the free fall time determines when the gas reaches the disk in either the SAMs or the simulations. However, SAMs do not yet model the impact of filaments on galaxy growth. Our two most massive galaxies, which are able to develop a hot halo at high redshifts \\citep[4 to 6; see also][]{derossi08}, are still capable of accreting cold gas that never approaches the virial temperature of the halo in filaments at high $z$. These filaments allow a significant amount of cold gas mass to accumulate inside of the galaxy, even for massive galaxies capable of supporting a shock at early times.  This cold gas cools quickly to the disk to form stars.  This rapid settling to the disk allows for the building of stellar disks at higher redshifts than predicted if all gas is heated to the virial temperature, as in the standard model.  We have shown that the discrepancy in temperatures between the shocked and unshocked gas components leads to significantly longer cooling times for the shocked gas (note that their densities will be discrepant as well, since most of the unshocked component are found in dense filaments, exacerbating the difference in cooling times).  If the unshocked gas had reached the virial temperature, the star formation rates in the disk would have been delayed (in the cases shown here, until after $z$ = 1).  A study of the angular momentum content of the accreted gas will be examined in future work, but the presence of cold flow gas accretion, particularly in filaments in massive halos, allows for the building of large, massive disks at high $z$, in agreement with observations."
    },
    "0812/0812.2002_arXiv.txt": {
        "abstract": "{The multifaceted role of the density dependent nuclear symmetry energy in the nuclear astrophysics involving neutron stars is highlighted.  Efforts toward a model independent determination of the dense matter equation state through a deconstruction of the neutron star structure equation utilizing the masses and radii of several individual neutron stars are described. The need for observational data of {\\em both} measurements for the same star is stressed.}  \\FullConference{10th Symposium on Nuclei in the Cosmos\\\\ July 27 - August 1 2008\\\\ Mackinac Island, Michigan, USA}      \\begin{document} ",
        "introduction": " ",
        "conclusions": "It is clear that  \\begin{enumerate}  \\item the masses and radii of several (say 5 to 7) of individual neutron stars can pin down (through deconstruction) the equation of state of neutron-star matter and shed light on the density dependence of the symmetry energy, i.e., the poorly known isospin dependence of strong interactions, and  \\item precise laboratory experiments, particularly those involving neutron-rich nuclei, are sorely needed to pin down the near-nuclear aspects of the symmetry energy which determines the masses, neutron skin thicknesses, collective excitations, etc., of nuclei.  \\end{enumerate}  It is also clear that continuing mutual interactions of astronomers, laboratory experimenters and theorists in the fields of astrophysics and strong interaction physics are necessary to fulfill the objective of understanding the physics of compact objects in which the ultimate energy density of observable cold-baryonic matter is realized~\\cite{LP05a}."
    },
    "0812/0812.5078_arXiv.txt": {
        "abstract": "Our ability to determine stellar ages from measurements of stellar rotation, hinges on how well we can measure the dependence of rotation on age for stars of different masses. Rotation periods for stars in open clusters are essential to determine the relations between stellar age, rotation, and mass (color). Until recently, ambiguities in $v\\sin i$ data and lack of cluster membership information, prevented a clear empirical definition of the dependence of rotation on color. Direct measurements of stellar rotation periods for members in young clusters have now revealed a well-defined period-color relation. We show new results for the open clusters M35 and M34. However, rotation periods based on ground-based observations are limited to young clusters. The Hyades represent the oldest coeval population of stars with measured rotation periods. Measurements of rotation periods for older stars are needed to properly constrain the dependence of stellar rotation on age. We present our plans to use the Kepler space telescope to measure rotation periods in clusters as old as and older than the Sun. ",
        "introduction": "Knowing stellar ages is fundamental to understanding the time-evolution of various astronomical phenomena related to stars and their companions. Accordingly, over the past decades much work has been focused on identifying the properties of a star that best reveal its age. For coeval populations of stars in clusters, the most reliable ages are determined by fitting model isochrones to single cluster members in the color-magnitude diagram. However, for the vast majority of stars not in clusters (unevolved late-type field stars), ages determined using the isochrone method are highly uncertain because the primary age indicators are nearly constant throughout their main-sequence lifetimes, and because their distances and thus luminosities are poorly known. Therefore, finding a distance-independent property of individual stars that can act as a reliable determinant of their ages will be of great value.  Stellar rotation (and the related measure of chromospheric activity - see paper by Mamajek this volume) has emerged as a promising and distance- independent indicator of age \\citep[e.g.][]{skumanich72,kawaler89,barnes03a,barnes07}. \\citet{skumanich72} first established stellar rotation as an astronomical clock by relating the average projected rotation velocity in young open clusters to their ages via the expression $\\overline{v\\sin i} \\propto t^{-0.5}$. The Skumanich relation is limited in mass to early G dwarfs and suffers from the ambiguity (due to the unknown inclination angle) of the $v\\sin i$ data. Furthermore, for ages beyond that of the Hyades cluster ($\\sim625$Myr), the Skumanich relationship is constrained only by a single G2 dwarf - the Sun.  Modern photometric time-series surveys in young open clusters can provide precisely measured stellar rotation periods (free of the $\\sin i$ ambiguity) for F, G, K, and M dwarfs. Based on such new data and emerging empirical relationships between stellar rotation, color, and age, a method was proposed by \\citet{barnes03a} to derive ages for late-type dwarfs from observations of their colors and rotation periods alone. We refer the reader to the paper in this volume by Barnes for a motivation and description of the method of {\\it gyrochronology}. However, our ability to determine stellar ages from stellar rotation, hinges on how well we can measure the dependence of rotation on age for stars of different masses. ",
        "conclusions": ""
    },
    "0812/0812.2975_arXiv.txt": {
        "abstract": "We report on a $\\sim$12 hr {\\em XMM-Newton} observation of the supergiant High-Mass X-ray Binary IGR~J16207--5129.  This is only the second soft X-ray (0.4--15~keV, in this case) study of the source since it was discovered by the {\\em INTEGRAL} satellite.  The average energy spectrum is very similar to those of neutron star HMXBs, being dominated by a highly absorbed power-law component with a photon index of $\\Gamma = 1.15^{+0.07}_{-0.05}$. The spectrum also exhibits a soft excess below $\\sim$2~keV and an iron K$\\alpha$ emission line at $6.39\\pm 0.03$~keV.  For the primary power-law component, the column density is $(1.19^{+0.06}_{-0.05})\\times 10^{23}$ cm$^{-2}$, indicating local absorption, likely from the stellar wind, and placing IGR~J16207--5129 in the category of obscured IGR HMXBs.  The source exhibits a very high level of variability with an rms noise level of 64\\%$\\pm$21\\% in the $10^{-4}$ to 0.05~Hz frequency range.  Although the energy spectrum suggests that the system may harbor a neutron star, no pulsations are detected with a 90\\% confidence upper limit of $\\sim$2\\% in a frequency range from $\\sim$$10^{-4}$ to 88~Hz.  We discuss similarities between IGR~J16207--5129 and other apparently non-pulsating HMXBs, including other IGR HMXBs as well as 4U~2206+54 and 4U~1700--377. ",
        "introduction": "The hard X-ray imaging by the {\\em INTEGRAL} satellite \\citep{winkler03} has been uncovering a large number of hard X-ray sources.  Since {\\em INTEGRAL}'s launch in 2002 October, 550 sources have been detected by the IBIS instrument in the $\\sim$20--50 keV band (based on version 29 of the ``General Reference Catalog''). Included in these sources are 236 ``IGR'' sources that were unknown or at least not well-studied prior to {\\em INTEGRAL}.  An important result from the {\\em INTEGRAL} mission has been the discovery of a relatively large number of High-Mass X-ray Binaries (HMXBs).  There are 37 IGR sources that have been classified as HMXBs (and it should be noted that about 1/3 of the IGR sources are still unclassified).  The IGR HMXBs are interesting both for the large number of new systems as well as the specific properties of these systems. These include a new class of ``Supergiant Fast X-ray Transients'' \\citep{negueruela06} that are HMXBs that can exhibit hard X-ray flares that only last for a few hours while the X-ray flux changes by orders of magnitude \\citep{intzand05,sguera06}.  Many of the IGR HMXBs are also extreme in having a high level of obscuration ($N_{\\rm H}$$\\sim$$10^{23}$--$10^{24}$ cm$^{-2}$) due to material local to the source \\citep{walter06,chaty08}.  For both the SFXTs and the obscured HMXBs, it is thought that a strong stellar wind is at least partially responsible for their extreme X-ray properties \\citep{fc04,walter06,wz07}.  The 37 IGR HMXBs presumably contain either neutron stars or black holes, but the nature of the compact object is only clear for the 12 systems for which X-ray pulsations from their neutron stars have been detected.  Eleven of these systems have pulse periods ranging from 5 to 1300~s, and the 12th system has an unusually long period of $\\sim$5900~s \\citep{patel04}.  Another X-ray property that can be taken as evidence for the presence of a neutron star is a very hard X-ray spectrum.  Typically, neutron star HMXBs have X-ray spectra that can be modeled with a power-law with a photon index of $\\Gamma$$\\sim$1 that is exponentially cutoff near 10--20~keV \\citep{nagase89,lutovinov05a}.  Although there are only a few known black hole HMXBs (e.g., Cygnus~X-1, LMC~X-1, LMC~X-3, M33~X-7), their X-ray spectral properties are similar to the general class of black hole X-ray binaries with power-law spectra with $\\Gamma$$\\sim$1.4--2.1 in their hardest spectral state \\citep{mr06}.  If black hole spectra show a cutoff, it is usually close to 100~keV \\citep{grove98}, significantly higher than seen for neutron star HMXBs.  However, it should be noted that evidence from the shape of the spectral continuum alone is usually taken only as an indication of the nature of the compact object. For example, HMXBs 4U~1700--37 and 4U~2206+54 both have continuum X-ray spectra similar to neutron star HMXBs, but they are considered to be only probable neutron star systems because pulsations have not been detected.  In this study, we focus on X-ray observations of IGR~J16207--5129, which was discovered in the Norma region of the Galaxy relatively early-on in the {\\em INTEGRAL} mission \\citep{walter04a,tomsick04_munich}.  Although it is a relatively faint hard X-ray source at $3.3\\pm 0.1$ millicrab in the 20--40~keV band \\citep{bird07}, it has been consistently detected by {\\em INTEGRAL} as well as in X-ray follow-up observations by {\\em Chandra} and {\\em XMM-Newton} (this work), indicating that it is a persistent source. The {\\em Chandra} observation provided a sub-arcsecond position that allowed for the identification of an optical counterpart with $R = 15.38\\pm 0.03$ and an IR counterpart with $K_{\\rm s} = 9.13\\pm 0.02$ \\citep{tomsick06}. Based on the optical/IR Spectral Energy Distribution, \\cite{tomsick06} found that the system has a massive O- or B-type optical companion.  Optical and IR observations confirmed this and indicate a supergiant nature for the companion \\citep{masetti06v,ns07,rahoui08}, and \\cite{rahoui08} estimated a source distance of $\\sim$4.1~kpc.  With further IR spectroscopy, the spectral type of the companion was narrowed down to B1~Ia and a source distance of $6.1^{+8.9}_{-3.5}$~kpc was estimated \\citep{nespoli08}.  Our current knowledge about the soft X-ray properties of IGR~J16207--5129 comes from the {\\em Chandra} observation that was made in 2005.  This showed that the source has a hard X-ray spectrum with a power-law photon index of $\\Gamma = 0.5^{+0.6}_{-0.5}$, and the source also exhibited significant variability over the 5~ks {\\em Chandra} observation \\citep{tomsick06}. Based on these properties, we selected this target for follow-up {\\em XMM-Newton} observations to obtain an improved X-ray spectrum and to search for pulsations that would provide information about the nature of the compact object.  Here, we present the results of the {\\em XMM-Newton} observation. ",
        "conclusions": "The {\\em XMM-Newton} observation of IGR~J16207--5129 confirms that it is a member of the group of obscured IGR HMXBs.  We measure a column density of $N_{\\rm H} = (1.19^{+0.06}_{-0.05})\\times 10^{23}$ cm$^{-2}$ for the average spectrum.  We find that the column density could have been that high during the only previous soft X-ray observation of this source with {\\em Chandra}, but a detailed spectral analysis of the {\\em XMM-Newton} data shows that $N_{\\rm H}$ can vary by a factor of $\\sim$2.  We detect an iron line in the energy spectrum, and its strength is consistent with what is expected for a spherical distribution of material with the measured $N_{\\rm H}$ around the compact object.  Although the X-ray spectrum is similar to those seen from other neutron star HMXBs with a hard primary power-law component and a soft excess, we do not detect the pulsations that might be expected if the compact object is a neutron star.  A detailed comparison between IGR~J16207--5129 and another apparently non-pulsating HMXB, 4U~1700--377, shows strong similarities in spectral and timing properties.  Most notably, the power spectra of both sources can be described as a single power-law down to $10^{-3}$-$10^{-4}$~Hz. Since pulsating HMXBs show power spectra that break near the pulsation frequency, it is possible that both of these sources harbor very slowly rotating neutron stars (although the possibility that the compact object is a black hole cannot be ruled out entirely).  We note that several of the IGR HMXBs are either known to harbor slowly rotating neutron stars or may harbor slowly rotating neutron stars in cases where pulsations have not been detected.  It is possible that in addition to uncovering obscured HMXBs, the HMXBs that {\\em INTEGRAL} is uncovering tend to contain slow rotators."
    },
    "0812/0812.0141_arXiv.txt": {
        "abstract": "We derive a generalised van Cittert-Zernike (vC-Z) theorem for radio astronomy that is valid for partially polarized sources over an arbitrarily wide field-of-view (FoV). The classical vC-Z theorem is the theoretical foundation of radio astronomical interferometry, and its application is the basis of interferometric imaging. Existing generalised vC-Z theorems in radio astronomy assume,  however, either paraxiality (narrow FoV) or scalar (unpolarized) sources. Our theorem uses neither of these assumptions, which are seldom fulfilled in practice in radio astronomy, and treats the full electromagnetic field. To handle wide, partially polarized fields, we extend the two-dimensional electric field (Jones vector) formalism of the standard ``Measurement Equation'' of radio astronomical interferometry to the full three-dimensional formalism developed in optical coherence theory. The resulting vC-Z theorem enables all-sky imaging in a single telescope pointing, and imaging using not only standard dual-polarized interferometers (that measure 2-D electric fields), but also electric tripoles and electromagnetic vector-sensor interferometers. We show that the standard 2-D Measurement Equation is easily obtained from our formalism in the case of dual-polarized antenna element interferometers. We find, however, that such dual-polarized interferometers can have polarimetric aberrations at the edges of the FoV that are often correctable. Our theorem is particularly relevant to proposed and recently developed wide FoV interferometers such as LOFAR and SKA, for which direction-dependent effects will be important. ",
        "introduction": "Polarimetric wide-field imaging is a recent and important trend in radio astronomy. Many new and planned radio telescopes such as LOFAR, LWA, MWA, and SKA have a wide field-of-view (FoV) as a major feature. Technologically, this has been made possible by the development of phased dipole arrays and focal plane arrays, both of which have an inherently wider FoV than traditional single-pixel radio telescopes. The motivation for wide FoV polarimetric radio telescopes in astronomy is that they facilitate the study of polarized phenomena not restricted to narrow fields such as the large, highly-structured, polarized features discovered in recent polarimetric galactic surveys \\citep{Taylor2006}. With LOFAR now producing its first \\emph{all-sky} images \\citep{LOFARteam}, a new era of polarimetric wide-field imaging is starting in radio astronomy.  Despite this trend, a complete theory for wide-field, polarimetric astronomical interferometry is lacking. Ultimately, classical interferometry is based on the far-zone form of the van Cittert-Zernike (vC-Z) theorem \\citep{TMS01,Mandel95} which, in its original form, is a scalar theory\\footnote{By scalar theory we mean a theory which only considers unpolarized radiation.} and only valid for narrow fields. Such restrictions are obviously unacceptable in astronomy where polarization often provides crucial astronomical information, and sources are distributed on the celestial sphere, not necessarily limited to small patches. In radio astronomy, a polarimetric, or vector (i.e. non-scalar), extension of the narrow-field vC-Z theorem was given by \\citet{Morris1964}, and more recently this was given a consistent mathematical foundation by \\citet{Hamaker2000} in the form of the so-called \\emph{Measurement Equation} (M.E.) of radio astronomy. On the other hand, a wide-field extension of the scalar vC-Z theorem for radio astronomy was given by \\citet{Brouw1971}, and more recently, imaging techniques for scalar, wide-fields have been an active field of research \\citep{Cornwell92,Sault1999b,Cornwell2005,McConnell06}. However, a generalised vC-Z theorem for radio astronomy that is both polarimetric and wide-field has not been derived. This may be because it has been assumed that such a theorem would be a trivial vector- or matrix-valued analogue of the wide-field, scalar theory \\citep{TMS01}. As we will see, however, the final result is not that simple.  The purpose of the present work is therefore to derive a vC-Z type relation by generalising the standard M.E. formalism to allow for arbitrarily wide fields. This is achieved by generalising the two-component Jones formalism to a three-component Wolf formalism \\citep{Wolf54}. In what follows, we will derive a vC-Z relation that is valid for the entire celestial sphere and is fully polarimetric. We will show that the standard M.E. can be recovered through a two-dimensional projection. We will also show that a dual-polarized interferometer of short electric dipoles (Hertzian dipoles) is inherently aberrated polarimetrically. We also extend our vC-Z relation to include not only the full electric field, but also the full magnetic field, and thereby establish the complete second-order statistical description of the electromagnetic relation between source terms and interferometer response.  Our wide-field vC-Z relations should be of relevance to the subject of \\emph{direction-dependent effects} in radio interferometric imaging, which has attracted recent attention \\citep{Bhatnagar2008}. ",
        "conclusions": "We have derived the full set of electromagnetic vC-Z relations, which are the basis of radio astronomical interferometry, without invoking the paraxial approximation. These relations allow all-sky imaging in a single telescope pointing. We have achieved these relations by generalising the usual 2-D Cartesian Jones vector based M.E. to a 3-D Wolf coherence matrix formulation on the (celestial) sphere. The derived wide-field vC-Z relations are not simply trivial matrix (or vector) analogues of the wide-field, scalar vC-Z, they exhibit direction dependent polarimetric effects. Indeed even in the scalar limit (that is, unpolarized radiation), our M.E. is not the same as the M.E. derived from scalar theory, in the case of Hertzian dipoles. Furthermore, we found that, for an arbitrary wide-field of sources, the electric visibilities generally have nine complex components for an arbitrary baseline. This implies that the standard Stokes (Cartesian) visibilities do not provide a full description of electric coherence for wide fields. We have also shown that our vC-Z relation (for the electric field) reduces to the standard 2-D M.E. after a 2-D projection. We showed that a consequence of this projection is that plane-polarized, Hertzian dipole interferometers are aberrated polarimetrically. Fortunately, these aberrations can be corrected for in the image plane for sources sufficiently far away from the plane of the dual-polarized antenna elements.  Besides its use in the derivation of the wide-field vC-Z relation, we believe that the 3-D Wolf formalism can be useful in constructing more general 3-D M.E. in cases requiring the full set of electric components. As examples, we mention the modelling of tri-polarized element arrays, the modelling of dual-polarized element arrays that are not strictly plane-polarized (due to manufacturing errors or the Earth's curvature), and the modelling of propagation that is not along the line-of-sight (due to refraction or diffraction in the ionosphere, e.g.)."
    },
    "0812/0812.4248_arXiv.txt": {
        "abstract": "We determine the Galactic production rate of strangelets as a canonical input to calculations of the measurable cosmic ray flux of strangelets by performing simulations of strange star mergers and combining the results with recent estimates of stellar binary populations. We find that the flux depends sensitively on the bag constant of the MIT bag model of QCD and disappears for high values of the bag constant and thus more compact strange stars.  In the latter case strange stars could coexist with ordinary neutron stars as they are not converted by the capture of cosmic ray strangelets. An unambiguous detection of an ordinary neutron star would then {\\em not} rule out the strange matter hypothesis. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.4624_arXiv.txt": {
        "abstract": "We study the details of preheating in an inflationary scenario in which the Standard Model Higgs, strongly non-minimally coupled to gravity, plays the role of the inflaton. We find that the Universe does not reheat immediately through perturbative decays, but rather initiate a complex process in which perturbative and non-perturbative effects are mixed. The Higgs condensate starts oscillating around the minimum of its potential, producing $W$ and $Z$ gauge bosons non-perturbatively, due to violation of the so called adiabaticity condition. However, during each semi-oscillation, the created gauge bosons partially decay (perturbatively) into fermions. The decay of the gauge bosons prevents the development of parametric resonance, since bosons cannot accumulate significantly at the beginning. However, the energy transferred to the decay products of the bosons is not enough to reheat the Universe, so after about a hundred oscillations, the resonance effects will eventually dominate over the perturbative decays. Around the same time (or slightly earlier), backreaction from the gauge bosons into the Higgs condensate will also start to be significant. Soon afterwards, the Universe is filled with the remnant condensate of the Higgs and a non-thermal distribution of fermions and bosons (those of the SM), which redshift as radiation and matter, respectively. We compute the distribution of the energy budget among all the species present at the time of backreaction. From there on until thermalization, the evolution of the system is highly non-linear and non-perturbative, and will require a careful study via numerical simulations. ",
        "introduction": "Inflation is nowadays a well established paradigm, consistent with all the observations, that solves most of the puzzles of the Hot Big Bang Model in a very simple and elegant way. It is able to explain not only the homogeneity and isotropy of the present Universe on large scales, but also the generation of almost scale invariant primordial perturbations that give rise to the structure formation \\cite{inflation}. However, the naturalness of inflation is directly related to the origin of the inflaton. Most of the inflationary models proposed so far require the introduction of new degrees of freedom  to drive inflation. The nature of the inflaton is completely unknown, and its role could be played by any candidate able to imitate a scalar condensate (tipically in the slow-roll regime), such as a fundamental scalar field, a fermionic or vector condensate, or even higher order terms of the curvature invariants. The number of particle physics motivated candidates is as big as the number of extensions of the Standard Model (Grand Unified Theories, supersymmetry, extra dimensions, etc.),  where it is not very difficult to find a field that could play the role of the inflaton~\\cite{Lindebook}.  In addition, given a model we must find a graceful exit to inflation and a mechanism to bring the Universe from a cold and empty post-inflationary state to the highly entropic and thermal Friedmann Universe~\\cite{reheating}. Unfortunately, the theory of reheating is also far from being complete, since not only the details, but even the overall picture, depend crucially on the different microphysical models. It seems difficult to study the details of reheating in each concrete model without the experimental knowledge of the strength of the interactions among the inflaton and the matter fields. Because of this, most of the work until now has focussed on models encoding the different mechanisms that could play a role in the process, with the strength of the couplings set essentially by hand. The relative importance of each one of these mechanisms can only be clarified in light of an underlying particle physics model, able to provide us with the couplings among the inflaton and matter fields. From this point of view it is very difficult to single out a given model of inflation, and even more difficult to understand the details of the reheating process via the experimental access to the couplings.  We may be very far away from understanding the microphysical mechanisms responsible for inflation, but maybe the natural candidate for being the inflaton was already there long ago. If we do not want to introduce new dynamical degrees of freedom in the theory, apart from those present in the Standard Model (SM) of particle physics, and at the same time we require Lorentz and gauge invariance, we are left with just one possibility: the Higgs field. Early models of inflation in terms of a Higgs-like scalar field $h$ with a  quartic self-interaction potential $\\frac{\\lambda}{4}h^4$  need an extremely small coupling constant $\\lambda\\sim 10^{-13}$ \\cite{Lindebook}, and are also nowadays excluded at around $3\\sigma$ by the present observational data \\cite{WMAPquarticinflation}. However, the SM-based inflation may be rescued from these difficulties replacing the usual  Einstein-Hilbert action by a non-minimal coupling of the Higgs field to the Ricci scalar \\cite{Mi77,Ad82,Smo79,Ze79,SalopekBardeenBond}, in a scalar-tensor theory fashion \\begin{equation} \\label{IG} S_{IG} =\\int d^4x \\sqrt{-g} \\; \\xi H^\\dagger H R\\;. \\end{equation} The induced gravity ($IG$) action is indeed the most natural generalization of the Standard Model in a curved spacetime, given the relevance of the non-minimal coupling to gravity for renormalizing the theory  \\cite{Birrell:1982ix}.  It belongs to a group of theories known as scalar-tensor theories, originally introduced by Brans and Dicke \\cite{BrDi61}  to explain the origin of the masses. In those theories  not only the active masses but also the gravitational constant $G$ ,  are determined by the distribution of matter and energy throughout the Universe, in a clear connection with the Mach principle. The interaction that gives rise to the masses should be the gravitational one, since gravity couples to all particles, {\\it i.e.} to their masses or energies. The gravitational constant $G$ in these theories is replaced by a scalar function that permeates the whole space-time and interacts with all the ordinary matter content, determining how the later moves through space and time. Any measurements of an object's mass depends therefore on the local value of this new field.  The successful Higgs mechanism lies precisely in the same direction of the original Mach's idea of producing mass by a gravitational-like  interaction. The  role of the Higgs field in the Standard Model is basically to provide the inertial mass of all matter fields through the local Spontaneous Symmetry Breaking mechanism\\footnote{We are not considering here Majorana masses.}. The Higgs boson couples to all the particles in the Standard model in a very specific way, with a strengh proportional to their masses \\cite{SModelref}, and mediates a scalar gravitational interaction of Yukawa type~\\cite{DeFrGh90,DeFr91}, between those particles which become massive as a consequence of the local Spontaneous Symmetry Breaking.  According to the Equivalence Principle, it seems natural to identify the gravitational and particle physics approaches to the origin of the masses. From this point of view, the induced gravity action (\\ref{IG}) would be an indication of a connection between the Higgs, gravity and inertia. Indeed, the action (\\ref{IG}) is, at least at the classical level, just a different representation of the Starobinsky's model of inflation \\cite{inflation, MuhkanovR2}, where inflation is entirely a property of the gravitational sector. Both representations of the same theory are simply related by a Legendre transformation. This fact, together with the possibility of having an inflationary expansion of the Universe,  makes the model extremely appealing. Unfortunately, the induced gravity model cannot be accepted as a completely satisfactory inflationary scenario, since the gauge bosons acquire a constant mass in the Einstein frame and totally decouple from the Higgs-inflaton field \\cite{SalopekBardeenBond}, which translates into an inefficient reheating of the Universe. Notice however that the action (\\ref{IG}) is not the most general one that can be written in a nontrivial background. As was shown in \\cite{SalopekBardeenBond,Shaposhnikov:2007} the simultaneous existence of a reduced \\textit{bare} Planck mass $M_P$ and a non-minimal coupling of a Symmetry Breaking field to the scalar curvature \\begin{equation}\\label{SHG} S_{HG} \\equiv \\int d^4x \\sqrt{-g} \\Bigg\\{ \\frac{M_P^2}{2}R+\\xi H^\\dagger H R \\Bigg\\}\\;, \\end{equation} avoid the decoupling of the gauge bosons and can give rise to an inflationary expansion of the Universe together with a potentially successful reheating.  In this paper we initiate the study of the reheating process in the model presented in Ref.~\\cite{Shaposhnikov:2007}, in which the Symmetry Breaking field is the Standard Model Higgs, strongly non-minimally coupled to gravity and playing the role of the inflaton. The novelty and great advantage of this model is its connection with a well-known microphysical mechanism, hopefully accessible in the near future accelerator experiments. The measurement of the Higgs mass will complete the list of the couplings of the Standard Model and, therefore, one should be able to study all the details of the reheating mechanism. This makes the model under consideration extremely interesting and, potentially, predictive. Reheating in the context of scalar-tensor theories has been studied, in the Hartree approximation by Ref.~\\cite{MaedaTorii2000}, and perturbatively by Ref.~\\cite{WatanabeKomatsu2007}, but without the Higgs boson playing the role of the inflaton, and therefore without an explicit coupling of the fundamental scalar field to matter.  Studying the details of reheating in this Higgs-Inflaton scenario, we have put special attention to the relative impact of the different mechanisms that can take place. We have found that the Universe does not reheat immediately through perturbative decays, but rather initiate a complex procces in which perturbative and non-perturbative effects are mixed. The Higgs condensate starts oscillating around the minimum of its potential, producing $Z$ and $W$ gauge bosons due to violation of the so called adiabaticity condition. During each semi-oscillation, the non-perturbatively created gauge bosons decay (perturbatively) into fermions. This decay prevents the development of the usual parametric resonance, since bosons do not accumulate significantly at the beginning. The energy transferred to the decay products of the bosons is not enough to reheat the Universe within a few oscillations, and therefore the resonance effects will eventually dominate over the perturbative decays. Around the same time, the backreaction from the gauge bosons into the Higgs condensate will also start to be significant. Soon afterwards, the Universe is filled with the remnant condensate of the Higgs and a non-thermal distribution of fermions and bosons (those of the SM), which redshift as radiation and matter, respectively. We end the paper computing the distribution of the energy budget among all the species present at the time of backreaction. From there on until thermalization, the evolution of the system is highly non-linear and non-perturbative, and will require a careful study via numerical simulations, to be described in a future publication.  The paper is organized as follows. In section~\\ref{Higgssector} we present the model with the Higgs field non-minimally coupled to the scalar curvature, transform it into a new frame where the action takes the usual Einstein-Hilbert form, and derive an approximate inflationary potential. In section~\\ref{mattersector} we study the effect of the conformal transformation in the matter sector, including the interaction among the Higgs, vector bosons and fermions. Section~\\ref{reheating} is devoted to the analysis of the different reheating mechanisms, both perturbative and non-perturbative, that can take place, leaving for Section~\\ref{stimulateddecays} the analysis of the combined effect of parametric resonance and perturbative decays. We then study the backreaction of the produced particles on the Higgs oscillations and the end of preheating in Section~\\ref{backreaction}. Finally the conclusions are presented in Section \\ref{conclusions}. ",
        "conclusions": "We have studied the different stages of reheating after inflation in a model where the role of the inflaton is played by the Higgs field of the Standard Model of particle physics with a non-minimal coupling to gravity. Inflation in this model takes place at the GUT scale, along the lines of the Starobinsky model of inflation since a conformal transformation makes these two models indistinguishable from the point of view of inflation. The usual difficulty with large self couplings of the Higgs is tamed here by the inclusion of a large non-minimal coupling to gravity, $\\xi\\sim10^5$, which nevertheless does not leave any signature at low (electroweak) scales due to the fact that the Higgs field acquires a vacuum expectation value and does not evolve at present, while the local spacetime curvature is negligible.  The advantage of this model of inflation for the study or reheating after inflation is that all the couplings of the Higgs-Inflaton to matter fields are known at the electroweak scale, and can be extrapolated to the GUT scale using the renormalization group equations, and therefore one can study in detail the process of reheating of the Universe, without having to impose ad hoc assumptions about their values. The surprise is that the process becomes more complicated than expected, and a series of subsequent stages take place, where essentially all different types of particle production mechanisms at preheating occur. Moreover, since the Standard Model couplings of the Higgs to gauge and matter fields are non-negligible, nor are their couplings among themselves, the process of non-perturbative decay via parametric resonance is mixed with the usual perturbative decays of the decay products, which complicates things significantly.  Inflation ends at values of the Higgs field of order the Planck scale and goes through a brief stage of tachyonic preheating soon after the end of inflation. The passage is so short that particle production is not significant at that stage. The same occurs with the production at the inflection point. Finally the Higgs-Inflaton field starts oscillating around the minimum of its potential with a curvature scale of order $10^{13}$ GeV. At this stage, particle production occurs whenever the Higgs passes through zero, creating mostly vector gauge bosons $W$ and $Z$. These gauge bosons acquire a large mass while the Higgs increases towards maximum amplitude and start to decay into all Standard Model leptons and quarks within half a Higgs oscillation, rapidly depleting the occupation numbers of gauge bosons, like in instant preheating. However, the fraction of energy of the Higgs that goes into SM particles is still very small compared with the energy in the oscillations, and therefore the non-perturbative decay is slow. This implies that a relatively large number of oscillations take place before a significant amount of energy is transferred to the gauge bosons and fermions.  The amplitude of Higgs boson oscillations decreases as the Universe expands in a matter-like dominated stage with zero pressure. Eventually, this amplitude is small enough that the gauge boson masses are not large enough for inducing a quick decay of the gauge bosons and these start to build up their occupation numbers very rapidly via parametric amplification. The question whether this effect can give rise to the production of a significant Gravitational Wave Background (GWB) potentially observable today remains to be addressed. Several papers have studied recently such an issue in the chaotic and hybrid models of inflation~\\cite{GarciaBellido:2007dg}, but in the present model, we don't have simply a parametric resonance phenomena but a combined preheating effect which, perhaps, could modify the properties of such a GWB. Similar arguments would affect also the production of magnetic fields at preheating~\\cite{DiazGil:2007dy} or even electroweak baryogenesis~\\cite{EWB}.  After about a hundred oscillations the gauge bosons produced backreact on the Higgs field and the resonant production of particles stops. The Higgs field acquires a large mass via its interaction with the gauge condensate and preheating ends. From there on, both Higgses and gauge fields decay perturbatively until their energy is transferred to SM particles. Since the stage after backreaction is very non-linear and non-perturbative, it cannot be solved analytically and we have to resort to numerical studies in the lattice. We leave the description of our numerical studies to a future publication.  \\  {\\bf Note added}: Upon completion of this paper, we received through the arXiv the preprint of Bezrukov et al.~\\cite{Bezrukov:2008}, where they also study preheating in the $\\nu$MSM. Although the formalism is common to both, our conclusions are somewhat different from those of Ref.~\\cite{Bezrukov:2008}. We find that the Higgs decay into gauge bosons is significantly faster, and that backreaction occurs much before thermalization. We thus think it is not possible to determine the reheating temperature without a careful numerical analysis with lattice simulations.  A few days after this work was completed, Ref.~\\cite{DeSimone:2008} was also posted in the arXiv, performing an analysis of the 2-loop quantum corrections to the running of all the parameters involved in the model. This paper allows for a different range of the Higgs self-coupling, which is compatible with the range~(\\ref{HiggsRange}) although more restricted. The main result of Ref.~\\cite{DeSimone:2008} is a relationship between the Higgs mass and the spectral index which, in principle, could be tested in the future against data from PLANCK and the LHC. Similar conclusions where found in Ref.~\\cite{Magnin:2008}.  \\subsection*{Acknowledgements}  We would like to thank Geneva University, SISSA-Trieste and MPI-Munich for hospitality during the development of parts of this research. DGF is supported by a FPU contract with Ref. AP2005-1092 and JR by an I3P contract. We also acknowledge financial support from the Madrid Regional Government (CAM) under the program HEPHACOS P-ESP-00346, and the Spanish Research Ministry (MEC) under contract FPA2006-05807. The authors participate in the Consolider-Ingenio 2010 CPAN (CSD2007-00042) and PAU (CSD2007-00060), as well as in the European Union 6th Framework Marie Curie Network ``UniverseNet\" under contract MRTN-CT-2006-035863."
    },
    "0812/0812.3004_arXiv.txt": {
        "abstract": "The effect of strong quantizing magnetic field on the equation of state of matter at the outer crust region of magnetars is studied. The density of such matter is low enough compared to the matter density at the inner crust or outer core region. Based on the relativistic version of semi-classical Thomas-Fermi-Dirac model in presence of strong quantizing magnetic field a formalism is developed to investigate this specific problem. The equation of state of such low density crustal matter is obtained by replacing the compressed atoms/ions by Wigner-Seitz cells with nonuniform electron density. The results are compared with other possible scenarios. The appearance of Thomas-Fermi induced electric charge within each Wigner-Seitz cell is also discussed. ",
        "introduction": "The theoretical investigation on the properties of compact stellar objects in presence of strong quantizing magnetic field have gotten a new life after the discovery of a large number of magnetars \\cite{R1,R2,R3,R4}. These exotic stellar objects are believed to be strongly magnetized young neutron stars. Their surface magnetic fields are observed to be $\\geq 10^{15}$G. Then it is quite possible that the fields at the core region may go up to $10^{18}$G (the theoretical estimate may easily be obtained from scalar virial theorem). The exact source of such strong magnetic field is of course yet to be known. These objects are also supposed to be the possible sources of anomalous X-ray and soft gamma emissions (AXP and SGR). Now, if the magnetic field is really so strong, specially at the core region, it must affect most of the important physical properties of these stellar objects and also some of the physical processes, e.g., the rates / cross-sections of elementary processes. In particular the weak and the electromagnetic decays / reactions taking place at the core region will change significantly.  The strong magnetic field affects the equation of state of dense neutron star matter. As a consequence the gross-properties of neutron stars \\cite{R5,R6,R7,R8}, e.g., mass-radius relation, moment of inertia, rotational frequency etc. should change significantly. In some recent work we have developed the relativistic version of Landau theory of Fermi liquid in presence of strong quantizing magnetic field to obtain the equation of state of dense neutron star matter \\cite{SCANN,SCLAND}. We have also shown that the nucleons acquire complex mass in the $\\sigma -\\omega -\\rho$ meson exchange type mean field in presence of strong magnetic field. It has been noticed that due to the complex nature of neutron and proton energies, there is a kind of relaxation or oscillation in the iso-spin space. In the case of compact neutron stars, the phase transition from neutron matter to quark matter, which may occur at the core region, is also affected by strong quantizing magnetic field. It has been shown that a first order phase transition initiated by the nucleation of quark matter droplets is absolutely forbidden if the magnetic field strength $\\sim 10^{15}$G at the core region \\cite{R9,R10}. However, a second order phase transition is allowed, provided the magnetic field strength $<10^{20}$G. This is of course too high to achieve at the core region. The study of time evolution of nascent quark matter, produced at the core region through some higher order phase transition, shows that in presence of strong magnetic field it is absolutely impossible to achieve chemical equilibrium ($\\beta$-equilibrium) configuration among the constituents of the quark phase if the magnetic field strength is as low as $ B \\sim 10^{14}G$.  The elementary processes, in particular, the weak and the electromagnetic decays/reactions taking place at the core region of a neutron star are strongly affected by such ultra-strong magnetic fields \\cite{R11,R12}. Since the cooling of neutron stars are mainly controlled by neutrino/anti-neutrino emissions, the presence of strong quantizing magnetic field should affect the thermal history of strongly magnetized neutron stars. Further, the electrical conductivity of neutron star matter, which mainly comes from free electron gas within the stars and directly controlls the evolution of neutron star magnetic field, should also change significantly \\cite{R12}.  In another kind of work, the structural stability of such strongly magnetized rotating objects are studied. It has been observed from the detailed general relativistic calculation that there are possibility of some form of geometrical deformation of these strongly magnetized objects from their usual spherical shapes \\cite{RR13,RR14,RR15}. In the extreme case such objects may either become black strings or black disks. It is quite possible to have gravity wave emission from these rotating magnetically deformed objects (see also \\cite{SOMA} for the magnetic deformation of nucleonic bags in neutron star matter).  In a recent study of the microscopic model of dense neutron star matter, it has been observed that if most of the electrons occupy the zeroth Landau level, with spin anti-parallel to the direction of magnetic field, and only a few are with spin along the direction of magnetic field  with non-zero Landau quantum number, then either such strongly magnetized system can not exist or such a strong magnetic field can not be realized at the core region of a neutron star \\cite{RR16}. We have observed that $10^{22}$G is the upper limit of magnetic field which the core of a neutron star can sustain. We have also noticed that since the electrical conductivity of the medium becomes extremely low in presence of ultra-strong magnetic field, the magnetic field at the core region must therefore decay very rapidly. Hence it may be argued that the magnetic field at the core of a magnetar can not be too high, it is only the strong surface field which has been observed.  Similar to the study of quark-hadron deconfinement transition inside neutron star core in presence of strong quantizing magnetic field, a lot of investigations have also been done on the effect of ultra-strong magnetic field on chiral properties of dense quark matter \\cite{SCCH,R13,R14,R15,R16, R17,R18,R19,Rina,VP2,SPK,WO} (see also \\cite{wang} for the chiral properties of strongly magnetized electron gas). The strong external magnetic field acts like a catalyst to generate mass dynamically. This is one of the most important effect of ultra-strong magnetic field on the properties of charged elementary particles. This effect is of course impossible to verify in the terrestrial laboratory. Only the cosmic laboratory offers us the opportunity to test this intrinsic phase transition in the world of elementary particles. Therefore, the presence of strong magnetic field in stellar matter causes a lot of interesting and significant changes in the properties of elementary particles. Also a lot of new phenomena are observed in such cosmic laboratory in presence of ultra-strong magnetic field.  The exotic process of magnetic photon splitting, i.e., the decay of a photon into two photons (or combination of two photons into a single one) in presence of a very strong magnetic field, has recently attracted renewed attention, mainly because of the great importance this process may have in the interpretation of the spectra of cosmic $\\gamma$-ray burst sources. The basic formulae for magnetic photon splitting had already been derived in the seventies \\cite{ADLER} (see also \\cite{ADLER1,WILKE,BAIER}). Again, since such strong magnetic field can not be generated, even in the pulse form, in the laboratory, it is absolutely impossible to verify this exotic phenomenon of photon splitting.  In this article we shall develop an exact (within the limitation of Thomas-Fermi-Dirac model) formalism of relativistic version of Thomas-Fermi-Dirac model in presence of strong quantizing magnetic field. From this model calculation, we shall obtain the equation of state of relatively low density crustal matter of magnetars, which is mainly dense iron crystal. We shall replace the iron ions / atoms by Wigner-Seitz cells with non-uniform electron density within each cell. We have organized the paper in the following manner: In the next section, we shall present the basic formalism of relativistic Thomas-Fermi-Dirac equation in presence of strong quantizing magnetic field. In section 3, we shall evaluate the equation of state of low density crustal matter of strongly magnetized neutron star. The kinetic energy part of the electron gas within each cell is obtained in section 4. Section 5 is dealt with the electron-nucleus interaction energy. In section 6, we evaluate the electron-electron direct interaction energy part. The exchange interaction part is obtained in section 7. The possibility of Thomas-Fermi induced charge within each cell is discussed in section 8 and finally in section 9, the last section, we shall discuss the results and future prospects of this model.  It is worth mentioning that the formalism we are presenting in this article is also applicable to strongly magnetized white dwarfs. For the sake of completeness, we have compared the results of the present model, with other possible physical scenarios, e.g., non-relativistic case with zero and non-zero values of Landau quantum number, with the non-relativistic and relativistic field free cases, etc.  Finally, we would like to mention, that the properties of low density magnetized crustal matter, mainly the electromagnetic properties, which also includes the transport properties have been studied thoroughly by Potekhin and Potekhin et. al. \\cite{RU}. ",
        "conclusions": "In this article we have developed the formalism for relativistic version of Thomas-Fermi-Dirac model in presence of strong quantizing magnetic field. The formalism is applicable to the outer crust of magnetars and also to strongly magnetized white dwarfs.  We have compared our results with several other cases, e.g., the well known non-relativistic model with zero magnetic field, field free relativistic case, non-relativistic model in presence of strong quantizing magnetic field for both $\\nu_{\\rm{max}}\\neq 0$ and $\\nu_{\\rm{max}}=0$.  We have noticed that in this formalism, to solve the Poisson equation numerically it is necessary to include a few more conditions, which were absent in the usual field free non-relativistic model or in presence of ultra-strong magnetic field ($\\nu_{\\rm{max}}=0$).  To remove singularity at the origin, we suggest, following \\cite{REMO}, to use finite dimension for the nuclei. It has also been noticed that unlike other scenario, one extra condition appears in the non-relativistic regime with $B\\neq 0$ and $\\nu_{\\rm{max}}\\neq 0$.  We have also given an approximate method to get an estimate of the induced charge within each cell and thereby obtain the variation of screening length with magnetic field strength.  In our model, the Wigner-Seitz cells are assumed to be spherical in nature and found that the radius of each cell decreases with the increase of magnetic field strength. The variation is given by $\\sim B^{-1/2}$.  The formalism is of course not applicable to the inner crust region, where the matter density is close to the neutron drip point, some of the neutrons may come out from the cells. In some future communication we shall present a modified version of this formalism appropriate for the inner crust region.  We have assumed that all the electrons within the cells are moving freely, i.e., they are not bound in any one of the atomic orbitals. In reality, it may happen that the electrons at the vicinity of the nucleus in a cell have negative energy. These electrons, therefore can not be treated as free. It is therefore absolutely necessary to get the total energy of an electron as a function of its position ($r$ or $x$) within the cell from the numerical solution of the Poisson's equation and the expressions for kinetic and various form of interaction energies. We expect that very close to the nucleus, the electron energy will be negative and for a particular value of $x$ ($=r/\\mu$) (which may be a function of $B$) it will become zero (quasi-free electrons) and then becomes positive. If it is found so, then we can not assume that all the $Z$-electrons in the cell are participating in statistical processes. On the other hand, if we consider the expression for electron energy as given in eqn.(86), then from the physics point of view all the electrons will become free (energy is always positive). Whereas, if we consider \\[ \\mu=~~~{\\rm{kinetic~ energy}}~~~ -e\\phi=~~~{\\rm{constant}}, \\] then we may have bound, quasi-free and free electrons within the cells. The presence of free electrons in the compressed cells in a dense medium is popularly known as {\\it{statistical ionization}} (see reference \\cite{SITE} for a detailed discussion).  To conclude our results, in the following we have given in tabular form the variation of Fermi momentum, Pressure, and various kinds of energy (except the exchange energy part) for electron gas within a typical Wigner-Seitz cell, with the strength of magnetic field.  \\noindent {\\bf{Table-II}} \\bigskip  \\begin{tabular} {|c|c|c|c|c|c|} \\hline\\hline $B/B_c$ & $p_F(x_s)$(MeV) & $P(x_s)$(MeV$^4$) & $E_{KE}(x_s)$(MeV) & $E_{en}(x_s)$(MeV) & $E_{ee}^{(d)}(x_s)$(MeV) \\\\ \\hline\\hline $10^5$ &  $2.57$  &$8.64\\times 10^4$&  $15.95$ & $46.71$&  $8.12$ \\\\ $5\\times 10^4$&  $2.58$&  $7153.02$&  $18.13$&  $27.15$&  $12.15$\\\\ $10^4$& $2.66$&  $4323.14$&  $19.58$&  $4.88$&  $19.80$\\\\ $5\\times 10^3$&  $2.76$&  $681.79$&  $24.53$& $-0.37$&  $23.82$\\\\ $10^3$& $3.56$&  $319.78$&  $27.30$& $-7.89$&  $36.18$\\\\ $5\\times 10^2$&  $4.82$&  $62.99$&  $37.67$& $-9.69$&  $44.55$\\\\ $10^2$& $6.95$&  $38.06$&  $44.81$& $-12.09$&  $78.04$\\\\ $50$&  $9.13$&  $31.49$&  $195.02$& $-13.13$&  $93.95$\\\\ $10$&  $13.61$&  $10.06$&  $781.47$& $-13.85$&  $104.24$\\\\ $1$&  $19.94$&  $6.28$&  $1.13\\times 10^5$& $-35.94$& $130.67$\\\\ \\hline\\hline \\end{tabular}  From the above tabular form of data one can see that the electron Fermi momentum and the corresponding kinetic energy decreases with the strength of magnetic field. Since the exchange energy has to be subtracted and its magnitude increases with the magnetic field strength, we may conclude that the system becomes more and more stable (total energy decreases) with the increase in magnetic field strength."
    },
    "0812/0812.2764_arXiv.txt": {
        "abstract": "Magnetic field is playing an important role at all stages of star evolution from star formation to the endpoints. The main effects are briefly reviewed. We also show that O--type stars have large convective envelopes, where convective dynamo could work. There, fields in magnetostatic balance  have intensities of the order of 100 G.  A few OB stars with strong polar  fields (\\cite[Henrichs et al. 2003a]{Henrichs03a})  show large N--enhancements indicating  a strong internal mixing. We suggest that the meridional circulation enhanced by an internal rotation law close to uniform in these magnetic stars  is responsible for the observed mixing. Thus, it is not the magnetic field itself which makes the mixing, but the strong thermal instability associated to solid body rotation.  A critical question for evolution is whether a dynamo is at work in radiative zones of rotating stars. The Tayler-Spruit (TS) dynamo is the best candidate.  We derive some basic relations for dynamos in radiative layers.  Evolutionary models with TS dynamo show important effects: internal rotation coupling and enhanced mixing, all model outputs being affected. ",
        "introduction": "The magnetic field of a star is, like the scent of a flower, subtle and invisible,  but it plays an essential role in evolution. Magnetic field  is often not accounted for in star models. However, the examples below show that from star formation to the endpoints as compact objects,  magnetic fields and rotation  strongly influence the course of evolution and all model outputs.  In Sect. \\ref{sectHR}, we give an overview where in evolution the fields are intervening. In Sect. \\ref{sectHen}, we emphasize some critical observations. In Sect. \\ref{sectdyn}, we focus on the general dynamo equations, with  examples in Sect. 5 and 6 for the Tayler--Spruit (TS) dynamo. ",
        "conclusions": ""
    },
    "0812/0812.3762_arXiv.txt": {
        "abstract": "{  The rapidly varying ($\\sim$10 minute timescale) non-thermal X-ray emission observed from Sgr A$^{\\star}$ implies that particle acceleration is occuring close to the event horizon of the supermassive black hole. The TeV $\\gamma$-ray source HESS\\,J1745$-$290 is coincident with Sgr A$^{\\star}$ and may be closely related to its X-ray emission. Simultaneous X-ray and TeV observations are required to elucidate the relationship between these objects. We report on joint H.E.S.S./Chandra observations performed in July 2005, during which an X-ray flare was detected.  Despite a factor of $\\approx$9 increase in the X-ray flux of Sgr~A$^{\\star}$, no evidence is found for an increase in the TeV $\\gamma$-ray flux from this region. We find that an increase in the $\\gamma$-ray flux of a factor of 2 or greater can be excluded at a confidence level of 99\\%. This finding disfavours scenarios in which the keV and TeV emission are associated with the same population of accelerated particles and in which the bulk of the $\\gamma$-ray emission is produced within $\\sim$$10^{14}$ cm ($\\sim$100$\\,R_{S}$) of the supermassive black hole.  } ",
        "introduction": "Measurements of stellar orbits in the central parsec of our galaxy have revealed the existence of a supermassive, $(3.6\\pm0.3)\\times 10^{6}$ solar mass, black hole coincident with the radio source Sgr~A$^{\\star}$~\\cite{GC:Eisenhauer05}. 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 \\citet{GC:Porquet08}, \\citet{GC:Eckart06} and \\citet{GC:Porquet03}). Variability on these timescales limits the emission region to within $<10$ Schwarzschild radii ($R_{S}$) of the black hole.  X-ray flares from \\astar\\ reach peak luminosities of $4\\times10^{35}$~erg~s$^{-1}$, two orders of magnitude brighter than the quiescent value~\\cite{GC:Porquet03, GC:Baganoff03}, and exhibit a range of spectral shapes~\\cite{GC:Porquet08}.  Several models of the origin of this variable emission exist, many of which invoke non-thermal processes close to the event horizon of the central black hole to produce a population of relativistic particles \\citep[see e.g.][]{GC:Markoff01,GC:Yuan03,GC:AharonianNeronov05apj, GC:Liu06,GC:Liu06HESS}.  Model-independent evidence that ultra-relativistic particles exist 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 measurements of this source, HESS\\,J1745$-$290, are those performed 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}, with work underway to reduce these uncertainties~\\cite{GC:ICRCHESSPos}.  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{GC:Levinson, GC:AharonianNeronov05apj}, TeV emission is produced in the immediate vicinity of the SMBH, and variability is expected. In alternative scenarios, particles are accelerated close to \\astar\\ but radiate within the central $\\sim$10~parsec region~\\cite{GC:AharonianNeronov05}, or are accelerated at the termination shock of a wind driven by the SMBH \\cite{GC:AtoyanDermer}.  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 presented by \\citet{HESS:gcprl}. Shell-type SNR are now well established TeV $\\gamma$-ray sources~\\cite{HESS:rxj1713p3,HESS:velajnr2} and several authors have suggested that Sgr A East is the origin of the TeV emission \\citep[see for example][]{GC:Crocker}. However, improvements in the uncertainty in the centroid position of HESS\\,J1745$-$290 \\cite{GC:ICRCHESSPos} effectively exclude the possibility of Sgr~A East being the dominant $\\gamma$-ray source in the region.  The pulsar wind nebula candidate G\\,359.95$-$0.04 discovered by \\citet{GC:Wang06} is located only 9$''$ from \\astar\\ and can plausibly account for 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, the possible origin of this source in the annihilation of WIMPs in a central dark matter cusp has been discussed extensively~\\cite{GC:Hooper04,GC:Profumo05,HESS:gcprl}.  Given the limited angular resolution of current VHE $\\gamma$-ray telescopes, the most promising tool in identifying the TeV source is the detection of \\emph{correlated variability} between the $\\gamma$-ray and X-ray, and/or NIR regimes.  A significant increase in the flux of HESS\\,J1745$-$290, occuring simultaneously with a flare detected in a waveband with sufficient angular resolution to isolate \\astar, would unambiguously identify the $\\gamma$-ray source.  Therefore, whilst not all models of the TeV emission from Sgr~A$^{\\star}$ predict variability in 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": "The absence of a significant increase in the $>$160~GeV $\\gamma$-ray flux of HESS\\,J1745-290 during a major X-ray flare (corresponding to an increase in flux by a factor of approximately 9 at maximum) suggests strongly that the keV and TeV emission cannot be attributed to the same parent population of relativistic particles. A possible component of the $\\gamma$-ray signal that has the same flaring behaviour as the X-rays is limited to a flux less than 100\\% of the quiescent state signal or $4.2\\times10^{-12}$ erg s$^{-1}$ cm$^{-2}$ (2--10 TeV), which should be compared with the $1.9\\times10^{-12}$ erg s$^{-1}$ cm$^{-2}$ 2--10 keV flux during the same period ($\\approx$$10\\times$ the quiescent flux).  The region of variable X-ray emission is limited by causality arguments to a size of $r_{X} < ct_{\\rm flare}$ or $\\sim$$10^{14}$~cm.  Following \\citet{GC:AtoyanDermer}, the radiation energy density at these distances from \\astar\\ is $U_{\\rm rad}\\,\\geq\\,3\\times10^{-4}$ erg cm$^{-3}$. If the X-ray emission is interpreted as synchrotron emission of $\\sim$TeV electrons, then the flux limit to inverse Compton (IC) emission at VHE energies during the flare implies that $U_{\\rm mag} > 0.5\\,U_{\\rm rad}$ and hence $B>50$ mG in this region. Since stronger magnetic fields are, in general, expected in this region \\citep[see e.g.][]{GC:Yuan03}, our result does not constrain models where the X-ray flares are assumed to be produced by synchrotron emission of relativistic electrons. We note that the arguments given above assume that the synchrotron flare is caused by an increase in the number of relativistic electrons. The alternative explanation that an increase in synchrotron emission occurs due to an impulsive increase in the magnetic field (with no direct effect on the IC flux) cannot be excluded.  The fact that HESS\\,J1745$-$290 does not appear to be associated with radiation processes within $10^{14}$~cm (or $\\sim$100$\\,R_{S}$) of the supermassive black hole does not exclude all scenarios in which \\astar\\ is the acceleration site for the particles responsible for the TeV emission. Scenarios in which the energy losses of the accelerated particles occur much farther from \\astar\\ (for example \\citet{GC:AharonianNeronov05}, \\citet{GC:AtoyanDermer} and \\citet{GC:Ballantyne07}) remain viable explanations for this $\\gamma$-ray source."
    },
    "0812/0812.1598_arXiv.txt": {
        "abstract": "We have found strong selective emission of the N~II 5000~\\AA\\ complex in the spectrum of the LMC hypergiant HDE~269896, ON9.7~Ia$^+$. Since this object also has anomalously strong He~II $\\lambda$4686 emission for its spectral type, an unusually wide range of ionization in its extended atmosphere is indicated.  The published model of this spectrum does not reproduce these emission features, but we show that increased nitrogen and helium abundances, together with small changes in other model parameters, can do so.  The morphological and possible evolutionary relationships of HDE~269896, as illuminated by the new spectral features, to other denizens of the OB Zoo are discussed.  This object may be in an immediate pre-WNVL (Very Late WN) state, which is in turn the quiescent state of at least some Luminous Blue Variables.  More generally, the N~II spectrum in HDE~269896 provides a striking demonstration of the occurrence of two distinctly different kinds of line behavior in O-type spectra: normal absorption lines that develop P~Cygni profiles at high wind densities, and selective emission lines from the same ions that do not.  Further analysis of these features will advance understanding of both atomic physics and extreme stellar atmospheres. ",
        "introduction": "Inverse N vs.\\ C, O anomalies in OB absorption-line spectra, denoted as OBN and OBC, have been described by Walborn (1976, 2003).  It is now generally accepted that the morphologically normal majority of OB supergiants display an admixture of CNO-cycled material in their atmospheres and winds, while the relatively rare OBC objects have physically normal (i.e., main-sequence) CNO abundances, and the OBN may have enhanced mixing as a result of additional effects such as binary interactions or rapid initial rotational velocities, with homogeneous evolution in extreme cases (Maeder \\& Meynet 2000).  For instrumental and historical reasons, as well as sometimes scientific ones, optical stellar spectroscopy has usually concentrated on certain wavelength regions, such as the blue-violet or the red.  As a result, new phenomena may be encountered when the orphan regions are examined. A recent example is the unexpected discovery of CNO anomalies in O2 spectra, from a survey of the 3400~\\AA\\ region (Walborn et~al.\\ 2004; Morrell et~al.\\ 2005).  Another example is the subject of this report, arising from a digital OB spectral-classification atlas extending somewhat beyond the traditional blue-violet limit around H$\\beta$ (http://www.fcaglp.unlp.edu.ar/$\\sim$mariela/atlas-mariela/index.htm).  The spectrum of HDE~269896 (= Radcliffe (R) 129, Sanduleak (Sk) $-68^{\\circ}135$), one of the brightest OB stars in the Large Magellanic Cloud, was first described in detail by Hyland \\& Bessell (1975) and Walborn (1977).  It displays two unusual characteristics: an emission line of He~II $\\lambda$4686 at its relatively late spectral type, and the CNO absorption-line intensity anomalies that define the ON supergiant class.  At earlier spectral types, the $\\lambda$4686 emission (Of) effect is a luminosity indicator (Walborn 1971, 1973, 2008; Walborn \\& Fitzpatrick 1990); hence, its presence in this spectrum was interpreted as an effect of superluminosity, which is supported by the very bright $M_V$ of $-$8.1, a magnitude brighter than typical Ia supergiants.  Thus, this spectrum was classified ON9.7~Ia$^+$ by Walborn (1977). ",
        "conclusions": "The ionization potentials of H$^0$, N$^+$, and He$^+$ are 13.6, 29.6, and 54.4~eV, respectively.  It is surprising to see both N~II and He~II features in the same spectrum; the former corresponds to later and the latter to earlier spectral types than that of HDE~269896.  Thus, this superluminous object must have an unusually extended atmosphere with a large range of ionization conditions.  These features together provide constraints on models of this atmosphere, as substantiated below.  Fitzpatrick (1991) discovered an interesting LMC relative of HDE~269896, namely Sk~$-66^{\\circ}169$, O9.7~Ia$^+$.  As implied by the spectral type, it has similar values of the classification criteria, including He~II $\\lambda$4686 emission (albeit somewhat weaker, consistent with its fainter but still superluminous $M_V$ of $-$7.5), and morphologically normal CNO features.  The optical spectra of both objects are illustrated by Fitzpatrick (1991), and their UV spectra from 900 through 1900~\\AA\\ by Walborn, Parker, \\& Nichols (1995) and Walborn et~al.\\ (2002a).  Comparative astrophysical analyses of these two stars are of considerable interest.  Indeed, they have been analyzed similarly by Crowther et~al.\\ (2002) and Evans et~al.\\ (2004).  Sections of the spectrum of HDE~269896 analyzed in the latter paper are reproduced here in Figures~2--4, and compared with the adopted model spectrum in green.  It is seen that, while fitting the He~I absorption lines well, the adopted model does not reproduce any of the N~II or He~II emission lines; in fact, it does not reproduce the N~III absorption-line strengths, either.  Fortunately, two other, unpublished trial models run at the same time are available; they are also shown here, in blue and red.  The blue model, with an increased N abundance (Table~1), does not produce emission lines, either.  However, the red model, with a similarly increased N abundance, but also somewhat lower effective temperature and increased mass-loss rate, does produce both N~II and He~II emission lines; actually, the N~II emission is somewhat overestimated (Table~2).  On the other hand, the He~I lines are badly underproduced, which is why this model was not favored at the time.  D.~Lennon (priv. comm.) has suggested that increasing the He abundance above the number ratio to H of 0.2 adopted by Evans et~al.\\ (2004) may remedy this discrepancy.  The results discussed here indicate that it may well be possible to reproduce most features in the optical spectrum of HDE~269896 with relatively small adjustments to the model parameters.  It is also very important to understand the atomic processes that produce the selective emission lines in O-type spectra, as emphasized by Walborn (2001).  These are lines that come into emission while others from the same ions remain in absorption; they are for the most part photospheric and respond to the temperature and luminosity.  A striking analogy to the N~II lines discussed here, but at higher ionization, is provided by the selective emission in N~IV $\\lambda$4058 vs. the absorption or P~Cygni profiles in the $\\lambda$3480 blend from the same ion, in early-O spectra (Walborn et~al.\\ 2004; Morrell et~al.\\ 2005).  It is remarkable that current models can produce these emission lines; evidently the essential physics has been correctly incorporated.  However, there has been no systematic effort to extract and elucidate the mechanisms involved, as was done for the Of N~III $\\lambda\\lambda$4634-4640-4642 triplet by Mihalas, Hummer, \\& Conti (1972) and Mihalas \\& Hummer (1973).  Such an effort will certainly provide interesting insights into the ionic level-population processes, and in turn further sensitive diagnostics for hot atmospheres.  Finally, it is interesting to compare HDE~269896 to other objects showing these N~II emission lines, particularly the WN10 and WN11 types defined by Crowther \\& Smith (1997); see also Smith, Crowther, \\& Prinja (1994). These objects have much more extensive emission-line spectra, including a strong He~I $\\lambda$5016 P Cygni profile adjacent to and blended with the N~II feature, which diagnose much denser winds.  Their typical He/H ratios range from 0.3 to 0.6.  Clearly HDE~269896 is a highly evolved object, but less so than the WN10/11 stars.  Thus, it is reasonable to propose that it may be in an immediate pre-WNVL (Very Late WN) stage, i.e., that it will develop a denser, slower wind, a more extensive emission-line spectrum, and a higher surface He/H ratio as its evolution proceeds.  Many (perhaps all) WNVL objects are now recognized as quiescent states of Luminous Blue Variables (Walborn et~al.\\ 2008 and references therein), and HDE~269896 has a comparable luminosity (Humphreys \\& Davidson 1994).  HDE~269896 is located in relative isolation, north of 30~Doradus, but near the O2~III(f*) star Sk~$-68^{\\circ}137$, also a very massive, likely WN progenitor (Walborn et~al.\\ 2002b)."
    },
    "0812/0812.4232_arXiv.txt": {
        "abstract": "We examine the distance of the two galactic microquasars \\gro\\ and \\aaa\\, which are potentially the two closest black holes to the Sun. We aim to provide a picture as wide and complete as possible of the problem of measuring the distance of microquasars in our Galaxy. The purpose of this work is to fairly and critically review in great detail every distance method used for these two microquasars in order to show that the distances of probably all microquasars in our galaxy are much more uncertain than currently admitted. Moreover, we show that many confirmations of quantitative results are often entangled and rely on very uncertain measurements.  We also present a new determination of the maximum distance of \\gro\\ using red clump giant stars, and show that it confirms our earlier result of a distance less than 2~kpc  instead of 3.2~kpc. Since it then becomes more likely that \\gro\\ could originate from the stellar cluster NGC~6242, located at 1.0 kpc, we review the distance estimations of \\aaa, which is so far the closest black hole with an average distance of about 1.0~kpc. We show that the distance methods used for \\aaa\\ are also problematic. Finally, we present a new analysis of spectroscopic and astrometric archival data on this microquasar, and apply the maximum-distance method of Foellmi et al. (2006). It appears that \\aaa\\ could indeed be even closer to the Sun than currently estimated, and consequently would be the closest known black hole to the Sun. ",
        "introduction": "Microquasars are stellar binaries in which one of the component is a black hole. The stellar companion is filling its Roche lobe, ejecting matter through the first Lagrangian point, which then organizes itself around the compact object as an accretion disk. The temperature of the disk increases strongly towards the center and large amounts of X-rays with powerful, persistent and collimated jets are produced. In some cases, the relativistic jets appear to be superluminal. Microquasars are galactic laboratories of high-energy phenomena, and they must be seen as part of a large paradigm where AGNs, microquasars and possibly gamma-ray bursts (GRBs) share similar physics \\citep{Mirabel-2004}.  There are various physical reasons why the distance is an important parameter in our understanding of microquasars. Firstly, the superluminal velocity effect of the jets \\citep[which is a relativistic time delay effect producing an apparent velocity on the sky larger than the speed of light; see e.g.][]{Rees-1966} depends on the distance. Although it is a pure geometrical effect, it has a minimum threshold of $\\beta = v/c \\geq 2^{-1/2}$ (with $c$ the speed of light) below which the jets are not superluminal, whatever their true spatial speed $v$. When one tries to model the jets, and more importantly, the jet's launching mechanism, it is important to know the true velocity (and hence power) of the jets. Moreover, this effect is not very common, and it will become more important to correctly identify superluminal objects when the statistics will grow with future X-ray missions.  Secondly, the possible misalignment of the jets with respect to the accretion disk axis. As a matter of fact, \\gro\\ is often quoted as the typical microquasar where the jets and the disk are misaligned \\citep[e.g.][]{Maccarone-2002}. This is directly related to the fact that the determination of the jet projection angle is linked to that of the distance, as explained below. The possible misalignment of the jet is an important clue of the formation and evolution of the system \\citep[see e.g.][]{Martin-etal-2008}.  Thirdly, the black holes' origin, and their orbit in our Galaxy. This is particularly true for \\gro, where a change in the distance by a factor of three makes a large difference to the orbit of the object in the Galaxy. See for instance \\citet{Mirabel-etal-2002} who have calculated the orbit for \\gro\\ with $D = 0.9$ and 3.2 kpc. This is important for our understanding of the origin of black holes in our Galaxy, and their distribution throughout the disk.  Finally, the true luminosity of the stellar companion can also be an important point in our understanding of the microquasars as dynamical objects, and in particular how the companion is affected by the filling of its Roche lobe and is certainly (partially) irradiated by X-rays from the accretion disk \\citep[e.g.][]{Dubus-etal-1999}.  The distance of microquasars is a simple yet central parameter in our description of these objects, conditioning a large part of our understanding both of the physics and the astronomical views of black hole stellar systems.  \\subsection{The origin of the disagreement on the distance of \\gro}  The origin of this work is to be found in the publication by \\citet{Mirabel-etal-2002} who note that the distance of \\gro\\ can be actually radically different from the accepted distance of 3.2~kpc determined by \\citet{Hjellming-Rupen-1995}. This later value has since then been apparently confirmed by many other studies. In fact, \\gro\\ is a runaway black hole and Mirabel et al. published its proper motion, obtained with the {\\it Hubble Space Telescope}. The opposite direction of the proper motion points almost perfectly (see their Fig. 1) to the center of a cluster (NGC~6242) located at 1.0 kpc from the Sun \\citep{Kharchenko-etal-2005}. It is then tempting to think that the system received a kick velocity at the moment of the primary's supernova explosion, and then moved away from its original cluster. This is the starting question: is the distance of \\gro\\ 3.2 or 1.0 kpc? Among other things, the jets are superluminal with $D=3.2$ kpc, but not at 1.0 kpc.  \\citet{Foellmi-etal-2006b} published a new method providing only an upper limit to the distance and giving $D\\lesssim 1.7$ kpc when applied to \\gro, thus challenging the accepted distance of 3.2~kpc. This method is also partially problematic, and it will be discussed critically below. With a maximum distance of 1.7~kpc, \\gro\\ becomes a likely candidate of being the closest (stellar) black hole to the Sun. The position is currently held by \\aaa, which has an average measured distance of about 1.0 kpc. We will show however that the distance of \\aaa\\ is also very uncertain, revealing different problems in each distance method.  In summary, none of the published distances of \\gro\\ and \\aaa\\ are so far reliable. This paper emphasizes the difficulties encountered when determining the distance of microquasars, and shows that the distances of probably many, if not all, microquasars in our Galaxy are much more uncertain than currently admitted. In this paper we aim to provide a fair and critical review on the distance methods used on these two galactic microquasars. Other microquasars will be the subject of a subsequent work. In addition to \\citet{Foellmi-etal-2006b}, partial material (incomplete and partially incorrect on the radio-jet distance of \\gro) has already been published in \\citet{Foellmi-2006e} and \\citet{Foellmi-2007c}. This paper addresses many more important details, and in particular the issues of the dynamical studies used to determine the absolute magnitude of the secondary star of \\gro. Moreover, we also provide a completely new estimate of the distance of \\gro. Finally, we apply the method of \\citet{Foellmi-etal-2006b} to \\aaa, along with archival astrometric data. We conclude that \\gro\\ is certainly closer than the current accepted distance, but that \\aaa\\ might be even closer still.   \\subsection{Distance methods: the basics}  The distance of a {\\it stellar} object is often measured by comparing its absolute and apparent magnitudes. There are other methods, such as astrometry and parallax, and jet speed measurements (see below) for the special case of microquasars. But most often, the core method is simply this one. It reads: \\begin{equation} m_{\\textrm{\\scriptsize true}} - M = 5 \\, \\log(D) - 5 \\end{equation} where $m_{\\textrm{\\scriptsize true}}$ is the true apparent magnitude of the object, affected only by the geometrical distance separating the object and the observer.  There are (at least) two major issues with this method: one is general, and one is specific to stellar binaries of short period. The first {\\it observational} difficulty is that $m_{\\textrm{\\scriptsize true}}$ is usually different from the observed apparent magnitude $m_{\\textrm{\\scriptsize obs}}$ since the light is going through some absorbing patchy interstellar medium, which makes the star dimmer and redder. We have: $m_{\\textrm{\\scriptsize true}} = m_{\\textrm{\\scriptsize obs}} - A$, where $A$ is the absorption, in the given passband, expressed in magnitudes. Therefore: \\begin{equation} m_{\\textrm{\\scriptsize obs}} - M - A = 5 \\, \\log(D) - 5 \\label{equ_magdist} \\end{equation} The determination of $A$ is critical. It is often not measured directly, but rather is the so-called color excess, or reddening: $E(B-V)$, which is the relative amount of additional red color between the $B$ and $V$ bands due to the absorption of bluer wavelengths by the gas and the dust. For a star, the color excess can be obtained by comparing the observed and the intrinsic $B$ and $V$ color indices (the latter being inferred from the spectral type): \\begin{equation} E(B-V) = (B-V)_0 - (B-V) \\label{equ_ebmv} \\end{equation} The absorption in the $V$ band follows: \\begin{equation} A_V = R \\times E(B-V) \\label{equ_AV} \\end{equation} where $R$ is the total-to-selective absorption. For other colors, this relation needs corrections \\citep[see e.g.][]{Fitzpatrick-1999}. The practice shows that what is often measured directly is $E(B-V)$, $R$ being often approximated\\footnote{The difference between $R=$3.0 and 3.7 implies a relative distance uncertainty of 1.38.} by a value between 3.0 and 3.7 with a more-or-less canonical value of 3.1. In fact, $R$ is a function of the interstellar reddening curve and the color of the stars, because the wide passbands of the photometric $B$ and $V$ filters makes the effective wavelengths of the filters shift with different stellar intensity gradients. \\citet{Olson-1975} gives an approximate relation for $R$: \\begin{equation} R = 3.25 + 0.25\\, (B-V)_0 + 0.05\\, E(B-V) \\label{equ_R} \\end{equation} where $(B-V)_0$ is the unreddened color index of the star. This relation is valid for normal stars with $(B-V)_0 < 0.4$ and $E(B-V) < 1.5$ to within an error of 0.05 in $R$ \\citep{Olson-1975}. See for instance \\citet{Crawford-Mandwewala-1976} for a comparison for various photometric systems, and \\citet{McCall-2004} for an updated discussion on this topic.  The other issue with the method consisting of comparing the magnitudes, specific to the stellar binaries with short period, is that the stellar companion (here the secondary star) is certainly not spherical anymore, since it completely fills its Roche lobe to feed the accretion disk. Moreover, its surface temperature might also not  be homogeneous because of irradiation \\citep[e.g.][]{Dubus-etal-1999,OBrien-etal-2002}, making its average observed temperature a function of the orbital phase. Hence uncertainties arise when one tries to estimate the true value of the absolute magnitude, $M$, in equation~\\ref{equ_magdist}, since it is often calibrated with \"normal\" (spherical) and isolated stars.  The absolute magnitude is normally obtained through the determination of the spectral type of the star, which gives the temperature $T_{\\textrm{{\\scriptsize eff}}}$, and the modelling of either the radial-velocity curve or the multi-color lightcurves of the binary system gives the size of its orbit. From the latter, one can compute the {\\it effective} Roche-lobe radius \\citep[see for instance][]{Eggleton-1983}, which is then identified to the radius of the star. Assuming a uniform temperature distribution across the star and that it is roughly spherical, the luminosity (and thus the absolute magnitude) can be estimated with: \\begin{equation} L = 4 \\pi \\sigma R^2 T^4 \\label{equ_lum} \\end{equation} where $\\sigma$ is the Stefan-Boltzmann constant. ",
        "conclusions": "It has been clearly shown that despite a large number of variants in distance methods, it is of prime importance to be extremely rigorous with assumptions and with the use of results made earlier and usually by others. The distance of microquasars, and in general of compact objects in our Galaxy, will benefit in a few years of the results of the satellite Gaia. Until then, we will certainly have to combine various types of information to infer good estimates to the distance of these objects. We mention here the work of \\citet{Lazorenko-etal-2007} who have obtained, through astrometric measurements using VLT/FORS1, a precision of 30 {\\it micro}arcseconds, similarly to what is expected for the VLTI instrument PRIMA, currently being commissioned at the ESO Cerro Paranal Observatory. This method requires a rather simple observational setup, a large field, and a large number of stars in the field. These requirements are easily fulfilled in the case of microquasars in the galactic plane like \\gro, although it might again be more difficult or impossible for \\aaa.  From the present work, we can conclude that: \\begin{itemize} \\item The use of the color excess measured using the sodium lines is risky. Not only one must have a spectroscopic resolution high enough, but its applicability range is limited. Moreover, one must be careful when choosing the relationship between the equivalent width and the color excess. \\item There is no proven systematic overestimation of the optical absorption as determined from X-ray data compared to that inferred from optical data. \\item The optical data also needs to be checked for being taken truly during optical quiescence or not. This is crucial for lightcurve models. \\item The quality of the data and the uncertainties associated to it must be assessed rigorously. It is meaningless to apply a sophisticated model on data of limited quality. \\item The maximum-distance method proposed by \\citet{Foellmi-etal-2006b} contains a systematic uncertainty, although its main hypothesis is sometimes implicitly used by others. \\end{itemize}  About \\gro: \\begin{itemize} \\item We have shown that although the upper limit of 3.5~kpc is a rather firm measurement, the value of 3.2 kpc has never been really measured. \\item We have also shown that the lower limit of 3.0~kpc is questionable and actually based on a questionable assumption on the interpretation of absorption lines in the radio spectrum. Moreover, we have challenged the relevance of comparing two radio spectra in a region where there HI clouds with anomalous velocities. \\item We have estimated that $E(B-V) \\sim$1.0 and $A_V \\sim 3.49$. \\item The peculiar absorption determined from X-rays towards \\gro\\ possibly indicate the presence of local dust close to the object. However, the X-ray flux in quiescence is also very variable. \\item The studies of the orbit and secondary star of \\gro\\ rely very much on the quality of the model of the disk, and on the assumed fixed parameters, often extracted from other incomplete studies. As matter of fact, no true model of the lightcurves of \\gro\\ combining the radial-velocity measurement and the multi-color lightcurves has been made while letting all the parameters truly vary at the same time. \\item A new distance method using red clump giant stars applied to \\gro\\ however confirms a distance less than 2 kpc. \\item Although not proven, the new estimation of a  smaller distance of \\gro\\ strengthen the idea of its possible origin in NGC~6242. It would make \\gro\\ one of the closest black holes to the Sun. \\item At 1.0 kpc, $\\beta = 0.28$ and there would be no misalignment between the disk and the jets, since $\\theta \\sim 71^{\\circ}$. \\end{itemize}  As for \\aaa: \\begin{itemize} \\item We have shown that the published distance and many confirmations of it are based (not always explicitely) on a single measurement of the color excess made 30 years ago on a spectrum made of 5 points, which is moreover in contradiction with an \\emph{HST/STIS} spectrum. \\item We have applied the maximum-distance method of \\citet{Foellmi-etal-2006b} to \\aaa\\ and found that it could be indeed located much closer, to a distance of $\\sim$0.4~kpc. \\item The example of \\aaa\\ illustrate that normal distance methods based on comparing the apparent and absolute magnitudes are difficult to apply to \\aaa\\ since the estimation of the extinction cannot rely on relationship established in the galactic plane. For this reason also, the red clump giant star distance method could not be applied to \\aaa. \\item We have also used archival images to infer an upper limit of the proper motion of \\aaa, which appears quite small. Although we have found a cluster of stars 2.8 degrees away, the origin of \\aaa\\ remains uncertain. \\end{itemize}  Finally, to the question \"What is the closest black hole to the Sun?\", our answer is the black hole X-ray binary \\aaa."
    },
    "0812/0812.3624_arXiv.txt": {
        "abstract": "{Observations suggest that low-mass stars condense out of dense, relatively isolated, molecular cloud cores, with each core spawning a small-$N$ cluster of stars.} {Our aim is to identify the physical processes shaping the collapse and fragmentation of a $5.4\\,{\\rm M}_\\odot$ core, and to understand how these processes influence the mass distribution, kinematics, and binary statistics of the resulting stars.} {We perform SPH simulations of the collapse and fragmentation of cores having different initial levels of turbulence ($\\alpha_{_{\\rm TURB}}=0.05,\\,0.10,\\,0.25$). We use a new treatment of the energy equation that captures (i) excitation of the rotational and vibrational degrees of freedom of H$_2$, dissociation of H$_2$, ionisation of H and He, and (ii) the transport of cooling radiation against opacity due to both dust and gas (including the effects of dust sublimation, molecules, and H$^-$ ions). We also perform comparison simulations using a standard barotropic equation of state.} {We find that -- when compared with the barotropic equation of state -- our more realistic treatment of the energy equation results in more protostellar objects being formed, and a higher proportion of brown dwarfs; the multiplicity frequency is essentially unchanged, but the multiple systems tend to have shorter periods (by a factor $\\sim 3$), higher eccentricities, and higher mass ratios. The reason for this is that small fragments are able to cool more effectively with the new treatment, as compared with the barotropic equation of state. We also note that in our simulations the process of fragmentation is often bimodal, in the following sense. The first protostar to form is usually, at the end, the most massive, i.e. the primary. However, frequently a disc-like structure subsequently forms round this primary, and then, once it has accumulated sufficient mass, quickly fragments to produce several secondaries.} {We believe that this delayed fragmentation of a disc-like structure is likely to be an important source of very low-mass stars in nature (both low-mass hydrogen-burning stars and brown dwarf stars); hence it may be fundamental to understanding the way in which the statistical properties of stars change -- continuously but monotonically -- with decreasing mass. However, in our simulations the individual cores probably produce too many stars, and hence too many single stars. We list the physical and numerical features that still need to be included in our simulations to make them more realistic; in particular, radiative and mechanical feedback, non-ideal magneto-hydrodynamic effects, and a more sophisticated implementation of sink particles.} ",
        "introduction": "{\\it The influence of thermodynamic effects on low-mass star formation in isolated, low-turbulence cores.} There is direct and indirect observational evidence to suggest that a significant fraction of low-mass stars form in small, relatively isolated, low-turbulence prestellar cores, with each core spawning only a few stars. This paper is concerned with exploring this mode of star formation, by means of numerical simulations, with a view to evaluating (a) how the statistical properties of the resulting protostars depend on the initial conditions, and (b) what role is played by thermal, chemical, and radiative processes. We adopt an analytic approach. That is to say, we do not seek to implement all the deterministic physical effects at once, and hence we do not expect at this stage to reproduce all the observed features of real star formation. Rather, we seek to establish whether particular physical effects, namely the thermal and radiative processes, influence the outcome in a systematic way. Given the complexity of star formation and the limitations of numerical codes, this incremental approach seems a more fruitful one to pursue.  {\\it The direct observational evidence for low-mass star formation in isolated, low-turbulence cores.} The direct evidence for this mode of star formation comes from a number of studies. First, Andr\\'e et al. (2007) have estimated the inter-core velocity dispersion in Ophiuchus (i.e. the dispersion in the bulk velocities of cores relative to their neighbours). Andr\\'e et al. infer that the frequency of interactions between cores is so low that a typical core is likely to have collapsed and fragmented, internally, before it interacts with a neighbouring core. Second, estimates of the level of turbulence in low-mass prestellar cores give subsonic values, i.e. $\\sigma_{_{\\rm TURB}} < \\sigma_{_{\\rm THERM}}$ (e.g. Myers 1983; Myers et al. 1991; Myers 1998; Andr\\'e et al. 2007). Indeed, Myers (1998) concludes that the decay of turbulence to subsonic levels may be a pre-requisite for the formation of low-mass protostars.  {\\it The indirect observational evidence for low-mass star formation in isolated, low-turbulence cores.} Indirect evidence for this mode of star formation comes from the binary statistics of young low-mass stars, which show that a high fraction are born in binary or higher-order multiple systems. The binary fraction decreases with decreasing primary mass, and with increasing age, but for a $1\\,{\\rm M}_\\odot$ primary it is still $\\sim 60\\%$ in the field (Duquennoy \\& Mayor 1991). Goodwin \\& Kroupa (2005) and Hubber \\& Whitworth (2005) have shown that this high multiplicity requires newly-born stars to complete their early dynamical evolution in small sub-clusters containing just a few stars (i.e. ${\\cal N}_{_{\\rm SUBCLUSTER}}\\sim 4\\pm 1\\;{\\rm stars}$). This is because, in a small sub-cluster, $N$-body interactions tend rather quickly to deliver a tight binary, usually comprising the two most massive stars, and to eject most of the remaining stars as singles (McDonald \\& Clarke 1993; Goodwin \\& Kroupa 2005; Hubber \\& Whitworth 2005). Therefore, if ${\\cal N}_{_{\\rm SUBCLUSTER}}$ is increased, a higher proportion of stars are ejected as singles, and therefore the primordial binary fraction is reduced, in contradiction with the observations.  {\\it The need for large ensembles of numerical simulations.} One advantage of this mode of star formation, involving a low-mass core fragmenting in relative isolation to form a small number of protostars, is that it can be simulated with quite high numerical resolution. A related advantage is that many realisations can be computed. Given the chaotic nature of turbulent, self-gravitating gas dynamics, this is essential if robust statistical inferences are to be made.  {\\it The effect of the initial level of turbulence in previous simulations.} The first comprehensive investigation of this mode of star formation was made by Goodwin et al. (2004a,b), who simulated the evolution of a large ensemble of cores, all having the same mass, initial density profile and turbulent power spectrum. They considered three different levels of turbulence, characterized by \\begin{eqnarray}\\label{EQN:TURB} \\alpha_{_{\\rm TURB}}&\\equiv&\\frac{E_{_{\\rm TURB}}}{|E_{_{\\rm GRAV}}|}\\;\\,=\\;\\,0.05\\,,\\;0.10\\,,\\;0.25\\,, \\end{eqnarray} where $E_{_{\\rm TURB}}$ and $E_{_{\\rm GRAV}}$ are the initial turbulent and gravitational energies of the core. For each value of $\\alpha_{_{\\rm TURB}}$, they performed many different simulations, each with a different realisation of the turbulent velocity field, in order to obtain reasonable statistics. They found that increasing the level of turbulence increased both the total number of stars formed, and the proportion of brown dwarfs.  {\\it The effect of the turbulent power spectrum in previous simulations.} In a subsequent paper, Goodwin et al. (2006) considered different turbulent power spectra, of the form $P_k\\propto k^{-x}$, and showed that increasing the exponent $x$ resulted in more fragments. In effect, increased $x$ means that there is more turbulent power on long wavelengths, and hence more large-scale fragmentation resulting in separate protostars. This is in accordance with the findings of Klessen \\& Burkert (2001), who note that turbulence has two effects. Turbulence on very short (sub-Jeans) length-scales can be viewed as a source of extra pressure, supplementing the thermal pressure. Conversely, turbulence on larger length-scales creates the structures which, if they become sufficiently massive and dense, are amplified by self-gravity to become prestellar cores.  {\\it The limitations of using a barotropic equation of state.} However, in their simulations Goodwin et al. (2004a,b, 2006) use a simple barotropic equation of state, which is designed to mimic the gross thermal behaviour of the gas at the centre of a collapsing, non-rotating $1\\,{\\rm M}_\\odot$ protostar, as determined by Masunaga \\& Inutsuka (2000). This is not realistic. First, it does not take proper account of the thermal history of protostellar gas, which depends on environment, metallicity, mass and geometry. Second, a barotropic equation of state is unable to capture thermal inertia effects, for example, the complex system of interacting pressure waves which dominates the dynamics when the gas becomes adiabatic and the thermal timescale suddenly becomes longer than the dynamical timescale. This is when fragmentation occurs (e.g. Boss et al. 2000), and thermal inertia effects are therefore critical.  {\\it Our new treatment of the energy equation.} We have therefore revisited the study of Goodwin et al. (2004a,b), but now using a significantly improved treatment of the energy equation due to Stamatellos et al. (2007a). This new treatment of the energy equation captures the critical thermal and radiative effects in protostellar gas much more realistically than a barotropic equation of state. For full details of the new treatment, and of the tests that have been performed to establish its fidelity, the reader is refered to Stamatellos et al. (2007). With the new treatment of the energy equation, we find  \\begin{itemize}  \\item{that the efficiency of fragmentation is increased, in the sense that on average more stars are formed, and in particular more brown dwarfs;}  \\item{that the mass distribution becomes bimodal, with a peak around $0.3\\;{\\rm to}\\;1.0\\,{\\rm M}_\\odot$, and a subsidiary peak in the brown dwarf regime around $0.03\\;{\\rm to}\\;0.06\\,{\\rm M}_\\odot$ (but note that we are only considering the fragmentation products of a single core mass -- we do not expect this bimodality to translate to the overall stellar mass function, when a distribution of core masses is considered);}  \\item{that fragmentation frequently involves the formation of a primary star (from material having low angular momentum) and the subsequent accumulation of a massive disc around this primary, with the disc then fragmenting ${\\sim\\!20,000}\\;{\\rm to}\\;{\\sim\\!100,000}\\;{\\rm years}$ later, to produce secondaries;}  \\item{that brown dwarfs result from disc fragments which {\\it either} are born in the {\\it outer} disc where there is not much material to assimilate, {\\it or} are quickly ejected from the disc, usually by mutual interactions, before they can assimilate much mass; disc fragments that form and remain in the {\\it inner} disc usually assimilate sufficient mass to become hydrogen-burning stars;}  \\item{that a significant fraction of the stars end up in binary systems, or higher multiples, and these systems tend to have large mass ratios (i.e. $q\\equiv M_{_2}/M_{_1}\\sim 1$);}  \\item{that there is little evidence for competitive accretion -- in fact, as noted above, there is a rather egalitarian process at work tending to produce binary systems with components of comparable mass.}  \\end{itemize}  {\\it The plan of the paper.} In Section 2, we describe the initial conditions, the constitutive physics, and the numerical method. In Section 3, we present the results of our simulations and the statistical distributions derived from them, and relate these distributions to the underlying physics. In Section 4 we summarise our conclusions.  \\begin{figure*} \\centering$ \\begin{array}{cccc} \\includegraphics[scale=0.2,angle=270]{N1.eps} & \\includegraphics[scale=0.2,angle=270]{N2.eps} & \\includegraphics[scale=0.2,angle=270]{N3.eps} & \\includegraphics[scale=0.2,angle=270]{N4.eps} \\\\ \\includegraphics[scale=0.2,angle=270]{T1.eps} & \\includegraphics[scale=0.2,angle=270]{T2.eps} & \\includegraphics[scale=0.2,angle=270]{T3.eps} & \\includegraphics[scale=0.2,angle=270]{T4.eps} \\\\ \\end{array}$ \\caption{Two simulations of the collapse and fragmentation of a $5.4\\,{\\rm M}_\\odot$ core having the same initial turbulent velocity field. In the first simulation (top row) the core is evolved using the barotropic equation of state. In the second simulation (bottom row) the core is evolved using the new treatment of the energy equation. Each frame shows the logarithm of the column density through the computational domain. For each simulation we show an image at, reading from left to right, $t=67\\,{\\rm kyr}\\,,\\;73\\,{\\rm kyr}\\,,\\;74\\,{\\rm kyr}\\,,\\;{\\rm and}\\;75\\,{\\rm kyr}\\,$. We see that with the new treatment of the energy equation, the disc is more unstable against fragmentation.} \\label{FIG:DISCS} \\end{figure*} ",
        "conclusions": "We have performed a large ensemble of SPH simulations of the collapse and fragmentation of an isolated, turbulent $5.4\\,{\\rm M}_\\odot$ core, with a view to establishing how the statistical properties of the resulting stars are influenced by (i) different initial levels of turbulence, and (ii) different treatments of the thermodynamics. We consider three initial levels of turbulence, characterised by $\\alpha_{_ {\\rm TURB}}=0.05,\\;0.10\\;{\\rm and}\\;0.25$. We treat the thermodynamics firstly with a standard barotropic equation of state, and secondly with a new treatment of the energy equation which captures all the important energy modes of the gas and takes account of radiation transport and variations in the opacity.  Increasing the initial level of turbulence tends to reduce the efficiency of star formation, $\\eta$ (i.e. the fraction of the core mass which is converted into stars after $300\\,{\\rm kyr}$), and to increase the number of stars formed by a single core, ${\\cal N}_{\\!\\star}$, but the effect is very small, and all the other statistical properties of the stars formed are essentially independent of $\\alpha_{_ {\\rm TURB}}$.  A common pattern is observed in which the low--angular-momentum material in the core collapses to form the first star (the primary) after $50\\;{\\rm to}\\;70\\,{\\rm kyr}$ (i.e. just a bit longer than the initial freefall time at the centre of the core), and then a massive disc builds up around the primary. As soon as this circumprimary disc becomes Toomre unstable (which may take from $20\\;{\\rm to}\\;100\\,{\\rm kyr}$), it rapidly breaks up into a clutch of secondary stars. Those secondaries that are quickly ejected from the disc normally end up as brown dwarf stars, whereas those secondaries that stay in the disc -- and particularly those that stay in the inner disc -- tend to grow by accretion, and sometimes they even grow larger than the primary.  Switching from the standard barotropic equation of state to our new more realistic treatment of the energy equation has several systematic effects: \\begin{itemize} \\item{the efficiency of star formation ($\\eta\\equiv\\sum\\left\\{M_{\\!\\star}\\right\\}/M_{_{\\rm CORE}}$) is reduced significantly (by $\\,\\sim\\!15\\%$);} \\item{the number of protostars formed (${\\cal N}_{\\!\\star}$) is greatly increased (by $\\,\\sim\\!40\\%$);} \\item{a higher proportion of brown dwarf stars is formed;} \\item{the mean period of multiple systems is reduced (by a factor $\\sim 3$);} \\item{the orbital eccentricities of multiple systems tend to be higher;} \\item{the mass ratios of multiple systems tend to be higher (i.e. more nearly equal components).} \\end{itemize}  All these trends can be attributed to the fact that the barotropic equation of state assumes that all gas is at the centre of a collapsing spherical $1\\,{\\rm M}_\\odot$ protostar, and therefore it becomes adiabatic at relatively low densities. In contrast, our new more realistic treatment of the energy equation allows the gas in low-mass proto-fragments to remain approximately isothermal to relatively high densities.  These simulations do not capture all the deterministic effects in star formation. In particular, we do not take account of radiative and mechanical feedback from stars, we do not include non-ideal magneto-hydrodynamics effects, and we invoke sink particles (which must prejudice interactions between stars, in favour of elastic dynamical ejections, and against dissipation and mergers). We plan to explore these additional effects in subsequent papers. In particular, we anticipate that by including feedback we can reduce the efficiency of star formation from the very high levels produced here ($\\,\\sim\\!60\\%$) to values more compatible with observation ($\\,\\lesssim\\!30\\%$; Alves et al. 2007; Nutter \\& Ward-Thompson 2007; Goodwin et al. 2008)."
    },
    "0812/0812.3412_arXiv.txt": {
        "abstract": "{Accurate measurement of the frequency-dependent shift of the self-absorbed radio core is required for multi-frequency analysis of VLBI data since absolute positional information is lost as a result of phase self-calibration. We use the cross-correlation technique of Croke \\& Gabuzda (2008) on the optically thin jet emission to align our VLBA images. Our results are consistent with those obtained from the phase-referencing method, as well as alignment by model-fitted optically thin jet components. Physical parameters of the compact jet regions, such as the magnetic field strength ($B$) and the distance of the radio core to the jet origin ($r$), can be calculated from these measurements. For the source Mrk 501, we find a magnetic field strength of $0.15\\pm0.04$ G in the 8.4-GHz core at a distance of $0.8\\pm0.2$ pc from the base of the jet. By extrapolating our 4.6 to 15.4 GHz results for BL Lac (2200+420), we estimate magnetic field strengths of the order of 1~G in the millimetre VLBI core. Using our core-shift measurement between 1.6 and 4.8 GHz for 1803+784, we find $B_{\\rm{core}}(4.8$~GHz$)=0.11\\pm0.02$ G and $r_{\\rm{core}}(4.8$ GHz$)=20\\pm5$ pc. The phase-referencing observations of this source at 8.4 and 43 GHz by Jim\\'enez-Monferrer et al.~(2008) imply $B_{\\rm{core}}(43$ GHz$)=1.0\\pm0.4$~G and $r_{\\rm{core}}(43$ GHz$)=2.0\\pm0.9$ pc.}  \\FullConference{The 9th European VLBI Network Symposium on The role of VLBI in the Golden Age for Radio Astronomy and EVN Users Meeting\\\\ September 23-26, 2008\\\\ Bologna, Italy}  \\begin{document} ",
        "introduction": "Highly collimated jets of relativistic plasma are considered to be generated and propelled outwards from the central regions of Active Galactic Nuclei (AGN) by magnetic forces surrounding the supermassive black hole, see \\cite{meierjapan} and references therein. High resolution VLBI observations can detect the synchrotron emission from these jets in the radio regime at distances of the order of $10^3-10^7$ $R_S$, where $R_S$ is the Schwarzchild radius of the black hole, corresponding to the sub-parsec to parsec scales of the jet \\cite{lobanovevn}. Current theoretical models \\cite{meierscience2001, vlahakiskonigl2004} invoke rotation of the magnetosphere around the compact central object which creates a stiff helical magnetic field that expels and collimates the jet flow. At a distance of several hundred gravitational radii from the central engine, the natural pinching action of the magnetic field can over-collimate the flow resulting in a ``collimation shock'' which can either partially or completely disrupt the helical field structure in the jet. Recent observational evidence suggests that this ``collimation shock'' may occur in the millimetre-wave VLBI core  \\cite{marschernature2008}. A significant number of VLBI observations at cm-wavelenghts support the presence of helical magnetic fields in the parsec scale radio jet (e.g., \\cite{osullivan2008} and references therein), indicating that the initially ordered magnetic field is not completely disrupted and that either remnants of the earlier field structure remain or possibly a current driven helical kink instability is generated \\cite{nakamurameier2004, nakamurajapan, careyhepro}.  The radio cores of AGN are generally flat-spectrum, self-absorbed regions. Since VLBI resolution is usually not sufficient to completely resolve the true optically thick radio core, the core region contains emission from around the optical depth ($\\tau$) $=1$ surface and some contribution from the optically thin inner jet. From the jet model of \\cite{blandfordkonigl1979}, the frequency dependence of the self-absorbed core is described by $r\\propto\\nu^{-1/k_{r}}$, where $r$ is the distance from the central engine and $k_r=((3-2\\alpha)m+2n-2)/(5-2\\alpha)$ with $m$ and $n$ describing the power-law fall-off in the magnetic field strength and particle number density, respectively, with distance from the central engine. For equipartition between the jet particle and magnetic field energy densities, the quantity $k_r=1$, with the choice of $m=1$ and $n=2$ being reasonable (e.g., \\cite{konigl1981, huttermufson1986, lobanov1998}), making $k_r$ independent of the spectral index. Therefore, in equipartition, the core-shift ($\\Delta r=|r_{\\nu_{1}}-r_{\\nu_{2}}|$) between two frequencies ($\\nu_{2}>\\nu_{1}$) is directly proportional to $(\\nu_{2}-\\nu_{1})/(\\nu_{2} \\nu_{1})$. Hence, for accurate analysis of multi-frequency VLBI maps, we must correct for the frequency-dependent position of the self-absorbed radio core. ",
        "conclusions": "Using a cross-correlation technique for aligning the optically-thin jet emission, we obtain accurate measurements of the frequency-dependent shift of the self-absorbed VLBI core of three AGN jets. We have tested our method against both phase-referencing observations and alignment by individual model-fitted steep-spectrum jet components. Using our measurements of the core-shift, we find magnetic field strengths in the compact centimetre jet regions of Mrk 501, BL Lac and 1803+784 of the order of 100's of mG; with values approaching 1 G in the mm-wave VLBI cores of BL Lac and 1803+784. We find that equipartition is a valid assumption for the compact inner jet regions of BL Lac, while we find some deviation from equipartition in the case of Mrk 501. We have also found estimates for distances of the observed VLBI core to the base of the radio jet for each source, providing the approximate location of the central supermassive black hole. Estimates such as these are essential in efforts to test the predictions of current theoretical models on the locations of regions of jet launching, acceleration and collimation.  \\newpage"
    },
    "0812/0812.3909_arXiv.txt": {
        "abstract": "{ We combine the catalogue of eclipsing binaries from the {\\it All Sky Automated Survey} ({\\it ASAS}) with the {\\it ROSAT} All Sky Survey ({\\it RASS}).  The combination results in 836 eclipsing binaries that display coronal activity and is the largest sample of active binary stars assembled to date. By using the ({\\it V-I} ) colors of the {\\it ASAS} eclipsing binary catalogue, we are able to determine the distances and thus bolometric luminosities for the majority of eclipsing binaries that display significant stellar activity.  A typical value for the ratio of soft X-ray to bolometric luminosity is $L_X/L_{\\rm bol}\\sim$ a few $\\times 10^{-4}$, similar to the ratio of soft X-ray to bolometric flux $F_X/F_{\\rm bol}$ in the most active regions of the Sun.  Unlike rapidly rotating isolated late-type dwarfs -- stars with significant outer convection zones -- a tight correlation between Rossby number and activity of eclipsing binaries is absent. We find evidence for the saturation effect and marginal evidence for the so-called ``super-saturation'' phenomena. Our work shows that wide-field stellar variability searches can produce a high yield of binary stars with strong coronal activity.  The combined {\\it ASAS} and {\\it RASS} catalogue, as well as the results of this work are available for download in a form of a file.}  {stars: eclipsing -- stars: binary -- stars: evolution -- stars: X-rays} ",
        "introduction": "In the last few decades, it has become clear that chromospheric and coronal activity of main sequence dwarfs is intimately tied to stellar rotation rate.  Main sequence stars are born rapidly rotating and quickly lose their angular momentum by some combination of stellar winds and magnetic torques.  It follows that stars with rapid rotation -- ideal for studies of stellar activity -- are quite rare.  In comparison to isolated main sequence stars, close eclipsing binaries are rare as well.  However, they are dramatically variable in comparison to rapidly rotating single stars. Close binaries that consist of two late-type dwarfs are thought to be synchronized as long as their orbital periods are shorter than $10$ days or so (Duquennoy \\& Mayor, 1991). This is true for the majority of close binaries, which then may be fruitfully employed in a study of stellar activity.  Contact binary stars were discovered to be strong X-ray emitters almost 30 years ago (Carroll et al. 1980). Shortly afterward Cruddace and Dupree (1984) identified 14 X-ray active W UMa systems and noticed a correlation between X-ray luminosity and rotational period, suggesting that for very fast rotators this relation may be reversed. Since then, the study of W UMa activity has broadened -- X-ray flux variations were discovered (McGale et al. 1996, Brickhouse and Dupree 1998), the color-activity and period-activity dependencies were studied on a bigger sample of 57 W UMa binaries (St\\c epie\\'n et al. 2001). But there is still no final theory describing the mechanism of X-ray emission.  Recently large databases of contact binaries were searched for coronal activity, offering still larger binary samples and thus better statistics. Geske et al. (2006) combined the Northern Sky Variability Survey (NSVS) catalogue with the ROSAT All-Sky Survey X-Ray catalogue and found 140 new binaries with active corona. A part of the All Sky Automated Survey (ASAS) catalogue (southern hemisphere) was also investigated (Chen et al. 2006), and 34 objects more were identified. In 2005 ASAS released the complete catalogue of variable stars south of declination +28deg (Pojma\\'nski et al. 2005), among which there are over 11,000 eclipsing binaries, thus providing the largest sample of galactic field binary stars.  In this short observational paper we combine the complete ASAS catalogue of variable stars (ACVS) with the ROSAT All-Sky Survey X-ray catalogue (RASS). We obtain the largest sample of X-ray active contact binary stars up to date, which constitutes 379 sources. Furthermore, we widen the study of stellar activity to include all types of eclipsing binaries i.e., including semi-detached and detached as well. The overall number of coincident sources amounts to 836. The overwhelming majority of the objects in our sample are close binaries that are expected to be in tidal synchronization. Therefore, we take the orbital period to be equal to the rotation period of each individual star.  Sections 2 and 3 describe both the ASAS and ROSAT catalogues, the process of preparing the combined sample and incidence of X-ray activity. In Section 4 we discuss the method of determining the bolometric and X-ray luminosities. The complete catalogue is presented in Section 5. A short analysis of the evolution of X-ray, or coronal, activity with rotation and colour is given in Section 6. We summarize the paper in Section 7. ",
        "conclusions": "We had produced a catalogue of the combined ASAS eclipsing binaries and RASS X-ray sources. All data had been put in a single file \"ASAS.ROSAT.cat\" which is available for download from the ASAS website (http://www.astrouw.edu.pl/asas/?page=rosat) in the form presented in Table 1 (see Section 5 for details).  The combination of the eclipsing binary catalogue with the RASS produces a total of 836 coincident sources ($\\sim$ a 7\\% yield). Among these are 379 contact binary stars (EC), which is the largest sample of X-ray active contact binary stars assembled up to date. Semi-detached (ESD) and detached (ED) binaries were also taken into account, resulting in 266 and 191 active variables, respectively). We observe similar incidence of X-ray activity in all 3 groups of variables, which is around ~7\\% in the distance unlimited sample.  An overwhelming majority of the coincident sources possesses orbital periods shorter than 10 days, so we expect tidal locking between the primary and the companion.  Therefore, the orbital period of the binary may be also thought of as the rotation period of the stellar components.  We analyzed the dependence of X-ray activity on Rossby number and colour. In comparison to typical isolated stars of similar primary mass, the ASAS eclipsing binaries display a higher level of coronal activity (see Fig.~5) % and  the maximum value $L_X/L_{\\rm bol}\\sim 3\\times 10^{-3}$, for the coincident sources is similar to that found in rapidly rotating isolated late-type dwarfs.  Both the activity - Rossby number (Fig.~6) % and the activity - colour (Fig.~7) % diagrams display a large scatter and they do not tighten any of the already known relations. The coronally-active binaries in our sample display the saturation effect, while there is no clear evidence of the so-called ``super-saturation'' effect.  However, for a given ({\\it V-I} ) colour ECs are rotating more rapidly than the other two classes. And at the same time, for a given colour the level of activity of ECs is lower in comparison to ESDs and EDs. This slight downturn in activity among the three classes may in fact be an indirect manifestation of the saturation effect in Fig.~7.   \\Acknow{We would like to thank K. St\\c epie\\'n for useful comments on the paper. This work was supported by the Polish MNiSW grants N203 007 31/1328 and N N203 304235. AS acknowledges support of a Hubble Fellowship administered by the Space Telescope Science Institute and of a Lyman Spitzer Jr. Fellowship at Astrophysical Sciences at Princeton University.}"
    },
    "0812/0812.1417_arXiv.txt": {
        "abstract": "{} {We study the connection between spatially resolved star formation and young star clusters across the disc of M~51. } {We combine star cluster data based on $B$, $V$, and $I$-band \\emph{Hubble Space Telescope} \\acs\\ imaging, together with new \\wfpc\\ $U$-band photometry to derive ages, masses, and extinctions of 1580 resolved star clusters using SSP models. This data is combined with data on the spatially resolved star formation rates and gas surface densities, as well as H$\\alpha$ and 20~cm radio-continuum (RC) emission, which allows us to study the spatial correlations between star formation and star clusters. Two-point autocorrelation functions are used to study the clustering of star clusters as a function of spatial scale and age.} {We find that the clustering of star clusters among themselves decreases both with spatial scale and age, consistent with hierarchical star formation. The slope of the autocorrelation functions are consistent with projected fractal dimensions in the range of 1.2--1.6, which is similar to other galaxies, therefore suggesting that the fractal dimension of hierarchical star formation is universal. Both star and cluster formation peak at a galactocentric radius of $\\sim$2.5 and $\\sim$5~kpc, which we tentatively attribute to the presence of the 4:1 resonance and the co-rotation radius. The positions of the youngest ($<10$~Myr) star clusters show the strongest correlation with the spiral arms, H$\\alpha$, and the RC emission, and these correlations decrease with age. The azimuthal distribution of clusters in terms of kinematic age away from the spiral arms indicates that the majority of the clusters formed $\\sim5$--20 Myr before their parental gas cloud reached the centre of the spiral arm.} {} ",
        "introduction": "\\label{sec:Introduction}  The increase in the amount of multi-wavelength data during the past few decades has been of tremendous importance for the field of star formation. On the largest scales, these multi-wavelength studies focussed primarily on the tight correlation between global radio-continuum (RC) and far-infrared (FIR) emission of galaxies (e.g. \\citealt{condon92, niklas97a, niklas97b, yun01}) and on the global Schmidt-law, which dictates that there is a powerlaw relation between the area-normalised star formation rate (SFR) and total gas surface density with an index of $1.4\\pm0.15$ \\citep{kennicutt98b}. These correlations have been proven useful in the derivation of global star formation rates from radio-continuum observations (e.g. \\citealt{condon92}), as well as in models of galaxy evolution \\citep{mihos94}.  More recently, these multi-wavelength correlations are exploited to study star formation \\emph{locally} across single galaxies. This is done by either utilizing the RC-FIR correlation \\citep{murphy06, schuster07, murphy08}, or by combining \\emph{Spitzer Space Telescope} FIR observations with H$\\alpha$ and Pa$\\alpha$ data, which allows for a more thorough extinction correction on the derived star formation rates \\citep{calzetti05, kennicutt07}. Combined with new CO and \\ion{H}{I} data, recent studies show that the Schmidt-law is also valid \\emph{locally} across the spiral galaxy M~51 \\citep{schuster07, kennicutt07}.  There is ample evidence that star formation leads to formation of star clusters (\\citealt{lada03}, and references therein). Therefore, star formation across spiral galaxies has been studied extensively through their star cluster populations, either by using ground-based data \\citep{larsen99b, larsen00a}, or by exploiting the higher resolution \\emph{Hubble Space Telescope} (\\hst) observations (e.g. \\citealt{holtzman92, meurer95, whitmore99b, degrijs03d, larsen04b}). Most of these studies, however, focus on the star cluster properties themselves, on the connection between the galaxy's star cluster population and SFR on \\emph{global} scales, or on their spatial correlation \\citep{zhang01}. There is still a lack of evidence, however, on the actual \\emph{fraction} of current star formation taking place in young star clusters (although \\citealt{degrijs03d} find a lower limit of $\\approx 35\\%$ for the Mice galaxies and \\citealt{fall05} find a lower limit of 20\\% in the Antennae). Although it is probably true that most star formation takes place in some sort of embedded, clustered form \\citep{lada03}, it is not clear yet which fraction of these clusters are actually observable as star clusters after they emerge from their natal gas cloud \\citep{bastian06b, goodwin06} and if, or how, this depends on environment. The answers to such questions could give us important clues on the process of star formation and the early evolution of young star clusters.  Before such a quantitative comparison can be made, however, we first need to determine the amount of \\emph{spatial} correlation between star and cluster formation, since any quantitative comparison is only meaningful for regions that are spatially correlated. Besides, determining the spatial correlation between clusters themselves allows us to study the amount and evolution of clustering of star clusters, which provide important clues as to which extent the formation of clusters is hierarchical \\citep[e.g.][]{efremov98, bastian07}.  For M~51 data with a wide wavelength range and good spatial resolution are available. This makes it possible to study, first of all, how well the positions of star clusters are tracing the regions of star formation across this spiral galaxy, and second of all, which fraction of the star formation is taking place in the form of visually observable star clusters. In this study we combine data on the \\emph{resolved} Schmidt-law in M~51 (SFR and gas surface densities, \\citealt{kennicutt07}) as well as archival radio-continuum, 2MASS and \\hst/\\wfpc\\ data with slightly resolved young star clusters from \\citet{scheepmaker07}. These datasets are covering a large fraction of the disc, allowing us to compare star and cluster formation  under a range of conditions, such a low/high gas and stellar surface densities, within a single galaxy consistently.  In this work we will introduce the new datasets used, and present new \\wfpc\\ \\uband\\ photometry. This photometry allows us to add ages, masses and extinctions to the star cluster data of \\citet{scheepmaker07}. In the remainder of this work we will then focus on the \\emph{spatial} correlations between star and cluster formation, for differently aged cluster populations. A \\emph{quantitative} comparison between star and cluster formation, as well as the cluster initial mass function (CIMF), will be the subject of a follow-up paper \\citep{scheepmaker09b}.  In Sect.~\\ref{sec:Observations and photometry} we describe the star cluster data we used, as well as the \\hst/\\wfpc\\ observations. In Sect.~\\ref{sec:Selection of the cluster samples} we estimate ages, masses and extinctions of the star clusters in order to select our cluster samples.  We describe the data on SFRs and gas densities in Sect.~\\ref{sec:Star formation rates, gas densities and spiral arms}, together with our method of tracing the spiral arms. Finally, in Sect.~\\ref{sec:Spatial distribution of star clusters and star formation} we study the spatial distribution of star clusters and star formation in M~51 and we summarize our main results in Sect.~\\ref{sec:Summary and conclusions}. ",
        "conclusions": "\\label{sec:Summary and conclusions}  We have studied the spatial correlations between star clusters and star formation across the spiral galaxy M~51. We combined published data of the resolved star formation law in M~51 (SFR and gas surface densities, \\citealt{kennicutt07}) with the star cluster data of \\citet{scheepmaker07}. We combined the star cluster data with new \\hst/\\wfpc\\ \\uband\\ data of 1580 clusters in order to estimate ages, masses and extinctions of these clusters using \\emph{SSP} models. We used radial and azimuthal distributions and cross-correlation functions to study the spatial correlation between star formation, star clusters and spiral structure. We used auto-correlation functions to study how the clustering of star clusters depends on age and spatial scale. The main conclusions of this work can be summarized as follows:  \\begin{enumerate} \\item We find that age determinations of young ($\\lesssim100$~Myr) star clusters based on broad-band UBVI photometry can depend strongly on the adopted stellar isochrones used in the SSP models. Using models based on the Padova isochrones, we generally find older, and therefore more massive, clusters compared to the same models based on the Geneva isochrones (Fig.~\\ref{fig:age_vs_age}).  \\item The radially averaged distributions of star formation and gas surface density in M~51 peak in the centre and subsequently at $\\sim2.5$ and $\\sim5$~kpc. These distributions are matched by the number surface density distribution of our youngest star cluster sample (ages $<10$ Myr, Fig.~\\ref{fig:radial distribution}), while for the older clusters the distributions are more flat. We note that the location of the peaks coincide approximately with the locations where the spiral arms meet the 4:1 resonance and the co-rotation radius.  \\item The distributions of azimuthal distances from the spiral arms show that most of the young \\emph{and old} star clusters are located at the centre of the spiral arms, with the youngest star clusters showing the most smooth distribution (Fig.~\\ref{fig:azimuthal}). In terms of kinematic age and for $R_{\\mathrm{G}} < 4$~kpc, the azimuthal distribution of the oldest clusters peaks at the largest separation downstream of the spiral arms (10--15~Myr), while the majority of the younger clusters are within 0--5~Myr downstream of the arms. This indicates that the majority of the clusters formed $\\sim5$--20~Myr before their parental gas cloud would have reached the centre of the spiral arm.  \\item We find that the amount of clustering of star clusters decreases both with spatial scale and age (Fig.~\\ref{fig:autocorrelation}), consistent with hierarchical star formation. The slope of the autocorrelation functions of the youngest clusters ($\\log(\\mathrm{age}) < 7.5$) is consistent with a projected fractal dimension of 1.6. We find a lower fractal dimension of 1.2 for the clusters with $\\log(\\mathrm{age}) < 7.0$, if we use a higher mass limit and thus, for this smaller mass range, a more complete sample. On spatial scales smaller than 400~pc, the amount of clustering of the \\emph{oldest} cluster samples is very uncertain due to uncertainties in the assumed SSP model parameters.  \\item The locations of the youngest star clusters are highly correlated with the flux from 20~cm radio-continuum emission and continuum subtracted H$\\alpha$ emission, but not with near-infrared $H$-band emission. The correlation with the radio emission is more diffuse compared to the correlation with H$\\alpha$, and both correlations decreases with the age of the cluster population. \\end{enumerate}  In a follow-up paper \\citep{scheepmaker09b} we will focus on the \\emph{quantitative} comparison between star and cluster formation, discussing the fraction of total star formation taking place in optically visible star clusters."
    },
    "0812/0812.2023_arXiv.txt": {
        "abstract": "We present statistical analysis of diffuse Galactic synchrotron emission and polarized thermal emission from dust. Both Galactic synchrotron emission and polarized thermal emission from dust reflect statistics of magnetic field fluctuations and, therefore, Galactic turbulence. We mainly focus on the relation between observed angular spectra and underlying turbulence statistics. Our major findings are as follows. First, we find that magnetohydrodynamic (MHD) turbulence in the Galaxy can indeed explain diffuse synchrotron emission from high galactic latitude. Our model calculation suggests that either a one-component extended halo model or a two-component model, an extended halo component (scale height $\\gtrsim 1kpc$) plus a local component, can explain the observed angular spectrum of the synchrotron emission. However, discrete sources seem to dominate the spectrum for regions near the Galactic plane. Second, we study how star-light polarization is related with polarized emission from thermal dust. We also discuss the expected angular spectrum of polarized emission {}from thermal dust. Our model calculations suggest that $C_l\\propto l^{-11/3}$ for $l\\gtrsim 1000$ and a shallower spectrum for $l\\lesssim 1000$. ",
        "introduction": "Diffuse Galactic synchrotron emission is an important foreground source of the cosmic microwave background (CMB) signals. Therefore proper understanding of Galactic synchrotron emission is essential for CMB studies. In this paper, we present statistical analysis of a synchrotron foreground emission map. In particular, we provide physical interpretation of the angular spectrum of the synchrotron intensity. In addition, we briefly discuss the properties of a model dust emission map, another important ingredient of the CMB foreground. As measurements of the polarized CMB signals become possible, more accurate  removal of Galactic polarized foreground emission is required. Synchrotron and dust emissions are known to be the most important sources of polarized foreground radiation. Therefore, measurements of angular power spectra of such foregrounds are of great interest. In this paper, we also provide estimation of angular spectrum of polarized emission by foreground dust.  The angular spectrum of the synchrotron emission delivers valuable information on the structure of the Galaxy. The observed spectra of synchrotron emission and synchrotron polarization (see papers in de Oliveira-Costa \\& Tegmark 1999) reveal a range of power-laws. Since the Galactic synchrotron emissivity is roughly proportional to the magnetic energy density \\footnote{ In fact, correct scaling is $\\epsilon({\\bf r}) \\propto n(e) |{B}_{\\bot}|^{\\gamma}$, where $\\epsilon({\\bf r})$ is the synchrotron emissivity, $n(e)$ is the cosmic-ray electron number density, and $B_{\\bot}$ is the strength of the magnetic field component perpendicular to the line of sight. In our Galaxy, we have $\\gamma \\sim 2$. Therefore, when $n(e)$ is more or less uniform, we have $\\epsilon({\\bf r}) \\propto  |{B}_{\\bot}|^{2}$.}, angular spectrum of synchrotron emission reflects statistics of magnetic field fluctuations in the Galaxy (see \\S\\ref{sect:synchem} for further discussions).  The interstellar medium (ISM) is turbulent and Kolmogorov-type spectra were reported on the scales from several AUs to several kpc (see Armstrong, Rickett, \\& Spangler 1995; Lazarian \\& Pogosyan 2000; Lazarian 2009). It is believed that magnetic field lines are twisted and bend by turbulent motions in the Galaxy. Therefore it is natural to think of the turbulence as the origin of the magnetic field fluctuations and thus the diffuse synchrotron foreground radiation. Indeed, several earlier studies addressed this issue. Tegmark et al.~(2000) suggested that the spectra may be relevant to Kolmogorov turbulence. Chepurnov (1999) and Cho \\& Lazarian (2002; hereinafter CL02) used different approaches, but both showed that the angular spectrum of synchrotron emission reveals Kolmogorov spectrum ($C_l\\propto l^{-11/3}$) for large values of multipole $l$. However, they noted that the spectrum can be shallower than the Kolmogorov one for intermediate values of multipole $l$, due to density stratification in the halo (Chepurnov 1999) or the Galactic disk geometry (CL02). In this paper, we further elucidate the relation between 3-dimensional turbulence spectrum and observed angular spectrum. We also investigate how structure of the Galactic halo affects observed angular spectrum.  Thermal emission from dust is also an important source of foreground emission. Dust emission is the most pronounced emission in far infrared (FIR) wavelengths. Schlegel, Finkbeiner, \\& Davis (1998) combined 100$\\mu m$ maps of IRAS (Infrared Astronomy Satellite) and DIRBE (Difuse Infrared Background Experiment on board the COBE satellite) and removed the zodiacal foreground and point sources to construct a full-sky map. Finkbeiner, Davis, \\& Schlegel (1999) extrapolated the 100$\\mu m$ emission map and 100/240$\\mu m$ flux ratio maps to sub-millimeter and microwave wavelengths. In this paper, we present statistical analysis of a model dust emission map above.  Thermal radiation from dust becomes polarized when dust grains are aligned. There is ample evidence that dust grains are aligned with respect to magnetic field. Therefore, polarization by dust is also related to the magnetic field fluctuations. Then how are they related? Polarized dust emission is difficult to observe directly. Therefore in this paper we first study the relation between star-light polarization and turbulence statistics. Then we describe how we infer the relation between star-light polarization and polarized thermal emission from dust.  In this paper, we show how MHD turbulence is related with the observed angular spectra of Galactic synchrotron emission  and starlight polarization. For this purpose, we use an analytical insight obtained in Lazarian (1992, 1995ab) and numerical results obtained in CL02. We also discuss  how we can estimate polarized microwave dust emission using star-light polarimetry. This problem is of great importance in view of recent interest to the foregrounds to the CMB polarization. In \\S2, we briefly describe the data sets we use in this paper, the 408MHz Haslam map and a model dust emission map at 94GHz. In \\S3, we review a simple model of the angular spectrum of synchrotron emission arising from homogeneous MHD turbulence. In \\S4, we present statistical analysis of the Haslam map, which is dominated by diffuse Galactic synchrotron emission. In \\S5, we investigate the dust emission map. In \\S6, we consider polarized emission from thermal dust. In \\S7, we discuss how to utilize our results to remove Galactic foregrounds. In \\S8, we calculate high-order structure functions of the synchrotron and the dust maps and we compare the results with those of turbulence. Finally in \\S9, we give summary. ",
        "conclusions": "In this paper, we have discussed the relation between 3D spatial spectrum and the observed angular spectrum (or the second-order angular structure function) of the CMB foreground emissions. We have focused on synchrotron total intensity and polarized thermal dust emission. Our current study, as well as earlier studies (Chepurnov 1999; CL02), predicts that $C_l$ will reveal true turbulence spectrum on small angular scales. Then, on what scales will we see the turbulence spectrum? Our model calculations that take into account stratification effects imply that \\begin{enumerate} \\item $\\theta <$ a few times $0.1^\\circ$ (or $l>$ a few times $100$) for synchrotron emission (see Fig.~\\ref{fig:model}), and \\item $\\theta \\lesssim 0.1^\\circ$ (or $l\\gtrsim 1000$) for polarized emission {}from thermal dust (see Fig.~\\ref{fig:starpol}). \\end{enumerate} On larger angular scales, spectra are expected to be shallower.  Then, how is the turbulence spectrum related to the angular spectrum (or the second-order angular structure function) on small scales? When the 3D spatial spectrum has a power-law spectrum, $k^{-m}$, the observed angular spectrum and the second-order angular structure function will be \\bea C_l & \\propto & l^{-m}   \\mbox{~~~ and} \\\\ D_2(\\theta) & \\propto & \\theta^{m-2} \\mbox{~~~~ if $\\theta \\ll L/d_{max}$}, \\eea where $L$ is the outer scale of turbulence and $d_{max}$ is the distance to the farthest eddy in case of homogeneous turbulence. We can interpret $d_{max} \\sim z_0$ if there is a stratification with scale height $z_0$ (see Fig.~\\ref{fig:model}). When the angular separation is large, we have a universal scaling for the angular correlation function: \\be K(\\theta) \\propto  \\frac{ \\pi -\\theta }{ \\sin\\theta } \\sim \\frac{1}{\\theta} \\mbox{ ~~~if $\\theta \\gtrsim L/d_{max}$} \\ee In this limit, the structure function is roughly constant and the angular spectrum is roughly proportional to $l^{-1}$.  In this paper, we have analyzed Haslam 408Mhz map and a model dust emission map and compared the results with model calculations. We have found that \\begin{enumerate} \\item The Haslam map for high galactic latitude ($b>30^{\\circ}$) can be explained by MHD turbulence in the Galactic halo. The measured second-order angular structure function is proportional to $\\theta^{1.2}$, which corresponds to an angular spectrum of $l^{-3.2}$. The high-order statistics for high galactic latitude ($b>30^{\\circ}$) is consistent with that of incompressible magnetohydrodynamic turbulence. Our model calculations show that a two-component model (see \\S\\ref{sect:3models} and Fig.~\\ref{fig:model}) can naturally explain the observed angular spectrum. The one-component model can also explain the observed slope. But, the slope of the spectrum shows a more abrupt change near $l\\sim 30$.  \\item The model dust emission map may not have anything to do with turbulence on large angular scales. That is, we do not find signatures of turbulence in the map.  \\item Both maps show flat high-order structure functions for the Galactic plane. This kind of behavior is expected when discrete structures dominate the map. \\end{enumerate}  We have described how we can obtain angular spectrum of polarized emission from thermal dust in high galactic latitude regions. Our model calculations show that starlight polarization arising from dust in high galactic latitude regions will have a Kolmogorov spectrum, $C_l\\propto l^{-11/3}$, for $l\\gtrsim 1000$ and a shallower spectrum for $l\\lesssim 1000$ (Fig.\\ref{fig:starpol}). We expect that polarized emission from the same dust also has a similar angular spectrum. That is, we expect that the angular spectrum of polarized emission from thermal dust is close to a Kolmogorov one for $l\\gtrsim 1000$."
    },
    "0812/0812.0817_arXiv.txt": {
        "abstract": "We present numerical simulations of driven magnetohydrodynamic (MHD) turbulence with weak/moderate imposed magnetic fields. The main goal is to clarify dynamics of magnetic field growth. We also investigate the effects of the imposed magnetic fields on the MHD turbulence, including, as a limit, the case of zero external field. Our findings are as follows. First, when we start off simulations with weak mean magnetic field only (or with small scale random field with zero imposed field), we observe that there is a stage at which magnetic energy density grows linearly with time. Runs with different numerical resolutions and/or different simulation parameters show consistent results for the growth rate at the linear stage. Second, we find that, when the strength of the external field increases, the equilibrium kinetic energy density drops by roughly the product of the rms velocity and the strength of the external field. The equilibrium magnetic energy density rises by roughly the same amount. Third, when the external magnetic field is not very strong (say, less than $\\sim 0.2$ times the rms velocity when measured in the units of Alfven speed), the turbulence at large scales remains statistically isotropic, i.e. there is no apparent {\\it global} anisotropy of order $B_0/v$. We discuss implications of our results on astrophysical fluids. ",
        "introduction": "Most astrophysical fluids are magnetized. Magnetic field in an astrophysical system can be divided into two components: large-scale regular and small-scale random components. The generation of magnetic field may involve with two separate issues: generation of the large-scale regular field (or mean field) and generation of small-scale random field.  The generation or growth of large-scale regular fields is an important topic in astrophysics. However, in this paper we assume fixed large-scale regular fields are already present and we only investigate how small-scale random fields are generated from the imposed large-scale fields. That is, we investigate the growth of magnetic energy in the presence of fixed mean fields. Therefore, the generation/growth of large-scale regular field itself is not a topic of this paper. Readers interested in the topic may refer to mean-field dynamo theories (see Moffatt 1978; Parker 1979; Krause \\& Radler 1980; Brandenburg \\& Subramanian 2005).  % The type of magnetic energy growth we deal with in this paper is so-called small-scale {\\it turbulence dynamo}\\footnote{ In this paper, turbulence dynamo means {\\it not} growth of mean field itself {\\it but} generation of random field on scales similar to or smaller than the driving scale of turbulence, which is sometimes referred to as fluctuation dynamo (see, for example, Schekochihin et al. 2007; Subramanian 2008). In all our simulations the strengths of the mean fields do not change with time. In this paper, by `amplification of magnetic field' or `growth of magnetic field', we actually mean amplification of fluctuating field(s) on scales similar to or smaller than the driving scale. Therefore, the small-scale turbulence dynamo is different from mean-field dynamo, which deals with growth of mean field itself. }, the origin of which can be traced back to the papers by Batchelor (1950) and Kazantsev (1968). We also study how properties of magnetohydrodynamic (MHD) turbulence (e.g. energy densities, power spectra, global anisotropy) change as the strength of the mean field changes.  When we introduce a mean field in a turbulent medium, the mean field interacts with turbulent motions. There are two distinct types of interaction, depending on the strength of the mean magnetic field. When the large-scale regular magnetic field is weak, turbulent motions stretch the magnetic field lines and, as a result, the magnetic energy density increases. MHD turbulence near the scale of the largest energy-containing eddies (i.e, the outer scale or the energy injection scale) will be more or less like ordinary hydrodynamic turbulence with small magnetic back reaction.  MHD turbulence in intracluster medium and intergalactic medium may fall in this type of turbulence. The origin of the seed magnetic fields is still uncertain. However, whatever the origin is, turbulence motions can produce small-scale field through the stretching of the seed fields. In this regime, the growth timescale will be of great importance. If the growth time scale is shorter than the age of the universe in a system, we expect that the system is strongly magnetized in the present time. Otherwise, we expect that the system is weakly magnetized.   On the other hand, when the imposed mean magnetic field is strong in the sense that the turbulent eddy turnover rate at the large scale (i.e. $L/v$) is  slower than the Alfvenic rate of the same scale (i.e. $L/B_0$)\\footnote{In what follows we measure the magnetic field in the units of Alfven velocity $V_A$.}, the resulting turbulence can be described through the nonlinear interaction of waves. The turbulence can be either {\\it strong}, meaning that the cascading happens within one eddy turnover time or {\\it weak}, meaning that the cascading takes more than one eddy turnover time. A classical study of Iroshnikov (1963) and Kraichnan (1965) presents an example of weak MHD turbulence. This is a hypothetical {\\it isotropic} MHD turbulence, while we know by now that the actual MHD turbulence is anisotropic (Shebalin et al 1983; Higdon 1986; Goldreich \\& Sridhar 1995; Matthaeus et al. 1998; see also Biskamp 2003 and references therein). An exact analytical treatment of weak MHD turbulence can be found in Galtier et al.~(2000). A successful\\footnote{The Goldreich \\& Sridhar model was successfully tested in incompressible 3D MHD simulations in Cho \\& Vishniac (2000b), Maron \\& Goldreich (2001), Cho, Lazarian \\& Vishniac (2002), as well as in compressible 3D MHD simulations in Cho \\& Lazarian (2002, 2003). While the particular points of the model, i.e. the particular slope of the spectrum are still the subject of debates (see Maron \\& Goldreich 2001; Boldyrev 2005, 2006; Beresnyak \\& Lazarian 2006; Mason, Cattaneo \\& Boldyrev 2006), the corner stone of the model, which is the critical balance, stays untouched.}  model of the strong MHD turbulence was presented in  Goldreich \\& Sridhar (1995).   In this paper we deal with the former case: turbulence with weak imposed magnetic field. In the interstellar medium community, this type of turbulence is called super-Alfvenic turbulence, which is favored by some researchers (see Padoan \\& Nordlund 1999; Padoan et al. 2004) as a model of turbulence in molecular clouds. In any case, such turbulence is expected to be present in any system with magnetic field below the equipartition value, which gets subject to intensive driving.  As the turbulence kinetic energy decreases with the scale, we expect the Goldreich-Sridhar (1995) model to be valid at some small scale when the equipartition is reached, while at larger scales we expect hydrodynamic motions to increase the energy of magnetic field.  It is a common knowledge that the effects of mean magnetic field have important astrophysical implications. In recent years, interest on the turbulent processes in tangled magnetic field has been growing. Relevant astrophysical problems include thermal diffusion in the intracluster medium (see Chandran \\& Cowley 1998; Narayan \\& Medvedev 2001; Cho et al. 2003; Lazarian 2006), cosmic ray propagation (see  Cassano \\& Brunetti 2005; Lazarian \\& Beresnyak 2006; Brunetti \\& Lazarian 2007) as well as star formation (Padoan et al. 2004; Li \\& Nakamura 2004; Vazquez-semadeni, Kim \\& Ballesteros-Paredes 2005).   Cho \\& Vishniac (2000a) numerically showed that magnetic energy grows until the magnetic energy density gets comparable to the kinetic energy density (see also Kulsrud \\& Anderson 1992; Kulsrud et al.~1997). In this paper, we present more comprehensive studies on the topic. Other aspects of the magnetic field generation also require further studies. For example, Haugen \\& Brandenburg (2004) discussed spectral change of MHD turbulence by mean field and showed that imposed magnetic field lowers the spectral magnetic energy in the inertial range. Mac Low (1999) demonstrated that mean magnetic field produces anisotropic structures along the mean field direction in strongly compressible MHD turbulence. Lee et al. (2003) discussed the behaviors of energy densities and spectral shapes for three different cases (very weak, weak, and strong mean magnetic field cases) but only for 2-dimensional  MHD turbulence. They showed the flow character in very weak field classes is similar to that of hydrodynamic turbulence, while the strong field cases show spectra shallower than the hydrodynamic one. Schekochihin et al. (2007) discussed the effects of mean fields for different values of magnetic Prandtl numbers. In all the papers above, the increase of magnetic field energy was noticed. In this paper, we present a comprehensive study on the effect of mean magnetic field on 3-dimensional MHD turbulence and turbulence dynamo.   Another issue that requires clarification is the effect of mean magnetic field on the decay of MHD turbulence. Stone, Ostriker, \\& Gammie (1998) and Mac Low et al. (1998) numerically showed that damping time-scales of compressible MHD turbulence are comparable to the large-scale eddy turnover time (see McKee \\& Ostriker 2007 for a collection of related results). Incompressible MHD turbulence also decays fast (see Cho et al. 2002). While the earlier works were mostly focused on two extreme limits - zero (see, for example, Biskamp \\& Muller 2000) and strong mean-field limits. In what follows we discuss how the decay time-scale changes as the mean field strength changes.  We will first consider the regime of very weak mean field. In this regime, we will mainly investigate the growth rate of magnetic field. Then, we will consider the effect of intermediate mean fields. We will investigate how magnetic and kinetic energy densities, anisotropy, and growth rate of magnetic field change with the increase of the strength of the mean field. In this paper, we deal with incompressible MHD turbulence. We describe our numerical methods in \\S 2, present our results for the very weak mean-field regime in \\S 3 and discuss the effects of the mean field in \\S 4. We compare MHD and hydrodynamic turbulence in \\S 5 and we give discussion in \\S 6. We provide our conclusions in \\S 7. ",
        "conclusions": "We have studied the growth of magnetic field in the presence of weak/moderate mean magnetic fields. We have considered only the cases of $\\nu=\\eta$. We have found that equipartition between the kinetic and magnetic energy densities occurs at a scale somewhat smaller than the kinetic energy peak, which is consistent with previous results with lower numerical resolution by Cho \\& Vishniac (2000a). We have found that runs with different numerical resolutions and simulation parameters show consistent results for the slope of the linear growth stage. In real astrophysical flows with very weak mean field and very large Reynolds number, the growth of magnetic energy will follow \\begin{equation} \\frac{ B^2(t) }{ 2 E_{turb} } \\sim 0.033 \\frac{ t }{ L/\\sqrt{v^2+B^2} }. \\end{equation}  We have also studied the effects of external magnetic fields on MHD turbulence. We have shown that magnetic energy density increases and kinetic energy density decreases as the external field becomes stronger. To be specific, we have shown that, when $B_0 \\leq 0.2$, \\begin{eqnarray} b^2 - b^{(0)2} \\propto vB_0 , \\label{conbsq}\\\\ v^2 - v^{(0)2} \\propto -vB_0, \\label{convsq} \\end{eqnarray} where $b^2$ is the magnetic energy density in the presence of the external field and $b^{(0)2}$ is the value when the external field is zero.  When the external magnetic field is not very strong ($\\sim$ less than 0.2 times the rms velocity), the turbulence remains statistically isotropic, i.e. there is no apparent anisotropy of order $B_0/v$. The slope of kinetic spectra at small k's is a function of $B_0$: when $B_0$ is small, kinetic spectra should be steeper than $k^{-5/3}$ and those for larger $B_0$ are similar to $k^{-5/3}$.  We have compared MHD and hydrodynamic turbulence. Spectra of total energy in MHD turbulence are similar to those of hydrodynamic turbulence. The Kolmogorov constant in MHD cases can be higher or lower than that of hydrodynamic turbulence, depending on the strength of the mean field."
    },
    "0812/0812.2812_arXiv.txt": {
        "abstract": "{% I will review some of the developments in studies of the host galaxy properties of Compact Steep Spectrum (CSS) and GigaHertz-Peaked Spectrum (GPS) radio sources. In contrast to previous reviews structured around observational technique, I will discuss the host galaxy properties in terms of  morphology, stellar content and warm gas properties and discuss how compact, young radio-loud AGN are key objects for understanding galaxy evolution. } ",
        "introduction": "In recent years it has become increasingly clear that AGN play a key role in galaxy evolution, in particular the way in which AGN and its host galaxy interact (e.g. Silk \\& Rees 1998; Fabian 1999; di Matteo et al 2005).  A key time to study the impact the AGN will have on its host galaxy, and how the properties of the host galaxy will influence the  AGN, is during the early stages of evolution. As discussed in the papers listed above, among others, after relocating to the centre of the galaxy after the merger, the central black hole will grow rapidly through merger-induced accretion, eventually `switching-on' and becoming a quasar. The merger will deposit large quantities of gas and dust into the nuclear regions, essential to fuel the young AGN, and this must eventually be shed through outflows and winds. However, much of this initial, rapid growth phase will be  obscured from view at optical wavelengths by the dense natal cocoon.  Cue the Compact Steep Spectrum (CSS) and GigaHertz-Peaked Spectrum (GPS) radio sources. Now believed to be small due to evolutionary stage, rather than confined and frustrated old sources (e.g. Fanti et al. 1990; Fanti et al. 1995; Owsianik et al. 1998; Murgia et al. 1999), the compact radio sources provide a unique tool for pinpointing young, recently  triggered AGN.  Until recently, most effort has focused on the AGN and it's radio jets, particularly in the radio band. However, the host galaxies can provide a number of key results, important for both understanding compact radio sources themselves and how they fit in to the bigger picture. Through deep optical spectroscopy with 4-m and 8-m class telescopes, it is possible to study both the emission lines and the near UV-optical continuum emission to search for signatures of AGN-host galaxy interaction and probe the stellar populations \\& star formation history.  In these proceedings, I give a brief overview of some of the recent developments in studies of compact radio source host galaxy properties. Instead of focussing on observation wavebands and techniques, we now know enough to discuss the various properties in terms of physically meaningful groups, namely the overall system morphology, the stellar and gas content and evidence for AGN-feedback. ",
        "conclusions": ""
    },
    "0812/0812.1306.txt": {
        "abstract": "We add the effect of turbulent viscosity via the $\\alpha-$prescription to models of the self-consistent formation and evolution of protostellar discs. Our models are non-axisymmetric and carried out using the thin-disc approximation. Self-gravity plays an important role in the early evolution of a disc, and the later evolution is determined by the relative importance of gravitational and viscous torques. In the absence of viscous torques, a protostellar disc evolves into a self-regulated state with disk-averaged Toomre parameter $Q \\sim 1.5-2.0$, non-axisymmetric structure diminishing with time, and maximum disc-to-star mass ratio $\\xi = 0.14$. We estimate an effective viscosity parameter $\\alpha_{\\rm eff}$ associated with gravitational torques at the inner boundary of our simulation to be in the range $10^{-4}-10^{-3}$ during the late evolution. Addition of viscous torques with a low value $\\alpha = 10^{-4}$ has little effect on the evolution, structure, and accretion properties of the disc, and the self-regulated state is largely preserved. A sequence of increasing values of $\\alpha$ results in the discs becoming more axisymmetric in structure, being more gravitationally stable, having greater accretion rates, larger sizes, shorter lifetimes, and lower disc-to-star mass ratios. For $\\alpha=10^{-2}$, the model is viscous-dominated and the self-regulated state largely disappears by late times. The axisymmetry and low surface density of this model may contrast with observations and pose problems for planet formation models. The use of $\\alpha=0.1$ leads to very high disc accretion rates and rapid (within 2 Myr) depletion of the disc, and seems even less viable observationally. Furthermore, only the non-viscous-dominated models with low values of $\\alpha = 10^{-4}-10^{-3}$ can account for an early phase of quiescent low accretion rate $\\dot{M} \\sim 10^{-8}~M_\\odot$~yr$^{-1}$ (interspersed with accretion bursts) that can explain the recently observed Very Low luminosity Objects (VeLLOs). We also find that a modest increase in disc temperature caused by a stiffer barotropic equation of state ($\\gamma=1.67$) has little effect on the disc accretion properties averaged over many disc orbital periods ($\\sim 10^4$~yr), but can substantially influence the instantaneous mass accretion rates, particularly in the early embedded phase of disc evolution. ",
        "introduction": "We have recently demonstrated that disc gravity plays an important role not only in the early embedded phase of disc evolution but also in the late accretion phase \\citep{VB3,VB4}. In the early embedded phase, when the infall of matter from the surrounding envelope is substantial, mass is transported inward by the gravitational torques from spiral arms that are a manifestation of the envelope-induced gravitational instability in the disc. In the late accretion phase, when the gas reservoir of the envelope is depleted, the distinct spiral structure is replaced by ongoing irregular nonaxisymmetric density perturbations. These perturbations produce a residual nonzero gravitational torque in the disc. In particular, the net gravitational torque in the inner disc tends to be negative during first several million years of the evolution, while the outer disc has a net positive gravitational torque.  This is a fundamental property of self-gravitating circumstellar discs around low-mass stars. There is also an overall net negative torque in the disc that is related to the removal of angular momentum by gas that is accreted in to the central object. Although we do not model the gas flow within a central sink of size 5 AU, the angular momentum that is carried in to this region may be lost to the system via additional processes including magnetic braking and outflows, thereby allowing the gas to reach the central star.  In this paper we seek to determine the effect of {\\it other} mechanisms of radial mass and angular momentum transport on the secular evolution of self-consistently formed circumstellar discs. It is well known that standard collisional viscosity (molecular viscosity) is negligible in application to circumstellar discs. The best candidate to date is turbulent viscosity induced by the magneto-rotational instability (MRI) \\citep{BH}, though other mechanisms such as nonlinear hydrodynamic turbulence cannot be completely eliminated due to the large Reynolds numbers involved \\citep[e.g.][]{Afshordi}. We make no specific assumptions about the source of turbulence and parameterize the magnitude of turbulent viscosity using the usual $\\alpha$-prescription \\citep{SS} \\begin{equation} \\nu=\\alpha \\,c_{s}\\,H, \\label{alpha} \\end{equation} where $c_{s}$ is the sound speed and $H$ is the disc scale height. The most appropriate value of the parameter $\\alpha$ in circumstellar discs is uncertain. Both the shearing box and global numerical simulations of the MRI tend to yield the values of $\\alpha$ that vary significantly from model to model and range between $10^{-4}$ and ${\\rm a~few} \\times 10^{-1}$ \\citep[e.g.][]{Hawley95, Fleming,Brandenburg96,Stone96,Armitage}. Motivated by these studies, we have adopted a similar range of values for $\\alpha$. We also assume that $\\alpha$ is spatially and temporally constant, i.e. it represents a mean value, time-averaged over many orbital periods of the disc. Radial variations in $\\alpha$ may (and should) be present in the disc but it requires a more %EV: eliminate \"(i.e. ionization balance)\" below thorough consideration of the disc physics and is left for a follow-up paper.  Our study of the self-consistent formation and long-term evolution of circumstellar discs has been preceded by numerical simulations of \\citet{Lin,Nakamoto95,Hueso05} and others. However, our work is different in one important aspect -- we employ a fully two-dimensional numerical hydrodynamics simulations in the thin-disc approximation (in contrast to earlier one-dimensional studies). As a result, we account for disc self-gravity self-consistently by solving the Poisson integral \\citep[see][]{VB2} and need not parameterize gravitational torques in terms of the $\\alpha$-prescription. Use of the thin-disc approximation is, unfortunately, a necessity. Fully three-dimensional numerical simulations are too computationally expensive to study the disc evolution on time scales of several Myr.   The paper is organized as follows. Section~\\ref{model} gives a brief description of model equations and initial conditions. The main results are presented in Section~\\ref{results}. The observational implications and comparison with earlier studies are discussed in Section~\\ref{discuss}. The main conclusions are summarized in Section~\\ref{summary}. ",
        "conclusions": "\\label{discuss}  \\subsection{Constraints on turbulent viscosity}  There are several known physical processes that can contribute to the radial transport of mass and angular momentum in circumstellar discs. These transport mechanisms include gravitational torques, either due to internal (self-gravity) or external (companion star) sources, and turbulent viscosity. The effect of self-gravity on the secular evolution of a circumstellar disc has been considered in our previous paper \\citep{VB3}. In this paper we added the effect of turbulent viscosity due to yet unspecified source of turbulence. Physical processes in circumstellar discs that can, under specific conditions, give rise to turbulence include the magneto-rotational instability and vertical convection. The latter may be heavily suppressed in the externally heated discs \\citep{Ruden}, which leaves the MRI as the most probable source of turbulence in circumstellar discs. \\cite{Afshordi} have recently argued on analytical grounds that hydrodynamic turbulence (not related to the MRI) can arise in cold Keplerian discs characterized by Reynolds numbers $\\ga10^4$ and having both the fine-tuned conditions and appropriate feedback mechanism, though this idea has not been proven so far by numerical simulations.  Many numerical efforts to characterize the MRI-induced turbulent viscosity qualitatively have been made in the recent years. These include local shearing box models \\citep[e.g.][]{Brandenburg96,Stone96,Fleming,Sato04} and global magnetohydrodynamic simulations of accretion discs \\citep[e.g.][]{Armitage}, though the latter are usually limited by a small number of orbital periods. Most authors calculate  the mean ratios of the Maxwell and Reynolds stresses versus midplane gas pressure, $T^{B}_{r\\phi}/P_0$ and $T^{Re}_{r\\phi}/P_0$, respectively, which, to a factor of unity, are proportional to $\\alpha$ in Keplerian discs\\footnote{In fact, this is true %EV: T^{B,R} - > T^{B} below, to agree with rest of text. Is this OK? Please check. only for axisymmetric discs, in which $T^{B}_{r\\phi}=(d \\log\\Omega /d \\log r) \\alpha P_0$ and the rate of strain $d \\log\\Omega /d \\log r$ is constant (e.g. -3/2 in Keplerian discs). In non-axisymmetric discs the relation between $T^{B}_{r\\phi}$ and $P_0$ is more complex.}.   The derived values of $\\alpha$ vary significantly between the studies. For instance, \\citet{Hartmann98} derived $\\alpha\\approx0.01$ using observed disc sizes and a simple model for the evolution of an axisymmetric viscous disc, in which viscosity is a power-law function of radius. \\citet{Fleming} have employed a stratified model of accretion discs in a shearing box approximation, in which the upper layers are MRI-active while the central regions are quiescent. They found that $T^{B}_{xy}/P_0$ has a mean value of about $10^{-3}$ in the MRI-active upper layers but drops to a negligible value in the midplane. On the other hand, $T^{Re}_{xy}/P_0$ in the vertical direction is roughly constant at a few $\\times 10^{-4}$. This implies mean values of $\\alpha$ in the range $10^{-4}-10^{-3}$. Low mean values of $\\alpha$ were also reported by Brandenburg et al. ($\\alpha=0.007$) and Stone et al. ($\\alpha\\la0.01$). On the other hand, global numerical simulations tend to yield larger values for $\\alpha$. For instance, \\citet{Armitage} found mean values between 0.05 and 0.1. %while \\citet{Hawley} argued that $\\alpha$ exhibits both systematic gradients %and large fluctuations, from $10^{-2}$ in the disc midplane at large radius to 10 (!) %at a few gas density scale heights above the midplane.  Large variations in $\\alpha$, both in space and time, imply that the development of the magneto-rotational instability is strongly dependent on the local conditions in the disc. Nevertheless, we can still learn about circumstellar discs from a simple model used in our paper, if we assume that the constant $\\alpha$ represents a mean value, time-averaged over many orbital periods of the disc. Radial variations in $\\alpha$ may (and should) be present in the disc but it requires a more thorough consideration of the disc physics (i.e. the ionization balance) and is left for a follow-up paper.   Our numerical simulations unambiguously demonstrate that circumstellar discs cannot sustain turbulent viscosity with a spatially and temporally averaged $\\alpha \\ga 0.1$. Such discs would have vanished during just one million year of evolution. The ubiquitous presence of older discs makes such large values of $\\alpha$ unlikely. On the other hand, low values of $\\alpha$ of order $10^{-4}$ make little effect on the secular disc evolution, which in this case is completely governed by gravitational torques rather than viscous ones. In the case of $\\alpha=10^{-3}$, viscosity does have some effect on the disc radial structure but the magnitude of these changes are modest -- the surface density profile becomes somewhat shallower in the inner disc and at the disc's outer boundary, and the disc size increases by a factor of 2 as compared to that of the non-viscous one. %EV: modified below The $\\alpha=0.01$ disc sees considerable changes in its radial structure in the late evolution, and the gravitational stabilization of such discs presents difficulty to account for non-axisymmetric structure and poses problems for the theoretical idea of giant planet formation via direct gravitational instability. Moreover, the gas surface density in the entire $\\alpha=0.01$ disc becomes lower than that of the MMSN after 1.0~Myr. This also poses problems for planet formation models, which often require discs with gas surface densities a few times greater than that of the MMSN \\citep[e.g.][]{Ida04}. We conclude that the mean value of $\\alpha$ (averaged over many orbital periods) should lie in the range $10^{-3}-10^{-2}$, although large transient variations around these values can still be present in real discs.   \\subsection{Effective viscosity due to gravitational torques}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{figure11.eps}} \\caption{Temporal evolution of the effective viscosity parameter $\\aleff$ that accounts for both viscous and gravitational torques, for each of our models.} \\label{fig11} \\end{figure}  An interesting way to account for the {\\it combined} effect of gravitational and viscous torques is to calculate an 'effective alpha' $\\aleff$ near the inner boundary of our simulation\\footnote{For various methods to define $\\aleff$ see e.g. \\citet{Lodato07}}. If we apply the steady-state mass accretion rate formula for thin viscous discs and also apply the $\\alpha-$prescription, then \\begin{equation} \\dot{M} = 3 \\pi \\nu \\Sigma = 3 \\pi \\aleff \\tilde{c}_{s} Z \\Sigma. \\end{equation} We calculate $\\dot{M}$ in each simulation at the boundary of the sink cell ($r=5$ AU), and $\\tilde{c}_{s}, Z,$ and $\\Sigma$ at some distance ($r \\approx 9$ AU) since $\\Sigma$ decreases somewhat near the sink cell. These numbers are used to generate $\\aleff$, which is plotted versus time in Fig. \\ref{fig11} for each of our models. The values of $\\aleff$ may be scaled down by a factor $\\approx 2$ if use also the values of $\\dot{M}$ at 9 AU instead of 5 AU. Clearly, gravitational torques in the absence of turbulent viscosity accounts for an $\\aleff$ in the range $10^{-4}-10^{-3}$ during the late evolution. This is why the addition of a viscous $\\alpha$ of at least $10^{-3}$ is required to see significant changes in the disc evolution.  It is also interesting to compare our results with previous numerical simulations of the secular evolution of viscous circumstellar discs. For instance, \\citet{Lin} have considered the formation and evolution of a circumstellar disc formed during the collapse of a rotating cloud core with initial mass $1.0~M_\\odot$. They use the usual diffusion equation describing the evolution of the surface density in a viscous {\\it axisymmetric} accretion disc \\citep{LBP} \\begin{equation} {\\partial \\Sigma \\over \\partial t} = - {1\\over r} {\\partial \\over \\partial r} \\left[ {1\\over (r^2\\Omega)^\\prime}  {\\partial \\over \\partial r} \\left(  \\nu \\Sigma \\,r^{3} \\Omega^\\prime \\right) \\right], \\end{equation} complemented by some form of the energy equation describing the internal energy balance in the disc due to viscous heating, energy input from the accretion process, and disc radiative cooling. Our model, though neglecting detailed thermodynamics, accounts for a possible disc asymmetry by directly solving the corresponding fluid dynamics equations for a thin disc. The effective kinematic viscosity in Lin \\& Pringle's model comes from turbulent viscosity and gravitational instability, the former is parameterized using a usual Shakura \\& Sunyaev $\\alpha$-prescription (eq.~\\ref{alpha}), while the latter is taken into account following \\citet{Lin87}: \\begin{equation} \\nu_{\\rm g} = \\left\\{ \\begin{array}{ll} {2\\mu \\over 3} \\left( {Q^2_c  \\over Q^2} -1 \\right) \\left( {c^2_s \\over \\Omega} \\right) & \\,\\, \\mbox{if $Q \\le Q_c$} \\\\ 0 & \\,\\, \\mbox{otherwise}.  \\end{array} \\right. \\label{alpha2} \\end{equation} We need to parameterize only the viscous torques, the effect of gravitational torques is taken into account self-consistently.  Both approaches yield circumstellar discs that share some common characteristics. For instance, Lin \\& Pringle's model A1 ($\\alpha,\\mu=0.01$) produces a rather cold disc ($T\\sim 10$~K) that features a near-flat surface density profile near the inner boundary and scales roughly as $r^{-1}$ at $100-1000$~AU. Their disc has a sharp outer edge upon formation but it spreads out in the %EV: cause -> course below course of evolution. These features are also seen in our modeling, though our disc in the corresponding model~3 has a smaller size ($\\sim 300$~AU) and a lower surface density. This brings about the most striking difference found between the two approaches -- the derived disc masses. \\citet{Lin} have %EV: eliminated exclamation mark below reported disc masses that are an order of magnitude larger than the corresponding stellar masses in the early evolution. Although this large difference reduces with time, the disc mass is still comparable to that of the star in the late evolution ($t\\ga 1.0$~Myr), irrespective of the value of $\\alpha$. Such massive discs are not observed. Our models predict maximum disc-to-star mass ratios of $\\xi=14\\%$ (model~1), and this value quickly reduces with time for the viscosity-dominated discs (e.g. model~3). %EV: added \"for non-viscous models\" nelow Although our obtained values of $\\xi$ for non-viscous models still seem to be greater than those usually inferred for T Tauri stars, $0.5\\%-1\\%$ \\citep[e.g.][]{Andrews}, they are not unfeasible given that the measured disc masses may be underestimated by conventional methods by as much as an order of magnitude \\citep[e.g.][]{Hartmann}. Figure~\\ref{fig5} indicates that gravitational torques are a dominant mechanism of mass and angular momentum transport in the disc in its early evolutionary phase.  The fact that \\citet{Lin} obtain overmassive discs in their numerical simulations suggests that either the parameterization given by equation~(\\ref{alpha2}) is inadequate, certainly in the early disc evolution, or the ratio of rotational-to-gravitational energies $\\beta$ in their model is too large, resulting in most of the initial cloud mass landing on to the disc and through the disc on to the star rather than directly on to the star. In other words, the phase of near constant accretion of matter from the envelope directly on to the star (see Fig.~\\ref{fig6}) is very short in the Lin \\& Pringle model and the disc is not capable of processing the infalling mass to the star fast enough to keep its mass low. Indeed, Lin \\& Pringle have adopted $\\beta$ in the range $0.25-0.64$, which is much larger than the values recently inferred for molecular cloud cores by \\citet{Caselli}, $\\beta=10^{-4}-0.07$. In our model, $\\beta$ is set to $1.4\\times 10^{-3}$."
    },
    "0812/0812.4096_arXiv.txt": {
        "abstract": "We present results obtained from near-infrared $JHK$ spectroscopic observations of  novae V2491~Cyg and V597~Pup in the early declining phases of their 2007 and 2008 outbursts respectively. In both objects, the spectra displayed  emission lines of HI, OI, HeI and NI. In V597 Pup, the HeI lines were found to strengthen rapidly with time.  Based on the observed spectral characteristics, both objects are classified as He/N novae. We have investigated the possibility of V2491 Cyg being a recurrent nova  as has been suggested. By studying the temporal evolution of the line widths in  V2491~Cyg  it appears unlikely that the binary companion is a giant star with heavy wind as in recurrent novae of the RS Oph type. Significant deviations from that of recombination case B conditions are observed in the  strengths of the HI lines. This indicates that the  HI lines, in both novae,  are optically thick during the span of our observations. The slope of the continuum spectra in both cases was found to have a  $\\lambda^{-(3-3.5)}$ dependence which deviates from a Rayleigh-Jeans spectral distribution. Both novae were detected in the post-outburst super-soft X-ray phase; V2491 Cyg being very bright in X-rays has  been the target of several observations. We discuss and correlate our infrared observations with the observed X-ray properties of these novae. ",
        "introduction": "We present here near-infrared studies of two classical novae V2491 Cygni and V597 Pup. The observations cover the early declining phase in both these novae. Both objects have been detected in X-rays but V2491 Cyg has specifically been the subject of considerable current interest. It is one of the rare novae to have been detected in X-rays even before its recorded optical outburst (Ibarra \\& Kuulkers, 2008). Apart from that, it has also shown very prominent and sustained super soft X-ray emission in recent times as  described later. Our infrared observations, in conjunction with observations from other wavelength regimes like the X-ray and optical, should give a more comprehensive understanding of these interesting objects.  The characteristics of a nova outburst depend on three principal parameters viz. the white dwarf mass, the core temperature of the white dwarf and the rate of mass accretion from the binary companion. The accreted mass which eventually ignites the explosion is ejected back into space at high velocities of the order of a few hundreds to thousands of kilometers per second. A white dwarf accreting mass below the critical rate undergoes a thermonuclear runaway in the hydrogen burning shell at the bottom of the accreted layer and becomes a classical nova when sufficient amount of mass has been accumulated. The brightness of a classical novae begins to decline once mass ejection subsides after the ejection of most of the envelope. As the remnant of the envelope collapses back onto the white dwarf, residual hydrogen burning occurs in the shell radiating most of the energy in the UV to X-ray band. The photosphere of the white dwarf shrinks which increases the effective temperature and the nova appears as a hot blackbody-like object. The temperature of the blackbody-like object (post nova) is predicted to be about 3$\\times$10$^5$ K for classical novae for a few years after optical maximum (Prialnik 1986). The post-outburst nova with this high temperature photosphere is a source of  super-soft X-rays which lasts for a duration depending on the leftover envelope mass, chemical composition of the envelope, mass of the white dwarf etc. (Kato 1997; Hachisu \\& Kato 2006). Seven such super-soft Galactic novae appear to have been observed with ROSAT and Swift till 2006 June 30 (Orio et al. 2001; Ness et al. 2007). Some of the novae also show hard X-ray emission, the origin of which is different from the emission from the white dwarf photosphere. X-ray emission from the white dwarf is not expected to be observed immediately after the nova outburst. X-ray detections in some of the novae during this phase is understood to  originate from some other processes such as shocks within the expanding shell or from a collision of the expanding shell and the atmosphere of the companion as seen in RS Opiuchi (Bode et al. 2006). Therefore, it is rather interesting to understand and relate the mechanism of classical nova outburst and the ensuing super-soft X-ray emission phase.  V2491~Cyg was discovered on 2008 April 10.728 UT at a visual magnitude of 7.7 (Nakano et al. 2008). Low resolution spectra of the source obtained on April 11.72 UT showed prominent broad Balmer emission lines indicating the object to be a nova in its early phase of outburst (Ayani \\& Matsumoto 2008). Subsequent infrared observations on April 12.56 UT confirmed the object as a nova with an unusual spectrum with extremely broad lines of complex profiles (Lynch et al. 2008). The presence of the helium line at 1.083 $\\mu$m and strong emission features of NI and NII in the near-infrared spectrum suggested the object to be a helium-nitrogen (He/N)nova. Balmer lines from H$\\alpha$ to H$\\delta$ in emission were seen in the optical spectrum taken on April 11.99 UT and April 13.95 UT (Tomov et al. 2008). The FWHM of the H$\\alpha$ line decreased from $\\sim$4800 km s$^{-1}$ to $\\sim$4600 km s$^{-1}$ within about 2 days. Another near-IR observation on April 17.6 UT (Rudy et al. 2008) found that the brightness of the nova had already declined by a factor of three  to  photometric magnitudes of J=7.7, H=7.8 and K=7.4; this was accompanied by  an  increase in the strengths of the emission lines. The OI lines at 0.8446 and 1.1287 $\\micron$, which are fluorescently excited by Lyman $\\beta$, were found to be the strongest emission features in the infrared spectrum. The relative strength of these lines imply  a reddening of E(B-V) =  0.43 (Rudy et al. 2008). Broad HI emission lines and prominent HeI and OI lines were detected in the near-infrared spectra of the nova on April 18 and 20 (Ashok et al. 2008). They have also pointed out the possibilities of the presence of NI lines at 1.1762 and 1.2462 $\\micron$s. A periodicity of 0.0958 day with  0.03-0.05 magnitude amplitude variations was suggested from the optical photometric observations (Baklanov et al. 2008).  Nova V597~Pup was discovered on 14 November 2007 at a visual magnitude of 7 (Pereira et al. 2007). Low resolution spectrum of V597~Pup on November 14.77 UT in the 410-670 nm range showed the presence of broad Balmer lines of Fe II with P-Cyg features along with the blue-shifted absorption lines of Fe II (49, 74), OI  and Si II (Naito \\& Tokimasa 2007). The expansion velocities on the same day were estimated to be 1800, 1700, and 1800 km s$^{-1}$, derived from the FWHM of the Balmer H$\\alpha$, H$\\beta$ and H$\\gamma$ emission lines respectively. The near-infrared spectrum (in the 0.8-2.5 $\\micron$ range) of the nova on November 30, 2007 showed that the He~I and O~I lines were very strong and the Br~$\\gamma$ line showed a double-peaked line profile, without any evidence of He~II lines and thermal emission from dust; a value of E(B-V) = 0.3 was estimated (Rudy et al. 2007). Similar IR observations on 2008 January 7 showed the presence of strong He~II lines and weak coronal lines of [S~VIII], [Si~VI] and [Ca~VIII] in the  spectrum which was accompanied by the appearance of  super-soft X-rays from the nova (Ness et al. 2008a).  Both  V2491~Cyg and V597~Pup were detected in the post-outburst, super-soft X-ray phase with various X-ray observatories. In case of V597~Pup, the X-ray emission was detected on 2008 January 8 and 17 from the Swift satellite (Ness et al. 2008a).  Most of the photons detected were below 0.7 keV indicating the onset of the super-soft X-ray phase. These were the only observations  reporting the X-ray detection in V597~Pup. However, V2491~Cyg was much brighter in X-ray during the super-soft X-ray phase and hence observed with Swift, Suzaku, and XMM-Newton observatories. Pre-outburst observations with Swift showed the presence of a variable soft X-ray source at the position of V2491~Cyg  even three months before the outburst. This makes the object specially interesting since it is only the second nova after V2487 Oph (Hernanz \\& Sala 2002) to be detected in X-rays even before eruption (Ibarra et al. 2008).   Nearly continuous monitoring of the nova from April 17 to May 14 with Swift found that the brightness of the source increased continuously since April 27, reaching $\\sim$200 count ks$^{-1}$ on May 14 with the significant increase below 1.5 keV (Page et al. 2008). The search for any periodicity in X-ray on the longest Swift observations on May 8 and 10, to corroborate the suggested 0.1 day orbital period (Baklanov al. 2008),  yielded a negative result. In addition to the Swift monitoring, Suzaku observation of the nova on May 9 showed the presence of strong low energy emission  from K-shell N, O, and Ne lines in the spectrum (Osborne et al. 2008). Swift observations on May 16 showed a very strongly increasing count rate which touched $\\sim$5 counts s$^{-1}$ on May 18. XMM-Newton observations of the nova, on 2008 May 20, detected several O and N absorption lines along with the broad Ne emission lines, suggesting the nova to be an ONe nova (Ness et al. 2008b).  \\begin{figure*} \\vskip 7.5 cm \\special{psfile=snaik_fig1a.eps vscale=45 hscale=45 hoffset=-65 voffset=220 angle=-90} \\special{psfile=snaik_fig1b.eps vscale=45 hscale=45 hoffset=185 voffset=220 angle=-90} \\caption{The light curves of V597~Pup (left panel; V band data from AAVSO shown in filled circles while the V band data from AFOEV is shown in open circles) and V2491~Cyg (right panel; using V band data from the AAVSO). The arrow marks in both the panels show the days of our near-infrared observations.} \\label{fig1} \\end{figure*} ",
        "conclusions": "While our IR observations do not throw direct light on the interesting X-ray behavior of V2491 Cyg, there are  some correlated aspects which could be discussed. But before proceeding to do so, it is appropriate to briefly discuss  the super soft X-ray phase in novae.  The discovery of X-ray emission from nova GQ~Mus, 463 days after the outburst (Ogelman et al. 1984), indicates that  novae in  post-outburst should go through an extended phase of residual hydrogen burning and appear as a super-soft X-ray source (SSS). The duration of the post-outburst super-soft X-ray phase depends on the nuclear burning of the residual mass on the surface of the white dwarf and also on the mass of the white dwarf. Massive white dwarfs show super-soft X-ray phases for short duration. This is because the stronger gravity in massive white dwarfs finishes the hydrogen burning in the residual envelope faster compared to that of less massive white dwarfs. The theoretical estimation of the duration of the residual hydrogen burning (super-soft X-ray phase) ranges from 5.6 years to 2$\\times$10$^4$ years for white dwarf masses from 1.35 to 0.6 M$_{\\odot}$ (Gehrz et al. 1998). However, observations showed a much shorter duration (a few hundred days) of the post-nova super-soft X-ray phase (Pietsch et al. 2005) implying that either the mass of the white dwarfs is much higher than the commonly assumed mass or the X-ray turn-off depends on some other parameters other than the white dwarf mass. Along with V597~Pup and V2491~Cyg, the super-soft X-ray emission has been detected, so it would appear, in fewer than a dozen novae (Orio et al. 2001; Drake et al. 2003; Ness et al. 2007 and references therein; Ness et al. 2008a, 2008b).  While the SSS phase is believed to arise from the hot photosphere of the white-dwarf, we investigate whether a fraction of the observed X-rays from V2491 Cyg could arise from an alternative mechanism viz. a shock in the ejecta. A plausible reason to expect a shock is the secondary outburst seen in V2491 Cyg. If this mini-outburst is due to a significant second mass-loss phase in which the ejected matter moves at a higher velocity than the primary ejecta, then this  matter  could catch up to and collide with the earlier expanding envelope leading to the formation of a shock. The hot shocked gas could then contribute to the observed X-rays. Hence it is necessary to check whether there is any significant increase in the velocity profiles during the second rebrightening.  However, from Figure 4, we do not see any evidence for any enhancement in velocity for the matter ejected around the peak of the mini-outburst ($\\sim$ 26 April). Therefore it is unlikely that a shock could have formed due to the mini-outburst. The optical light curve of V2491~Cyg resembles  that of  two other novae, V1493~Aql and V2362~Cyg, in terms of the presence of a secondary peak. The secondary peaks in V2491~Cyg and V1493~Aql (Venturini et al. 2004)  were observed about 16 and 48 days after the primary outburst respectively, where as in case of V2362~Cyg it was after $\\sim$240 days (Kimeswenger et al. 2008). Among these three novae, the SSS phase has been detected in V2491~Cyg (present work) and V2362~Cyg (Hernanz et al. 2007). Because of only three novae showing the secondary peak in the light curve, out of which only two are detected in X-ray, it is premature to say whether the secondary peak is linked to the detected X-ray emission from the novae.  The beginning of the SSS phase (X-ray turn-on) i.e. the residual nuclear burning in various novae is generally detected after more than a hundred days of outburst. Some of the examples are, 250 days after the outburst in V1974~Cyg (Krautter et al. 1996), 200 days in Nova LMC 1995 (Orio et al. 2003), 180 and 222 days  in V382~Vel (Orio et al. 2002), about 250 days  in V1494~Aql (Drake et al. 2003), about 180 days in V4743~Sgr (Ness et al. 2003) and about 175 days in V574~Pup (Ness et al. 2007). However, in case of V597~Pup and V2491~Cyg, the super-soft X-ray state was detected about 55 and 30 days after outburst respectively, which is  rather early compared to the other novae discussed above. As per the model of Hachisu \\& Kato (2006) the relatively early onset of the SSS phase in these novae suggests  the white dwarfs in these systems are very massive and close to the Chandrasekhar limit."
    },
    "0812/0812.1020_arXiv.txt": {
        "abstract": "We estimate the sensitivity of various experiments detecting ultra-high-energy cosmic rays to primary photons with energies above $10^{19}$~eV. We demonstrate that the energy of a primary photon may be significantly (up to a factor of $\\sim 10$) under- or overestimated for particular primary energies and arrival directions. We consider distortion of the reconstructed cosmic-ray spectrum for the photonic component. As an example, we use these results to constrain the parameter space of models of superheavy dark matter by means of both the observed spectra and available limits on the photon content. We find that a significant contribution of ultra-high-energy particles (photons and protons) from decays of superheavy dark matter is allowed by all these constraints. ",
        "introduction": "Recent studies~\\cite{AGASA_Risse,A+Y,Ylim,Auger_sdlim} put strong limits on the presence of photons among primary particles of ultra-high-energy (UHE, $\\gtrsim 10^{19}$~eV) cosmic rays (CR). However, while at energies $\\sim 10^{19}$~eV current gamma-ray limits ($<2\\%$ of the integral flux of cosmic particles at 95\\% CL~\\cite{Auger_sdlim}) are very restrictive, the best limit at $10^{20}$~eV allows (at 95\\% CL) as much as 36\\% of primary gamma rays~\\cite{A+Y}. At the same time, the reconstruction of the UHECR spectrum often relies on a general assumption of hadronic primaries. This assumption is explicit in Monte-Carlo simulations (AGASA~\\cite{AGASA_Eest}, Telescope Array~\\cite{TAsurface}) and implicit in methods which use calibration relations obtained for the bulk of the lower-energy events which are mostly hadronic (Yakutsk~\\cite{PravdinICRC2005}, Pierre Auger~\\cite{Auger_sd_eest}). Well justified at $10^{19}$~eV, this approach may lead to systematic distortions of the spectrum in the very interesting energy range $\\gtrsim 10^{20}$~eV, where a significant fraction of gamma rays is allowed. Because of this systematics, models which predict primary photons of these energies should be tested with great care.   The purpose of this paper is twofold. First, we allow for a hypothetical contribution of primary photons which is consistent with all experimental limits and study its possible effect on the derivation of the spectrum. Second, we consider particular theoretical models which predict such a contribution and constrain them with the simultaneous account of the spectrum and of photon limits.  Though ultra-high energy photons have not been observed by now, they are expected to be seen among secondary products of the Greisen--Zatsepin--Kuzmin \\cite{g,zk} reaction (see e.g.\\ Ref.~\\cite{GZKphotons}). They are also predicted in exotic hypothetical top-down models of UHECR origin (see Refs.~\\cite{siglreview,KachDMrev} for reviews), notably in the superheavy dark-matter (SHDM) models.  In this paper, we give a quantitative analysis of reconstruction of the spectrum by various experiments in the presence of a fraction of gamma-ray primaries. This analysis is obligatory when exotic scenarios of UHECR origin are constrained: theoretical predictions for the photon fraction depend on the normalization of the exotic contribution to the spectrum.  Air showers induced by primary photons differ significantly from the bulk of hadron-induced events (see e.g.\\ Ref.~\\cite{RisseRev} for a recent review). There are two competitive effects responsible for the diversity of showers induced by primary photons.  First, due to the Landau, Pomeranchuk \\cite{LP} and Migdal \\cite{M}~(LPM) effect the electromagnetic cross-section is suppressed at energies $E~>~10^{19}$~eV. The LPM effect leads to the delay of the first interaction and the shower arrives to the ground level under-attenuated. Another effect is the $e^\\pm$ pair production due to photon interaction with the geomagnetic field above the atmosphere. Secondary electrons produce gamma rays by synchrotron radiation generating a cascade in the geomagnetic field. The probability of this effect is proportional to the square of the product of photon energy and perpendicular component of geomagnetic field. The shower development therefore depends on both zenith and azimuthal angles of photon arrival direction. If the effect is strong enough, the particles enter the atmosphere with energies below the LPM threshold. As a result, not only the shower development differs from that of an average hadronic shower but also this difference is strongly direction-dependent.  The energy reconstructed by an experiment may therefore differ significantly from the true energy of the primary photon. In addition, acceptance of fluorescence detectors for photons may differ from that for hadronic primaries which is assumed in the spectral calculation. The difference in the reconstructed energy and acceptance should be accounted for individually for each experiment. In this note, we estimate quantitatively this difference in the energy reconstruction and discuss its possible effect on the spectrum and implications for constraining the SHDM models.  UHECR spectra measured by different experiments are not in the mutual agreement. The disagreement is sometimes attributed to systematic errors in energy determination. Both the normalization of AGASA~\\cite{AGASA_eest}, HiRes~\\cite{HiRes_spec} and Yakutsk~\\cite{Yakutsk_spec} spectra and the position of the astrophysically motivated dip agree within this approach \\cite{Berezinsky_escale}. The spectrum observed by the Pierre Auger Observatory (PAO) \\cite{Auger_spec} agrees with the others in the region above $10^{19}$~eV in the same assumption, see Fig.~\\ref{fig:4spec}. \\begin{figure} \\begin{center} \\includegraphics[width=0.7 \\columnwidth]{4spec.eps} \\end{center} \\caption{ \\label{fig:4spec} Spectra of AGASA~\\cite{AGASA_eest} (red triangles), HiRes~\\cite{HiRes_spec} (blue diamonds), Yakutsk~\\cite{Yakutsk_spec} (grey boxes) and PAO \\cite{Auger_spec} (green stars) for energies scaled according to the best fit with protons from extragalactic sources described in Sec.~\\ref{sec:shdm}. } \\end{figure} The energy rescaling is motivated by discrepancies in different methods of energy estimation for hadronic showers: for instance, the energy estimated by the surface detector of PAO alone is about 30\\% larger than the energy estimated in the standard reconstruction procedure based on calibration to fluorescence-detector data \\cite{Engel_cic}. For AGASA and Yakutsk the use of CORSIKA for energy estimation leads to systematic shifts of energies downwards by about 10-15\\% and 40\\%, respectively~\\cite{Teshima_ichep,Dedenko_yak_en,Dedenko_yak_confirm}. Note that the rescaled spectra do not coincide at the highest energies ($E \\gtrsim 10^{20}$~eV); the discrepancy may be attributed either to insufficient statistics or to the presence of energy-dependent systematics, for instance of a non-standard component at the highest energies. This question should be addressed in the future with more data available.  Several limits on the fraction $\\epsilon_\\gamma$  of UHE photons in the integral cosmic-ray flux have been set by several independent experiments. They are summarized in Fig.~\\ref{fig:gamma-limits}. \\begin{figure} \\begin{center} \\includegraphics[width=0.8 \\columnwidth]{limitsV.eps} \\end{center} \\caption{ \\label{fig:gamma-limits} Upper limits (95\\% C.L.) for the fraction $\\epsilon_\\gamma$ of gamma rays in the integral flux of cosmic rays with energies higher than $E$: Haverah Park~\\cite{HP_lim} (HP), AGASA \\cite{AGASA_1stlim} (A), \\cite{AGASA_Risse} (AH), AGASA and Yakutsk \\cite{A+Y} (AY), Yakutsk \\cite{Ylim} (Y), Auger fluorescence detector \\cite{Auger_fdlim} (PF), Auger surface detector \\cite{Auger_sdlim} (PA). } \\end{figure} The most restrictive limits (95\\% C.L.) are $\\epsilon _\\gamma <0.36$ for $E>10^{20}$~eV from the AGASA and Yakutsk joint dataset~\\cite{A+Y}, $\\epsilon _\\gamma <0.22$ for $E>4\\times 10^{19}$~eV from Yakutsk \\cite{Ylim}, $\\epsilon _\\gamma <0.05$ for $E>2\\times 10^{19}$~eV and $\\epsilon _\\gamma <0.02$ for $E> 10^{19}$~eV from the Auger surface detector \\cite{Auger_sdlim}. Even when the energy reconstruction of photons is properly taken into account (which was done in the calculation of these most restrictive limits), the limits on $\\epsilon _\\gamma$ may depend on the uncertainty in the energy reconstruction of non-photonic primaries, notably in the case of low statistics (see discussion and Fig.~2 in Ref.~\\cite{Ylim}). A more stable quantity is the flux of gamma rays; the Pierre Auger upper limit on the integral flux of photons above $10^{19}$~eV is $3.8\\times10^{-3}$~km$^{-2}$sr$^{-1}$yr$^{-1}$ at the 95\\% CL. These photon limits may be used to constrain top-down models provided a theoretical model for the top-down photon flux is given. In a self-consistent analysis, the latter should be normalized to the observed spectrum. This normalization requires in turn the account of the energy reconstruction of photons which constitute a significant fraction of the top-down flux. Below, we perform a joint analysis of the spectrum and of photon limits for the SHDM models and constrain the space of two SHDM parameters (mass and lifetime of the superheavy particle). Contrary to the previous studies, most of which used either a naive AGASA normalization or an overall (independent from the energy and arrival direction) multiplicative correction for the reconstructed photon energies, our results suggest that a significant fraction of cosmic rays from the SHDM decays is allowed by all experimental constraints.  The rest of the paper is organized as follows. In Sec.~\\ref{sec:sens} we estimate sensitivity of four experiments (AGASA, HiRes, Pierre Auger and Yakutsk) to the primary photon component. In Sec.~\\ref{sec:shdm} we consider an example of constraining SHDM parameters using primary spectra and photon limits. Section~\\ref{sec:concl} summarizes our results. ",
        "conclusions": "\\label{sec:concl} Modern experiments have different sensitivities to the photon component and this should be taken into account when testing particular models. The AGASA experiment overestimated the energy of primary photons with energies $E \\gtrsim 10^{19}$~eV by a factor of 2 in average, though the overestimation reaches a factor of $\\sim 10$ for particular energies and arrival directions. Contrary, the surface detector of the Pierre Auger Observatory underestimates the energy of primary photons in this energy range by a factor of $\\sim 4$ in average, while underestimation by a factor of $\\sim 10$ happens for particular energies and arrival directions. The HiRes detector overestimated the photon energies by a factor of $\\sim 1.1$, uniformly over arrival directions and well within the systematic uncertainties. However, it had a significantly lower exposure for primary photons than for primary hadrons.  One of the scenarios predicting a significant amount of primary UHE photons is the superheavy-dark-matter model. We analyzed constraints on its parameters from the observed spectra and limits on the photon content. While the most restrictive photon limits~\\cite{A+Y,Ylim,Auger_sdlim} account for peculiarities in the energy reconstruction for photons, a dedicated study was required and performed for the spectral fits. A significant (more than one half of the flux at $E>10^{20}$~eV) SHDM component is still allowed by all limits. Though there seems no present need for the SHDM to explain the UHECR spectrum, a large contribution of SHDM at the highest energies is not excluded and may be further constrained (or validated) by future experiments. Among the tests are measurements of the spectrum, studies of anisotropy and searches for primary photons and neutrinos. Our study indicates that a significant SHDM contribution is allowed for masses of dark-matter particles $M_X \\gtrsim 4\\times 10^{22}$~eV, so one needs experiments with large aperture at $E>10^{20}$~eV (e.g.\\ JEM-EUSO~\\cite{JEM-EUSO}) to test the shape of the spectrum. At lower energies ($10^{19}$~eV to $10^{20}$~eV) the model may be constrained by improving the gamma-ray limits and by precise studies of the Galactic anisotropy, e.g.\\ with the help of fluorescent detectors (which treat uniformly both photon and hadron primaries).  As a final remark, we note that the example of photons should warn one against naive tests of models predicting ``exotic'' primaries with the experimental data. For instance, the correlations with BL Lac type objects observed by HiRes \\cite{BL:HiRes,HiRes:BL} require neutral primary particles. If the latters were photons, apparent absence of correlations in the preliminary data of the PAO surface detector \\cite{PAO:BL} is easily explained \\cite{index} by underestimation of their energies as compared to the bulk of hadronic primaries. With more exotic primary particles, the analysis becomes even less trivial.  We are indebted to V.~Berezinsky, D.~Gorbunov, J.-M.~Fr\\'ere, M.~Libanov, M.~Pravdin, V.~Rubakov, D.~Semikoz, P.~Tinyakov and I.~Tkachev for valuable discussions. This work was supported in part by the Russian Foundation of Basic Research grants 07-02-00820 and 09-07-00388, by Federal Agency for Science and Innovation under state contracts 02.740.11.0244 and 02.740.11.5092, by the grants of the President of the Russian Federation NS-1616.2008.2, MK-1966.2008.2 (OK), MK-61.2008.2 (GR), by the Dynasty Foundation (GR), by the Russian Science Support Foundation (GR), by IISN and by the Belgian Science Policy (IAP VI-11).  Numerical part of the work was performed at the computer cluster of the Theoretical Division of INR RAS. GR and ST thank Service of Theoretical Physics, ULB, for kind hospitality."
    },
    "0812/0812.4605_arXiv.txt": {
        "abstract": "Cyclical wind variability is an ubiquitous but as yet unexplained feature among OB stars. The O7.5 III(n)((f)) star $\\xi$ Persei is the brightest representative of this class on the Northern hemisphere. As its prominent cyclical wind properties vary on a rotational time scale (2 or 4 days) the star has been already for a long time a serious magnetic candidate. As the cause of this enigmatic behavior non-radial pulsations and/or a surface magnetic field  are suggested. We present a preliminary report on our attempts to detect a magnetic field in this star with high-resolution measurements obtained with the spectropolarimeter Narval at TBL, France during 2 observing runs of 5 nights in 2006 and 5 nights in 2007. Only upper limits could be obtained, even with the longest possible exposure times. If the star hosts a magnetic field, its surface strength should be less than about 300 G. This would still be enough to disturb the stellar wind significantly. From our new data it seems that the amplitude of the known non-radial pulsations has changed within less than a year, which needs further investigation. ",
        "introduction": "Like many O and B stars, the O7.5III(n)((f)) star $\\xi$ Per shows very prominent cyclical wind variability in the UV resonance lines (Fig.~1a), manifested by discrete absorption components (DACs) which migrate from red to blue, and narrow when they approach (but not reach) the terminal velocity as measured from saturated wind profiles (\\textit{e.g.} \\cite[Kaper \\etal\\ 1999)]{kaper99}. In $\\xi$ Per the DAC period is 2.09 d.  Multiwavelength observations of a number of OB stars, including $\\xi$ Per, have shown that the cyclic behavior is present down to the surface of the star (\\cite[de Jong \\etal\\ 1997)]{dejong97}, and that the typical timescale varies with the (estimated) rotational timescale. This strongly argues in favor of a surface phenomenon which perturbs the base of the flow.  Two scenarios are proposed: in the so-called Corotating Interaction Region (CIR) model a perturbation at the surface of the star causes a local increase (or decrease) of the radiative force driving the stellar wind, or surface magnetic fields may disturb the outflow, both resulting in a rotationally modulated stellar wind (see \\cite[Cranmer and Owocki 1996)]{cranmer96}. A number of coordinated UV and optical observations have confirmed the CIR model, including for the case of $\\xi$ Per, but the origin of the perturbations is not known, which is one of the most challenging problems in stellar wind research of the last decades. We present here our most recent efforts to measure the magnetic field and the pulsation properties.  \\begin{figure}[htp] \\centering \\leftline{\\includegraphics[bb=0 13 596 842,height=4.5cm,keepaspectratio]{fig1a}}  \\vspace{-4.7cm} \\centerline{\\includegraphics[bb=38 218 551 770,height=4.5cm,keepaspectratio]{fig1b}}  \\vspace{-4.6cm} \\rightline{\\includegraphics[bb=38 203 551 770,height=4.7cm,keepaspectratio]{fig1c}}  \\vspace{0.1cm} \\caption{(a) Progressing DACs in the Si IV wind lines in 1994, (b) Non radial pulsation (period 3.5 h) in the He I 5411 line, December 2006. (c)  September 2007; the amplitude of the moving features seems to have changed after 9 months. } \\end{figure}  \\begin{figure}[htp] \\centering \\leftline{\\includegraphics[bb=30 378 507 742,height=3.2cm,keepaspectratio]{fig2a}} \\vspace{-3.2cm} \\centerline{\\includegraphics[bb=30 378 507 742,height=3.2cm,keepaspectratio]{fig2b}} \\vspace{-3.2cm} \\rightline{\\includegraphics[bb=8 240 503 632,height=3.2cm,keepaspectratio]{fig2c}}  \\caption{Preliminary magnetic results of Narval spectropolarimetry as a function of time. (a) December 2006. (b) September 2007. (c) LSD Stokes V profile of the last point in September 2007.  No significant Zeeman signature was found. } \\end{figure} ",
        "conclusions": "Previous Musicos magnetic measurements of $\\xi$ Per were presented by \\cite[de Jong \\etal\\ (2001)]{dejong01}, with no detection. For Narval data the magnetic analysis is essentially the same, applying the least-squares deconvolution method (\\cite[Donati \\etal\\ 1997)]{donati97} to the spectral lines sensitive to magnetic effects. This yields the longitudinal component of the field, averaged over the facing hemisphere of the star. In Fig.~2a and 2b the 45 results are plotted as a function of time. No Zeeman signature was found. As an example, Fig.~2c shows the LSD profile of the best exposed spectrum of September 2007 with a S/N = 1960. The magnetic values are preliminary, as the used spectral linelist can still be optimized.  We also analysed a number of spectral lines for non-radial pulsations with known period of 3.5 h found in a previous Musicos campaign (\\cite[de Jong \\etal\\ 1999]{dejong99}). The amplitude may have changed during the 9 months between our Narval runs, see Fig.~1b and 1c. This obviously needs further investigation, as the impact on the stellar wind may have changed.  Observations with a higher efficiency should be able to detect a possible field, unless the geometry is such that the magnetic properties of the two hemispheres cancel out."
    },
    "0812/0812.1216_arXiv.txt": {
        "abstract": "{% The final inspiral and coalescence of a black hole binary can produce highly beamed gravitational wave radiation. To conserve linear momentum, the black hole remnant can recoil with ``kick\" velocity $v_\\mathrm{kick}$ $\\leq$ 4000 km/s. We present two sets of full N-body simulations of recoiling massive black holes (MBH) in high-resolution, non-axisymmetric potentials. The host to the first set of simulations is the main halo of the Via Lactea I simulation (Diemand et al. 2007). The nature of the resulting orbits is investigated through a numerical model where orbits are integrated assuming an evolving, triaxial NFW potential, and dynamical friction is calculated directly from the velocity dispersion along the major axes of the main halo of Via Lactea I. By comparing the triaxial case to a spherical model, we find that the wandering time spent by the MBH is significantly increased due to the asphericity of the halo. For kicks larger than 200 km/s, the remnant MBH does not return to the inner 200 pc within 1 Gyr, a timescale an order of magnitude larger than the upper limit of the estimated QSO lifetime. The second set of simulations is run using the outcome of a high-resolution gas-rich merger (Mayer et al. 2007) as host potential. In this case, a recoil velocity of 500 km/s cannot remove the MBH from the nuclear region.} ",
        "introduction": "In the context of the currently favored  $\\Lambda$CDM cosmogony, where large halos are assembled through the hierarchical merging and accretion of small progenitors, MBH mergers should be common. While dynamical friction against the dark matter and baryons can only form binaries with separations of $\\sim$1 pc, other dynamical process such as stellar ejection and gas drag can efficiently reduce the binary orbital separation down to $\\sim$0.001 pc, the regime where gravitational wave radiation dominates the orbital energy loss. This radiation is typically anisotropic due to asymmetries associated with the black holes' mass and spin, causing the center of mass of the system to recoil in order to balance the linear momentum carried away by the gravitational wave radiation (\\cite{bekenstein73}; \\cite{fitchett84}; \\cite{favata04}).  The ``kick'' velocity of the remnant depends on the mass ratio ($m_2/m_1 < 1$) and the spin parameters ($a_1,\\, a_2$) of the binary, but not on the total mass of the system.  Early estimates of the recoil velocity, framed in the post-Newtonian regime for non-spinning, unequal mass black holes (\\cite{fitchett83}; \\cite{redmount89}), yielded velocities in the range $100< v_\\mathrm{kick}< 500$ km/s which were successfully reproduced by numerical data (\\cite{baker06b}; \\cite{gonzalez07b}; \\cite{herrmann07}).  Simulations with varying (arbitrary) spin orientations (\\cite{campanelli07b}) and mass ratios (\\cite{baker08}) show that recoil velocities can be significantly larger, reaching ${v_\\mathrm{kick} \\sim 2000}$ km/s (\\cite{campanelli07a}; \\cite{gonzalez07b}) and are predicted to be as large as $v_\\mathrm{kick} \\sim 4000$ km/s for unequal mass ($m_2/m_1 = 1/3$), maximally spinning black holes (\\cite{campanelli07a}). In order to address the issue of whether off-nuclear QSO/AGN are common and detectable, we ought to understand the dynamics of their orbits. The radial orbit of a recoiling MBH in a spherically symmetric potential  was studied analytically by \\cite{madau04}. They showed that large kicks (400 km/s) can displace MBHs a few tens of kiloparsecs away from the center of a Milky-Way size stellar bulge. After the kick, the MBH undergoes several oscillations before decaying back to the bottom of the potential. Most of the orbital energy is lost during the MBH passages through the center, where dynamical friction is most efficient. \\cite{gualandris07} substantiated these results by performing direct summation N-body simulations of MBH recoil in a spherically symmetric galaxy, modeled as a binary-depleted core-S\\'{e}rsic profile (\\cite{graham03}). They find that after reaching the core radius, the MBH and the core experiences damped oscillations with bulge's center of mass, which decelerate the remnant MBH until it reaches thermal equilibrium with the surrounding stars.  They estimate that this final stage may take up to 1 Gyr in large galaxies. As we show in Sect.~\\ref{tri_model}, the MBH wandering time can be largely increased if the host potential is triaxial, while the presence of gas can significantly damp the overall MBH motion.  After a major merger, however, the remnant dark matter halo tends to be prolate (e.g. \\cite{novak06}), a factor that can play a major role for large kicks.  In the following we carry out N-body simulations of recoiling MBH in two high-resolution non-axisymmetric potentials: the Via Lactea I simulation's main halo, a triaxial dark-matter-only host with known triaxiality (\\cite{kuhlen07}) and evolution parameters (\\cite{diemand07}), and the gas-rich galaxy merger remnant described in \\cite{mayer07}. In the dark-matter only case, we present a numerical model that successfully reproduces the main features of the MBH orbits, and can therefore be used to estimate the fallback time and maximum displacement for any given kick velocity (\\cite{guedes08}, in preparation).  \\begin{figure}[t]% \\hskip-3mm \\scalebox{1.06}{ \\includegraphics[scale=0.45,clip]{orb3d_vl01.eps} } \\vskip-4mm \\caption{(online colour at: www.an-journal.org) Resulting 3D orbit corresponding to simulation VL080. The first 0.5 Gyr are plotted in yellow, the following 0.5 Gyr in orange and the remaining 0.1 Gyr is plotted in dark red. The box has side length $L=1.2$ kpc.} \\label{orb3dvl01} \\end{figure} ",
        "conclusions": "If MBH mergers are common in the context of $\\Lambda$CDM, then the ``rocket effect\" should also be common, especially at high-redshift where the bulk of the mass assembly occurs. Because recoiling MBHs carry gas and stars with them, the detection of off-nuclear QSOs is possible by finding velocity shifts between the narrow line region associated with gas accreting onto the MBH, and the broad line emission associated with the galaxy left behind (\\cite{bonning07}). The best off-nuclear QSO candidate today is thought to be powered by a recoiling SMBH of mass $M_{\\bullet} =6 \\times\\! 10^8 \\,\\rm M_{\\odot}$ which travels at a velocity of $2650$ km/s with respect to the host galaxy (\\cite{komossa08}).  In order to assess the detectability of so-called naked QSOs, we ought to understand the dynamics of the  recoiling MBH orbits and the environment provided by the host potential. While detectability requires the presence of an accretion disk around the remnant MBH, the medium of its host galaxy can damp its motion significantly, disfavoring its detection. However, if the kick is large enough that, while keeping an accretion disk and stellar sphere on influence, the MBH can reach as far as the (generally) prolate dark matter halo, dynamical friction is much less efficient and therefore the detection probability is mostly determined by the lifetime of the QSO."
    },
    "0812/0812.1570_arXiv.txt": {
        "abstract": "Torsion oscillations of the neutron star crust are Landau damped by the Alfven continuum in the bulk. For strong magnetic fields (in magnetars), undamped Alfven eigenmodes appear. ",
        "introduction": "It is thought that Israel et al (2005) have detected torsion oscillations of the neutron star crust. Here we show that crustal modes are strongly affected by the Alfven continuum in the bulk. For week magnetic fields, there is Landau damping of crustal eigenmodes. At stronger magnetar-like magnetic fields undamped Alfven eigenmodes appear. In this regime, the Alfven waves in the bulk control the torsion oscillations of the crust.   It might be that Israel et al (2005) have actually measured the Alfven speed in the bulk $c_A$, rather than the torsion speed in the crust $c_t$. If so, the measured $\\nu \\approx 20$Hz frequency gives $c_A \\approx 2\\nu R \\approx 4\\times 10^7$cm/s (assuming that $\\nu$ is the fundamental Alfven eigenmode).   We  calculate the Landau damping of the crustal torsion modes and Alfven eigenmode (undamped) frequencies. The thin-crust uniform-field constant-density model is used, and we only consider the antisymmetric axisymmetric modes. This is obviously not good enough for serious neutron star seismology. One must include: (i) the non-axisymmetric modes, (ii) the toroidal magnetic field, (iii) the realistic density profile. But the qualitative results -- Landau damping and emergence of undamped Alfven eigenmodes -- should be robust.   This paper is a formal mathematical development of the ideas of Levin (2007). ",
        "conclusions": "It is thought that Israel et al (2005) have detected torsion oscillations of the neutron star crust. Here we show that crustal modes are strongly affected by the Alfven continuum in the bulk. For week magnetic fields, there is Landau damping of crustal eigenmodes. At stronger magnetar-like magnetic fields undamped Alfven eigenmodes appear. In this regime, the Alfven waves in the bulk control the torsion oscillations of the crust.   It might be that Israel et al (2005) have actually measured the Alfven speed in the bulk $c_A$, rather than the torsion speed in the crust $c_t$. If so, the measured $\\nu \\approx 20$Hz frequency gives $c_A \\approx 2\\nu R \\approx 4\\times 10^7$cm/s (assuming that $\\nu$ is the fundamental Alfven eigenmode).   We  calculate the Landau damping of the crustal torsion modes and Alfven eigenmode (undamped) frequencies. The thin-crust uniform-field constant-density model is used, and we only consider the antisymmetric axisymmetric modes. This is obviously not good enough for serious neutron star seismology. One must include: (i) the non-axisymmetric modes, (ii) the toroidal magnetic field, (iii) the realistic density profile. But the qualitative results -- Landau damping and emergence of undamped Alfven eigenmodes -- should be robust.   This paper is a formal mathematical development of the ideas of Levin (2007)."
    },
    "0812/0812.3743_arXiv.txt": {
        "abstract": "The bulge is a region of the Galaxy of tremendous interest for understanding galaxy formation. However measuring photometry and kinematics in it raises several inherent issues, such as severe crowding and high extinction in the visible. Using the Besan\\c{c}on Galaxy model and a 3D extinction map, we estimate the stellar density as a function of longitude, latitude and apparent magnitude and we deduce the  possibility of  reaching and measuring bulge stars with Gaia. We also present an ongoing analysis of the bulge using the Canada-France-Hawaii Telescope. ",
        "introduction": "Observing towards the bulge gives measures for a large number of individual stars that are the tracers of the Galactic formation and evolution. A detailed study of the bulge is necessary to understand its structure, dynamics and formation. The bulge of the Milky Way is the only one where the parameters in the six dimensions of phase space can be determined, as well as elemental abundances, on a star by star basis. Its detailed observation is thus very useful to test different scenarii of galactic bulges formation.  The bulge has been explored by a few ground based surveys, for example DENIS, 2MASS and microlensing surveys (OGLE, MACHO, DUO, EROS-2 and MOA) in the visible and the infrared, and from space by Spitzer in the mid infrared. The interstellar matter distributed in the plane is a major obstacle to the observation of the stars in the direction of the bulge and the Galactic center, particularly in the visible.  Gaia will give unprecedented view of the Galaxy, but it is often stated that, due to the wavelength coverage of instruments, it will be hardly usable for bulge studies. However, estimations of what Gaia will effectively see in the bulge is worthwhile to do before launch. We have investigated this question using a population synthesis model and a realistic 3D map of the extinction in the inner Galaxy. Studies of the bulge from ground before Gaia are also possible. We present an on-going survey of the bulge undertaken with Megacam aimed at producing good photometry and proper motions for a substancial number of clump stars. ",
        "conclusions": "Thanks to new photometric and astrometric surveys, the bulge structure can be revealed, bringing important constraints on the galaxy formation and evolution scenarii. We presented here several studies using such data set. Other available data are worthwhile to be analysed, such as the 49 OGLE-II fields for which proper motions have been measured \\cite{Sumi}. A preliminary analysis of these data with the Galaxy model shows that the bulge rotation can be constrained with these data.  Furthermore, radial velocities and metallicities data are available in the bulge directions \\citep{Ibata1995,Minniti1996,Tiede1999,Rich}.The analysis of these data will allow us to refine the kinematical parameters of the bulge, as well as bulge metallicity. Finally, we showed that Gaia should give a detailed survey of bulge stars, down to the main sequence, in terms of photometry but also kinematics, including regions very close to the Galactic plane.."
    },
    "0812/0812.4213_arXiv.txt": {
        "abstract": "V838~Mon erupted at the beginning of 2002 becoming an extremely luminous star with $L \\simeq 10^6$~L$_{\\sun}$. Among various scenarios proposed to explain the nature of the outburst the most promising is a stellar merger event. In this paper we investigate the observational properties of the star and its surroundings in the post outburst phase. We have obtained a high resolution optical spectrum of V838~Mon in October 2005 using the Keck~I telescope. We have identified numerous atomic features and molecular bands present in the spectrum and provided an atlas of those features. In order to improve the spectrum interpretation we have performed simple modeling of the molecular bands. Our analysis indicates that the spectrum is dominated by molecular absorption features arising in photospheric regions with temperatures of $\\sim$2400~K and in colder outer layers, where the temperature decreases to $\\sim$500~K. A number of resonance lines of neutral alkali metals are observed to show P-Cyg profiles. Particularly interesting are numerous prominent emission lines of [Fe~{\\small II}]. All of them show practically the same profile, which can be well described by a Lorentzian profile. In the blue part of the spectrum photospheric signatures of the B-type companion are easily seen. We have fitted the observed spectrum with a synthetic one and the obtained parameters are consistent with the B3V type. We have also estimated radial and rotational velocities of the companion. ",
        "introduction": "}  The eruption 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 \\citep{muna02,kimes02,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. At the same time a B3V companion to the erupted object was discovered in the optical \\citep{mundes}. A detailed analysis of the evolution of the object in the outburst and decline can be found, e.g. in \\cite{tyl05}.  Several mechanisms have been proposed to explain the eruption of V838~Mon, including an unusual nova \\citep{it92}, a late He-shell flash \\citep{law05}, and a stellar merger \\citep{soktyl03}. They have critically been discussed in \\cite{tylsok06}. These 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 a $\\sim 8\\,M_\\odot$ main-sequence star.  In October/November 2006 the B3V companion significantly faded and a strong H$\\alpha$ emission appeared in the spectrum of V838~Mon \\citep{gor06,bond06}. Late in 2004 an emission-line spectrum, composed mainly of [\\ion{Fe}{2}] lines, started to develop \\citep{barsuk06}, reaching its maximum around the 2006 eclipse-like event \\citep{mun07}.  In the present paper we present and discuss a high resolution spectrum of V838~Mon acquired with the Keck~I telescope in October 2005. At that time the [\\ion{Fe}{2}] emission-line spectrum was already well developed and the object was a year before the eclipse of the B3V companion. ",
        "conclusions": "We have obtained a spectrum of V838~Mon on 2005 October 13 using the Keck/HIRES instrument. The spectrum of a resolving power of $\\sim 34\\,000$ covered a wavelength range of 3720~--~7960~\\AA. The spectrum is very complex and shows numerous spectroscopic features which can be identified as arising from different components or environments. V838~Mon itself, i.e. the object which erupted in 2002, is seen mainly in the green and red parts of the spectrum as a cool continuum with strong, broad and complex molecular absorptions. Numerous atomic lines showing P-Cyg profiles or broad absorption features are also present in this component. The blue part of the spectrum is dominated by a photospheric spectrum of the B3V companion, showing broad absorptions lines of \\ion{H}{1} and \\ion{He}{1}. In the blue and green parts of the spectrum, numerous emission lines can be identified. These mostly arise from forbidden transitions in singly ionized iron. Finally, several absorption features of interstellar origin have been identified. These include atomic resonance lines as well as DIBs.  The numerous [\\ion{Fe}{2}] emission lines observed in our spectrum evidently arise from rarefied and (partly) ionized matter. All the lines show the same profile, which can be well fitted with a Lorentzian profile. The lines are centred at a heliocentric radial velocity of $\\sim$13~km\\,s$^{-1}$ and have a typical FWHM of $\\sim$80~km\\,s$^{-1}$.  In the cores of the photospheric absorption lines of the B3V companion we see weak emission features. The double feature in the H$\\beta$ line can be interpreted as a single emission with a narrow absorption superimposed on it. The derived parameters of the emission component are close to those of the [\\ion{Fe}{2}] lines. This suggests that the Balmer emission features arise from the same environment as the [\\ion{Fe}{2}] emission lines.  The spectrum displays a number of strong atomic, mostly resonance, lines with P-Cyg profiles. Evidently V838~Mon, being almost 4 years after its outburst, continues losing mass. The emission component typically peaks at a heliocentric velocity of $\\sim$80~km\\,s$^{-1}$, while the absorption component usually extends from $\\sim$50~km\\,s$^{-1}$ to $\\sim -100$~km\\,s$^{-1}$. The absorption component often shows a complex profile. In particular, in a few lines a strong and narrow absorption component is clearly seen at a velocity of $\\sim -80$~km\\,s$^{-1}$. A few atomic lines were observed to show pure absorption profiles.  The spectrum from the cool component is dominated by numerous, often very strong, molecular bands in absorption. Some bands are so deep that the flux from the cool component practically goes to zero. In these cases we have a sort of spectral windows, through which we can see the B3V companion, even in the red part of the spectrum. All the identified bands are from oxides. They include TiO, VO, AlO, ScO, and YO.  Our analysis shows that the bands are formed in different regions. Some of them, especially those from highly excited states, are formed, at least partially, in the photospheric regions. The temperature in these layers is as high as $\\sim$2400~K. These bands show a heliocentric radial velocity of $\\sim$58~km\\,s$^{-1}$, which we identify as a radial velocity of V838~Mon itself.  Most of the bands are formed in regions significantly cooler than the photospheric ones. The vibrational temperatures derived from some bands go down to $\\sim$500~K. These are usually very strong bands arising from ground states. The radial velocity of these bands is usually much lower than that of the photospheric regions and goes down to $\\sim -125$~km\\,s$^{-1}$. Overall, our analysis of the molecular bands shows that they are mostly formed in outflowing matter with a significant velocity gradient.  We have detected bands arising from forbidden transitions in TiO. As far as we know, these are first observations of TiO forbidden transitions in astrophysical objects.  The observed spectrum of the B-type companion, seen primarily in the blue part of our spectrum, can be fitted with a model atmosphere spectrum with solar abundances, effective temperature of $\\sim$18\\,000~K, and a gravity, log\\,$g \\simeq 4.0$. These values agree well with those expected for a B3V star. The absorption lines are significantly broadened, which we interpret as a rotational broadening with a velocity, $V\\,{\\rm sin}\\,i \\simeq 250$~km\\,s$^{-1}$. We have derived a heliocentric radial velocity of the B3V companion of $37\\pm17$~km\\,s$^{-1}$.  Several interstellar absorption features are seen in our spectrum. They include lines of \\ion{Ca}{2}, \\ion{Na}{1}, and \\ion{K}{1}, as well as DIBs. The lines of \\ion{Ca}{2} and \\ion{Na}{1} are observed in the spectrum of the B3V companion. Their profiles are very similar to those observed in the spectrum of V838~Mon during the 2002 outburst. This result is consistent with the idea that V838~Mon and the B3V companion are related objects. From equivalent widths of diffuse interstellar bands we have estimated an interstellar reddening of $E_{B-V} \\simeq 0.9 - 1.0$, which agrees well with interstellar extinction values obtained by other authors using other methods.  A detailed analysis and interpretation of the data presented in the present paper, as well as of data from other sources, will be done in a forthcoming paper."
    },
    "0812/0812.3611_arXiv.txt": {
        "abstract": "The standard model of cosmology suggests the existence of two components, ``dark matter'' and ``dark energy'', which determine the fate of the Universe. Their nature is still under investigation, and no direct proof of their existences has emerged yet. There exist alternative models which reinterpret the cosmological observations, for example by replacing the dark energy/dark matter hypothesis by the existence of a unique dark component, the {\\it dark fluid}, which is able to mimic the behaviour of both components. After a quick review of the cosmological constraints on this unifying dark fluid, we will present a model of dark fluid based on a complex scalar field and discuss the problem of the choice of the potential. ",
        "introduction": "The standard model of cosmology suggests that the total energy density of the Universe is presently dominated by the densities of two components: the {\\it dark matter} component, which is a pressureless matter fluid having an attractive behaviour, and the {\\it dark energy}, whose main properties are a negative pressure and a nearly constant energy density today, which evoke the idea of vacuum energy. The nature of both components remains unknown, and in the near future we can hope that the Large Hadron Collider (LHC) will be able to provide hints on the nature of dark matter. In spite of the mysteries of the dark components, it is generally considered that dark matter can be modeled as a system of collisionless particles, whereas the most usual models of dark energy are the scalar field based quintessence models. However, many difficulties in usual dark energy and dark matter models still question the cosmological standard model bases, leaving room for other models to be investigated. In this paper, we consider a unifying model in which the dark matter and dark energy components can be considered as two different aspects of a same component, the {\\it dark fluid}. We will first review how such a unifying model can be constrained by the cosmological observations. Then we will consider a dark fluid model based on a complex scalar field, and we will discuss the question of the choice of the scalar field potential. ",
        "conclusions": ""
    },
    "0812/0812.1422_arXiv.txt": {
        "abstract": "We study the quantum remnant of a scalar field protected by the uncertainty principle. The quantum remnant that survived the later stage of evolution of the universe may provide dark energy and dark matter depending on the potential. Though the quantum remnant shares some useful property of complex scalar field (spintessence) dark energy  model, quantum fluctuations are still unstable to the linear perturbations for $V \\sim \\phi^q$ with $q<1$ as in the spintessence model. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.4946_arXiv.txt": {
        "abstract": "An interacting scalar field with largish coupling to curvature can support a distinctive inflationary universe scenario.   Previously this has been discussed for the Standard Model Higgs field, treated classically or in a leading log approximation.  Here we investigate the quantum theory using renormalization group methods.  In this model the running of both the effective Planck mass and the couplings is important.  The cosmological predictions are consistent with existing WMAP5 data, with $0.967\\lesssim n_s\\lesssim0.98$ (for $N_e=60$) and negligible gravity waves.   We find a relationship between the spectral index and the Higgs mass that is sharply varying for $m_h\\sim 120-135$\\,GeV (depending on the top mass); in the future, that relationship could be tested against data from PLANCK and LHC.   We also comment briefly on how similar dynamics might arise in more general settings, and discuss our assumptions from the effective field theory point of view. ",
        "introduction": "\\label{sec:introduction}  The hypothesis that there was a period in the early history of the universe during which a local Lorentz invariant energy density -- i.e., an effective cosmological term -- dominated the equation of state, causing exponential expansion, explains several otherwise puzzling features of the present universe (flatness, isotropy, homogeneity) \\cite{Guth,Linde,AlbrechtSteinhardt82,Linde83}.  It also suggests a mechanism whereby primordial density fluctuations arise through intrinsic fluctuations of quantum fields, leading to qualitative and semi-quantitative predictions that are consistent with recent observations. However the physics behind inflation remains mysterious.  What, specifically, is the source of the energy density?  Ideas ranging from fields associated with supersymmetry, string moduli, ghosts, branes, and others abound \\cite{ModLinde, ModLiddle,ModDvali, ModMcAllister, ModNima, ModKachru}.    One (or more) of them might be correct, but all are highly speculative, and none is obviously compelling.   Alternatively, we can look for inflationary dynamics based on degrees of freedom already present in the Standard Model.   We can also attempt to maintain the guiding philosophy of the Standard Model, including gravity, to allow only local interactions which are gauge invariant and have mass dimension $\\leq 4$.   Within this very restrictive framework, there remains the possibility to include the non-minimal gravitational coupling  $\\xi H^\\dagger H\\mathcal{R}$.  Here $H$ is the Higgs field, $\\mathcal{R}$ is the Ricci scalar, and $\\xi$ is a dimensionless coupling constant, whose value is unknown and largely unconstrained by experiment.\\footnote{For $\\xi=-1/6$ the Higgs is conformally coupled to gravity.} Indeed renormalization of the divergences arising in a self-interacting scalar theory in curved spacetime requires a term of this form \\cite{BirrellDavies}. The Higgs sector is then described, classically, by the Lagrangian \\beq \\mathcal{L}_h= -|\\partial H|^2 + \\mu^2H^\\dagger H-\\lambda (H^\\dagger H)^2+\\xi H^\\dagger H\\,\\mathcal{R} \\label{lagrangianH} \\eeq where $\\lambda$ is the Higgs self coupling and $\\mu$  is the Higgs mass parameter.    It has been known for some time that such minimal classical Lagrangians can support inflation driven by an interesting interplay between the quartic term and the non-minimal coupling term \\cite{Salopek,Fakir,Kaiser,Komatsu}.  For ease of reference, we will call this general set-up ``running inflation''; the name seems appropriate, since evolution of the effective Planck mass and the effective scalar mass is central to the dynamics.\\footnote{The term ``running inflation\" was used in a different context in \\cite{runningpaper}.} This quasi-renormalizable set-up allows use of renormalization group methods, as will be illustrated here. By quasi-renormalizable, we mean that the theory is renormalizable when gravity is treated classically; in particular, we ignore quantum corrections from graviton exchange (see Appendix \\ref{NewApp}). In the investigation of (non-gravitational) quantum effects, it is appropriate to focus specifically on the Standard Model, for two reasons. First, because (as we will see) it illustrates important qualitative issues in a very concrete, familiar setting.  Second, because it -- or something close to it -- might actually contain the degrees of freedom relevant to real-world inflation, in which case the specific predictions we derive could help describe reality.    Recently, the idea that the Standard Model Higgs field, non-minimally coupled to gravity, can lead to inflation was proposed in Ref.~\\cite{Bezrukov}. Those authors argued that the radiative corrections to the potential are negligible and hence the inflationary parameters can be computed using the classical Lagrangian. They found that the cosmological predictions are in good agreement with cosmological data, independent of the Standard Model parameters, such as $\\lambda$. On the other hand the authors of Ref.~\\cite{Barvinsky} criticized their approach, suggesting that the quantum corrections to the potential can be very important. They concluded that a Higgs lighter than $230$\\,GeV cannot serve as the inflaton, because the predicted spectral index is ruled out by  WMAP5 data \\cite{wmap5}. Ref.~\\cite{Barvinsky} only incorporated  quantum corrections at leading log order, extrapolated from low energies. Here, in contrast, we will compute the full renormalization group improved effective action at 2-loops. We conclude that running inflation based upon a Standard Model Higgs makes predictions that are consistent with current cosmological data, and leads to firm predictions for the PLANCK satellite and the LHC. Our main result is a correlation between the spectral index and the Higgs mass, see Fig.~\\ref{nsplot}. This correlation is absent in the classical theory. The origin of the correlation lies in the interactions of the Standard Model, which dictate the form of the effective action.  In Section \\ref{sec:nonmin} we review inflation with non-minimally coupled scalars. In Section \\ref{sec:classical} we investigate the classical theory of the Higgs non-minimally coupled to gravity. In Section \\ref{sec:quantum} we describe our method for obtaining the quantum corrected effective action. We compute all the inflationary observables numerically and present results in Section \\ref{sec:results}. Finally, we review our results and discuss their significance in Section \\ref{sec:discussion}.    \\begin{figure}[t] \\begin{center} \\includegraphics[scale=1.0]{nsVSmhSmallmh2.pdf} \\end{center} \\caption{The spectral index $n_s$ as a function of the Higgs mass $m_h$ for a range of light Higgs masses. The 3 curves correspond to 3 different values of the top mass: $m_t=169$\\,GeV (red curve), $m_t=171$\\,GeV (blue curve), and $m_t=173$\\,GeV (orange curve). The solid curves are for $\\alpha_s(m_Z)=0.1176$, while for $m_t=171$\\,GeV (blue curve) we have also indicated the 2-sigma spread in $\\alpha_s(m_Z)=0.1176\\pm0.0020$, where the dotted (dot-dashed) curve corresponds to smaller (larger) $\\alpha_s$. The horizontal dashed green curve, with $n_s\\simeq 0.968$, is the classical result. The yellow rectangle indicates the expected accuracy of PLANCK in measuring $n_s$ ($\\Delta n_s\\approx 0.004$) and the LHC in measuring $m_h$ ($\\Delta m_h\\approx 0.2$\\,GeV).  In this plot we have set $N_e=60$.} \\label{nsplot} \\end{figure} ",
        "conclusions": "\\label{sec:discussion} A number of papers have discussed bounds on the Higgs mass coming from demanding stability of the vacuum, e.g., see \\cite{Sher,Giudice,Nimabds,Isidori}. Cosmological constraints only require metastability on the lifetime of the universe, which places the constraint $m_h\\gtrsim 105$\\,GeV \\cite{Giudice}. However, if we further demand that the Higgs drive inflation, we find that heavier Higgs are required: $m_h\\gtrsim 126$\\,GeV (depending on the top mass, see eq.~(\\ref{boundsmh})), which essentially coincides with the bounds from absolute stability. Furthermore, by demanding that the theory remains perturbative to high energies ($m_h\\lesssim 190$\\,GeV), we establish a correlation between both stability and triviality bounds, and inflation.  More precisely, {\\em we have established a  mapping between the renormalization group flow and the cosmological spectral index}.  Over a substantial range of parameter space the classical value $n_s\\simeq 0.968$ (for $N_e=60$) emerges as a good approximation, but there are corrections.  Given a detailed microphysical theory, such as the Standard Model, we can explicitly calculate such corrections, as summarized in Fig.~\\ref{nsplot}. This plot displays a sharp rise in the spectral index towards 0.98, or so, as we approach vacuum instability for a light Higgs.     It is likely that the Standard Model  is only  the low-energy limit of a more complete theory,  accommodating the facts that do not find explanation within the Standard Model, such as neutrino masses, dark matter, baryon asymmetry, etc. Our methodology is still applicable, so long as we can control the relevant $\\beta$-functions.   In principle some quite different scalar field, not connected to the Standard Model Higgs, could drive running inflation.   The central requirement is a large coefficient for the $\\phi^2  \\mathcal{R}$ term.   It is possible that such a coefficient could emerge as some sort of Clebsch-Gordan coefficient, or from the coherent addition of several smaller terms (involving more basic scalars $\\phi_j$).   It is also possible to consider, in the same spirit, the dimension 3 interaction  $\\phi  \\mathcal{R}$, which arises for generic scalar fields, though not of course for the Standard Model Higgs. Furthermore, as discussed in Appendix \\ref{NewApp} and Refs.~\\cite{Burgess,Espinosa}, the inclusion of higher dimension operators may significantly affect the predictions of the original theory and even spoil its validity (as is the case in many inflationary models). These possibilities, and their possible embedding in unified field theory or string theory, deserve further investigation.   If the Higgs boson exists and it is in the  mass  range considered in this Letter, the LHC will discover it and will determine the Higgs mass with a precision of about $0.1\\%$ \\cite{tdr}, which means an uncertainty $\\Delta m_h\\approx 0.2$ GeV. In order to extract accurate correlations between the inflationary observables and the Higgs mass it is crucial to improve the precision with which we know the other parameters of the Standard Model, in particular the top quark mass and the strong coupling. The current value of the top mass from direct observation of events is $m_t=(171.2\\pm 2.1)$ GeV \\cite{PDG}. In the near future, the LHC will improve the determination of the top mass, but relatively large systematic uncertainties will prevent a top mass determination to better than $1$ GeV; more conservatively, the top mass will be determined at LHC with an error $\\Delta m_t=1 \\div  2$ GeV. Looking further ahead, the ILC is expected to be able to measure the top mass to $\\sim100$\\,MeV. So together with the measured Higgs mass from the LHC and improved precision on the strong coupling, as well as calculating higher order effects and reheating details, running inflation in the Standard Model will predict a rather precise value for the spectral index.     \\bigskip  We would like to thank M.~Amin, F.~D'~Eramo, A.~Guth, and C.~Santana for useful discussions, and A.~Riotto for comments on the manuscript. We thank F.~L.~Bezrukov, A.~Magnin, and M.~Shaposhnikov for correspondence. The work of ADS is supported in part by the INFN ``Bruno Rossi'' Fellowship. The work of ADS, MPH, and FW is supported in part by the U.S. Department of Energy (DoE) under contract No. DE-FG02-05ER41360."
    },
    "0812/0812.0424_arXiv.txt": {
        "abstract": "We investigate correlation between the arrival directions of ultra-high-energy cosmic rays (UHECRs) and the large-scale structure (LSS) of the Universe by using statistical quantities which can find the angular scale of the correlation. The Infrared Astronomical Satellite Point Source Redshift Survey (IRAS PSCz) catalog of galaxies is adopted for LSS. We find a positive correlation of the highest energy events detected by the Pierre Auger Observatory (PAO) with the IRAS galaxies inside $z=0.018$ within the angular scale of $\\sim 15^{\\circ}$. This positive correlation observed in the southern sky implies that a significant fraction of the highest energy events comes from nearby extragalactic objects. We also analyze the data of the Akeno Giant Air Shower Array (AGASA) which observed the northern hemisphere, but the obvious signals of positive correlation with the galaxy distribution are not found. Since the exposure of the AGASA is smaller than the PAO, the cross-correlation in the northern sky should be tested using a larger number of events detected in the future. We also discuss the correlation using the all-sky combined data sets of both the PAO and AGASA, and find a significant correlation within $\\sim 8^{\\circ}$. These angular scales can constrain several models of intergalactic magnetic field. These cross-correlation signals can be well reproduced by a source model in which the distribution of UHECR sources is related to the IRAS galaxies. ",
        "introduction": "\\label{introduction}  The origin of ultra-high-energy cosmic rays (UHECRs) above $10^{19}$ eV has been an open problem in astroparticle physics. The highly isotropic distribution of their arrival directions and the deflection angles of UHECRs estimated in the Galaxy have indicated the extragalactic origin of them. In 1999, Akeno Giant Air Shower Array (AGASA) found a small-scale anisotropy of the arrival distribution of UHECRs above $4 \\times 10^{19}$ eV within $2.5^{\\circ}$, which is comparable with its angular resolution \\cite{takeda99}. The small-scale anisotropy can be interpreted as a signal of point-like sources.   Motivated by the AGASA result, many researchers have tested correlation between the arrival directions of UHECRs and the positions of several classes of astrophysical objects. The correlation with BL Lac objects was discussed by using the AGASA and Yakutsk data \\cite{tinyakov01,tinyakov02,gorbunov02}. The authors of Refs. \\cite{smialkowski02,singh04} studied the correlation with infrared galaxies using {\\it Infrared Astronomical Satellite Point Source Redshift Survey} (IRAS PSCz) catalog \\cite{saunders00}, and discussed the possibility that the AGASA events below $\\sim 10^{20}$eV come from luminous infrared galaxies. In Ref. \\cite{hague07}, the correlation with nearby active galactic nuclei (AGNs) listed in the {\\it Rossi X-ray Timing Explorer} (RXTE) catalog of AGNs was investigated and no significant correlation was found. The authors of Ref. \\cite{gorbunov05} made a comprehensive study of correlation with various classes of powerful extragalactic sources and found that BL Lac objects and unidentified gamma-ray sources might be correlated with the AGASA and Yakutsk data. The spatial correlation with the supergalactic plane has also been discussed and positive signals have been obtained \\cite{stanev95,uchihori00,stanev08}. On the other hand, High Resolution Fly's Eye (HiRes) reported no significant signal for any point-like sources \\cite{abbasi05,abbasi07} and the correlation with BL Lac objects studied by Refs. \\cite{tinyakov01,tinyakov02,gorbunov02} was not supported by the HiRes data \\cite{abbasi06}. Theoretically, several authors predicted the anisotropy of UHECR arrival distribution in future observations when UHECR source distribution follows the large-scale structure (LSS) of the Universe \\cite{waxman97,yoshiguchi03b,takami08,kashti08}. Cosmic rays above $8 \\times 10^{19}$ eV lose their energies rapidly by interactions with cosmic microwave background (CMB) and therefore they cannot reach the earth from distant sources, typically above 100 Mpc (Greisen-Zatsepin-Kuz'min (GZK) mechanism \\cite{greisen66,zatsepin66}). In other words, the GZK mechanism predicts positional correlation between such highest energy cosmic rays and nearby UHECR sources if the Universe is not magnetized. We predicted such positional correlation within a few degree scale even taking into account a structured intergalactic magnetic field (IGMF), which reproduces the observed structure of local Universe, and 5 yrs observation of the Pierre Auger Observatory (PAO) will unveil the local UHECR source distribution \\cite{takami08}.   Recently, the PAO has reported correlation between the arrival directions of its 27 events above $5.7 \\times 10^{19}$eV and the positions of nearby active galactic nuclei (AGNs) with 3$\\sigma$ confidence level \\cite{pao07a,pao08}. This report showed that the isotropic distribution of the arrival directions of the highest energy cosmic rays is disfavored. This was also a clear evidence which indicated directly that UHECR sources are extragalactic objects and can be regarded as the start of UHECR astronomy.   However, the interpretation of the PAO results is problematic. All of the PAO-correlated AGNs except Centaurus A are classified into Seyfert galaxies and Low-Ionization Nuclear Emission Regions (LINERs), which have much weaker activities than radio-loud AGNs, one of plausible candidates of UHECR sources \\cite{moskalenko08}. Thus, it is an open question whether the PAO-correlated AGNs are really UHECR sources.   The PAO adopted the 12th edition of Veron-Cetty \\& Veron Catalog (VC catalog) of AGNs \\cite{veron06} for a correlation analysis. This catalog is a compilation of many astronomical catalogs of AGNs available in literature and therefore AGNs in the VC catalog are selected based on different criteria. In order to investigate the nature of UHECR sources, the authors of Ref. \\cite{george08} analyzed the spatial correlation of the PAO events with hard X-ray selected AGNs compiled by the Swift satellite, an AGN catalog complied based on a criterion. They found a significant correlation by using the 2 dimensional generalization of the Kolmogorov-Smilnov test (2D KS test).   More generally, the arrival directions of UHECRs are expected to correlate with a galaxy distribution, which is a representative of LSS of the Universe, if UHECR sources are some astrophysical objects. When a correlation with LSS is investigated, a catalog as homogeneous as possible is required. Ref. \\cite{kashti08} studied correlation between the PAO events and LSS using the IRAS PSCz catalog. They showed that the PAO data is inconsistent with random distribution of UHECR sources at 98\\% confidence level and favors a modeled source distribution following LSS. They analyzed the similarity between the observed arrival distribution of UHECRs and mock one from their cosmic ray intensity map constructed based on the IRAS catalog using $6^{\\circ} \\times 6^{\\circ}$ angular bins. The bin size was determined by considering possible deflections of UHECRs by the IGMF in order to avoid the effect of the IGMF. Ref. \\cite{ghisellini08} discussed the correlation with HI-selected galaxies by using 2D KS test and a significant correlation was found. The authors interpreted the result as an evidence that spiral galaxies are the hosts of UHECR sources and could suggest that newly born magnetars are the best candidates of UHECR sources. Note that 2D KS test cannot estimate the angular scale of a spatial correlation. These works are important steps toward the understanding of UHECR sources. However, if the angular scale is naturally found by another analysis without fixing the scale, we can obtain a better evidence of the correlation and information on intervening magnetic fields and the composition of UHECRs.   In this study, we investigate spatial correlation between the arrival directions of the observed highest energy cosmic rays and a galaxy distribution as a representative of LSS by using statistical quantities which need not to set an artificial angular scale and can estimate the angular scale of the correlation. We adopt the PAO data for cosmic rays observed in the southern hemisphere and the AGASA data for cosmic rays detected in the northern hemisphere. The HiRes is a UHECR observatory with the largest exposure in the northern hemisphere and recently, the HiRes reported no significant correlation with AGNs listed in the VC catalog in their data by using the same method as the PAO \\cite{abbasi08}. However, HiRes has not published the detailed data and, moreover, the analytical estimation of the apertures of fluorescence detectors is more complicated than that of ground-based detectors. Thus, we do not consider the HiRes in this paper. We also check whether a source model associated with LSS can reproduce the cross-correlation signal and try to constrain the IGMF strength taking into account UHECR propagation in magnetized intergalactic space. This approach allows us to analyze the correlation taking account of the radial information of the positions of UHECR sources.   This paper is laid out as follows. In Section \\ref{methods}, we explain a galaxy catalog which represents LSS and UHECR samples used in this study, and our methods of the statistical analysis. In Section \\ref{correlation}, we calculate cross-correlation function introduced in Section \\ref{methods} between the arrival directions of the observed highest energy cosmic rays and the positions of nearby galaxy distribution, and discuss correlation with its angular scale. In Section \\ref{modelp}, we construct a simple model of UHECR sources, check the reproducibility of the observed cross-correlation and try to constrain the IGMF strength. Finally, in Section \\ref{conclusion}, we discuss our results and make conclusions. ",
        "conclusions": "\\label{conclusion}  In this paper, we have investigated spatial correlation between the arrival directions of the highest energy cosmic rays observed by the PAO and AGASA, and the IRAS galaxies as a representative of the local structure of the universe. For statistical analysis, we adopted a cross-correlation function which can find the angular scale of the spatial correlation. This statistic need not to be set an artificial angular scale on correlation studies. We confirmed that the arrival directions of UHECRs above $5.7 \\times 10^{19}$eV detected by the PAO are inconsistent with isotropic distribution and found that these have significant correlation with the galaxy distribution in the local universe inside $z=0.018$ within the angular scale of $\\sim 15^{\\circ}$. On the other hand, the AGASA data did not have obvious correlation with the IRAS galaxies. We also performed the same cross-correlation analysis on the whole sky using the combined data set of both the PAO and AGASA, and found the evidence of positive correlation within the angular scale of $\\sim 8^{\\circ}$. These results could be well reproduced by a source model that UHECR sources are distributed related to galaxies and the source number density is $\\sim 10^{-4}$ Mpc$^{-3}$, taking propagation process in magnetized intergalactic space into account. These indicate that the distribution of UHECR sources are related to that of galaxies.   An important point of the results in this study is the angular scale in which the correlation is positive. We showed that the combined data has positive correlation compared to random event distributions within $\\sim 8^{\\circ}$ in Section \\ref{modelp}. This scale is comparable with the angular scale of the auto-correlation of galaxies within $z=0.018$. Thus, we cannot distinguish the origin of the cross-correlation scale; the deflection angles of UHECRs in the IGMF or the spread of UHECR source distribution, but the result provides us with an upper limit of $\\theta_{\\rm obs}$. It enables us to constrain several structured IGMF models. Several simulations of the LSS formation with magnetic field have been performed and have applied to the studies of UHECR propagation in the local universe \\cite{sigl04,dolag05,das08}, but the resultant magnetic structures are different except the centers of clusters. Among the three simulation results, 70\\% of protons reach the Earth with $\\theta_{\\rm obs} \\geq 15^{\\circ}$ even for the energies above $10^{20}$ eV in an IGMF model of Ref. \\cite{sigl04}, in which only sources are strongly magnetized. Such large deflection angles are inconsistent with the angular scale of the spatial correlation estimated in this study. Thus, the magnetic structure in local Universe proposed by Ref. \\cite{sigl04} is not consistent with the PAO data. Note that the authors of Ref. \\cite{sigl04} pointed out the uncertainty of observer positions in their simulated universe (in other words, the magnetic structure of local Universe). The angular scale gives information on a magnetic field in local Universe   The angular scale of positive correlation against random event distribution also have information on the composition of UHECRs at the Earth, though interpretation depends on the results of the LSS formation simulations. Since UHECR sources are expected to be associated with dense structures, like clusters of galaxies and filamentary structures, propagating UHECRs are affected by magnetic fields in these structures. A LSS simulation by Ref. \\cite{das08} showed that the filamentary structure has relatively strong magnetic field up to 10 nG. A simple but structured IGMF model developed in Ref. \\cite{takami06} also reproduce such structures. Propagating protons above $6 \\times 10^{19}$ eV in the magnetic structure by Ref. \\cite{das08} is deflected by less than $\\sim 10^{\\circ}$, which is comparable with the angular correlation scale. When we assume heavy composition of UHECRs, like irons, the deflection angles of UHECRs are a few tens times larger than protons. This conflicts with our estimated angular correlation scale. Thus, protons or light nuclei are favored as the composition of the highest energy cosmic rays if the IGMF models proposed by Refs. \\cite{das08,takami06} are real. On the other hand, a simulation by Ref. \\cite{dolag05} predicted the filamentary structure with the IGMF strength of 0.1-1 nG and the volume fraction with the strength of magnetic field larger than 1nG is much smaller than the other 2 simulations. Thus, heavy nuclei-dominated composition is allowed if the IGMF model of Ref. \\cite{dolag05} is real.   Another approach to estimate the strength of intervening magnetic field is the measurement of the separation angle between real UHECR sources and the arrival directions of UHECRs from them. The authors of Ref. \\cite{angelis08} estimated the IGMF strength by using the correlation scale proposed by PAO ($\\sim 3.1^{\\circ}$) as sub-nG assuming the PAO-correlated AGNs are real UHECR sources, though the assumption itself is controversial at present. When several sources are identified in the future, more detailed approaches to the magnetic field can be performed like Ref. \\cite{angelis08}.   In fact, it has shown that the GMF significantly contributes to the deflection angles of UHECRs \\cite{takami06,stanev97,medinatanco98,alvarez02,kachelriess07,yoshiguchi04,takami07b}. Ref. \\cite{takami07b} showed that the trajectories of UHE protons with $\\sim 6 \\times 10^{19}$ eV are deflected by $\\sim 4^{\\circ}$ on average in BS models except for protons arriving from the direction of the Galactic center. The deflection pattern is complicated and dependent on the arrival directions of protons. Also, GMF could affect the arrival directions of protons not only arriving near the Galactic plane but also arriving far from the Galactic plane \\cite{takami07b}. In this situation, we completely neglected GMF and adopted all data without removing events with arrival directions near the Galactic plane as a first step in this study. The statistical analysis taking GMF into account is a next target of our study.   Finally, we comment on the calibration of the energy-scale. When we discussed the cross-correlation using the combined data, the energies of the AGASA data were shifted by 10\\% to lower energies and those of the PAO data were shifted by 50\\% to higher energies on the assumption of the dip-calibration, a physically motivated calibration method \\cite{berezinsky08,aloisio07}. However, the energy-shift of the PAO data is larger than its systematic uncertainty of 22\\% \\cite{pao08,dawson08} while the shift of the AGASA is within its systematic uncertainty \\cite{hayashida00,takeda03}. If the spectral dip at around $10^{19}$eV is generated by pair-creation interaction with the CMB, either/both experiments possibly underestimate the uncertainty on the energy determination. On the other hand, if we assume that the energy-scale of the HiRes is correct, the energy spectra of all UHECR observatories can be consistent by shifting the energies of each experimental data within each systematic uncertainty \\cite{berezinsky08b}. In this standpoint, we can interpret the spectral ankle not as the pair-creation dip but as the transition point from Galactic to extragalactic cosmic rays. The precise determination of the position of the dip enables us to distinguish the two interpretations on the ankle. When the correlation with extragalactic objects is investigated with combined data sets of observations of both the northern and southern hemispheres, as we have done, the calibration of the energy-scale is important because the GZK radius of protons in energy range from $6 \\times 10^{19}$ to $10^{20}$ eV is sensitive to their energy. Precise energy determination is required in order to develop charged particle astronomy. A consensus between the energy-scales of the experiments is required.   Now the PAO and Telescope Array \\cite{fukushima07} are operating and Extreme Universe Space Observatory (JEM-EUSO) \\cite{ebisuzaki07} and the northern site of the PAO \\cite{nitz07} are proposed. These experiments will not only accumulate much more statistics of UHECR events, but also determine the energy-scale precisely by more understanding physics in extensive air shower. We will be able to discuss UHECR sources more precisely in the near future.   \\ack We are grateful to Susumu Inoue and Tokonatsu Yamamoto for useful discussions. The works of H.T., T.N, and K.Y. are 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, Culture, Sports, Science and Technology (MEXT) of Japan through Research Grants S19104006. This work was also supported by World Premier International Research Center Initiative (WPI Initiative), MEXT, Japan.  \\clearpage"
    },
    "0812/0812.1034_arXiv.txt": {
        "abstract": "We report our complete database of X-ray eclipse timings of the low mass X-ray binary \\exo\\ observed by the {\\it Rossi X-Ray Timing Explorer} ({\\it RXTE}) satellite. As of this writing we have accumulated 443 full X-ray eclipses, 392 of which have been observed with the Proportional Counter Array on {\\it RXTE}. These include both observations where an eclipse was specifically targeted and those eclipses found in the RXTE data archive. Eclipse cycle count has been maintained since the discovery of the \\exo\\ system in February 1985. We describe our observing and analysis techniques for each eclipse and describe improvements we have made since the last compilation by \\citeauthor{whw+02} The  principal result of this paper is the database containing the timing results from a seven-parameter fit to the X-ray light curve for each observed eclipse along with the associated errors in the fitted parameters. Based on the standard \\OMC\\ analysis, \\exo\\ has undergone four distinct orbital period epochs since its discovery. In addition, \\exo\\ shows small-scale events in the \\OMC\\ curve that are likely due to short-lived changes in the secondary star. ",
        "introduction": "\\label{section:introduction}  The orbital period is one of the most fundamental parameters characterizing binary star systems. Understanding the evolution of binary systems requires understanding how the orbital period might change on both short and long time scales, both from a theoretical and observational standpoint. As a binary system evolves mass and angular momentum are exchanged between the binary components, angular momentum is carried out of the system via mass loss, and each component can evolve via either nuclear compositional and stellar structural changes. Eclipse transitions provide good fiducial timing markers allowing the precise timing of the orbit which, in turn, make possible the long-term characterization of period changes and the magnitude of the effects these physical processes have on the system parameters.  To date, five low mass X-ray binaries (LMXBs) showing full X-ray eclipses are known (\\exo, \\xonesixfiveeight, XTE$\\,$J1710$-$281, AX$\\,$J1745.6$-$2901, and GRS$\\,$J1747$-$312). Of these sources, only \\exo\\ has been persistently X-ray active over the past two decades and thus available for continuous monitoring of its orbital period behavior \\citep{pwgg86,whw+02}. \\exo\\ is a transient X-ray system with a 3.82 hour orbit period first observed by the European Space Agency's X-ray Observatory ({\\it EXOSAT}) satellite when it became X-ray active in February of 1985 \\citep{pwgg86}. It has remained X-ray active at a level of at least a few mcrab since its discovery. The eclipsing LMXB \\xonesixfiveeight\\ was discovered by the SAS-3 satellite in October 1976 \\citep{lhdl76} and was eventually determined to be a X-ray bursting, eclipsing LMXB with an orbital period of 7.12 hours \\citep{cw89}. \\xonesixfiveeight\\ became undetectable in X-rays in 1979 but then returned to X-ray activity in April 1999, repeatedly being observed by BeppoSAX \\citep{ops+01} and the {\\it Rossi X-Ray Timing Explorer} ({\\it RXTE}) \\citep{wsb00} until it became X-ray inactive again in October 2000. The observations reported by \\citet{wsb00} appeared to indicate that the orbital period was changing with a time scale of $\\tau_{orb} = | \\Porb / \\Porbdot | \\sim 10^7$ years even though little mass transfer had occurred between epochs of X-ray activity. \\xonesixfiveeight\\ has remained X-ray inactive since October 2000 making long term monitoring of its X-ray orbital period impossible.  GRS$\\,$J1747$-$312 in globular cluster Terzan 6, and XTEJ1710$-$281 are both regularly recurring transient sources showing full X-ray eclipses and have orbital periods 12.4 hours and 3.28 hours, respectively. Their ephemerides can be established via X-ray observations \\citep{zhm+03}. However, in the case of GRS$\\,$J1747$-$312, timing observations based on the mid-eclipse timings are difficult because of the relatively long eclipse totality ($\\sim43$ minutes). Eclipse observations must wait through a long duration of totality when virtually no X-ray activity from the source can be seen in order to accurately time both the ingress and egress transitions of any single eclipse. Finally, the transient faint source AX$\\,$J1745.6$-$2901 is very close to \\sgrastar\\ and thus very difficult to observe because of source confusion in that crowded field. Nevertheless, it has been detected by both the Advanced Satellite for Cosmology and Astrophysics ({\\it ASCA}) \\citep{mks+96,skm+02} and by the X-ray Multi-Mirror Mission ({\\it XMM-Newton}) with the 23-minute eclipse variations being visible and the orbital period measured at 8.36 hours \\citep{pgg+07}.  For all of these reasons, \\exo\\ is a uniquely good system for long-term study of LMXB orbital evolution. Thus, early in the {\\it RXTE} mission, we began a program of monitoring the X-ray eclipses in \\exo\\ in an effort to detect orbital period variations. The ultimate goal at that time was to measure the orbital period derivative \\Porbdot\\ and compare this with estimates of magnitude of orbital period changes in theoretical calculations of LMXB evolution. \\exo\\ was observed intensely by the EXOSAT satellite \\citep{pwgg86} resulting in the first estimate of its orbital period and also revealing that, because the source showed X-ray bursts, the compact object in the binary was a neutron star. It should be noted, however, that even the 23-year time baseline reported on here is short compared to the time steps in most evolutionary calculations of binary systems which are in the range $\\Delta t > 10^{4-5}$ years. Theoretical calculations of LMXB evolution must cover a time span of 1-10 Gyrs and time-steps less than $10^4$ years would make such computations prohibitively expensive. Thus, we still have yet to access a sufficiently long baseline to allow robust comparison with most theoretical calculations of LMXB evolution. Conversely, most theoretical calculations do not resolve binary system effects, such as magnetic field generation and cycling in the secondary star, that can potentially influence the system \\OMC\\ behavior on our 23-year observational time scale.  With these monitoring observations we expected to find smooth variations in the \\OMC\\ residuals (observed minus calculated residuals are the delay or advance in the mid-eclipse timings compared to that expected for a specified ephemeris) that could be compared to the estimates of LMXB evolutionary time scales predicted by theory \\citep[e.g.,][]{rjw82}. What we found was considerably more complicated. The first paper in this series, \\citet{hwc95}, concluded that the variable eclipse durations and profiles observed in \\exo\\ implied there is an additional source of uncertainty in timing measurements and that this uncertainty is intrinsic to the binary system, correlated from observation to observation, and had variance that increased as a function of binary cycles between observations. \\citeauthor{hwc95} also suggested that spurious trends in the \\OMC\\ residuals could be misinterpreted as changes in the orbital period. \\citet{hwc97} went further and identified a component of ``intrinsic scatter'' in the \\OMC\\ residuals of $\\sim 0.15$ s per orbit cycle. Finally, \\citet{whw+02} found that neither a constant orbital period derivative or any reasonable polynomial expansion in \\Porb, \\Porbdot, etc., provided an acceptable fit to the entire historical set of \\OMC\\ residuals. Rather, the system appeared to go through ``epochs'' in which the orbital period would change on the order of milliseconds from epoch to epoch. A model for the observed \\OMC\\ variations that included observational measurement error, cumulative orbital period jitter intrinsic to the binary system, with some underlying period evolution was found to be consistent with the observed mid-eclipse timing data and implied a rapid timescale for orbit evolution, $\\tau_{orb} \\sim 2 \\times 10^7$ years. ",
        "conclusions": "\\label{section:conclusions}  We have analyzed 392 full X-ray eclipses from the \\exo\\ system observed by the {\\it RXTE} satellite. We have carefully fitted a seven-parameter model to each eclipse light curve profile in order to determine the mid-eclipse timings and the errors on those timings. The observed \\OMC\\ behavior of the \\exo\\ system we have found is much different than was expected when this project was started. The evolution of the binary system after more than 12 years of {\\it RXTE} observations and eclipse timings is one of multiple orbital period epochs where the orbital period takes on distinct values that change rapidly. We have identified magnetic field cycling of the secondary star as the most likely cause of these \\OMC\\ residual variations and given a brief description of how variations in the secondary's magnetic field might bring about changes in the system orbital period. Also, the \\OMC\\ diagram shows significant intrinsic jitter of various magnitudes during each separate epoch. This jitter most likely reflects subtle changes in the occulting edge of the secondary star caused by magnetic eruptions in the secondary star's chromosphere."
    },
    "0812/0812.3031_arXiv.txt": {
        "abstract": "Neutrino and antineutrino fluxes from a core-collapse galactic supernova are studied, within a representative three-flavor scenario with inverted mass hierarchy and tiny 1-3 mixing. The initial flavor evolution is dominated by collective self-interaction effects, which are computed in a full three-family framework along an averaged radial trajectory. During the whole time span considered ($t=1$--20~s), neutrino and antineutrino spectral splits emerge as dominant features in the energy domain for the final, observable fluxes. The main results can be useful for SN event rate simulations in specific detectors. Some minor or unobservable three-family features (e.g, related to the muonic-tauonic flavor sector), as well as observable effects due to variations in the spectral input, are also discussed for completeness. ",
        "introduction": "%   The fluxes of $\\nu$ and $\\overline\\nu$ from a core-collapse supernova (SN) can encode very interesting information about both the (anti)neutrino properties \\cite{Di08} and the SN explosion mechanism \\cite{Expl}.  Assuming that the latter is understood, one may  try to detect distinctive features of flavor change in either the energy ($E$) or time ($t$) spectra of the observable fluxes $F_\\alpha (E,\\,t)$, as compared with the ``featureless,'' unoscillated fluxes $F_\\alpha^0(E,\\,t)$, for one or more neutrino species $\\nu_\\alpha$.  After the recent, seminal work in \\cite{Coll,Semi}, it has been fully realized that flavor change phenomena driven by (anti)neutrino self-interactions \\cite{Pant} may induce dramatic, observable effects in a variety of SN scenarios, especially for inverted mass hierarchy \\cite{Pend}; see \\cite{Di08,Nonl} for reviews of this rapidly growing research field. In particular, the so-called spectral split \\cite{Split} or swap \\cite{Swap} emerges as a distinctive feature in the neutrino \\cite{Miri} and possibly the antineutrino \\cite{Miri,Tamb} energy spectra. If this feature dominates, then a  stepwise flavor conversion of $\\nu_e$ develops across a certain critical energy $E_c$ \\cite{Analysis}, namely, \\begin{equation} \\label{P} P_{ee}(E)\\simeq \\left\\{ \\begin{array}{ll} 1 & (\\mathrm{for\\ }E<E_c)\\ ,\\\\ 0 & (\\mathrm{for\\ }E>E_c)\\ , \\end{array} \\right. \\end{equation} in terms of the survival probability $P_{ee}=P({\\nu_e\\to\\nu_e})$ at the end of collective effects. A somewhat similar feature has been observed in the $\\overline\\nu_e$ sector \\cite{Miri,Tamb}, although with at different (generally smaller) split energy $\\overline E_c$, \\begin{equation} \\overline P_{ee}(E)\\simeq \\left\\{ \\begin{array}{ll} 1 & (E<\\overline E_c)\\ ,\\\\ 0 & (E>\\overline E_c)\\ . \\end{array} \\right. \\label{Pbar} \\end{equation} In a strictly $2\\nu$ framework with flavors ``$e$'' and ``$x$'', the net result of collective effects would be a complete interchange, or ``swap,'' of the fluxes $F_\\alpha$ between ($\\nu_e,\\,\\nu_x$) above $E_c$ \\cite{Split,Swap,Miri}, as well as between $(\\overline\\nu_e,\\,\\overline\\nu_x)$  above $\\overline E_{c}$ \\cite{Tamb}.    Analogously, in a three-flavor scenario with both $\\nu_\\mu$ and $\\nu_\\tau$  acting as a flavor ``$x$,'' \\begin{eqnarray} 2F_x \t\t\t&=& F_\\mu + F_\\tau\\ ,\\\\ 2\\overline F_x \t&=& \\overline F_\\mu + \\overline F_\\tau\\ , \\end{eqnarray} and with usual initial conditions \\begin{equation} F^0_\\mu=F^0_\\tau= F^0_x= \\overline F^0_x=\\overline F^0_\\tau=\\overline F^0_\\mu\\ , \\end{equation} a swap is expected to occur between the $e$ flavor and one of the two $x$ flavors, namely, \\begin{eqnarray} \\label{je} F'_e &=& F_e^0 P_{ee} + F^0_x(1-P_{ee}) \\simeq \\left\\{\\begin{array}{ll}F^0_{e} & (E<E_c)\\ ,\\\\ F^0_x & (E>E_c)\\ ,\\end{array} \\right.\\\\ \\label{jx} \\overline F'_{ e} &=& \\overline F_e^0 \\overline P_{ee} + \\overline F^0_x(1-\\overline P_{ee})\\simeq  \\left\\{\\begin{array}{ll}\\overline F^0_{e} & (E<\\overline E_c)\\ , \\\\ \\overline F^0_{ x} & (E>\\overline E_c)\\ ,\\end{array} \\right. \\label{jebar}\\\\ 2F'_x &=& [F_e^0(1-P_{ee})+F_x^0P_{ee}]+F_x^0 \\simeq  \\left\\{\\begin{array}{ll}2 F^0_{x} & (E<E_c)\\ ,\\\\ F^0_e + F^0_x & (E>E_c)\\ ,\\end{array} \\right.\\\\ 2\\overline  F'_{x}  &=& [\\overline F_e^0(1-\\overline P_{ee})+\\overline F_x^0\\overline P_{ee}]+\\overline F_x^0 \\simeq  \\left\\{\\begin{array}{ll} 2 \\overline F^0_{x} & (E<\\overline E_c)\\ , \\\\ \\overline F^0_{e} + \\overline F^0_{ x} & (E>\\overline E_c)\\ .\\end{array} \\right. \\label{jxbar} \\end{eqnarray} In the above equations, the primes denote fluxes at the end of collective effects within the SN, to be further evolved up to the exit from the SN and to the final detector.  It is important to test the above $3\\nu$ expectations within comprehensive three-flavor calculations, for various reasons. First, one has to show that the assumed dominance of spectral split phenomena can indeed be manifest (without being disrupted by ordinary matter effects) in a sufficiently general, uncontrived and interesting supernova $3\\nu$ scenario. Second, the effective factorization of one family out of  three {\\em self-interacting families\\/} (leading to nonlinear equations) is never totally obvious, and ``effective $2\\nu$'' expectations like those in Eqs.~(\\ref{je})--(\\ref{jxbar}) need to be explicitly checked. Third, in the self-interaction context, the probabilities $P_{\\alpha\\beta}$ depend, among other things, on the initial conditions and on the absolute fluxes---which do change if the SN energy luminosity (which also varies in time) is distributed over $3\\nu$ rather than $2\\nu$. Finally, while the $\\nu$ spectral split phenomenon is robust and largely understood in terms of lepton number conservation and adiabatic flavor evolution \\cite{Simple,Smirnov}, the $\\overline \\nu$ split seems to be more fragile and related to (not completely understood) nonadiabatic aspects of the evolution \\cite{Tamb}; it is thus worth checking its appearance in a full $3\\nu$ calculation.  Several recent works have focussed on collective effects in $3\\nu$ scenarios \\cite{Neutroniz,Esteban,Das3,Full3,Basu,Earth,SNDB,Gava} and have successfully recovered spectral split features in the neutrino sector (and perhaps also in the antineutrino sector \\cite{Das3,Gava}) in inverted hierarchy. We think it useful to add an independent contribution to this very recent research field, by discussing a SN $3\\nu$ scenario where both the $\\nu$ and $\\overline\\nu$ split features are shown to emerge clearly for a relatively long time of $\\sim\\! 20$ seconds after SN explosion. Such a scenario is thus particularly suited to prospective experimental tests in high-statistics, time-integrated energy spectra of events from the next galactic core-collapse SN.  Our work is structured as follows. In Sec.~2 we describe a representative SN neutrino framework, where collective effects (in the form of spectral splits) are expected to dominate over ordinary matter effects. In Sec.~3 we discuss the $3\\nu$ formalism used to compute the flavor evolution of the initial fluxes $F^0_\\alpha$. In Sec.~4 we present our results for the fluxes $F'_\\alpha$ at the end of collective effects, and show that they confirm the simple expectations in Eqs.~(\\ref{je})--(\\ref{jxbar}); we also discuss minor or unobservable features of our calculations. In Sec.~5 we complete the flavor evolution of the  observable fluxes $F_\\alpha$, by accounting for final  phase-averaging and matter effects.  In Sec.~6 we discuss the effects of some variations in the input SN neutrino spectra, with respect to our default scenario. Section~7 summarizes our results.  The reader is reminded that our results, as well as many other observable features discussed in the growing literature of self-interaction neutrino effects in supernovae, must be taken with a grain of salt. Indeed, many open questions are still open from the viewpoint of the theory (validity of the mean-field approximation in the neutrino evolution equations), of the SN explosion energetics (neutrino luminosities and energy spectra) and geometry (asymmetries, turbulence), and of the neutrino evolution numerics (convergence and robustness of calculations). Results which currently appear to be rather generic (such as the spectral split phenomena in inverse hierarchy) might be unpredictably challenged by improvements and further understanding in any of the above issues. Hopefully, the next galactic supernova explosion(s) will help to reduce the current level of uncertainty in the physics and astrophysics hidden in the expected neutrino signal. ",
        "conclusions": "Building upon recent literature on $3\\nu$ collective effects in core-collapse supernovae \\cite{Neutroniz,Esteban,Das3,Full3,Basu,Earth,SNDB,Gava}, we have performed an independent study of three-flavor evolution of neutrinos and antineutrinos within a rather ``standard'' SN model (Figs.~1 and 2). Then, assuming inverted neutrino mass hierarchy and tiny $\\theta_{13}$ (i.e., strongly nonadiabatic $H$-resonance), self-interaction effects are expected to provide the dominant spectral features, in the form of ``splits''---or ``stepwise swaps''---for both $\\nu$ and $\\overline\\nu$.  We have explored this SN $3\\nu$ scenario by solving the evolution equations for the $3\\nu$ density matrix in vector form and single-angle approximation, after discretization in energy. The numerical results for the evolved fluxes at the end of collective effects (Fig.~3) confirm basic expectations at low and high energy [Eqs.~(\\ref{je})--(\\ref{jxbar})] in terms of unevolved fluxes (Fig.~4). Effects of the ``solar'' mass splitting and of the $\\nu_\\mu$-$\\nu_\\tau$ interaction energy difference in matter are found to be negligible; effects due to variations of $\\theta_{23}$ are large (Fig.~5) but unobservable. The final fluxes at the detector (Fig.~6) are obtained by standard application of adiabatic $L$-resonance effects.  Both $\\nu$ and $\\overline\\nu$ spectral split features tipically emerge with similar characteristics in the whole time interval considered ($t=1$--20~s). Observations of such features (if realized in nature) and of their possible variations require low-energy thresholds, where the signal is suppressed by low cross sections; however, time integration can partly overcome the suppression, without necessarily canceling the persistent spectral features. If the collected statistics and the energy resolution allow, spectra collected at different times might even reveal migrations of the split energies due to time variations of neutrino temperatures. The low-energy frontier in SN neutrino physics may thus be the key to access the physics of inverted mass hierarchy, especially if $\\theta_{13}$ is very small.     \\medskip \\medskip \\medskip"
    },
    "0812/0812.3388_arXiv.txt": {
        "abstract": "In cyclic universe models based on a single scalar field ({\\it e.g.}, the radion determining the distance between branes in M-theory), virtually the entire universe makes it through the ekpyrotic smoothing and flattening phase, bounces, and enters a new epoch of expansion and cooling. This stable evolution cannot occur, however, if scale-invariant curvature perturbations are produced by the entropic mechanism  because it requires two  scalar fields ({\\it e.g.}, the radion and the Calabi-Yau dilaton) evolving along an unstable classical trajectory. In fact, we show here that an overwhelming fraction of the universe fails to make it through the ekpyrotic phase; nevertheless, a sufficient volume survives and cycling continues forever provided the dark energy phase of the cycle lasts long enough, of order a trillion years.   Two consequences are a new role for dark energy and a global structure of the universe radically different from that of eternal inflation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.3494_arXiv.txt": {
        "abstract": "Polarized light provides the most reliable source of information at our disposal for diagnosing the physical properties of astrophysical plasmas, including the three-dimensional (3D) structure of the solar atmosphere. Here we formulate and solve the 3D radiative transfer problem of the linear polarization of the solar continuous radiation, which is principally produced by Rayleigh and Thomson scattering. Our approach takes into account not only the anisotropy of the solar continuum radiation, but also the symmetry-breaking effects caused by the horizontal atmospheric inhomogeneities produced by the solar surface convection. We show that such symmetry-breaking effects do produce observable signatures in $Q/I$ and $U/I$, even at the very center of the solar disk where we observe the forward scattering case, but their detection would require obtaining very-high-resolution linear polarization images of the solar surface. Without spatial and/or temporal resolution $U/I{\\approx}0$ and the only observable quantity is $Q/I$, whose wavelength variation at a solar disk position close to the limb has been recently determined semi-empirically. Interestingly, our 3D radiative transfer modeling of the polarization of the Sun's continuous spectrum in a well-known 3D hydrodynamical model of the solar photosphere shows remarkable agreement with the semi-empirical determination, significantly better than that obtained via the use of one-dimensional (1D) atmospheric models. Although this result confirms that the above-mentioned 3D model was indeed a suitable choice for our Hanle-effect estimation of the substantial amount of ``hidden\" magnetic energy that is stored in the quiet solar photosphere, we have found however some small discrepancies whose origin may be due to uncertainties in the semi-empirical data and/or in the thermal and density structure of the 3D model. For this reason, we have paid some attention also to other (more familiar) observables, like the center-limb variation of the continuum intensity, which we have calculated taking into account the scattering contribution to the continuum source function. The overall agreement with the observed center-limb variation turns out to be impressive, but we find a hint that the model's temperature gradients in the continuum forming layers could be {\\em slightly} too steep, perhaps because all current simulations of solar surface convection and magnetoconvection compute the radiative flux divergence ignoring the fact that the effective polarizability is not completely negligible, especially in the downward-moving intergranular lane plasma. ",
        "introduction": "The main aim of this investigation is to formulate and solve the 3D radiative transfer problem of the linear polarization of the Sun's continuous spectrum, which results principally from Rayleigh scattering at neutral hydrogen and Thomson scattering at free electrons. Our scientific motivation is twofold. On the one hand, we want to use the polarization of the Sun's continuous radiation as a diagnostic tool of the thermal and density structure of the solar photosphere. On the other hand, a realistic modeling of the continuum polarization is important for quantitative analysis of the linearly-polarized solar limb spectrum.  As we shall see below, the polarization of the solar continuum radiation is largely determined by the effective polarizability and by the radiation field's anisotropy, which in turn depend on the density and thermal structure of the stellar atmosphere under consideration. Collisional and/or magnetic depolarization do not play any role on the polarization of the continuum radiation of the Sun's visible spectrum. Therefore, the more realistic is the thermal and density structure of a given solar atmospheric model, the closer to the empirical data will be the calculated linear polarization of the solar continuum radiation.  While our present telescopes and polarimeters can reach very high accuracy in determining the relative polarization variations between line and continuum features in the solar spectrum (e.g., Stenflo \\& Keller 1997), they are not suitable for  determining the zero-point of the polarization scale with a similar polarization precision. Nevertheless, Stenflo (2005) could extract the polarization of the Sun's continuous spectrum from Gandorfer's (2000, 2002, 2005) atlas of the linearly-polarized solar limb spectrum. As emphasized by Ivanov (1991), this so-called second solar spectrum contains multitude of fractional linear polarization ($Q/I$) features seen both ``in emission\" and ``in absorption\", i.e., with correspondingly larger and smaller polarization than in the adjacent continuum\\footnote{We recall that the Stokes $I$ parameter represents the intensity at a given wavelength, while Stokes $Q$ is the intensity difference between linear polarization parallel and perpendicular to a given reference direction (e.g., the direction perpendicular to the straight line joining the disk center with the observed point). Stokes $U$ would be then the intensity difference between linear polarization at $+45^{\\circ}$ and $-45^{\\circ}$ with respect to the chosen reference direction, where positive angles are measured counterclockwise by an observer facing the radiation.}. While many of such $Q/I$ signals are dominated by the {\\em selective emission} and {\\em selective absorption} processes caused by the radiatively induced quantum coherences in the energy levels of the atoms and molecules of the solar atmosphere (e.g., Trujillo Bueno 2001), many of the ``absorption\" features result from depolarizing line contributions which depress the continuum polarization level. Stenflo's (2005) determination of the polarization of the Sun's continuous spectrum is {\\em semi-empirical} because it was based on a model for the behavior of these spectral lines that depolarize the continuum polarization. His strategy consisted in choosing spectral windows dominated by purely depolarizing lines and in modeling the observed relative line depth in $Q/I$ in terms of the relative line depth in Stokes $I$. He pointed out that his semi-empirically determined continuum polarization lies systematically lower than the values predicted by Fluri \\& Stenflo (1999) for $\\lambda{>}4000$ \\AA\\ from radiative transfer modeling in 1D  semi-empirical models of the solar atmosphere.  Probably, the most interesting result of Stenflo's (2005) semi-empirical determination of the wavelength dependence of the continuum polarization at $\\mu={\\rm cos}{\\theta}\\,{\\approx}\\,0.1$ (with $\\theta$ the heliocentric angle) is the confirmation of the 1979 theoretical expectation by L. Auer (see Stenflo, Twerenbold \\& Harvey 1983; Stenflo 2003) that it must show a prominent jump around the Balmer limit at 3646 \\AA. At larger  wavelengths the dominant contribution to the polarization of the solar continuous spectrum is Rayleigh scattering at neutral hydrogen (Chandrasekhar 1960), which in turn is dominated by scattering in the distant line wings of the Lyman series lines (Stenflo 2005). At visible wavelengths this Rayleigh scattering at neutral hydrogen is the second most important opacity source after that of the (non-polarizing) H$^{-}$ absorption. As it happens with other observables of the Sun's continuous spectrum (e.g., the center-limb variation and the contrast of solar granulation) a prominent jump should be seen also in the scattering polarization when the extra contribution to the continuum opacity due to Balmer bound-free transitions becomes operative, because this contribution is larger than the Rayleigh scattering one. Interestingly enough, Stenflo's (2005) determination shows that this Balmer jump is indeed produced in the polarization of the Sun's continuous spectrum, but approximately at a wavelength 80 \\AA\\ larger than the nominal position of the Balmer series limit. According to Stenflo (2005) this can be understood in terms of Stark broadening of the Balmer bound-bound transitions, which become very crowded when approaching the Balmer limit and merge into a quasi-continuum whose ensuing opacity becomes larger than the Rayleigh scattering opacity well before the Balmer series limit is reached.  In this paper we take the opportunity of contrasting Stenflo's (2005) determination of the polarization of the Sun's continuous spectrum between 3000 \\AA\\ and 7000 \\AA\\ with detailed radiative transfer simulations in a ``realistic\"  3D snapshot model of the solar photosphere, which resulted from Asplund et al.'s (2000) hydrodynamical simulations of solar surface convection. This is the same 3D snapshot model we used in our investigations of the Hanle effect in the Sr {\\sc i} line at 4607 \\AA, which allowed us to conclude that in the ``quiet\" regions of the solar photosphere there is a substantial amount of hidden magnetic energy carried by tangled magnetic fields at subresolution scales (Trujillo Bueno, Shchukina \\& Asensio Ramos 2004; Trujillo Bueno, Asensio Ramos \\& Shchukina 2006; Trujillo Bueno \\& Shchukina 2007). As we shall see below, our 3D radiative transfer modeling of the polarization of the Sun's continuous spectrum is in very good agreement with Stenflo's (2005) semi-empirical determination. However, there are some small discrepancies, whose origin may be due to various possible reasons (e.g., uncertainties in Stenflo's  semi-empirical determination and/or in the thermal and density structure of the 3D model). In order to investigate this issue from a complementary perspective, we have considered also other (more familiar) observables, like the center-limb variation of the continuum intensity and histograms of the intensity fluctuations produced by the solar granulation pattern, which we have calculated taking into account the scattering contribution to the continuum source function.  The outline of this paper is the following. The formulation of the continuum polarization problem in a 3D medium like that of the solar atmosphere is presented in \\S2, where we establish the relevant equations and mention how we have solved them. After illustrating and discussing in \\S3 the atmospheric models we have chosen for this investigation, we show in \\S4 the spatial variation of the radiation field tensors (which quantify the symmetry properties of the radiation field), while in \\S5 we show both the spatial and the wavelength variation of the effective polarizability. The center-limb variation of the Sun's continuous spectrum is carefully considered in \\S6, while \\S7 focuses on confronting the calculated histogram of the intensity distribution of the model's granulation with that observed on the Sun with the {\\em Hinode} space telescope. \\S8 presents a careful investigation of the wavelength dependence of the anisotropy factor, which is one of the fundamental quantities for scattering polarization. Our calculations of the polarization of the Sun's continuous spectrum are presented in \\S9, where we compare them with Stenflo's (2005) semi-empirical determination. The interesting polarization effects produced by the breaking of the axial symmetry of the photospheric radiation field are discussed in \\S10. Finally, \\S11 summarizes our main conclusions with an outlook to future research. ",
        "conclusions": "We have formulated and solved the 3D radiative transfer problem of the polarization of the stellar continuous radiation, which in solar-like atmospheres is principally produced by the interaction between radiation and hydrogen plus free electrons (Rayleigh and Thomson scattering). Of particular interest is that we have taken into account not only the anisotropy of the stellar continuum radiation, but also the symmetry-breaking effects caused by the horizontal atmospheric inhomogeneities produced by the stellar surface convection. As shown in the last figure of this paper, detection of the observational signatures of such symmetry-breaking effects would require observing the linear polarization of the Sun's continuum radiation at high spatial and temporal resolution, so as to be able to see clearly the details of the solar granulation pattern. In particular, at the very center of the solar disk, where we observe the forward scattering case, the predicted $Q/I$ and $U/I$ signals have a subgranular pattern which is solely due to the symmetry breaking effects we have discussed in \\S10 (see also Trujillo Bueno \\& Shchukina 2007, for the impact of such effects on the scattering polarization of the Sr {\\sc i} 4607 \\AA\\ line).  Without spatial and/or temporal resolution $U/I{\\approx}0$ and the only observable quantity for the solar case is of course $Q/I$, whose wavelength variation at a disk position close to the limb (i.e., at $\\mu{\\approx}0.1$) was extracted semi-empirically by Stenflo (2005) from Gandorfer's (2000; 2002, 2005) atlas of the second solar spectrum. The fact that collisional and/or magnetic depolarization do not play any role on the linear polarization of the continuum radiation of the Sun's visible spectrum implies that we can use the continuum polarization to evaluate the degree of realism of any given solar photospheric model, such as the 3D snaphot model we have taken from Asplund's et al. (2000) hydrodynamical simulations. Overall, our 3D radiative transfer modeling of the polarization of the solar continuum radiation in such a 3D model shows a notable agreement with Stenflo's (2005) semi-empirical data, significantly better than that obtained via the use of 1D atmospheric models. Of especial interest is to note that the amplitude of the Balmer jump in $Q/I$ is larger in 3D than in 1D, which is in line with the sizable amplitude of the Balmer jump seen in Stenflo's (2005) semi-empirical values.  Given that our 3D modeling of the polarization of the Sun's continuous radiation gives us Stokes $I$, in addition to Stokes $Q$ and $U$, we have taken the opportunity of investigating also the reliability of the above-mentioned 3D snapshot model by considering other (more familiar) quantities, like the absolute continuum intensities (Fig. 2), the center-limb variation of the continuum intensity (Fig. 7), the RMS contrast of solar granulation at the 6301 \\AA\\ continuum wavelength observed with the spectropolarimeter of the {\\em Hinode} space telescope (Fig. 9), and histograms of the ensuing continuum intensity fluctuations (Fig. 10). In particular, Fig. 7 contrasts the observed and the calculated wavelength variation of the center-limb variation of the solar continuum intensity. Although the overall agreement is indeed impressive, there are however some small but significant discrepancies in the center-limb curves with $0.2 {\\le} {\\mu} {\\le} 0.5$ and for wavelengths around $\\lambda=5000$ \\AA, which support the possibility that the temperature gradients in the continuum forming layers of Asplund et al's (2000) 3D model could be {\\em slightly} too steep. We emphasize that we have done this type of investigations by using Eq. (5) for $S_I$ --that is, taking into account that the continuum source function is not exactly equal to the Planck function, because the scattering contribution is not completely negligible, especially for wavelengths sensibly smaller than 5000 \\AA. In this respect, it is of interest to note that around the visible ``surface\" layers the effective polarizability in the downward-moving intergranular lane plasma is significantly larger than in the upward-moving granule plasma, with values of the order of $1\\%$ at 4000 \\AA\\ (see Fig. 4).  The reported small discrepancies between the measured and the calculated observables might be due to the fact that the radiative flux divergence at each time step in Asplund et al's (2000) simulations of solar surface convection was computed assuming $S_I=B_{\\nu}(T)$ (i.e., neglecting the scattering contribution to the continuum source function). Another possibility could be a plausible impact of the unresolved magnetic field that Trujillo Bueno et al. (2004; 2006) inferred via a joint analysis of the Hanle effect in the Sr {\\sc i} 4607 \\AA\\ line and in the C$_2$ lines of the Swan system. These authors concluded that the bulk of the ``quiet\" solar photosphere is teeming with tangled magnetic fields at subresolution scales, with a mean field strength of the order of 100 gauss, which might be important for the overall energy balance of the solar atmosphere. As shown in this paper, the agreements between the observed and the calculated center-limb variation (see Fig. 7) and between the inferred and the calculated continuum polarization (see the left panel of Fig. 12) are so good that we think that our conclusion on the presence of a substantial amount of ``hidden\" magnetic energy in the quiet solar photosphere will remain valid after similar Hanle-effect investigations are carried out using instead the forthcoming, new generation of 3D models (which will probably go beyond the current approach for calculating the radiative flux divergence). Such future models of stellar surface convection and magnetoconvection will be useful for clarifying whether or not the Sun has a sub-solar metallicity, but for further progress in our understanding of the small-scale magnetic activity of our nearest star it is important to perform increasingly realistic solar surface dynamo simulations and to use them as models of the solar photosphere.  \\newpage"
    },
    "0812/0812.0945_arXiv.txt": {
        "abstract": "{The observations of the proplyds in the Orion Nebula Cluster, showing clear evidence of ongoing photoevaporation, have provided a clear proof about the role of the externally induced photoevaporation in the evolution of circumstellar disks. NGC~6611 is an open cluster suitable to study disk photoevaporation, thanks to its large population of massive members and of stars with disk. In a previous work, we obtained evidence of the influence of the strong UV field generated by the massive cluster members on the evolution of disks around low-mass Pre-Main Sequence members. That work was based on a multi-band BVIJHK and X-ray catalog purposely compiled to select the cluster members with and without disk.} {In this paper we complete the list of candidate cluster members, using data at longer wavelengths obtained with Spitzer/IRAC, and we revisit the issue of the effects of UV radiation on the evolution of disks in NGC~6611.} {We select the candidate members with disks of NGC~6611, in a field of view of $33^{\\prime}\\times34^{\\prime}$ centered on the cluster, using IRAC color-color diagrams and suitable reddening-free color indices. Besides, using the X-ray data to select Class~III cluster members, we estimate the disks frequency vs. the intensity of the incident radiation emitted by massive members.} {We identify 458 candidate members with circumstellar disks, among which 146 had not been revealed in our previous work. Comparing of the various color indices we used to select the cluster members with disk, we claim that they detect the excesses due to the emission of the same physical region of the disk: the inner rim at the dust sublimation radius. Our new results confirm that UV radiation from massive stars affects the evolution of nearby circumstellar disks.} {} ",
        "introduction": "\\label{intro} Circumstellar disks around young Pre-Main Sequence (PMS) stars have been deeply studied since the discovery of their infrared emission \\citep{Mendo66,Mendo68}. Until now, several disk models have been developed and a large number of star-forming regions have been observed, in order to understand the physical properties and the evolution of circumstellar disks. Several authors (e.g. \\citealt{Ha01}) claim that the emission from the inner region of the disks, responsible for NIR excess in T-Tauri stars, declines on a timescale between $\\sim$1 and $\\sim$10 Myr. Important clues about the typical timescale of disk evolution come from the studies of PMS stars that have cleared the inner region of their circumstellar disks, thus showing excesses in mid- and far-infrared but not in near-infrared bands. These stars with disk are usually considered to be in a transitional phase between Class~II and disk-less Class~III PMS stars.\\par The evolution of circumstellar disks can be influenced by nearby massive stars, as it was shown by the observations of the evaporating protoplanetary disks in the Trapezium, near the massive star $\\Theta^1$ Ori (see, for example, \\citealp{Stor99} and references therein). In this case, the photoevaporation process, induced by UV radiation emitted by $\\Theta^1$ Ori, dissipates the nearby disks on timescales shorter than $10^6$ years. On the other hand, several authors (for example \\citealp{Eis06}) claim that the externally induced photoevaporation is not an efficient mechanism for the truncation of circumstellar disks around low mass PMS stars. The debate about this topic is still open.\\par Comparing the spatial variations of disk frequencies respect to the positions of massive stars in young open clusters it is possible to study how the radiation from massive stars affects the evolution of circumstellar disks and the formation of new generations of stars. \\citet{Bal07} and \\citet{io07}, hereafter GPM07, applied this approach to two young clusters: NGC~2244 and NGC~6611 respectively. In their Spitzer/IRAC-MIPS study of NGC~2244 (distant 1.4-1.7 Kpc from the Sun and with an age of about 2-3 Myr), \\citet{Bal07} found that disk frequency drops quickly at a distance of 0.5 parsec from the massive cluster members. For larger distances, however, stars with disk do not experience effects induced by massive stars, since the disk frequency is not spatially correlated with their positions. \\par Also the young open cluster NGC~6611, in the Eagle Nebula, is an ideal target for this kind of study, thanks to its rich PMS population and its large number of massive members (56 with spectral classes earlier than B5, \\citealp{Hil93}) that are distributed irregularly in the central region of the cluster. In order to obtain a reliable list of cluster members GPM07 used a membership criterion based on infrared excesses (to select members with disk) and X-ray emission (to select members without disk). In fact, thanks to the high X-ray luminosity of PMs stars, this criterion allows to obtain an unbiased sample of PMS cluster members. With this method, a total of 1122 candidate PMS stars associated to NGC~6611 or to the outer part of the Eagle Nebula were identified. In NGC~6611, GPM07 found evidences that the UV radiation from massive stars have dissipated the disks around nearby PMS members on short timescales, since the disks frequency declines for small distances from massive members. To obtain this result, GPM07 calculated the flux emitted by all the massive stars incident on each cluster members and then obtained the disk frequency for various bins of incident flux. To this calculation, the projected distances were used, instead the real ones. However, GPM07 showed that this approximation have not affected their result. \\par The effects of the energetic radiation from massive stars on the evolution of nearby circumstellar disks are also studied with simulations of the evolution of open clusters. For example \\citet{Ada06} claim that externally induced photoevaporation does no affect the evolution of disks in clusters with less than 1000 members. Instead, more populated clusters, as shown by \\citet{Fatu08}, can be environments in which disks are efficiently evaporated in small timescales, since they have intense UV fields practically independent from the number of stellar members. \\par This paper is the follow up of the GPM07 study, enriching the list of cluster members by using Spitzer/IRAC data and confirming their principal results; both works, therefore, will be the base of a subsequent paper, focused on the analysis of the Spectral Energy Distributions of selected cluster members. \\par The Spitzer/IRAC observations of NGC~6611 used in this work are obtained from the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire (GLIMPSE; \\citealt{Ben03}). The {\\it Infrared Array Camera} (IRAC), mounted on Spitzer Space Telescope \\citep{Faz04}, allows studying stars with circumstellar disk in four infrared bands, centered on 3.6~$\\mu m$, 4.5~$\\mu m$, 5.8~$\\mu m$ and 8.0~$\\mu m$. In this region of the spectrum, the contribution from the stellar photosphere is usually small compared to the emission from disk and envelope, and the effects of extinction by interstellar medium is smaller than at shorter wavelengths. These facts permit a more reliably selection of stars with circumstellar disks than using 2MASS data alone. \\par Our paper is organized as follow: in Sect. \\ref{catal} we show the results of the cross identification between GLIMPSE catalog and that compiled in GPM07; in Sect. \\ref{colcolsp} and \\ref{Qpar} we will describe the identification of stars with disk respectively using IRAC color-color diagram and suitable reddening-free color indices that combine optical and NIR colors; in Sect. \\ref{compat} we will compare the two diagnostics used in this paper; in Sect. \\ref{memcat} we describe the catalog of the cluster members; in Sect. \\ref{final} we will re-examine the results obtained in GPM07 using the updated member list.\\par ",
        "conclusions": "\\label{thatsallfolks}  In this paper, we analyze Spitzer/IRAC data of the young open cluster NGC~6611, in the Eagle Nebula, in a large field wide approximately $30^{\\prime}\\times30^{\\prime}$. In a previous paper, we have already selected the members of this cluster using our published BVIJHK and X-ray multi band catalog. The obtained list of cluster members, both with and without disk, is used to verify that the former are more frequent at larger distances from the massive members of the cluster. \\par In this work, in order to select new members with disk, we use and compare two different disk diagnostics: the IRAC color-color diagram, that use all the IRAC bands simultaneously, and four suitable reddening-free color indices ($Q$ indices), each defined to detect the excesses in a specific IRAC band. \\par Using the IRAC color-color diagram, we identify 182 Young Stellar Objects (147 Class~II; 13 Class~I and 22 with colors compatible with both classifications). This diagnostic is very efficient to select largely embedded young sources, but it can be used only for stars with good photometry in all IRAC bands. We discuss how significantly this affects the selection of faint YSOs. With the use of $Q$ indices we partially overcome this problem, since they allow us to select stars with excesses in each IRAC band taken alone; with this diagnostic, we select 174 YSOs with excesses in IRAC bands. Among them, 66 stars cannot be detected with the IRAC color-color plane, since they are not well measured in all the IRAC bands.\\par Combining the outcome of both diagnostics, we identify 146 new cluster members with disk with respect to our previous work, which was based on 2MASS photometry alone (116 from the IRAC color-color diagram and 30 from $Q$ indices). At the end, our catalog includes a total of 474 candidate Pre-Main Sequence members with disk and 790 without disk (the latter only in the central $17^{\\prime}\\times17^{\\prime}$ field, where we found 118 X-ray sources with disk). \\par Comparing all the $Q$ indices defined with 2MASS and IRAC bands (with the exception of [8.0] band), we claim that they are all sensitive to the emission from the same physical region of the disk, namely the inner rim at the dusts sublimation radius. These indices will be efficiently used, then, to select Class~II stars that are not largely embedded (i.e. for which it is possible to detect the photospheric emission) and whose inner disk is not evacuated. The embedded YSOs associated to the cluster can be selected only thank to the IRAC color-color diagram, and not with the $Q$ indices.\\par We discuss evidence suggesting that star formation activity is ongoing in the outer region of the Eagle Nebula, and not only in the center of NGC~6611 as showed by previous works. For example, the diagnostic based on IRAC color-color diagram allow us to identify a probable new star formation site, rich of Class~I and embedded Class~II objects, at North with respect to the center of the cluster. \\par In the central $17^{\\prime}\\times17^{\\prime}$ field we find an average disk fraction equal to $24\\% \\pm 2\\%$. In this field, the disk frequency also has a strong dependence on the incident radiation emitted by massive members, varying from $31\\% \\pm 4\\%$ to $16\\% \\pm 3\\%$ across the whole range of values of incident flux. This is an evidence of the influence of massive stars radiation on the evolution of circumstellar disks and star formation process, as already shown in GPM07."
    },
    "0812/0812.2246_arXiv.txt": {
        "abstract": "We calculate the probability distribution for the volume of the Universe after slow-roll inflation both in the eternal and in the non-eternal regime. Far from the eternal regime the probability distribution for the number of $e$-foldings, defined as one third of the logarithm of the volume, is sharply peaked around the number of $e$-foldings of the classical inflaton trajectory. At the transition to the eternal regime  this probability is still peaked (with the width of order one $e$-folding) around the  average, which gets twice larger at the transition point. As one enters the eternal regime the probability for the volume to be finite rapidly becomes exponentially small. In addition to developing techniques to study eternal inflation, our results allow us to establish the quantum generalization of a recently proposed bound on the number of $e$-foldings in the non-eternal regime: the probability for slow-roll inflation to produce a finite volume larger than $e^{S_{dS}/2}$, where $S_{dS}$ is the de~Sitter entropy at the end of the inflationary stage, is smaller than the uncertainty due to non-perturbative quantum gravity effects. The existence of such a bound provides a consistency check for the idea of de Sitter complementarity. ",
        "introduction": "The Universe is accelerating today~\\cite{SSTC,SCPC,WMAP}, and it is extremely likely that it was experiencing a period of accelerated expansion (inflation) back in the past~\\cite{WMAP,Guth:1980zm,Linde:1981mu,Albrecht:1982wi}. In both cases the pressure to density ratio is very close to $-1$ and the local geometry is very close to that of de~Sitter (dS) space.  It is plausible that both these periods of inflation are eternal, {\\it i.e.}  some space-time regions keep inflating forever. Indeed the most economical explanation for the cosmic acceleration observed now is that we are stuck in a metastable vacuum \\cite{Guth:1982pn}, and this results in eternal inflation unless the vacuum decay rate $\\Gamma$ is faster than the expansion rate of the Universe, $\\Gamma\\gtrsim H_0$~\\footnote{ This latter possibility appears rather unlikely and fine-tuned given the non-perturbative nature of the decay, although it may receive support from future particle physics data \\cite{ArkaniHamed:2008ym}.}. Also there is strong evidence that in the past we underwent a phase of inflation driven by a rolling scalar field, which could have been preceeded by a period of eternal inflation, as predicted in many field theoretical models~\\cite{Vilenkin:1983xq,Linde:1986fd,Linde:1986fc}. Further, the current picture of the string landscape~\\cite{Douglas:2006es} suggests that the observed part of the Universe was created as a result of tunneling from some higher-scale eternally-inflating  vacuum \\cite{Bousso:2000xa}. All these arguments make the study of eternal inflation very important.  Moreover eternal inflation provides a natural framework to implement Weinberg's solution of the cosmological constant problem~\\cite{Weinberg:1987dv}, which is the most plausible so far in spite of the many efforts made to find an alternative. According to this solution the choice of the vacuum is made, at least partially, by anthropic reasons such as the requirement that structures were able to form in our Universe. These arguments raise the notoriously difficult and puzzling question of making predictions in an eternally inflating Universe.  On a purely practical side we possess a well developed machinery of quantum field theory in curved space-time that proved to be very successful in calculating properties of the primordial density perturbations with many fine details in the case of non-eternal inflation (see e.g. \\cite{Maldacena:2002vr}). The applications of these techniques to the case  of eternal inflation is instead much more challenging, as in this latter regime the size of the quantum fluctuations is large, a non-perturbative treatment is required and calculations become in general much more difficult (see e.g. \\cite{Creminelli:2008es} for a recent discussion). On the other hand, without a clear understanding of the eternally inflating geometry, it might be hopeless to solve the issues raised by eternal inflation such as the measure problem in the landscape; this is why we find it very important to make explicit and precise calculations in this regime.  On a more theoretical side, dS space appears to share many properties with the black hole geometry---most importantly, in both space-times the most natural sets of observers (the asymptotic observers in the black hole case and the comoving ones in the accelerating Universe) see a gravitational horizon with the associated thermodynamic properties, such as Hawking temperature and Bekenstein entropy, the latter, in the case of de~Sitter, being equal to \\cite{Gibbons:1977mu} \\[ S_{dS}=\\pi{M_{\\rm Pl}^2\\over H^2}\\;, \\] where $M_{\\rm Pl}^2\\equiv 1/G_N$ with $G_N$ the Newton constant. A finite entropy suggests that the system is described by a finite number of degrees of freedom, or more formally, the Hilbert space describing the system at the quantum level has a finite dimensionality equal to $e^{S_{dS}}$. It is widely believed that both black holes and de~Sitter space always arise as a subsector of a larger theory with an infinite dimensional Hilbert space. This is obvious for a black hole in asymptotically Minkowski or adS space-times, but less clear for dS space. Indeed, in principle one could imagine a quantum gravity theory with, for example, a single positive energy vacuum. However, this is not the case in the known string theory landscape~\\cite{Kachru:2003aw} and there are general arguments strongly suggesting that dS vacua are always metastable with respect to decay to either the Minkowski or the adS minima of the potential~\\cite{Goheer:2002vf}. This will be the point of view adopted in the current paper and by dimensionality of the Hilbert space describing black hole or dS space we always understand the dimensionality of the corresponding subsector in a larger, most likely infinite-dimensional, Hilbert space (although, it is worth noting that an alternative line of thought is also being pursued \\cite{Banks:2005bm}).  Taking seriously the similarity between the causal structures of de~Sitter and Schwarzschild causes a serious doubt on the validity of the global semiclassical picture of the eternally inflating Universe\\footnote{This idea is being pursued by a number of authors, see {\\it e.g.} \\cite{Bousso:2006ge}.  Our discussion  mainly follows that of  \\cite{ArkaniHamed:2007ky}.}. Indeed, the remarkable fact about black holes is that the global effective field theory description of the space-time claiming to describe both the interior and the exterior of the horizon eventually breaks down~(see e.g. \\cite{Lowe:1995ac}).  More concretely, if one considers the set of space-like slices covering both the exterior and the interior of a black hole, as shown in fig.~\\ref{fig:penrose}a, and insists that the field theory description is valid on this set of slices, one comes to the conclusion that the information is lost in the course of the black hole evaporation~\\cite{Hawking:1976ra}. This conclusion was proven to be wrong \\cite{Maldacena:2001kr,Hawking:2005kf} by the adS/CFT correspondence  that provides a description of the system involving black holes (gravity in the bulk adS space) in terms of an unitary boundary CFT. \\begin{figure}[t] \\begin{center} \\includegraphics[height=4cm]{bh-pen.eps} \\hspace{2cm} \\includegraphics[height=4cm]{dS-pen.eps} \\caption{\\label{fig:penrose} \\small \\it Set of space-like slices covering both the exterior and the interior of: a) a black hole (left) and b) an eternal inflating universe (right).} \\end{center} \\end{figure} Consequently, one is forced to conclude that the global effective field theory description breaks down at the time scale of order the evaporation time, $t\\sim R_s^3/G_N$, where $R_s$ is the Schwarzschild radius and $G_N$ the Newton constant (for a recent review see for example \\cite{ArkaniHamed:2007ky}). After this time the information about the inside observer gets reprocessed in the evaporating Hawking quanta, and by insisting on a simultaneous local description of the exterior and the interior on longer time-scales one would run into a contradiction with the ``no quantum xerox\" principle (or equivalently with the linearity of quantum mechanics) \\cite{'tHooft:1990fr,Susskind:1993if}. This conclusion is really surprising given that one can always choose a set of slices that avoid the region close to the singularity, so that naively one would expect the effective field theory to hold.   Given the similarity between the causal structures of de~Sitter space and that of Schwarzschild geometry, one may suggest that also in dS space the global description in terms of a similar set of slices (shown in fig.~\\ref{fig:penrose}b) eventually breaks down.  Note that these slices are just the conventional FRW slices that are commonly used to describe the inflationary Universe. Going on with the analogy with the black holes one expects this breakdown to happen at a time-scale of order $t\\sim H^{-3}/G_N$ or, equivalently, after a period of order $S_{dS}$ $e$-foldings.  One may expect that space-time events outside the region containing $e^{S_{dS}}$ Hubble patches get encoded in de~Sitter fluctuations,  similarly to how the information inside the black hole gets released in the Hawking quanta after the evaporation time.  Note that, unlike in the black hole case, we are not running into any paradox if this does not happen, therefore it may well be that pushing the analogy this far is too naive. However, if one takes it seriously there is an immediate test to pass. Indeed, in the black hole case, due to the presence of a curvature singularity, it is impossible to read the information in the Hawking quanta and then jump into the black hole and read the same information again. Similarly,  it should not be possible to get the same information twice in the de~Sitter case as well.  However, if inflation could last for an arbitrarily long time without becoming eternal, an observer would be able to first read the information in the dS fluctuations and later, as more and more space-time becomes visible after inflation, he would be able to access it directly. Consequently, for the above ideas about de~Sitter complementarity to be consistent, there should be a limit on how long inflation can last without becoming eternal \\cite{ArkaniHamed:2007ky}, \\be \\label{vague_limit} 3N\\leq cS_{dS}\\;, \\ee where $N$ is the number of $e$-foldings (defined as one third of the logarithm of the total volume after inflation) and $c$ is a coefficient of order one. On the other hand, no limit on the number of $e$-foldings is required on those realization that are eternal, because in those cases the observer is not able to access all the volume after reheating.  It was proven in \\cite{ArkaniHamed:2007ky} that the bound (\\ref{vague_limit})  indeed holds for the classical inflaton trajectory in any theory of inflation that does not allow violations of the null energy condition, \\[ \\rho+p\\geq 0\\;, \\] where $\\rho$ and $p$ are the energy density and the pressure. On the other hand, violation of the bound (\\ref{vague_limit}) is possible in theories able to violate the null-energy condition, such as ghost inflation \\cite{ArkaniHamed:2003uz}. This is actually encouraging and supports  arguments establishing the link between horizon complementarity and duration of inflation, as also the conventional black hole thermodynamics  breaks down in theories where the null energy condition can be violated \\cite{Dubovsky:2006vk,Eling:2007qd}.  There are two reasons leading to an uncertainty in the numerical value of the coefficient $c$ in the bound (\\ref{vague_limit}) as proven in \\cite{ArkaniHamed:2007ky}. First, at the time when the bound was proposed the exact condition for inflation to become eternal was unknown, and it was not even clear whether there is a sharp distinction between eternal and non-eternal regimes. This issue was addressed in \\cite{Creminelli:2008es} and the conclusion is that there is a sharp transition between these two regimes, with the condition not to have eternal inflation being \\be \\label{no_eternal} \\Omega\\equiv {2\\pi^2\\over 3}{\\dot{\\phi}^2\\over H^4}\\geq 1\\;, \\ee where $\\dot\\phi$ is the classical velocity of the inflaton field.  Now it is straightforward to find a bound on the number of $e$-foldings for the classical inflaton trajectory in single-field slow-roll inflation  in the non-eternal regime. Namely, one writes \\be \\label{precise_classical} {dS_{dS}\\over dN_c}\\equiv {M_{\\rm Pl}^2dH^{-2}\\over Hdt}=-{2M_{\\rm Pl}^2\\dot{H}\\over H^4}=12\\Omega\\;, \\ee where at the last step we made use of the second Friedmann equation. By integrating (\\ref{precise_classical}) and using the condition (\\ref{no_eternal}) for the absence of eternal inflation we obtain (\\ref{vague_limit}) with the value $c=1/4$, \\be \\label{vague_limit1/4} 3N_c\\leq {S_{dS}\\over 4}\\;, \\ee where $N_c$ is the number of $e$-foldings on the classical inflaton trajectory.   However, this does not establish the sharp version of the bound (\\ref{vague_limit}) yet. There is  another reason for the uncertainty in the bound (\\ref{vague_limit}). Namely,  the above analysis  (and that of  \\cite{ArkaniHamed:2007ky}) is restricted to the classical inflaton trajectory. This approximation clearly breaks down when $\\Omega$ is of order (but still larger than) one, so that inflation is close to be eternal. In this case, even though inflation is not eternal, the typical inflaton trajectory is very different from the classical one and can be much longer. More generally, at the quantum level for any value of $\\Omega$ there is always a non-vanishing probability for the actual  inflaton trajectory to be long enough to violate the bound (\\ref{vague_limit}) for any value of $c$.  In this situation it is natural to study what is the probability distribution for inflaton trajectories of different lengths---in other words, what is the probability distribution for the volume of the Universe $\\rho(V)$ after inflation. It is not clear a priori what the natural generalization of the bound (\\ref{vague_limit}) should look like at the quantum level. One might expect that there exists  a value of $c$ such that the probability to violate (\\ref{vague_limit}) is suppressed, for example as non-perturbative quantum gravity effects $e^{-S_{dS}}$ (which would correspond to $\\rho(V)\\sim 1/V^\\alpha$) or even more, for example exponentially with the volume $\\rho(V)\\sim e^{-V}$ (which would correspond to order $e^{-e^{S_{dS}}}$ effects). What we find from our analysis is that such value of $c$ exists (it is $c=1/2$) and that the probability associated to the violation of the bound is actually super-exponentially small, \\emph{i.e.} $\\sim e^{-e^{ S_{dS}}}$.  To achieve this goal we obtained another result of independent interest. Namely, we calculated in an explicit form the probability distribution for the volume of the Universe after slow-roll inflation $\\rho(V)$ both in eternal and non-eternal regimes. This offers further insight in the actual geometry of the eternally inflating spacetime. While, unlike the density perturbation spectrum, this quantity is not of much interest for current observations, it still appears to be one of the natural ``theoretical\" observables to look at in the study of eternally inflating Universes. In particular, according to \\cite{Creminelli:2008es}, the order parameter for the transition to eternal inflation is the normalization of $\\rho(V)$, \\be P_{\\rm ext}=\\int_0^\\infty dV\\rho(V)\\;. \\ee At $\\Omega>1$ this quantity is equal to 1, in agreement with the naive expectation. However, at $\\Omega<1$ the normalization $P_{\\rm ext}$ becomes smaller then 1, indicating that there is a non-zero probability $(1-P_{\\rm ext})$ for the reheating volume to be infinite, {\\it i.e.} for inflation to last forever. In this paper we will rederive this result in yet another, somewhat more explicit, way. It is also worth noting that recently the far exponential tail of the probability $\\rho(V)$ was calculated in \\cite{Winitzki:2008ph} in the eternal regime. This result was used there to define a ``reheating-volume\" measure for observables after eternal inflation. It appears that the explicit expression for $\\rho(V)$ has good chances to be useful in further theoretical studies of de~Sitter space and eternal inflation.  The rest of the paper is organized as follows. We start section~\\ref{sec:bact2infl} with a review of the discrete stochastic branching process introduced in  \\cite{Creminelli:2008es} to describe inflation. Then we use this model to derive a differential equation (\\ref{eq:phi}) for the Laplace transform of the probability distribution $\\rho(V)$. Similar discrete models were used in \\cite{Winitzki:2005ya} and they are essentially equivalent to the stochastic description of inflation by Starobinsky \\cite{starobinsky}. Within the stochastic  approach equation (\\ref{eq:phi}) is known as the ``non-linear Fokker--Planck equation\" \\cite{Winitzki:2001np}.  In section~\\ref{sec:solutions} we provide  approximate solutions for the equation (\\ref{eq:phi}) and calculate the probability distribution in different regimes. We start with a discussion of the general properties of the solutions of (\\ref{eq:phi}), and rederive in a new and very explicit way that $\\Omega=1$ is the transition point to the eternally inflating regime, {\\it i.e} that $P_{\\rm ext}=1$ for $\\Omega>1$ and $P_{\\rm ext}<1$ for $\\Omega<1$. Then in section~\\ref{sec:moments}  we study the moments of the volume distribution. We show that there is a simple way of calculating them without actually solving  equation (\\ref{eq:phi}) and performing the Laplace transform. In agreement with the results of \\cite{Creminelli:2008es} we prove that at $\\Omega>1$  sufficiently high moments diverge for any value of $\\Omega$ if the inflaton field is allowed to take arbitrarily high values. We find the values  $\\Omega_n$ such that  the $n$-th moment diverges at $\\Omega<\\Omega_n$, and illustrate our method by explicitly calculating the average and the variance.  In section~\\ref{sec:classical} we start analyzing the properties of the probability distribution by performing the Laplace transform. First, we consider the semiclassical limit $\\Omega\\gg1$, and find an approximate solution for eq.~(\\ref{eq:phi}) in this limit. It turns out that because of the non-commutativity of the large-volume and the large-$\\Omega$ limits this solution does not capture correctly the large volume tail of the probability distribution where the probability becomes smaller than $\\sim e^{-\\Omega}$. The study of this solution is also instructive for developing an intuition on how to perform the Laplace transform of the solutions of eq.~(\\ref{eq:phi}). In section~\\ref{sec:Wgen} we apply this intuition for general $\\Omega>1$, when  it is not possible to find an approximate solution to (\\ref{eq:phi}) in closed form. By solving this equation in different regimes one can obtain enough information to reconstruct the probability distribution in the physically relevant case $N\\gg1$. $\\rho(V)$ turns out to be peaked around the average value \\begin{eqnarray} \\label{Nav} {\\overline N}={2N_c\\over 1+\\sqrt{1-\\Omega^{-1}}}\\;. \\end{eqnarray} For $N<{\\overline N}$ it takes the following Gaussian form \\be \\label{gaussian} \\rho(N)\\simeq {\\cal N}e^{-{(3N-3{\\overline N})^2\\over 2\\sigma^2}}\\;,\\qquad \\Omega > 1 \\ee where  the width $\\sigma$ is equal to \\begin{eqnarray} \\sigma^2={2\\over\\Omega(1+\\sqrt{1-\\Omega^{-1}})^2}\\;. \\end{eqnarray} As $\\Omega$ approaches the transition point, $\\Omega=1$, the width $\\sigma$ becomes of order one, which is still narrow in the regime of large number of $e$-foldings, $N\\gg 1$, the one we are interested in. So the most important consequence of the change of $\\Omega$ is that the average number of $e$-foldings ${\\overline N}$ changes. In agreement with the naive expectation it increases as $\\Omega$ approaches the transition point, but it does not grow a lot: at $\\Omega=1$  the average number of $e$-foldings is twice as large as the classical one.  At large volumes, $N\\gtrsim {\\overline N}$, the probability distribution becomes exponential in $N$ (or, equivalently, power-law in the volume $V$), \\be \\label{exp} \\rho(N)\\propto e^{-6\\Omega N\\l(1+\\sqrt{1-{1\\over \\Omega}}\\r)}=V^{-2\\Omega \\l(1+\\sqrt{1-{1\\over \\Omega}}\\r)}\\;. \\ee  \\begin{figure}[t] \\begin{center} \\includegraphics[width=14cm]{rhoV-proto.eps} \\caption{\\label{fig:rhoV-proto} \\small \\it Typical shape for the probability distribution of the volume $\\rho(V)$. For small volumes the behavior is gaussian with the number of $e$-foldings ($\\rho\\sim e^{-c(N-\\overline N)^2}$); for volumes larger than the average value $\\overline V$, $\\rho(V)$ follows a power law in the volume ($\\rho\\sim 1/V^\\alpha$) that eventually turns into an exponential law ($\\rho \\sim e^{-{\\rm const}\\cdot V}$) at large enough volumes ($V\\gtrsim V_b$). When $\\Omega<1$ the exponential tail starts earlier at $V\\simeq V_\\epsilon=e^{\\pi/(2\\sqrt{1-\\Omega})}$.} \\end{center} \\end{figure}   Then we proceed with the eternal inflation regime. First, we study what happens in the vicinity of the transition point, {\\it i.e.} when $\\Omega=1-\\epsilon$ with $0<\\epsilon\\ll 1$. We find that the probability distribution $\\rho(N)$ is not changed until $N\\sim\\pi/(6\\sqrt{\\epsilon})$. At large volumes, $N\\gtrsim\\pi/(6\\sqrt{\\epsilon})$, it becomes much more strongly suppressed, $\\rho(N)\\propto e^{-{\\rm const} \\cdot e^{3N}}$. This behavior is easier to interprete  in terms of the volume distribution $\\rho(V)$ (rather than the $e$-folding distribution $\\rho(N)$). It  indicates that if the volume gets large enough, $V\\gtrsim e^{\\pi/(2\\sqrt{\\epsilon})}$, the probability for inflation to terminate is exponentially small $P_{\\rm ext}\\propto e^{-{\\rm const}\\cdot V}$. Related to that, we find that the total probability for inflation to terminate is of order $e^{-\\Omega(3N_c)^2}$ (this also applies for $\\epsilon\\sim{\\cal O}(1)$), indicating that it is saturated by the small-$N$ tail of the Gaussian distribution (\\ref{gaussian}). This behavior smoothly matches with yet another regime where one can find the probability distribution explicitly---$\\Omega\\simeq 0$. Here the probability distribution $\\rho(V)$ is exponentially small $\\rho(V)\\propto e^{-V/2}$ for all volumes of interest, $V\\gg 1$.  All these considerations were made in the approximation where the inflaton potential goes up to arbitrary high values of the inflaton field. This is clearly unrealistic and we conclude in section~\\ref{sec:solutions} by discussing what happens in the presence of a ``barrier\" at large values of the inflaton field. As expected, the presence of a barrier affects only the far tail of the volume distribution by making it exponentially suppressed at large volumes, $\\rho(V) \\propto e^{-{\\rm const}\\cdot V}$, for any value of $\\Omega$. Note that if the initial value of the inflaton field is not too close to the barrier, this effect is relevant only for inflaton trajectories whose probability is smaller than the uncertainty coming from non-perturbative quantum gravity effects ($\\sim e^{-S_{\\rm dS}}$). The various behaviors of the probability distribution are summarized in fig.~\\ref{fig:rhoV-proto} that shows the shape of $\\rho(V)$ in the different regimes of the volume.  These results imply that the quantum version of the bound (\\ref{vague_limit}) does hold with $c=1/2$. In the concluding section~\\ref{conclusions} we give the physical explanation of the behavior of the probability distribution and show that an inflaton trajectory  with more than $S_{dS}/6$ $e$-foldings and such that inflation terminates globally in the entire space is super-exponentially improbable. We also speculate on the possibility that the value $c=1/2$ for the coefficient in (\\ref{vague_limit}) that we obtained in our analysis might have a natural physical interpretation. In the appendix we cross check our results by calculating the average volume directly from the inflaton stochastic equations. ",
        "conclusions": "\\label{conclusions}   To summarize, in this paper we calculated explicitly the probability distribution $\\rho(N)$ for the volume of the Universe after a period of slow-roll inflation (as before, by the number of $e$-foldings $N$ we understand  one third of the logarithm of the total volume produced during inflation, $N={1\\over 3}\\log V$). Our results cover both the eternal and the non-eternal regime. Let us start this concluding section by summarizing these results and then by explaining how all different kinds of behavior we found for $\\rho(N)$ have an intuitive physical explanation.  In general, the function $\\rho(N)$ exhibits three qualitatively different regions. Namely, in the non-eternal regime $\\Omega>1$ it is peaked at $N\\sim {\\overline N}$, where the average number of $e$-foldings ${\\overline N}$ is given by \\be \\label{sumNav} {\\overline N}={2N_c\\over 1+\\sqrt{1-\\Omega^{-1}}}\\;. \\ee For $N\\lesssim {\\overline N}$ it has a gaussian form, \\be \\label{sumgaussian} \\rho(N)\\propto e^{-{(3N-3{\\overline N})^2\\over 2\\sigma^2}}\\;, \\ee with a width $\\sigma$ given by \\begin{eqnarray} \\label{sumwidthNav} \\sigma^2={2\\over\\Omega(1+\\sqrt{1-\\Omega^{-1}})^2}\\;. \\end{eqnarray} The behavior changes at $N \\gtrsim {\\overline N}$ where the probability distribution becomes exponential in $N$ (or, equivalently, power-law in the volume $V$), \\be \\label{sumexp} \\rho(N)\\propto e^{-6\\Omega N\\l(1+\\sqrt{1-{1\\over \\Omega}}\\r)}=V^{-2\\Omega \\l(1+\\sqrt{1-{1\\over \\Omega}}\\r)}\\;. \\ee Finally, if a barrier preventing the inflaton to take arbitrary large values is present, this power-law tail becomes further suppressed at larger volumes and turns into an exponential in the  {\\it volume} \\be \\label{sumexpexp} \\rho(N)\\propto e^{-c e^{3N}}=e^{-c V}\\;. \\ee This behavior sets in at $N\\sim N_b$, where $N_b$ is the number of $e$-foldings on the classical inflaton trajectory from the barrier  till reheating.  What changes in the eternal regime, $\\Omega<1$, is that the exponential behavior (\\ref{sumexpexp}) sets in at $N\\simeq{\\pi/ 6 \\sqrt{1-\\Omega}}$ even if the barrier is absent. If  $(1-\\Omega)$ is not too small this happens at $N<{\\overline N}$ so that the gaussian regime (\\ref{sumgaussian}) interpolates directly to the superexponential one (\\ref{sumexpexp}) without intermediate exponential behavior (\\ref{sumexp}).  It is rather straightforward to understand the origin of the above three different types of behavior for $\\rho(N)$ at the intuitive level. First, to produce a small number of $e$-foldings $N\\ll N_{c}$, the inflaton field during the first few $e$-foldings needs to perform a jump in the whole volume to a value of $\\phi$ corresponding to an average number of $e$-folding equal to $N$. The  inflaton fluctuations away from the classical trajectory follow, at early times, a gaussian distribution (because of the random walk) and the size of the jump in field space is directly proportional to $(N-{\\overline N})$, so the probability of such a jump is suppressed by $e^{-c (N-{\\overline N})^2}$ ($c$ is a constant) in agreement with our result (\\ref{sumgaussian}) at small volumes.  On the other hand, the least expensive way to produce a large number of $e$-foldings $N\\gg N_{c}$ is for one Hubble patch to go high up the potential, till values of the field corresponding to have a classical trajectory with $N$ $e$-foldings. If such a fluctuation happened, the probability to produce at least $N$ $e$-foldings becomes of order one. To estimate the probability of such a fluctuation, note again that the probability distribution for the inflaton fluctuations $\\Delta\\phi$ around the classical trajectory is gaussian. The crucial difference with the small volume case is that now  it is not necessary for the fluctuation to happen in a short period of time, and the variance of the inflaton distribution grows linearly as a function of time $t$. As a result, the probability $p$ of the fluctuation is maximized for times $t$ corresponding to order $N$ $e$-foldings, and depends exponentially on $N$, $p\\propto e^{-c_1 \\Delta\\phi^2/t}\\propto e^{-c_2 N}$ in agreement  with (\\ref{sumexp}). This argument is essentially identical to the one provided in \\cite{Creminelli:2008es} to explain why high enough moments of the volume distribution diverge in the non-eternal regime if no barrier is introduced.  The last argument fails at sufficiently large volumes both in the presence of a barrier at high values of the inflaton field and in the eternal regime. In the first case it fails because there is a limit on the maximum length of the classical trajectory, while in the second because even if a fluctuation to high values of the inflaton field happened, one is not guaranteed to end up with a {\\it finite} number of $e$-foldings (in fact, the higher the fluctuation is the smaller is the probability for inflation to terminate). In both cases one ends up with  getting a superexponential suppression of the probability distribution at large volumes, (\\ref{sumexpexp}). This can roughly be explained by the necessity for an exponential number of independent and rather improbable events (corresponding to the individual Hubble volumes produced during inflation) to happen. Namely, in the eternal regime, inflation should terminate for $\\sim e^{3N }$ Hubble patches and all their children (while in the non-eternal regime the probability for this  to happen is equal to one). In the non-eternal case with a barrier, for $N\\gg N_b$, an order $\\sim e^{3N }$  Hubble patches have to live an anomalously long time, corresponding to a number of $e$-foldings much larger than $N_b$.  To summarize, we see that our results on the shape of the probability distribution of the reheating volume not only pass a number of consistency checks with results obtained by different methods, but also can be understood in a rather intuitive level. We believe making the explicit form of the volume probability available  may help to understand the geometry of eternal inflation and offers also a natural ``theoretical\" observable that can be useful in further studies of eternal inflation.  A further motivation for doing our calculation was to establish whether a quantum version of the bound \\be \\label{vague_limit_2} 3N \\leq cS_{dS}\\;, \\ee exists and, if so, whether this bound can be made sharp, {\\it i.e.} whether there is a concrete value for the coefficient $c$ in (\\ref{vague_limit_2}). Our results provide a positive answer to both questions. Namely, we find that the quantum version of the bound (\\ref{vague_limit_2})  can be formulated as follows:  \\emph{The probability for slow-roll inflation to produce a finite volume larger than $e^{S_{dS}/2}$, where $S_{dS}$ is de~Sitter entropy at the end of the inflationary stage, is suppressed below the uncertainty due to non-perturbative quantum gravity effects.}  Indeed, our results imply that an inflaton trajectory  with more than $2N_c$ $e$-foldings and such that inflation terminates globally in the entire space is exponentially unprobable, $\\sim e^{-N}$. Given that in the non-eternal regime we have the bound (\\ref{vague_limit1/4}), we conclude that the probability for slow-roll inflation to last more than $S_{dS}/6$ $e$-foldings and to terminate globally in the entire space is smaller than the uncertainty coming from non-perturbative quantum gravitational effects ($\\sim e^{-S_{\\rm dS}}$). In fact, there is an even stronger statement that follows from the observation of section~\\ref{sec:barrier} about the bound on the barrier position $N_b$. Indeed, as long as $\\Omega\\gtrsim 1$, $N_b$ is smaller  than $S_{dS}/12$, which implies that the transition to the regime where $\\rho(V)$  drops exponentially with the volume ($e^{-{\\rm const}\\cdot V}$) starts exactly when the produced volume  starts violating the bound, $V>e^{S_{dS}/2}>e^{6N_b}$. Consequently, the probability to produce much larger volumes $V\\gg e^{S_{dS}/2}$ is super-exponentially suppressed, $\\sim e^{-e^{S_{dS}}}$. In turn this produces a super-exponentially small probability to produce volumes $V\\gg e^{S_{dS}/2}$.  We believe this provides a  non-trivial test of the de~Sitter complementarity idea. However, one may wish to go further and ask whether the particular value $c=1/2$ that we found provides any insights into how the de~Sitter complementarity works. Of course, all our calculations were done in the limit where gravity is non-dynamical, therefore it is not clear whether the value $c=1/2$ really tells us something fundamental about the properties of de~Sitter space. Let us however assume that it does and speculate on what would be the interpretation of this particular value. Recall that the original motivation for the bound (\\ref{vague_limit}) is coming from the idea that the black hole complementarity applies in the de~Sitter case as well, so that the global effective field theory description of the FRW slices breaks down and information about the outside observers gets released in de~Sitter fluctuations. Then, if inflation ends, there is the danger of violating the linearity of quantum mechanics by creating a ``quantum xerox\"---one can see the information twice, first holographically in the de~Sitter fluctuations and later in a direct way when the corresponding mode comes in.   Recall that the analogous paradox could potentially arise in the black hole case after one  measures more than $S_{bh}/2$ Hawking quanta ($S_{bh}$ is the black hole entropy). In this case, the factor $1/2$ arises because, as pointed out by Page \\cite{Page:1993wv}, if one measures $k$ degrees of freedom of a larger system described by $n$ degrees of freedom, generically, the resulting density matrix looks thermal and carries less than a single bit of information as long as $k<n/2$.  This does not give rise however to any paradox: indeed, if  one waits a long enough time outside the black hole, so that $\\sim S_{bh}/2$ Hawking quanta are emitted, it is impossible to observe the same information a second time inside the black hole, since it gets destroyed by the curvature singularity.   It is tempting to interprete our result by saying that the de~Sitter analogue of one Hawking quanta is produced every time a new Hubble patch is created from the volume where the observer lives ({\\it{i.e}} every $1/3$ of $e$-folding). Then the quantum-xerox problem would arise after $S_{dS}/6$ $e$-foldings but it does not because of the bound.  Although  it may sound quite natural, this interpretation raises a number of issues. First, as argued in \\cite{Parikh:2008iu}, a single observer in {\\it eternal} de~Sitter can never run into the xerox paradox, as the largest amount of entropy that can be stored within a single causal patch is bounded by the Bekenstein area of the largest black hole that fits in a single patch. The latter is equal to $S_{dS}/3$ and smaller than the minimal amount ($S_{dS}/2$) required to extract the first bit of information.  However we are in a rather different situation here, as inflation eventually terminates. Therefore should the bound be not there, one would have the possibility to measure $S_{dS}/4$ quanta, give them to a friend who leaves for a different Hubble patch, then measure other $S_{dS}/4$ quanta and, after  inflation ends, the two could  meet  and in this way collect $S_{dS}/2$ quanta in total. Then the bound suggests that for the purpose of extracting information it is legitimate to collect quanta in one Hubble patch and keep them until the end of inflation in another. However, this reasoning also implies that it should not be possible to use quanta collected in {\\it different} Hubble patches for extracting information---otherwise one would be forced to conclude that the actual bound on the number of $e$-folding is much stronger, $N\\lesssim \\log{S}$.  Notice also that if inflation does not terminate everywhere, then there is no problem of duplication of information if one observer lasts in the inflationary phase for longer than $S_{dS}/6$ $e$-foldings: when the observer undergoes reheating and enters the Minkowski phase she is never able to see all the rest of the de~Sitter space.  Another puzzle with drawing the parallel between the factor 1/2 in our bound and the one in the Page argument is that, if there are $n$ light species around then, likely, the rate of how fast the information gets released is proportional to the number of species (at least this is so in the black hole case). On the other hand our bound does not depend on $n$ (it could if the light species are scalars but does not for other fields).  Finally, it is worth mentioning yet another reason why one should be cautious in taking the value $c=1/2$ too seriously. Namely, here we focused on slow-roll models of inflation. In the effective field theory language of \\cite{Cheung:2007st} these correspond to the models where the inflaton perturbations have sound velocity $c_s$ equal to one. The classical version of the bound (\\ref{vague_limit}) was proven in \\cite{ArkaniHamed:2007ky} for  general inflationary models, including those with small $c_s$. In fact at small $c_s$ the classical bound becomes even stronger, $N_c\\lesssim c_s^5 S_{dS}$. However the reason it arises is somewhat different. At small sound velocities the strong coupling regime, where the effective field theory breaks down, sets in always before the eternal inflation regime. So the bound (\\ref{vague_limit}) in this case follows from the requirement that the system is weakly coupled and the null energy condition is not violated. Consequently, at the present stage we cannot exclude the possibility of violating the bound in the strongly coupled regime, although this possibility appears highly unlikely.   To summarize, we proved the quantum version of the bound on how long slow-roll inflation can last without becoming eternal. The existence of such a bound (eq.~(\\ref{vague_limit_2})) provides non-trivial support to de~Sitter complementarity ideas, while it is still unclear whether the value of the coefficient $c=1/2$ appearing in the bound that we found has a definite physical meaning. Before finishing, it is worth stressing that we believe that, independently of the answer to the last question, explicit calculations providing a detailed quantitative understanding of the transition to the eternal inflation regime, like those we preformed in the current paper, have an independent value. Indeed,  it appears to us that a detailed understanding of the geometry and the dynamics of the eternally inflating Universe might be an important step to reach a final verdict on the puzzling issues raised by the observation of the cosmic acceleration and the string landscape."
    },
    "0812/0812.0844_arXiv.txt": {
        "abstract": "By performing local three-dimensional MHD % simulations of stratified accretion disks, we investigate disk winds driven by MHD turbulence. Initially weak vertical magnetic fields are effectively amplified by magnetorotational instability and winding due to differential rotation. % Large-scale channel flows develop most effectively at 1.5 - 2 times the scale heights where the magnetic pressure is comparable to but slightly smaller than the gas pressure. The breakup of these channel flows drives structured disk winds by transporting the Poynting flux to the gas. These features are universally observed in the simulations of various initial fields. This disk wind process should play an essential role in the dynamical evaporation of protoplanetary disks. The breakup of channel flows also excites the momentum fluxes associated with Alfv\\'{e}nic and (magneto-)sonic waves toward the midplane, which possibly contribute to the sedimentation of small dust grains in protoplanetary disks. ",
        "introduction": "Magnetorotational instability \\citep[MRI;][]{bh91} is regarded as a robust mechanism to provide turbulence for an efficient outward transport of angular momentum in accretion disks. MHD simulations in a local shearing box have been carried out \\citep[e.g.,][]{hgb95,bra95,san04} to study the properties of MRI-driven turbulence. \\citet{ms00} studied vertically stratified local disks with free boundaries to allow leaks of mass and magnetic field. While their main purpose is to study general properties of stratified disks such as disk coronae rather than disk winds, they concluded that the mass flux of the outflows is small in the cases of initially toroidal and zero-net vertical flux magnetic fields.  On the other hand, protoplanetary disks around young stars should have net vertical magnetic fields that are connected to their parental molecular clouds. In this case, the physical conditions of the surface of the disk are analogous to the open coronal holes of the Sun where the solar wind is driven by turbulent footpoint motions of the magnetic field lines \\citep{sak07,tsu08}. Obviously, MHD turbulence excited by MRI in the disk is also expected to drive winds from the surfaces of the accretion disk. Although such a disk wind mechanism may play a significant role in the evolution of accretion disks \\citep{fdc06}, quantitative studies have not been carried out so far because of difficulties in the numerical treatment: a long-term calculation of the wind process requires accurate description of outgoing boundary conditions for various types of waves, rather than simple free boundaries \\footnote{Simple free boundaries means setting the derivatives of variables to be zero. In this case, however, unphysical reflections of waves usually occur. For the real outgoing boundary, only the derivatives of incoming characteristics should vanish \\citep{tho87}.}.   In this Letter, we investigate disk winds driven by MRI with initially vertical magnetic fields utilizing the rigorous outgoing boundary condition that was originally developed for the simulations of solar \\citep{si05,si06} and stellar \\citep{suz07} winds. ",
        "conclusions": "We have shown that the gas is lifted up from the injection regions at $|z|/H_0\\approx 1.5$-$2$ to the surfaces by the Poynting flux and streams out from the upper and lower boundaries of the simulation box. In effect, our results determine the condition of ``mass loading'' in various models of global disk wind \\citep[e.g.,][Ferreira et al 2006]{bp82,ks97}, in which one usually fixes the mass flux in advance by setting the densities at the `bases' of winds. With our outgoing boundary condition, we implicitly assume that once the gas goes out of the $z$-boundaries of the simulation box it does not return. % The validity of this treatment needs to be examined by the global modeling of accretion disks.  Although in the shearing box treatment the time-averaged net mass flow in the $x$-direction is zero, we can estimate the accretion velocity, $-v_r\\approx\\alpha c_s^2/r\\Omega_0$, from the angular momentum balance under steady states, where $r$ is a cylindrical distance from a central star \\citep{ss73}. The ratio of the mass-loss rate, $\\dot{M}_z$, from the simulation box by the disk winds to the mass accretion rate, $\\dot{M}_r$, passing through the $y$-$z$ plane becomes \\begin{eqnarray} \\frac{\\langle\\dot{M}_z\\rangle}{\\langle\\dot{M}_r\\rangle}&\\approx& \\frac{\\int\\int dx dy \\langle\\rho v_z\\rangle}{\\int dy\\langle\\Sigma\\bar{\\alpha}\\rangle c_s^2/(r\\Omega_0)}=\\frac{\\langle\\langle\\rho v_z\\rangle\\rangle{L_x}r\\Omega_0}{\\langle\\langle\\Sigma\\bar{\\alpha}\\rangle\\rangle c_s^2}\\nonumber\\\\ &\\approx&0.05\\left(\\frac{r}{10H_0}\\right)\\left(\\frac{L_x}{H_0}\\right), \\label{eq:msfx} \\end{eqnarray} where $\\Sigma=\\int dz\\rho$ and % $\\bar{\\alpha}=\\int dz\\rho\\alpha/\\Sigma$; $\\langle\\rangle$ denotes the time-average and $\\langle\\langle\\rangle\\rangle$ is the average of time and $x$-$y$ planes. The final value is estimated from $\\langle\\langle\\rho v_z\\rangle\\rangle\\approx 8\\times 10^{-5}\\rho_{0}c_s$ and $\\bar{\\alpha}\\approx 0.012$ in our fiducial run for a moderately thin disk with $H_0/r=0.1$. We should take {\\it the above estimate as an upper limit} because the calculation with a larger vertical box size may give a smaller mass flux at the top and bottom boundaries. The angular momentum loss rates in vertical and radial directions give the same scaling, \\begin{equation} \\frac{\\langle\\dot{\\cal{L}}_z\\rangle}{\\langle\\dot{\\cal{L}}_r\\rangle}=\\frac{\\int\\int dx dy\\langle\\rho v_z\\rangle r^2\\Omega_0} {\\int dy r\\langle\\Sigma\\bar{\\alpha}\\rangle c_s^2}=\\frac{\\langle\\langle\\rho v_z\\rangle\\rangle L_x r\\Omega_0}{\\langle\\langle\\Sigma\\bar{\\alpha}\\rangle\\rangle c_s^2}\\approx\\frac{\\langle\\dot{M}_z\\rangle}{\\langle\\dot{M}_r\\rangle}, \\end{equation} because the time-averaged specific angular momentum carried by the winds is approximately the same as that in the disk material at the same radial position in the shearing box treatment. In a realistic situation, however, winds possibly carry a larger specific angular momentum than the disk material. For such studies, we need to model global accretion disks with disk winds, which is also important from the viewpoint of angular momentum evolution of the star-disk system \\citep[e.g.][]{mp05}.  Hereafter we discuss the evolution of protoplanetary disks, as an application of our results. As a reference model, we use the minimum-mass solar nebula (MMSN) of Hayashi (1981), % which gives the midplane density, $\\rho_{0}=1.4\\times 10^{-9}\\left(\\frac{r}{\\rm 1AU}\\right)^{-11/4}{\\rm g\\;cm^{-3}}$. % Then, the initial vertical magnetic field of $\\beta_0=10^6$ in our fiducial run corresponds to $B_{z,0}\\approx 0.01$ G, and the saturated field strength is $B\\approx 1$ G at 1 AU.  First, we examine how much the disk wind contributes to the evaporation of protoplanetary disks \\citep[see e.g.][for other mechanisms]{dul07}. After the saturation of the magnetic fields, $\\approx 5$\\% of the total disk mass is lost from the simulation box from 200 to 400 rotations by the disk winds in our fiducial case. Assuming a disk around a central star with the solar mass (1 rotation = 1 yr at 1 AU), we have the timescales of the evaporation, $\\tau_{\\rm ev}\\approx 4000$ yr at 1 AU, and $6\\times 10^5$ yr at 30 AU. Although this is rather short in comparison with recent observational results \\citep[typically $\\tau_{\\rm ev}\\sim 10^{6-7}$ yrs, e.g., ][]{hll01}, this is not a severe contradiction because we have not yet taken into account the global radial accretion of the disk mass, which continuously supplies the mass from the outer region. Another important issue that affects, and might reduce, the mass flux of disk winds is the effect of resistivity, which requires an additional detailed analysis of the ionization structure \\citep{Sano00,InutsukaSano2005} and will be the scope of our next paper. Here, the estimated $\\tau_{\\rm ev}$ should be taken as a lower limit.  The disk scale height has a relation of $H_0/r\\propto r^{1/4}$ for the MMSN. Combining with Equation (\\ref{eq:msfx}), we infer that the dynamical evaporation by disk winds, in comparison with accretion, becomes relatively more important in the inner parts of protoplanetary disks than in the outer regions for a constant initial $\\beta_0$ structure ($B_{z,0}^2\\propto r^{-11/4}$ for the MMSN).     Finally, we should point out the effects of waves on dusts in protoplanetary disks. We have shown that the momentum flux of Alfv\\`{e}nic and sound-like waves directs to the midplane from the injection regions. The momentum flux of the sound-like waves ($\\delta\\rho\\delta v_z$) can push dust grains to the midplane by gas-dust collisions. Dusts are usually weakly charged; in this case Alfv\\'{e}nic waves also contribute to the sedimentation of dusts to the midplane through ponderomotive force or dust-cyclotron resonance \\citep{vj06}.  This work was supported in part by Grants-in-Aid for Scientific Research from the MEXT of Japan (T.K.S.: 19015004 and 20740100, S.I.: 15740118, 16077202, and 18540238), and Inamori Foundation (T.K.S.). Numerical computations were in part performed on Cray XT4 at Center for Computational Astrophysics, CfCA, of National Astronomical Observatory of Japan. The page charge of this paper is supported by CfCA."
    },
    "0812/0812.0419_arXiv.txt": {
        "abstract": "21cm emission from residual neutral hydrogen after the epoch of reionization can be used to trace the cosmological power spectrum of density fluctuations.  Using a Fisher matrix formulation, we provide a detailed forecast of the constraints on cosmological parameters that are achievable with this probe.  We consider two designs: a scaled-up version of the MWA observatory as well as a Fast Fourier Transform Telescope.  We find that 21cm observations dedicated to post-reionization redshifts may yield significantly better constraints than next generation Cosmic Microwave Background (CMB) experiments. We find the constraints on $\\Omega_\\Lambda$, $\\Omega_{\\rm m}h^2$, and $\\Omega_\\nu h^2$ to be the strongest, each improved by at least an order of magnitude over the Planck CMB satellite alone for both designs. Our results do not depend as strongly on uncertainties in the astrophysics associated with the ionization of hydrogen as similar 21cm surveys during the epoch of reionization.  However, we find that modulation of the 21cm power spectrum from the ionizing background could potentially degrade constraints on the spectral index of the primordial power spectrum  and its running by more than an order of magnitude.  Our results also depend strongly on the maximum wavenumber of the power spectrum which can be used due to non-linearities. ",
        "introduction": "Recently, there has been much interest in the feasibility of mapping the three-dimensional distribution of cosmic hydrogen through its spin-flip transition at a resonant rest frame wavelength of 21cm \\cite{Furlanetto:2006jb,2007RPPh...70..627B}.  Several first generation experiments are being constructed to probe the epoch of reionization (MWA \\cite{MWAref}, LOFAR \\cite{LOFARref}, PAPER \\cite{PAPERref}, 21CMA \\cite{21CMAref}) and more ambitious designs are being planned (SKA \\cite{SKAref}).  One driver for mapping hydrogen through 21cm emission is to measure cosmological parameters from the underlying cosmic power spectrum. During the epoch of reionization (EoR), the 21cm power spectrum is shaped mainly by the structure of ionized regions.  Even without precise knowledge of the ionization power spectrum it is possible to isolate the cosmological power spectrum by exploiting anisotropies in redshift space due to peculiar velocities \\cite{Barkana:2004zy,2006ApJ...653..815M}.  Recent work \\cite{2008PhRvD..78b3529M,2008arXiv0805.1920P} has shown that the 21cm power spectrum accessible during the EoR has the potential to put tight constraints on cosmological parameters; however, these constraints depend on model-dependent uncertainties, the most important of which involves the ionization power spectrum.  In this paper, we explore in detail the constraints achievable from mapping the residual cosmic hydrogen after reionization \\cite{2007arXiv0708.3392W,2007arXiv0709.2955W,Chang:2007xk,2008PhRvL.100p1301L,Pen:2008fw,2008arXiv0808.2323W}. The hydrogen resides in pockets of dense galactic regions which are self-shielded from the UV background, also known as damped Lyman-$\\alpha$ absorbers (DLAs) in quasar spectra \\cite{Wolfe:2005jd}. In difference from traditional galaxy redshift surveys, 21cm surveys do not need to resolve individual galaxies but rather are able to monitor the smooth variation in their cumulative 21cm emission owing to their clustering on large scales \\cite{2008PhRvL.100p1301L}.  Measuring the 21cm power spectrum after reionization offers several key advantages as compared to the EoR.  First, the near uniformity of the UV radiation field guarantees that the 21cm power spectrum would reliably trace the underlying matter power spectrum. However, the UV background could introduce a scale dependent modulation to the 21cm power spectrum as large as one percent \\cite{2008arXiv0808.2323W}.  This modulation could introduce degeneracies that would significantly weaken constraints on the spectral index of the primordial power spectrum and its running. Another advantage of probing low redshifts is that the brightness temperature of the galactic synchrotron emission scales as $(1+z)^{2.6}$.  However, this advantage is offset by the fact that the mass weighted neutral fraction of hydrogen is only a few percent at redshifts $z\\lesssim 6$ \\cite{Prochaska:2005wy}.  In this paper we use the Fisher matrix formalism to quantify how effectively futuristic surveys dedicated to post-reionization redshifts ($z\\lesssim6$) can constrain cosmology.  We consider both a survey with ten times the number of cross dipoles of the MWA \\cite{MWAref} ($5000\\times 16=8\\times 10^4$) but operating at higher frequencies, which we term MWA5000, and a Fast Fourier Transform Telescope (FFTT) with $10^6$ dipoles over a square kilometer area \\cite{2008arXiv0805.4414T}. We also show results for the combination of these surveys and a next generation Cosmic Microwave Background (CMB) experiment (Planck).  This paper is structured as follows. In \\S II we discuss the detectability of the 21cm signal after reionization and in \\S III we discuss its power spectrum. In \\S IV we consider the details of our Fisher matrix calculation and in \\S V we present its results. Finally, we discuss and summarize our conclusions in \\S VI and \\S VII. ",
        "conclusions": "A 21cm survey after reionization provides a promising probe of cosmology and fundamental physics \\cite{2008PhRvL.100p1301L}.  We find the simple form of the 21cm power spectrum after reionization provides less uncertainty in forecasting constraints than 21cm during the EoR.  However, we find that errors on estimates of the spectral index and the running of the primordial power spectrum could be increased by an order of magnitude if the ionizing background introduces scale dependent effects into the post reionization 21cm power spectrum.  Aside from its greater simplicity, the measurement sensitivity after reionization is improved over EoR surveys due to lower foreground brightness temperature, a bias that is greater than unity, and a larger growth factor of density perturbations. These advantages are somewhat offset by the fact that there is less neutral hydrogen after reionization.  The constraints we derive are competitive with those from the EoR and for the FFTT are significantly better than next generation CMB experiments on their own.  We find errors on $\\Omega_\\Lambda$, $\\Omega_{\\rm m}h^2$, and $\\Omega_\\nu h^2$ to be improved the most, each lowering the errors from Planck alone by at least an order of magnitude for both FFTT and the survey similar in design to MWA5000.    \\appendix*"
    },
    "0812/0812.3901_arXiv.txt": {
        "abstract": "If the accelerated expansion of the Universe at the present epoch is driven by a dark energy scalar field, there may well be a non-trivial coupling between the dark energy and the cold dark matter (CDM) fluid.  Such interactions give rise to new features in cosmological structure growth, like an additional long-range attractive force between CDM particles, or variations of the dark matter particle mass with time.  We have implemented these effects in the N-body code {\\small GADGET-2} and present results of a series of high-resolution N-body simulations where the dark energy component is directly interacting with the cold dark matter.  As a consequence of the new physics, CDM and baryon distributions evolve differently both in the linear and in the nonlinear regime of structure formation.  Already on large scales a linear bias develops between these two components, which is further enhanced by the nonlinear evolution.  We also find, in contrast with previous work, that the density profiles of CDM halos are less concentrated in coupled dark energy cosmologies compared with $\\Lambda $CDM, and that this feature does not depend on the initial conditions setup, but is a specific consequence of the extra physics induced by the coupling. Also, the baryon fraction in halos in the coupled models is significantly reduced below the universal baryon fraction. These features alleviate tensions between observations and the $\\Lambda $CDM model on small scales. Our methodology is ideally suited to explore the predictions of coupled dark energy models in the fully non-linear regime, which can provide powerful constraints for the viable parameter space of such scenarios. ",
        "introduction": "\\label{i}  The last decade has seen an astonishing amount of new cosmological data from many different experiments, ranging from large-scale structure surveys \\citep[e.g.]{Percival_etal_2001} to Cosmic Microwave Background (CMB) \\citep{wmap5} and Type Ia Supernovae \\citep{Riess_etal_1998, Perlmutter_etal_1999, SNLS} observations. These experiments all consistently show that the Universe is almost spatially flat with a current expansion rate of about $70\\,{\\rm km\\, s^{-1} Mpc^{-1}}$, and contains a total amount of matter that accounts for only $\\sim 24\\%$ of the total energy density.  From a theoretical point of view, understanding the remaining $76\\%$ -- which must be in the form of some dark energy (DE) component able to drive an accelerated expansion -- is a serious challenge.  A simple cosmological constant would be in agreement with a large number of observational datasets, but would also raise two fundamental questions concerning its fine-tuned value and the coincidence of its domination over cold dark matter (CDM) only at a relatively recent cosmological epoch.  An alternative consists of identifying the dark energy component with a dynamic scalar field \\citep{Wetterich_1988, Ratra_Peebles_1988}, thereby seeking an explanation of the fundamental problems challenging the cosmological constant in the properties of the dynamic evolution of such scalar field.  An interesting suggestion recently developed concerns the investigation of possible interactions between the dark energy scalar field and other matter species in the Universe \\citep{Wetterich_1995, Amendola_2000, Farrar2004, Gubser2004, Farrar2007}. The existence of such a coupling could provide a unique handle for a deeper understanding of the DE problem \\citep{Quartin:2008px, Brookfield:2007au, Gromov:2002ek, Mangano:2002gg, Anderson:1997un}. It is then crucial to understand in detail what effects such a coupling imprints on observable features like, for example, the CMB and structure formation \\citep{LaVacca_etal_2009, Bean:2008ac, Bertolami:2007zm, Matarrese_etal_2003, Wang:2006qw, Guo:2007zk, Mainini:2007ft, Lee:2006za}.  In this paper, we perform the first fully self-consistent high-resolution hydrodynamic N-body simulations of cosmic structure formation for a selected family of coupled DE cosmologies.  The interaction between DE and CDM is expected to imprint characteristic features in linear and nonlinear structures, and could possibly open up new ways to overcome a series of observational challenges for the $\\Lambda $CDM concordance cosmology, ranging from the satellite abundance in CDM halos \\citep{Navarro_Frenk_White_1995}, to the observed low baryon fraction in large galaxy clusters \\citep{Allen_etal_2004, Vikhlinin_etal_2006, LaRoque_etal_2006}, and to the so-called  ``cusp-core'' problem for the density profiles of CDM halos, that was first raised in relation to observations of dwarf and low-surface-brightness galaxies \\citep{Moore_1994, Flores_Primack_1994, Simon_etal_2003}, but that more recent studies seem to extend to a wider range of objects including Milky-Way type spiral galaxies \\citep{Navarro_Steinmetz_2000, Binney_Evans_2001} up to galaxy groups and clusters \\citep{Sand_etal_2002, Sand_etal_2004, Newman_etal_2009}.   In the following we will consider DE as the energy associated with a scalar field $\\phi$ \\citep{Wetterich_1988, Ratra_Peebles_1988} whose energy density and pressure are defined respectively as the $(0,0)$ and $(0,i)$ components of its stress energy tensor, so that\\footnote{In this work we always denote with a prime the derivative with respect to the conformal time $\\tau $, and with an overdot the derivative with respect to the cosmic time $t$. The two time variables are related via the equation: ${\\rm d}\\tau = {\\rm d}t / a$.}: \\be \\label{rho_phi} \\rho_\\phi \\equiv \\frac{1}{2}\\frac{{\\phi '}^2}{a^2} + U(\\phi) \\,, \\ee \\be \\label{p_phi} p_\\phi \\equiv \\frac{1}{2}\\frac{{\\phi '}^2}{a^2} - U(\\phi) \\,, \\ee with $U(\\phi)$ being the self interaction potential. The models we investigate in this work by means of detailed high-resolution N-body simulations are physically identical to the ones already studied by \\citet{Maccio_etal_2004}. With respect to this previous work, however, our analysis significantly improves on the statistics of the number of cosmic structures analyzed in each of our different cosmological models and on the dynamic range of the simulations.  We also extend the analysis to a wider range of observable effects arising from the new physics of the DE-CDM interaction, and we include hydrodynamics. Despite the physical models investigated be identical, our outcomes are starkly different from the ones found in \\citet{Maccio_etal_2004}.   This work is organized as follows. In Section~{\\ref{cde}} we detail the coupled DE cosmologies under investigation, both regarding the background evolution and the perturbations, and we introduce a numerical package to integrate the full set of background and perturbation equations up to linear order.  In Section~\\ref{sim}, we describe the numerical methods we use, and the particular set of simulations we have run.  We then present and discuss our results for N-body simulations for a series of coupled DE models in Section~\\ref{results}.  Finally, in Section~\\ref{concl} we draw our conclusions. ",
        "conclusions": "\\label{concl}  We have investigated coupled DE cosmologies, both with respect to their expected background and linear perturbation evolution, as well as in their predictions for the nonlinear regime of structure formation.  To do so we have developed and tested in this work a modified version of the cosmological N-body code {\\small GADGET-2} suitable for evolving these kinds of cosmological models.  The numerical implementation we have developed is in fact quite general and not restricted to the simple specific models of coupled quintessence that we have investigated in this paper. Instead, it should be well suited for a much wider range of DE models. We also note that the ability to selectively enable or disable each of the modifications discussed above, makes the code suitable for cosmological models that are unrelated to the coupled DE scenario but require similar new degrees of freedom that our implementation allows. These are: \\begin{enumerate} \\item the expansion history of the Universe can be specified according to any desired evolution of the Hubble rate as a function of the scale factor $a$; \\item a global variation in time of the gravitational constant and/or a variation in time of the gravitational strength for each individual matter species. This includes the possibility to have a long range repulsive interaction between different particle species; \\item variation in time of the particle mass of each individual matter species; \\item extra velocity-dependent terms in the equation of motion for each individual matter species. \\end{enumerate}   With this implementation we have investigated the effects of coupled DE models with a constant coupling $\\beta _{c}$ to the CDM fluid on structure formation. We have shown that the halo mass function is modified in coupled DE models, but can still be well fitted at different redshifts by the {\\em Jenkins et al.} \\citep{Jenkins_etal_2000} fitting formula, or by the {\\em Sheth \\& Tormen} \\citep{Sheth_Tormen_1999} formula, which yields a moderately better agreement, especially at $z > 0$.  We have confirmed the analytic prediction that density fluctuations in baryons and CDM will develop a bias on all scales due to the presence of a fifth-force acting only between CDM particles.  We have also shown that in addition to this the  bias is enhanced when moving from the linear regime of very large scales to smaller and progressively more non-linear scales.  We have investigated the evolution of the bias between baryons and CDM overdensities down to the very nonlinear regime found in the inner part of collapsed objects, in the same fashion as described in \\citet{Maccio_etal_2004}.  We found here similar results with this previous work, namely an enhancement of the bias in the nonlinear region within and around massive halos. We also recover from this analysis the large scale value of the linear bias computed from the power spectrum when integrating the bias function up to very large radii from the centre of CDM halos.  The enhancement of the bias in highly nonlinear structures has an impact on the determination of the baryon fraction from cluster measurements, and we have computed for all our halos, without including non-adiabatic processes, the evolution of this fraction within the virial radius $r_{200}$ with coupling, finding that the baryon fraction is reduced with increasing coupling by up to $\\sim 8-10\\%$ with respect to $\\Lambda $CDM for the largest coupling value.    We have also investigated the effect of the coupling on the halo density profiles. We find that they are remarkably well fit over the resolved range by the NFW formula for any value of the coupling. There is a clear trend of a decrease of the inner halo overdensity with respect to $\\Lambda $CDM with increasing coupling (or, equivalently, an increase of the scale radius $r_{s}$ for increasing coupling). This result conflicts with previous claims for the same class of coupled DE models \\citep{Maccio_etal_2004}.  Using a number of special test simulations, we have identified the origin of this effect of reduced halo concentrations for increasing coupling.  It actually arises from a combination of two peculiar features that the coupling introduces in the Newtonian limit of gravitational dynamics.  The first of these is the decrease of CDM particle mass with time, which causes the total potential energy of a halo to decrease, and hence effectively moves the system to a configuration where an excess of kinetic energy is present relative to virial equilibrium.  The second one is the additional velocity-dependent term, which directly raises the total kinetic energy density of halos by accelerating their particles in the direction of their peculiar velocity. Both of these effects cause a halo to slightly expand in order to restore virial equilibrium, and this reduces the halo concentration.  In conclusion, we have developed a general numerical implementation of coupled dark energy models in the {\\small GADGET-2} code. We have then performed the first fully self-consistent high-resolution hydrodynamic N-body simulations of interacting dark energy models with constant coupling, and carried out a basic analysis of the non-linear structures that formed.  Interestingly, we found that a larger coupling leads to a lower average halo concentration. Furthermore, both the baryon fraction in massive halos and the inner overdensity of CDM halos decrease with increasing coupling.  These effects alleviate the present tensions between observations and the $\\Lambda$CDM model on small scales, implying that the coupled DE models are viable alternatives to the cosmological constant included in standard $\\Lambda$CDM."
    },
    "0812/0812.1559_arXiv.txt": {
        "abstract": "{Using Monte Carlo simulations we analyze the potential of the upcoming transit survey Pan-Planets. The analysis covers the simulation of realistic light curves (including the effects of ingress/egress and limb-darkening) with both correlated and uncorrelated noise as well as the application of a box-fitting-least-squares detection algorithm. In this work we show how simulations can be a powerful tool in defining and optimizing the survey strategy of a transiting planet survey. We find the Pan-Planets project to be competitive with all other existing and planned transit surveys with the main power being the large 7 square degree field of view. In the first year we expect to find up to 25 Jupiter-sized planets with periods below 5 days around stars brighter than V = 16.5 mag. The survey will also be sensitive to planets with longer periods and planets with smaller radii. After the second year of the survey, we expect to find up to 9 Warm Jupiters with periods between 5 and 9 days and 7 Very Hot Saturns around stars brighter than V = 16.5 mag as well as 9 Very Hot Neptunes with periods from 1 to 3 days around stars brighter than i' = 18.0 mag.} ",
        "introduction": "Thirteen years after the first discovery of a planet revolving around a main-sequence star \\citep{1995Natur.378..355M} more than 300 extra solar planets are known. The majority of these have been detected by looking for periodic small amplitude variations in the radial velocity of the planet's host star. This method reveals the period, semi-major axis and eccentricity of the planetary orbit, but due to a generally unknown inclination only the minimum mass of the planet can be derived. The situation is different in the case of a planet that transits its host star. The planet blocks a small fraction of the stellar surface resulting in a periodic drop in brightness. In combination with radial velocity measurements the light curve provides many additional parameters like inclination, true mass and radius of the planet.\\\\ Transiting planets are subject of many detailed follow-up studies such as measurement of thermal emission using the secondary transit \\citep{2008arXiv0802.0845C,2008ApJ...673..526K} or measurement of the spin-orbit alignment using the Rossiter-McLaughlin effect \\citep[e.g.][]{2007ApJ...665L.167W}.\\\\ In the past years many transit projects have monitored hundreds of thousands of stars looking for periodic drops in the light curves. A total of about 50 transiting planets are known to date. Remarkably, more than half of the transiting planets have been found in the past year making the transit method equally successful in that period compared to the radial velocity method.\\\\ The majority of the recently detected transiting planets have been found by wide-angle surveys targeting bright stars such as WASP, HAT, TrES or XO \\citep{2006PASP..118.1407P,2008ApJ...673L..79N,2007ApJ...663L..37O,2005PASP..117..783M}. Also the space mission Corot has contributed by adding four new discoveries \\citep{2008ASPC..384..270A}. Deep surveys like OGLE targeting highly crowded regions of the Milky Way disk have not been able to keep up with the increased detection rate of all-sky monitoring programs mainly due to limited amount of observation time and a lower number of target stars.\\\\ In 2009 Pan-Planets - a new deep transit survey - will start taking first observations. This project will be more powerful than all existing deep surveys because of its by far larger field of view, bigger telescope and faster readout.\\\\\\\\ With the first detections of transiting extra-solar planets, several groups have started to predict the number of planets that could be found by existing and planned surveys. First estimates were based on optimistic assumptions and have been mostly over-predictions \\citep[e.g.][]{2003ASPC..294..361H}. For example the frequency of very close-in planets had been extrapolated from the metallicity biased results of the radial velocity surveys. Late type dwarfs with higher metallicity turned out to have a higher frequencies of close-in planets \\citep{2005ApJ...622.1102F} and therefore transit surveys find less planets compared to radial velocity surveys due to a lower average metallicity \\citep{2006AcA....56....1G}. It was further assumed that planets could be found around all stars in the target fields whereas planets transiting giants show a much too faint photometric signal due to the larger radius of the star. In addition, the efficiency of the detection algorithm was not taken into account and all light curves with 2 visible transits were assumed to lead to a detection which is not the case.\\\\ More realistic methods have been introduced by \\citet{2003AcA....53..213P} and \\citet{2006AcA....56....1G}. Both groups use an analytical approach assuming a stellar and a planetary distribution and integrating over period, stellar mass, planetary radius and volume probed taking into account the detection probability. Similarly, \\citet{2007A&A...475..729F} modeled the OGLE survey and compared the predicted distributions to the parameters actually found by OGLE. In a recent study, \\citet{2008arXiv0804.1150B} generalized the formalism of \\citet{2006AcA....56....1G} in order to provide a method that can be used to calculate planet yields for any photometric survey given the survey parameters like number of nights observed, bandpass, exposure time, telescope aperture, etc. They applied their method to a number of different planned surveys like SDSS-II and the Pan-STARSS 3$\\pi$ survey.\\\\ In this work we use Monte Carlo simulations to predict the number of planets of the Pan-Planets survey. Our approach is quite general and applicable to any transit survey. Based on stellar and planetary populations we model the survey by constructing realistic light curves and running a detection algorithm on them. In this way we are able to directly include the effects of limb darkening, ingress/egress and observational window functions which have not been included in most previous studies. In addition we introduce a model for correlated noise and study its impact on the efficiency of the detection algorithm.\\\\ To optimize the survey strategy of Pan-Planets we want to address the following questions: What is the best observing block size (1h or 3h) and how many fields (3 to 7) should we observe? Given the optimized survey strategy, we study how many Very Hot Juptiters (VHJ) and Hot Jupiters (HJ) are expected in the first year and what is the potential of Pan-Planets to find planets with longer periods, such as Warm Jupiters (WJ) or planets with smaller radii, such as Very Hot Saturns (VHS) and Very Hot Neptunes (VHN). We further study whether it will be more efficient to observe the same target fields in the second year of the Pan-Planets survey or to choose new ones.\\\\\\\\ In \\S\\ref{overview} we give a brief overview of the Pan-Planets survey. \\S\\ref{simulations} describes in detail the simulations we performed.  We present our results in \\S\\ref{results}.  In order to verify our results we perform a consistency check with the OGLE-III survey by comparing our predicted yield with the actual number of planets found (\\S\\ref{consistency}). Finally we draw our conclusions in \\S\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} The aim of this work was to study the influence of the survey strategy on the efficiency of the Pan-Planets project and to predict the number of detections for an optimized strategy.\\\\ Our calculations are based on the simulation of realistic light curves including the effects of limb-darkening, ingress/egress and observational window functions. In addition we have introduced a model to simulate correlated (red) noise which allows us to include the effects of correlated noise on the efficiency of the BLS detection algorithm. Our approach can be applied to any transit survey as well.\\\\\\\\ Below we summarize the caveats and assumptions that were made in our simulations : \\begin{itemize} \\item{Our results depend on the spectral type and magnitude distribution of the Besan\\c{c}on model. The model does not include second order substructure such as spiral arms.} \\item{We neglect the effects of blending. Due to crowding into the direction of the Galactic disk some stars are blended by neighboring sources.} \\item{We assume all planets that are detected by the BLS-algorithm to be followed-up and confirmed spectroscopically. In particular, we assume that no true candidate is rejected by any candidate selection process. The detailed follow-up strategy of the Pan-Planets survey will be presented in Afonso et al. (in prep.).} \\item{Our simulations are done for 1 sq.deg. and the results are scaled to the actual survey area. We assume that all fields (3, 4, 5, 6 or 7 case) have homogeneous densities and non-varying (or similar) stellar populations. Simple number counts on the USNO-catalog showed that we can find up to 7 fields with similar total number of stars (see \\S\\ref{results}). For a larger number of fields, the assumption of a constant density might be too optimistic since we are restricted to fields that are close to each other in order to keep the observational overhead low.} \\item{Our results directly scale with the assumed planet frequencies. The values of 0.14\\% and 0.31\\% we use for VHJ and HJ have uncertainties of a factor of 2. For WJ, VHS and VHN we have used hypothetical values of 0.31\\%, 0.14\\% and 5\\% respectively.  After completion of the Pan-Planets survey we will be able to derive more accurate absolute frequencies for all five planet populations.} \\item{The quality of the data is assumed to be homogeneously good over the whole detector area. Bad pixel regions and gaps between the individual CCDs are not taken into account and result in an effective field of view that is smaller than 7 sq.deg.} \\end{itemize} Comparing different observing strategies we found that observing more fields is more efficient. Concerning the observation time per night, we compared 1h blocks to 3h blocks and found the shorter ones to be more efficient. This is still the case for a 2 yr campaign.\\\\\\\\ For an RRN level of 2 mmag we expect to find up to 15 VHJ and 10 HJ in the first year around stars brighter than V = 16.5 mag. The survey will also be sensitive to planets with longer periods (WJ) and smaller radii (VHS and VHN). Assuming that the frequencies of stars with WJ and VHS is 0.31\\% and 0.14\\% respectively, we expect to find up to 2 WJ and 3 VHS in the same magnitude range.\\\\ We found that observing the same fields in the second year of the 3.5 yr lifetime of the survey is more efficient than choosing new fields. We expect to find up to 24 VHJ, 23 HJ, 9 WJ and 7 VHS. In particular for longer periods (HJ and WJ) and smaller radii (VHS) we will more than double the number of detections of the first year if we continue to observe the same targets.\\\\ We have investigated the potential of the Pan-Planets survey to detect VHN transiting M dwarfs brighter than i' = 18 mag. Assuming the frequency of these objects is 5\\%, we expect to find up to 3 detections in the first year and up to 9 detections observing the same fields in the second year.\\\\ As a consistency check we modeled the OGLE-III Carina survey and found 2.18 VHJ, 1.46 HJ, 0.45 WJ, 0.12 VHS and zero VHN which is in agreement with the 2 VHJ and 1 HJ and the zero WJ, VHS and VHN that have been actually detected."
    },
    "0812/0812.2650_arXiv.txt": {
        "abstract": "Torsional (shear) oscillations of neutron stars may have been observed in quasiperiodic oscillations of Magnetar Giant Flares. The frequencies of these modes depend on the shear modulus of neutron star crust.   We calculate the shear modulus of Coulomb crystals from molecular dynamics simulations.  We find that electron screening reduces the shear modulus by about 10\\% compared to previous Ogata et al.   results.  Our MD simulations can be extended to calculate the effects of impurities and or polycrystalline structures on the shear modulus. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.0978_arXiv.txt": {
        "abstract": "The high-frequency-peaked BL-Lacertae object \\objectname{1ES 0806+524}, at redshift z=0.138, was observed in the very-high-energy (VHE) gamma-ray regime by VERITAS between November 2006 and April 2008. These data encompass the two-, and three-telescope commissioning phases, as well as observations with the full four-telescope array. \\objectname{1ES 0806+524} is detected with a statistical significance of 6.3 standard deviations from 245 excess events. Little or no measurable variability on monthly time scales is found. The photon spectrum for the period November 2007 to April 2008 can be characterized by a power law with photon index $3.6 \\pm 1.0_{\\mathrm{stat}} \\pm 0.3_{\\mathrm{sys}}$ between $\\sim$300 GeV and $\\sim$700 GeV. The integral flux above 300 GeV is $(2.2\\pm0.5_{\\mathrm{stat}}\\pm0.4_{\\mathrm{sys}})\\times10^{-12}\\:\\mathrm{cm}^{-2}\\:\\mathrm{s}^{-1}$ which corresponds to 1.8\\% of the Crab Nebula flux. Non-contemporaneous multiwavelength observations are combined with the VHE data to produce a broadband spectral energy distribution that can be reasonably described using a synchrotron-self-Compton model. ",
        "introduction": "Active galactic nuclei (AGN) are galactic cores with a luminosity that dominates the host galaxy. These central engines are believed to be powered by supermassive black holes which draw surrounding matter into an accretion disc. The unified model \\citep{1995PASP..107..803U} of AGN implies that the observed emission from AGN is a strong function of the observation angle in relation to the relativistic jet, with the highest energy observed emission beamed along the direction of the jet.   \\objectname{1ES 0806+524} was identified as a BL Lacertae object \\citep{1993ApJ...412..541S} using radio observations from the Green Bank 91-m telescope \\citep{becker91} and X-ray observations from the Einstein Slew Survey \\citep{elvis92}. The redshift of the host galaxy is z=0.138 \\citep{bade98}. The blazar \\objectname{1ES 0806+524} was suggested as a candidate VHE gamma-ray source based on the presence of both high-energy electrons and sufficient seed photons \\citep{costamante2002}. That work predicts an intrinsic flux in the VHE regime of $\\mathcal{F}_{\\mathrm{E}>0.3\\:\\mathrm{TeV}}=1.36 \\times 10^{-11}\\:\\mathrm{cm}^{-2}\\:\\mathrm{s}^{-1}$.  Very high-energy gamma-ray observations of \\objectname{1ES 0806+524} reported by the Whipple Collaboration \\citep{horan04a,delacalleperez03} indicated flux upper limits of $\\mathcal{F}_{\\mathrm{E}>0.3\\:\\mathrm{TeV}}<1.4 \\times 10^{-11}\\:\\mathrm{cm}^{-2}\\:\\mathrm{s}^{-1}$ and $\\mathcal{F}_{\\mathrm{E}>0.3\\:\\mathrm{TeV}}<16.8 \\times 10^{-11}\\:\\mathrm{cm}^{-2}\\:\\mathrm{s}^{-1}$ from the 1996 and 2000 observing seasons and $\\mathcal{F}_{\\mathrm{E}>0.3\\:\\mathrm{TeV}}<1.47 \\times 10^{-11}\\:\\mathrm{cm}^{-2}\\:\\mathrm{s}^{-1}$ from the 2000 to 2002 observing seasons. Observations reported by the HEGRA Collaboration \\citep{2004A&A...421..529A} indicated an upper limit of $\\mathcal{F}_{\\mathrm{E}>1.09\\:\\mathrm{TeV}}<43 \\times 10^{-11}\\:\\mathrm{cm}^{-2}\\:\\mathrm{s}^{-1}$ based on hour of observations. Observations of \\objectname{1ES 0806+524} in late 2006 and early 2007 with VERITAS \\citep{Cogan2007a} yielded evidence for weak emission with a statistical significance of 2.5 standard deviations in 35 hours. The data reported in \\cite{Cogan2007a} are combined with the data from the 2007/2008 observing season here. ",
        "conclusions": "VERITAS has observed the blazar \\objectname{1ES 0806+524} for a total 65 hours from November 2006 to April 2008 resulting in the discovery of very-high-energy gamma rays with a statistical significance of $6.3\\:\\sigma$. The differential energy spectrum between $\\sim$300 GeV and $\\sim$700 GeV can be fitted by a relatively soft power law with index $\\Gamma=3.6 \\pm 1.0_{\\mathrm{stat}}\\pm 0.3_{\\mathrm{sys}}$. The integral flux above 300 GeV is $(2.2\\pm0.5_{\\mathrm{stat}}\\pm0.4_{\\mathrm{sys}})\\times10^{-12}\\:\\mathrm{cm}^{-2}\\:\\mathrm{s}^{-1}$ which corresponds to $\\left(1.8\\pm0.5\\right)$\\% of the Crab Nebula flux as measured by VERITAS above 300 GeV. Assuming absorption on the infrared component of the extragalactic background light according to \\cite{2008A&A...487..837F}, the intrinsic integral flux above 300 GeV is $(4.4\\pm0.6_{\\mathrm{stat}}\\pm0.5_{\\mathrm{sys}})\\times10^{-12}\\:\\mathrm{cm}^{-2}\\:\\mathrm{s}^{-1}$ which is approximately one order of magnitude less than the flux predicted by \\cite{costamante2002}. The broadband spectral energy distribution can be fitted using a one-zone SSC model with standard parameters.  Observations by the EGRET gamma-ray space telescope reveal gamma-ray emission from the region around \\objectname{1ES 0806+524}, however the large error box indicates that the emission could be associated with either \\objectname{B 0803+5126} (z=1.14) or \\objectname{1ES 0806+524}, with the former being favored \\citep{2003ApJ...590..109S}. Observations of this region with the Large Area Telescope on the Fermi Gamma-Ray Space Telescope should resolve the source(s) of gamma-ray emission in this region.   VERITAS is the most sensitive instrument of its kind in the Northern Hemisphere, and it is ideally suited to observations of extragalactic objects. This has been demonstrated by the detection of three new BL Lac objects by VERITAS in 2008 (\\cite{2008ATel.1415....1S}, \\cite{2008ATel.1753....1S}, \\cite{2008ApJ...684L..73A})."
    },
    "0812/0812.0388.txt": {
        "abstract": "%Context: { %{\\bf \\Large NOTE: Article Draft, please do not spread!}\\\\ The use of model atmospheres for deriving stellar fundamental parameters, such as $T_{\\rm eff}$, $\\log g$ and [Fe/H], will increase as we find and explore extreme stellar populations where empirical calibrations are not yet available. Moreover calibrations for upcoming large satellite missions of new spectrophotometric indices, similar to the $uvby-$H$\\beta$ system, will be needed. } %Aim: {We aim to test the power of theoretical calibrations based on a new generation of MARCS models by comparisons with observational photomteric data. } %Methods {We calculate synthetic $uvby-$H$\\beta$ colour indices from synthetic spectra. A sample of 388 field stars as well as stars in globular clusters is used for a direct comparison of the synthetic indices versus empirical data and for scrutinizing the possibilities of theoretical calibrations for temperature, metallicity and gravity. } %Results: {We show that the temperature sensitivity of the synthetic $(b-y)$ colour is very close to its empirical counterpart, whereas the temperature scale based upon H$\\beta$ shows a slight offset. The theoretical metallicity sensitivity of the $m_{1}$ index (and for G-type stars its combination with $c_1$) is somewhat larger than the empirical one, based upon spectroscopic determinations. The gravity sensitivity of the synthetic $c_1$ index shows a satisfactory behaviour when compared to obervations of F stars. For stars cooler than the sun a deviation is significant in the $c_{1}$--$(b-y)$ diagram. The theoretical calibrations of $(b-y)$, $(v-y)$ and $c_{1}$ seem to work well for Pop II stars and lead to effective temperatures for globular cluster stars supporting recent claims by Korn et al. (2007) that atomic diffusion occurs in stars near the turnoff point of NGC 6397. }",
        "introduction": "The $uvby$ photometric intermediate-band system of Str\\\"omgren (1963) and the H$\\beta$ narrow-band system (Crawford 1958, 1966) were combined early and soon proved to be a most powerful means of determining fundamental parameters of stars. Calibrations of the systems were also developed early by Str\\\"omgren, Crawford and collaborators. Generally, semi-empirical methods were used with the calibrations made by means of sets of stars with parameters determined in more fundamental ways. The first metallicity calibration of the $m_1$ index for solar-type stars by Str\\\"omgren (1964) was thus based on spectroscopical [Fe/H] values of Wallerstein (1962) and the first luminosity calibration of the $c_1$ index by Str\\\"omgren (see Crawford 1966) used cluster stars in the main-sequence band. This general semiempirical approach continued with calibrations like those of the $m_1$ index by Nissen (1970), Gustafsson $\\&$ Nissen (1972), Nissen $\\&$ Gustafsson (1978) and Nissen (1981) where very-narrow-band spectrophotometry of groups of weak lines was calibrated with synthetic spectra. These stars were then used as calibration stars for $uvby$-H$\\beta$ photometry. Other calibrations follow this semiempirical approach e.g. those of Crawford (1975), Ardeberg $\\&$ Lindgren (1981) and Olsen (1988) for G and K-type stars, as well as by Schuster \\& Nissen (1989) for metal-poor stars. In the calibration by Alonso et al. (1996) the effective temperature calibration was based on stellar temperatures from the Infrared Flux Method. The calibrations by Nordstr\\\"om et al. (2004) and Holmberg et al. (2007) additionally utilised Hipparcos parallaxes for the surface gravity calibration.   Another development towards calibration of the systems also started early: the direct calculation of photometric indices by means of model atmospheres. Such theoretical calibrations were attempted for early-type stars with relatively line-free spectra. For late type stars, a statistical correction for the effects of spectral lines was made by Baschek (1960) in his calibration of the Str\\\"omgren $m$ index (a predecessor to $m_1$).  A first systematic and detailed calculation of $uvby-$H$\\beta$ indices for a grid of F and G dwarf model atmospheres was published by Bell (1970), using scaled solar model atmospheres. Bell $\\&$ Parsons (1974) calculated $uvby$ colours for flux-constant model atmospheres of F and G supergiants, while Gustafsson $\\&$ Bell (1979) produced theoretical colours in a number of systems, including the $uvby$ system, for a grid of giant-star model atmospheres. Relya \\& Kurucz (1978) calculated $uvby$ and $UBV$ colours from early ATLAS models, and discussed their shortcomings for late-type stars. $uvby$ colours for new sets of Kurucz models were published by Lester, Gray $\\&$ Kurucz (1986). Castelli $\\&$ Kurucz (2006) published H$\\beta$ indices. Sometimes, semiempirically corrected fluxes from model atmospheres have also been used for calibrations of Str\\\"omgren photometry, see e.g. Lejeune et al. (1999) and Clem et al. (2004).  The need for reliable calibrations of $uvby-$H$\\beta$ photometry has increased in the last decade, not the least for estimating parameters of new and more \"exotic\" stars, such as very metal-poor and super-metal-rich stars which are not found at great abundance in the solar neigbourhood so that relatively complete sets of calibration stars cannot easily be established. Furthermore, the preparation for the Gaia satellite includes a careful analysis of the power of model atmospheres to provide a detailed astrophysical calibration of the photometric system of the satellite. This analysis needs support by a detailed test of the problems and possibilities to make a detailed theoretical calibration of, e.g., the $uvby-$H$\\beta$ photometry.  Subsequently, we shall present the theoretical models and colours (Sect. \\ref{sec:models} and \\ref{sec:colours}). Stellar samples for empirical comparisons are discussed in Section \\ref{sec:stellsamp}. Next, the discussion will be focused on the determination of effective temperature, metallicity and surface gravity of the stars, by discussing the calibration of $(b-y)$ and H$\\beta$ indices (Sect. \\ref{sec:teffcalib}, effective temperature), of $m_1$ (Sect. \\ref{sec:metcalib}, metallicity) and $c_1$ (Sect. \\ref{sec:gravcalib}, gravity) indices, devoting more limited interest to \"secondary\" effects such as the metallicity sensitivity of $(b-y)$ and $c_1$, or the gravity sensitivity of $m_1$. In each section the results will be compared with empirical and semi-empirical data and calibrations. Finally, in the last Section some comments will be made on the success and the problems of the theoretical calibrations, conclusions will be drawn and recommendations given. ",
        "conclusions": "We have explored the possibilities and shortcomings of synthetic $uvby-$H$\\beta$ photometry based on new MARCS model atmospheres. In general a good agreement with empirical calibrations of this photometric system is found. However, a number of systematic deviations between theory and observations also become apparent. The temperature sensitivity of the $(b-y)$ colour (i.e. $\\delta (b-y)/\\delta T_{\\rm eff}$) seems marginally larger for the calculated colours than is found when using infrared-flux-method determinations of temperatures. %, although this might possibly be due to problems with the latter empirical temperature scale %instead. A similar, and even somewhat larger, difference occurs for temperatures based on the H$\\beta$ index, when compared with the empirical scale. The $(b-y)$ calibration for Pop II stars supports Korn et al. (2007), who claim the signatures of atomic diffusion in the metal-poor globular cluster NGC 6397.  For the metallicity index $m_1$ the theoretical sensitivity $\\delta m_1/\\delta $[Me/H] is somewhat larger than the empirical one that is based on spectroscopic [Fe/H]-determinations. For the gravity sensitivity of the $c_1$ Balmer discontinuity measure, we find a reasonably good agreement with observations for stars hotter than the Sun, where the Balmer discontinuity is significant. Considerable problems remain for the cooler stars, although the model calibration works well for Pop II giants.  One might ask whether these problems may be solved when even more detailed atomic and molecular line data become available to feed into model atmospheres. This is possible but it is also possible that thermal inhomogenities as generated by convection and possibly also non-LTE effects, contribute to the difference between synthetic and observed values of $(b-y)$ and $m_1$. Concerning the problems in reproducing the observed $c_1$ indices of cooler stars, these may well be due to failures in line-data, although thermal inhomogeneities in the atmospheres may also here turn out to be the important effect. One should also realize that there are still residual problems associated with differences and uncertainties in the observed and adopted Str\\\"omgren passbands (Manfroid \\& Sterken, 1987). The current study illustrates the problems in synthetic photometry in the visual wavelength regions. %such as it is currently applied for the calibration of the photometric system of the Gaia %satellite. Continued efforts along such lines must be complemented with more detailed studies of the shortcoming of classical model atmospheres and the replacement of those, also in large-scale calibration efforts, by physically more realistic models."
    },
    "0812/0812.4306_arXiv.txt": {
        "abstract": "{} { In this study we investigate the evolution of shape and kinematics of elliptical galaxies in a cosmological framework. } {We use a set of hydrodynamic, self-consistent simulations operating in the context of a concordance cosmological model where relaxed elliptical-like objects (ELOs) were identified at redshifts $z=0$, $z=0.5$, $z=1$ and $z=1.5$.} {The population of elliptical systems analysed here present a systematic change through time, i.e. evolution, by becoming rounder in general at $z=0$ and, at the same time more velocity dispersion supported. This is found to be primarily due to major dry mergers where only a modest amount of angular momentum is involved into the merger event. Despite the general trend, in a significant amount of cases the merger event involves a higher specific angular momentum, which in general causes the system to acquire a higher rotational support and/or a more oblate shape. These evolutionary patterns are still present when we study our systems in projection,  mimicking real observations, and thus they should become apparent in future observations.} {} ",
        "introduction": "Recent surveys of high-redshift galaxies suggest that some relaxed, massive elliptical galaxies could already be in place at $z \\sim 2-1.5$ (Cimatti et al. 2004, Conselice et al. 2007), or even by $z \\sim 5$ (Wiklind et al.~2008). These results notwithstanding, according to the current formation scenarios, merging has played an important role in the mass assembly of most of the local massive elliptical galaxies. In fact, observations and theory suggest that, first, violent mergers at high $z$ have transformed most of the available gas into stars, and, later on, these systems might have evolved through gas-free (the so-called dry) mergers (Conselice 2006; Scarlata et al.~2007).  Since the study by Bertola \\& Cappacioli (1975) we know that elliptical galaxies are supported against gravity by random motions as well as rotation. Davies et al. (1983) studied the now classical $V_{max}/\\sigma_o$ vs. $\\epsilon$ diagram for spheroids (Illingworth 1977; Binney 1978), where $V_{max}$ is the maximum of the line-of-sight rotation curve, $\\sigma_o$ is the central l.o.s. velocity dispersion and $\\epsilon$ is the mean ellipticity inside a given radius. They found that luminous (and massive) elliptical galaxies were characterised by low $V_{max}/\\sigma_o$ and a fairly round aspect (low $\\epsilon$), while ellipticals with intermediate luminosity tend to have larger values of $V_{max}/\\sigma_o$ and $\\epsilon$. Several later observations in 1D spectroscopy of near-by elliptical galaxies (Lauer 1985, Bender 1988, Nieto et al. 1989, Bender et al. 1994, Pellegrini 2005, Lauer et al. 2005) and 2D spectroscopy show such behaviour for ellipticals (see Emsellem et al. 2007, Cappellari et al. 2007, hereafter CAP07). Recently van der Marel \\& van Dokkum (2007) presented evidences of evolution of the rotation support of elliptical systems since z=0.5. Present formation schemes should explain these observations on the kinematics and shape of elliptical galaxies and their possible evolution.  A number of recent N-body simulations of isolated galaxy mergers have dealt with the population of the classical diagram and the formation of boxy and disky objects (Naab \\& Burkert 2003, Gonz\\'alez-Garc\\'{\\i}a \\& Balcells 2005, Gonz\\'alez-Garc\\'{\\i}a \\& van Albada 2005, Naab, Khochfar \\& Burkert et al. 2006, Bournaud et al. 2005, Robertson et al. 2006, Cox et al. 2006, Gonz\\'alez-Garc\\'{\\i}a et al. 2006, Jesseit et al. 2007, Naab \\& Ostriker 2007). These studies indicate that mergers between disk galaxies tend to produce too large rotation when compared with present day massive elliptical galaxies. Besides, mergers between elliptical galaxies do reproduce the observed characteristics of massive ellipticals. Khochfar \\& Burkert (2003), Kang et al. (2007) (and references there in) present first attempts of semi-analytical modelling to address the issue of kinematics and shape in early type galaxies. Naab et al. (2007) study the formation of three massive galaxies from cosmological initial conditions. However, a detailed and statistical analysis of the internal kinematics and shape of objects formed in fully self-consistent cosmological simulations and their possible evolution is, to the best of our knowledge, still missing and mandatory to provide a clear picture of the mechanisms at play in the formation and evolution of present day E's.  In this paper we present results from self-consistent cosmological simulations where we have investigated the shape and kinematic evolution of elliptical-like objects (here after ELOs) for several redshifts. For the present analysis we consider both three dimensional (hereafter, 3D) and projected onto the sky data. Other kinematic parameters recently proposed will be investigated in a fore-coming paper, but are out of the present study for the sake of simplicity. ",
        "conclusions": "We have  compared the merger characteristics with the changes measured in shape and rotational support. We have found that there is a close correlation between the shape and the amount of rotational support of a given ELO at a given time and the characteristics of the last merger event it has suffered. Specifically: MMs always produce changes in both shape and rotational support in the outer and the inner parts. MMs with little angular momentum will most often decrease the rotational support inside $r_{e}$ producing mostly rounder (i.e., a larger value of $c/a$) prolate spheroids, while MMs with a high (intermediate) amount of angular momentum increase rotational support producing oblate or triaxial systems. This picture  about the role of angular momentum is consistent with the transformation found in binary  mergers of spheroidal systems by Gonz\\'alez-Garc\\'{\\i}a \\& van Albada (2005). mMs most often produce triaxial or oblate spheroids and increase the rotational support, although the effect is most likely to affect the outer parts of the ELO (see Balcells \\& Quinn 1990, Eliche-Moral et al. 2006), except for penetrating mMs. Finally, aggregation processes are also important and may affect the shape and kinematics of the final object. We find that around $85\\%$ of the changes in shape and in rotational support are associated to merging (as opposite to aggregation).  Combining these findings on the trends of merger characteristics with time and the different effects a merger causes according to these characteristics, we can understand that prolate ELOs are the most frequent at high $z$, and then they are transformed most often into rounder triaxial objects with less rotation support. An interesting conclusion from this study is that, due to dry merging, the most luminous (i.e. massive) ellipticals must be rounder and less rotationally supported in general. Indeed, this is what is found in observational data (see e.g. Davies et al. 1983).  The results presented  here do not intend to fully represent at a quantitative level the  characteristics of the shape and kinematic evolution of elliptical galaxies, but rather to unveil some qualitative trends in dense environments and the mechanisms causing them in a cosmological context. It is reassuring to find that all the simulations analysed here show in general an evolution of the population of E's towards roundish and less rotationally supported systems, mostly driven by dry merging and that the mechanisms at play described above, causing either the general trend or the exceptions, are consistent in all kind of simulations, run with different box size, particle number, input $\\sigma_8$ and code.  Van der Marel \\& van Dokkum (2007) report a similar evolution for two samples of $\\sim 40$ galaxies at $z=0.5$ and $z=0$. Although the evolution we observe in our simulations seems milder than the one reported by this study, it remains open the amount of evolution to larger redshifts. Thus, it would be highly desirable to perform larger statistical studies of the shape and kinematics of real E's at low $z$ and to aim at observations of kinematics at higher redshifts to confirm the evolutionary trend of ellipticals.  This work was partially supported by the MCyT (Spain) grants: AYA-07468-C03-03 and AYA2006-15492-C03-01 from the PNAyA, and the regional government of Madrid through the ASTROCAM Astrophysics network (S\u00ad0505/ESP\u00ad0237). We thank the Centro de Computaci\\'on Cient\\'ifica (UAM, Spain) for computing facilities."
    },
    "0812/0812.4289_arXiv.txt": {
        "abstract": "We exploit the accumulating, high-quality, multi-wavelength imaging data of nearby supernova (SN) hosts to explore the relationship between SN production and host galaxy evolution.  The Galaxy Evolution Explorer ({\\it GALEX}, \\citep{Martin:05:L1}) provides ultraviolet (UV) imaging in two bands, complementing data in the optical and infra-red (IR).  We compare host properties, derived from spectral energy distribution (SED) fitting, with nearby, well-observed SN~Ia light curve properties.  We also explore where the hosts of different types of SNe fall relative to the red and blue sequences on the galaxy UV-optical color-magnitude diagram (CMD, \\citep{Wyder:07:293}).  We conclude that further exploration and larger samples will provide useful results for constraining the progenitors of SNe. ",
        "introduction": "The diagnostic power of {\\it GALEX} imaging for galaxy evolution is revealed in a plot of simple stellar population age versus wavelength shown at the top of Figure~1 in \\citet{Martin:05:L1}. This plot covers the UV to the IR and demonstrates that the UV is most sensitive to star formation on time scales of $\\sim$10$^8$ yr.  On the same figure, we see the flux of stellar populations with differing Scalo b ($\\equiv$ \\.M/$\\langle$\\.M$\\rangle$, where \\.M is the star formation rate in M$_{\\odot}$ yr$^{-1}$) plotted against wavelength.  The UV shows more discriminating power than other wavelengths.  {\\it GALEX} imaging is sensitive over at least five orders of magnitude of Scalo b, from b $= 10$ down to b $= 10^{-3}$.  Another illustration of the leverage of UV imaging on galaxy evolution is shown in the galaxy CMD in Figure~9 of \\citet{Wyder:07:293}.  Plotting galaxy near-UV minus r-band (NUV-r) color versus absolute r-band magnitude (M$_r$) clearly separates the blue and red sequences of galaxies, affording an exploration of the transition galaxies between these sequences.  Cosmic downsizing of galaxies hosting active star formation is also illustrated using {\\it GALEX} data, by plotting the surface density of star formation versus redshift (Figure~18 of \\citet{Martin:07:415}).  These evolution diagnostics are further enhanced when comparing a large local galaxy sample \\citep{Paz:07:185} in the restframe UV with deep optical high-z galaxy surveys that observe the younger universe also in the restframe UV.  What do these diagnostics tell us about galaxies hosting various types of SNe?  To begin to answer this, we analyze a preliminary set of $\\sim100$ galaxies hosting both core-collapse (CC) and thermonuclear (type Ia) SNe. We derive their integrated UV photometry using a preliminary version of the {\\it GALEX} Large Galaxy Atlas (GLGA, \\citep{Seibert:09}).  To this we add integrated optical and IR photometry extracted from online databases\\footnote{{\\tt http://nedwww.ipac.caltech.edu/} and {\\tt http://irsa.ipac.caltech.edu/}}.  We fit the resulting SEDs with star formation history models \\citep{Fioc:97:950, Borgne:02:446, Borgne:04:881}, and plot the photometry on diagnostic diagrams to explore the link between host evolution and SN production.  \\begin{figure}[t] \\includegraphics[angle=90.,height=0.3\\textheight]{neill_1_fig1} \\caption{Derived host stellar mass versus SN~Ia light curve stretch for the hosts of low-redshift, well-observed SNe~Ia.  Note that for hosts of SNe~Ia with stretch $ < 0.8$ (open red circles), the mass range is much narrower than hosts of higher-stretch SNe~Ia (open black circles).\\label{fig_str_mstar}} \\end{figure} ",
        "conclusions": "We have seen that SED fitting and diagnostic plots derived from integrated photometry of SN host galaxies have the potential to tell us about the host evolutionary state and hence constrain the progenitors of SNe of different types.  We will continue to accumulate high quality SN host photometry so we can improve the sample size and depth for each host and SN type.  We will also continue to refine our SED fitting technique and apply this and galaxy evolution diagnostic techniques to the host photometry.  Our first efforts will concentrate on the nearby set of SNe~Ia that have well determined stretch values \\citep{Neill:09}, but future efforts will explore all types of SNe.    \\begin{theacknowledgments} We acknowledge the use of the NASA Extragalactic Database and the Infra-Red Science Archive and their vital contribution to this work. \\end{theacknowledgments}"
    },
    "0812/0812.3656_arXiv.txt": {
        "abstract": "A white dwarf (WD) approaching the Chandrasekhar mass may in several circumstances undergo accretion-induced collapse (AIC) to a neutron star (NS) before a thermonuclear explosion ensues.  It has generally been assumed that such an accretion-induced collapse (AIC) does not produce a detectable supernova (SN).  If, however, the progenitor WD is rapidly rotating (as may be expected due to its prior accretion), a centrifugally-supported disk forms around the NS upon collapse.  We calculate the subsequent evolution of this accretion disk and its nuclear composition using time-dependent height-integrated simulations with initial conditions taken from the AIC calculations of Dessart et al.~(2006).  Soon after its formation, the disk is cooled by neutrinos and its composition is driven neutron-rich (electron fraction $Y_{e} \\sim 0.1$) by electron captures.  However, as the disk viscously spreads, it is irradiated by neutrinos from the central proto-NS, which dramatically alters its neutron-to-proton ratio.  We find that electron neutrino captures increase $Y_{e}$ to $\\sim 0.5$ by the time that weak interactions in the disk freeze out.  Because the disk becomes radiatively inefficient and begins forming $\\alpha$-particles soon after freeze out, powerful winds blow away most of the disk's remaining mass.  These $Y_{e} \\sim 0.5$ outflows synthesize up to a few times $10^{-2}M_{\\sun}$ in $^{56}$Ni.  As a result, AIC may be accompanied by a radioactively-powered SN-like transient that peaks on a timescale of $\\sim 1$ day.  Since few intermediate mass elements are likely synthesized, these Nickel-rich explosions should be spectroscopically distinct from other SNe.  The  timescale, velocity, and composition of the AIC transient can be modified if the disk wind sweeps up a $\\sim 0.1 M_{\\sun}$ remnant disk created by a WD-WD merger; such an ``enshrouded'' AIC may account for sub-luminous, sub-Chandrasekhar Type I SNe.  Optical transient surveys such as PanSTARRS and the Palomar Transient Factory should detect a few AIC transients per year if their true rate is $\\sim 10^{-2}$ of the Type Ia rate, and the Large Synoptic Survey Telescope (LSST) should detect several hundred per year.  High cadence observations ($\\lesssim 1$ day) are optimal for the detection and follow-up of AIC. ",
        "introduction": "\\voffset=-2cm \\vspace{-0.2 cm} \\label{sec:int} \\label{sec:intro}  It is generally thought that when an accreting C-O white dwarf (WD) approaches the Chandrasekhar mass $M_{\\rm Ch}$, the carbon ignites in its core, causing the WD to explode and produce a Type Ia supernova (SN).  If, however, the WD's composition is O-Ne instead of C-O, the outcome may be qualitatively different: electron captures on nuclei will cause the WD to collapse to a neutron star (NS) before an explosion can ensue (Miyaji et al.~1980; Canal et al.~1990; Gutierrez et al.~2005; Poelarends et al.~2008).  Such an accretion-induced collapse (AIC) may also occur for C-O WDs with large initial masses ($M > 1.1M_{\\sun}$) in systems with mass-transfer rates $\\dot{M} \\gtrsim 10^{-8}M_{\\sun}$ yr$^{-1}$ (Nomoto $\\&$ Kondo 1991) or following the merger of two WDs in a binary if their total mass exceeds $M_{\\rm ch}$ (e.g., Mochkovitch $\\&$ Livio 1989, 1990; Yoon et al.~2007).  Indeed, spectroscopic surveys of WD-WD binaries (e.g., SPY; Napiwotzki et al.~2004) have discovered at least one system that will merge within a few Hubble times with a total mass tantalizingly close to $M_{\\rm Ch}$ (Napiwotzki et al.~2002).  Despite its likely occurrence in nature, AIC has not yet been observationally identified, probably because it is less common and produces less $^{56}$Ni than a normal SN (e.g., Woosley $\\&$ Baron 1992).  Identifying AIC or constraining its rate would, however, provide unique insights into the paths of degenerate binary evolution, the true progenitors of Type Ia SNe (e.g., Yungelson $\\&$ Livio 1998), and the formation channels of globular cluster NSs (Grindlay 1987), low mass X-ray binaries (van den Heuvel 1984; Michel 1987), and millisecond pulsars (Grindlay $\\&$ Bailyn 1988; Bailyn $\\&$ Grindlay 1990).  Identifying the optical signature of AIC is also important because these events may be a source of strong gravitational wave emission (Fryer et al.~2002; see Ott 2008 and references therein).  In this Letter we show how AIC may itself produce a SN-like transient which, although less luminous and shorter-lived than a typical SN, may nonetheless be detectable with upcoming optical transient surveys.  Although little $^{56}$Ni is ejected by the newly-formed proto-NS itself, we show that powerful winds from the accretion disks formed during AIC produce up to a few times $10^{-2}M_{\\sun}$ in $^{56}$Ni.  \\vspace{-0.6 cm} \\subsection{Nickel-Rich Winds from AIC Disks} \\label{sec:bigpicture} \\vspace{-0.2 cm} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{cartoon.eps}} \\caption{Stages in the accretion-induced collapse of a WD.  (A) WD accretes up to near the Chandrasekhar mass ($M_{\\rm Ch}$) from either a binary companion or following the disruption of another WD in a merger.  Due to the accretion of angular momentum, the WD is rapidly rotating prior to collapse. (B) Electron captures cause the WD to collapse to a proto-NS surrounded by a compact accretion disk with a mass $\\sim 0.1 M_{\\sun}$ and a size $\\sim 100$ km.  Electron captures drive the composition neutron-rich (i.e., to an electron fraction $Y_{e} \\sim 0.1$).  As matter begins to accrete onto the proto-NS, the bulk of the disk spreads to larger radii. (C)  As the disk spreads, the midplane is irradiated by neutrinos from the central proto-NS.  Electron neutrinos are absorbed by neutrons, raising the electron fraction to $Y_{e} \\sim 0.5$.  (D) The disk's electron fraction freezes out.  Soon thereafter, powerful outflows blow the disk apart.  As the ejecta expands adiabatically, material with $Y_{e} \\gtrsim 0.5$ is efficiently synthesized into $^{56}$Ni, the subsequent decay of which powers a SN-like optical transient.} \\label{fig:cartoon} \\end{figure}  Figure $\\ref{fig:cartoon}$ summarizes the stages of AIC.  Due to their preceding accretion, WDs that undergo AIC are probably rotating rapidly prior to collapse (Fig.~$\\ref{fig:cartoon}$A).  As a result, a substantial fraction of the WD mass must be ejected into a disk during the collapse in order to conserve total angular momentum (Shapiro $\\&$ Lightman 1976; Michel 1987).  Dessart et al.~(2006, 2007; hereafter D06,07) perform 2D axisymmetric AIC calculations, in which they find that a quasi-Keplerian accretion disk with mass $M_{\\rm d,0} \\sim 0.1-0.5M_{\\sun}$ forms around the newly-formed proto-NS.  In Metzger, Piro, $\\&$ Quataert (2008a,b; hereafter MPQ08a,b) we studied the time-dependent evolution of accretion disks with similar initial properties that are formed from black hole (BH)-NS or NS-NS mergers.  Many of these results also apply to proto-NS accretion following AIC.  Although the disk is initially concentrated just outside the proto-NS surface, the bulk of the disk's mass must spread to larger radii as matter accretes in order to conserve angular momentum.  Early in their evolution, such ``hyper-accreting'' disks are neutrino-cooled and geometrically thin, and their composition is driven neutron-rich by an equilibrium between electron and positron captures under degenerate conditions (e.g., Pruet et al.~2003; Fig.~$\\ref{fig:cartoon}$B): the electron fraction $Y_{e} \\sim 0.1$, where  $Y_{e} \\equiv n_{p}/(n_{p}+n_{n})$ and $n_{p}/n_{n}$ are the proton/neutron number densities.  As the disk viscously spreads and its temperature decreases, however, weak interactions eventually become slow compared to the evolution timescale of the disk and the value of $Y_{e}$ freezes out.  Soon following freeze-out, neutrino cooling becomes inefficient and the disk becomes advective and geometrically thick (MPQ08a,b).  Since advective disks are only marginally bound, powerful winds likely begin to unbind most of the disk's remaining mass (Blandford $\\&$ Begelman 1999), aided by the nuclear energy released as $\\alpha-$particles begin to form (Lee $\\&$ Ramirez-Ruiz 2007; MPQ08a,b).  Because $\\sim 20\\%$ of the initial disk mass $M_{d,0}$ still remains when the disk becomes advective (or when $\\alpha-$particles form), a robust consequence of disk formation during AIC is the ejection of up to a few times $10^{-2}M_{\\sun}$.  Since the ejecta are hot and dense, heavy isotopes are synthesized as they expand away from the midplane and cool.  Because weak interactions are already slow by this point, which heavy isotopes are produced depends on $Y_{e}$ in the disk at freeze out.  In the case of BH accretion following BH-NS or NS-NS mergers, the disk freezes out neutron-rich with an electron fraction $Y_{e} \\sim 0.2-0.4$ (MPQ08a,b).  A crucial difference in the case of AIC, however, is the presence of the central proto-NS, which radiates a substantial flux of electron neutrinos as it deleptonizes during the first few seconds following core bounce (e.g., Burrows $\\&$ Lattimer 1986).  This $\\nu_{e}$ flux irradiates the disk, which acts to raise $Y_{e}$ via absorptions on free neutrons ($\\nu_{e} + n \\rightarrow e^{-} + p$; Fig.~\\ref{fig:cartoon}C).  In $\\S\\ref{sec:disk_evo}$ we present calculations of the evolution of the accretion disks formed from AIC using time-dependent height-integrated simulations with initial conditions and neutrino irradiation taken from the AIC simulations of D06.  We show that due to $\\nu_{e}$ absorptions, a substantial fraction of the disk freezes out with $Y_{e} \\gtrsim 0.5$.  As a result, late time outflows from the disk primarily synthesize $^{56}$Ni (Fig.~\\ref{fig:cartoon}D).  We discuss the observational signature of these Ni-rich outflows in $\\S\\ref{sec:miniSN}$, and we conclude in $\\S\\ref{sec:rates}$ by evaluating the prospects for detecting AIC with upcoming optical transient surveys and as electromagnetic counterparts to gravitational wave sources. \\vspace{-1 cm} ",
        "conclusions": ""
    },
    "0812/0812.3018_arXiv.txt": {
        "abstract": "{The presence of magnetic fields in the crust of neutron stars causes a non-spherically symmetric temperature distribution. The strong temperature dependence of the magnetic diffusivity and thermal conductivity, together with the heat generated by magnetic dissipation, couple the magnetic and thermal evolution of NSs, that cannot be formulated as separated one--dimensional problems. } {We study the mutual influence of thermal and magnetic evolution in a neutron star's crust in axial symmetry. Taking into account realistic microphysical inputs, we find the heat released by Joule effect consistent with the circulation of currents in the crust, and we incorporate its effects in 2--dimensional cooling calculations. } {We solve the induction equation numerically using a hybrid method (spectral in angles, but a finite--differences scheme in the radial direction), coupled to the thermal diffusion equation. To improve the boundary conditions, we also revisit the envelope stationary solutions updating the well known $T_{\\rm b}-T_{\\rm s}$--relations to include the effect of 2--D heat transfer calculations and new microphysical inputs. } {We present the first long term 2--dimensional simulations of the coupled magneto-thermal evolution of neutron stars. This substantially improves previous works in which a very crude approximation in at least one of the parts (thermal or magnetic diffusion) has been adopted. Our results show that the feedback between Joule heating and magnetic diffusion is strong, resulting in a faster dissipation of the stronger fields during the first $10^5-10^6$ years of a NS's life. As a consequence, all neutron stars born with fields larger than a critical value ($> 5 \\times 10^{13} G$) reach similar field strengths ($\\approx 2-3 \\times 10^{13} G$) at late times. Irrespectively of the initial magnetic field strength, after $10^6$ years the temperature becomes so low that the magnetic diffusion timescale becomes longer than the typical ages of radio--pulsars, thus resulting in apparently no dissipation of the field in old NS. We also confirm the strong correlation between the magnetic field and the surface temperature of relatively young NSs discussed in preliminary works. The effective temperature of models with strong internal toroidal components are systematically higher than those of models with purely poloidal fields, due to the additional energy reservoir stored in the toroidal field that is gradually released as the field dissipates. } {} ",
        "introduction": "The neutron star (NS) magnetic field (MF) maintained by electric currents circulating in the crust modifies the crustal temperature distribution by means of two mechanisms. The first one is due to the anisotropy of thermal conductivity in presence of a strong MF and causes important changes on how heat flux flows from the core through the crust and envelope up to the surface. The second mechanism is the generation of heat due to MF dissipation which results in a non spherically symmetric allocation of heating sources. Therefore, the magnetic diffusivity, besides its tensorial character due to the presence of the MF (which eventually results in the Hall drift of the field), becomes a non-spherically symmetric quantity. Under these conditions, an initially purely dipolar field evolves very quickly in a magnetized NS to generate complex structures including toroidal fields and/or higher order multipoles. Different studies on magneto-hydrostatic equilibrium configurations indicate that stable configurations require both toroidal and poloidal components (see e.g. \\cite{Reis2008} and references therein).   The complexity of the problem has limited previous works to partial studies of the complete problem. The anisotropic heat flux and its consequences for the - in principle observable - surface temperature ($T_{\\rm s}$) distribution has been considered by \\cite{Geppert2004,Geppert2006} and \\cite{Azorin2006a,Azorin2006b}, but for fixed, prescribed MF. The long--term effect of Joule heating has also been recently discussed by \\cite{Aguilera2008a} or \\cite{UK2008} (see also references therein), although in a very crude approach. In \\cite{Aguilera2008b} a full 2-D cooling code has been used to describe the temperature evolution, but the Joule heating rate is estimated by an analytical approximation and considered to be uniform. In \\cite{UK2008} the former results of \\cite{Miralles1998} are revisited in the context of high field radio--pulsars, but with the simplified approach of considering the diffusion of a one--mode (dipolar) poloidal field, without actually performing consistent simulations including the temperature evolution. On the other hand, the evolution of the MF including the influence of the Hall term at early times has been studied in \\cite{PonsGeppert2007}. The Hall drift causes a somewhat faster dissipation for strong initial fields due to the reorganization of the field in smaller scales. In this latter work the temporal evolution of the temperature has been prescribed according to a generally accepted cooling law and assumed to be the same in the whole crust. The effect of ambipolar diffusion in the liquid core can also be relevant, as recently studied for example in \\cite{Hoyos2008}.  One of the reasons that leads to complex MF geometries in the crust is due to the non-spherically symmetric magnetic diffusivity. It is effective even if the Hall drift is negligible. Since the temperature within the crust is no longer uniform, the magnetic diffusivity in this approach becomes a function of the radial and of the polar coordinate, thus rendering the one--mode approximation inappropriate. Up to now, the effect of the dependence of the magnetic diffusivity on the polar angle $\\theta$ (if axial symmetry is assumed) has not been considered. In some studies about cooling of NSs, the source term in the thermal diffusion equation includes an angular--averaged Joule heating rate, where the heat production is uniform in spherical shells at a given radius $r$ \\citep{Page2000,Aguilera2008a,UK2008}. These are not fully consistent calculations, because the currents can locally be very intense and release heat in a highly non-spherically symmetric way. Therefore, it is important to attack the problem of the full magneto-thermal evolution of NSs, by solving simultaneously the induction equation and the heat transfer equation taking into account the anisotropy of the thermal conductivity and the electrical resistivity tensors in a consistent way. This is the main goal of this paper, in which we solve this problem (in axial symmetry) for the first time. The aim of this study is to provide a self-consistent model for the coupled evolution of the MF and temperature in axially symmetric NSs, in general relativity, and including the effect of angular variations of temperature in the magnetic diffusivity and the thermal conductivity. We also revisit the envelope stationary solutions updating the well known $T_{\\rm b}-T_{\\rm s}$--relations to include the effect of 2--D heat transfer calculations and new advances in microphysical ingredients (phonons, ion--ion interactions). We provide new fits to solutions of equilibrium magnetic envelope that can be used as boundary conditions in multidimensional cooling simulations.  The paper is organized as follows. In the next Section we derive and present the basic equations together with the boundary conditions. Next we describe the input microphysics and discuss the effect of superfluid neutrons on the heat flux anisotropy. The following Section is devoted to the presentation of the magneto--thermal evolution. Finally we discuss the results by comparing them with observable consequences.\\\\ ",
        "conclusions": "We have performed consistent 2D simulations of the coupled magneto-thermal evolution of NSs for the first time, including realistic microphysical input and general relativistic corrections. By properly taking into account the interplay of the MF and temperature evolution in cooling simulations we have found that their mutual influence is important and affects the outcome when the initial field strength is of the order or larger than a critical value of $B_{\\rm crit} = 5 \\times 10^{13}$ G. It has effects on both sides: the temperature and the field strength.  The average effective temperature of a NS born with $B_p <  B_{\\rm crit}$ is barely affected, while those born as magnetars are subject to significant heating by the dissipation of currents in the crust. In addition, since heating is locally important in the regions where currents are more intense, the surface temperature distribution can be very different depending on the field geometry. While the pole (or other regions in which the field lines are nearly radial through the crust) is in thermal equilibrium with the core and has its same temperature, regions in which the field lines are tangential remain essentially thermally isolated. This can produce cooler areas if no heating process is considered or, conversely, hotter regions if heating is important. This is the case of poloidal fields, in which currents are located near the equatorial region that remains warmer than the rest of the crust during a long time. The effective temperature of models with strong internal toroidal components are systematically higher than that of models with purely poloidal fields, due to the additional energy reservoir stored in the toroidal field that is gradually released as the field dissipates.  In the models with stronger fields, as a result of the average higher temperatures, the crustal electrical resistivity is enhanced and magnetic diffusion proceeds faster during the first $10^5-10^6$ years of a NS's life. As a consequence, all NSs born with fields larger than a critical value ($> 5 \\times 10^{13} G$) reach similar field strengths ($\\approx 2-3 \\times 10^{13} G$) at late times, irrespectively of the initial strength. After $10^6$ years the temperature is so low that the magnetic diffusion timescale becomes longer than the typical ages of radio--pulsars, resulting in apparently no dissipation of the field in old NSs. We confirm the strong correlation between the MF and the surface temperature of relatively young NSs discussed in preliminary works. Notice that if the MF of magnetars is caused by superconducting currents in the core (as opposed to crustal currents) the longer diffusion timescale in the core would allow magnetars to live much longer with their original large fields. Thus, observations of magnetars can help to discern between models with currents located in the crust or in the core.  It should be mentioned that magnetic field evolution by ambipolar diffusion in the core may produce qualitatively similar effects to those obtained in this work: keeping the temperature high while the field is strong and stopping field decay when the temperature drops \\citep[see e.g.][]{Reis2008}. In this work we have not considered the MF evolution in the core because ambipolar diffusion is usually considered under the assumption that the core is in a non--superfluid state and therefore is important during the very early stages of evolution. We are more interested in the long-term evolution, after the temperature rapidly drops below the critical temperature for nucleon superfluidity. In order to quantify the relative importance of both effects one would need to consider ambipolar diffusion in a superconducting fluid coupled to the dissipation of crustal currents studied here.  The detection of magnetars with {\\it true} ages (the spin-down age can be seriously overestimated) $t \\approx 10^6$ yr, or the detection of a young highly magnetized NS with $T<10^6$ K would be a serious challenge for crustal field models. At present our results are in agreement with the known population of high field NSs and magnetars and support the idea of the existence of a strong crustal MF component in magnetars.  Given the complexity of the feed-back between temperature and MF, it seems necessary to extend this work in two main lines that can shed new light on our knowledge of the cooling theory of NS: engaging 3D simulations and including the Hall term in the induction equation. The complex geometry that may arise in a realistic case with hot spots and irregular fields is certainly not treatable with our present code and needs further investigation. Similarly, the influence of the Hall term at late times (when the temperature is so low that the magnetic diffusivity is negligible) or a consistent treatment of ambipolar diffusion in a superconducting fluid coupled to the dissipation of crustal currents, can produce new interesting effects and are issues worth to explore in future works."
    },
    "0812/0812.2472_arXiv.txt": {
        "abstract": "We continue our presentation of an alternative cosmology based on conformal gravity, following our kinematical approach introduced in a recent paper. In line with the assumptions of our model, which proposes a closed-form expression for the cosmic scale factor $R(t)$, we first revise the Hubble and deceleration parameters and also introduce modified cosmological distances, analyzing in particular the case of the luminosity distance.  Our kinematical conformal cosmology is then able to explain the anomalous acceleration of the Pioneer spacecraft, as due to a local region of gravitational blueshift. From the reported values of the Pioneer anomaly we also compute the current value of our first fundamental parameter, $\\gamma _{0}=1.94\\times10^{-28}\\ cm^{-1}$, in line with the original estimate by P. Mannheim of this quantity.  Our second fundamental parameter, $\\delta_{0}=3.83\\times10^{-5}$, interpreted as the current value of a cosmological time variable, is derived from a detailed fitting of type Ia Supernovae \\textquotedblleft gold-silver\\textquotedblright\\ data, producing Hubble plots of the same quality of those obtained by standard cosmology, but without requiring any dark matter or dark energy contribution.  If further experiments will confirm the presence of an anomalous frequency blueshift in the outer region of the Solar System, as described by our model, kinematical conformal cosmology might become a viable alternative to standard cosmological theories. ",
        "introduction": "Introduction}  This paper is the second part of a project aimed at introducing an alternative cosmology based on Conformal Gravity (CG), as originally proposed by H. Weyl (\\cite{Weyl:1918aa}, \\cite{Weyl:1918ib}, \\cite{Weyl:1919fi}) and recently revisited by P. Mannheim and D. Kazanas (\\cite{Mannheim:1988dj}, \\cite{Kazanas:1988qa}, \\cite{Mannheim:2005bf}). In the first paper on the subject \\cite{Varieschi:2008fc} (paper I, in the following) we presented the mathematical foundations of our new kinematical approach to Conformal Cosmology. This was based on a critical re-analysis of fundamental astrophysical observations, starting with the cosmological redshift, and on the fact that modern metrology defines our common units of length and time using non-gravitational physics, i.e., through emission/propagation/absorption of electromagnetic waves or similar phenomena.  Since the laws of electromagnetism are notoriously invariant under a conformal transformation, we argued that on a cosmological scale a conformal \\textquotedblleft stretching\\textquotedblright\\ of the space-time might be present and might yield to an effective change in wavelength/frequency of electromagnetic radiation, equivalent to the observed cosmological redshift (or blueshift, if any). As the origin of this presumed conformal stretching of the metric in the Universe can only be gravitational, we searched its connection with existing theories of gravity that allow this possible conformal symmetry. Our attention was focused on Weyl's Conformal Gravity, since this is the simplest known conformal generalization of Einstein's General Relativity (GR). Weyl's theory is also based on the same principles and assumptions of GR, such as the equivalence principle and other foundational concepts.  The complexity of CG, in particular its fourth order field equations, as opposed to Einstein's second order equations, have rendered this theory quite intractable until Mannheim-Kazanas (MK) found the first complete solutions, such as the exterior solution for a static, spherically symmetric source (\\cite{Mannheim:1988dj}, \\cite{Kazanas:1988qa}) and they also showed that it reduces to the classic Schwartzschild solution in the limit of no conformal stretching. In addition, the MK solution is able to interpolate smoothly between the classic Schwartzschild solution and the Robertson-Walker (RW) metric, through a series of coordinate transformations based on the conformal structure of the theory.  It was precisely this ability to transform from the conformal MK\\ solution to the standard RW metric, used to describe the cosmological expansion, which convinced us that a universal conformal stretching might be able to mimic the expansion of the Universe. The gravitational origin of this conformal stretching should also lead to the change of observed wavelength/frequency of cosmic radiation. The principles of General Relativity, which still apply to its conformal extension, naturally propose such a mechanism: the gravitational redshift.  We have shown in our first paper how the original MK potential can support this explanation and how the chain of transformations, from Static Standard Coordinates used in the MK solution to the Robertson Walker coordinates, can lead to a unique expression of the cosmic scale factor $R$. In this way, the conformal symmetry of the Universe is \\textquotedblleft kinematically\\textquotedblright\\ broken and the precise amount of stretching at each space-time point can be determined, once certain parameters of the original MK potential are measured.  In this second paper we will show how we can determine these fundamental parameters using astrophysical data, such as the luminosities of type Ia Supernovae (SNe Ia) and others. In this way our kinematical conformal cosmology might become a viable alternative model for the description of the Universe, with the advantage of avoiding most of the controversial features of the standard model, such as dark energy, dark matter, inflationary phases, etc.  In the next section we will review our cosmological solutions from paper I, then we will obtain expressions for the Hubble constant and deceleration parameter and also revise the definitions of the standard cosmological distances. In Sect. \\ref{sect:cosmological_parameters} we will fit current astrophysical data in order to compute our cosmological parameters and check the consistency of our model. Finally, in Sect. \\ref{sect:consequences}, we will explore the immediate consequences of our model, in terms of the behavior of fundamental constants and other physical quantities. ",
        "conclusions": "Conclusions}  We have introduced experimental evidence in support of our kinematical conformal cosmology and determined the values of its fundamental parameters. In particular, we have focused our analysis on reproducing the Hubble plots for type Ia Supernovae, with the same level of accuracy obtained by standard cosmology.  To achieve this goal we critically reconsidered all the standard distances commonly used in cosmology and we added a new scaling property for the masses (or the energies) as a function of the redshift parameter. Our new expression for the distance modulus as a function of redshift can effectively fit the \\textquotedblleft gold-silver\\textquotedblright\\ SNe data with the required accuracy and also yield a current-time value for our fundamental parameter $\\delta_{0}$, which is small and positive as expected.  Since type Ia Supernovae, or other astrophysical candles, are distant cosmological objects, our second point of focus was to consider more local effects due to our kinematical conformal cosmology, which might be more suitable for the determination of the parameters. In particular, a local blueshift region was expected, given the estimated values of the quantities in our model, and is possibly evidenced by the recently discovered Pioneer anomaly.  We have seen how our model can account for this effect and can be used to estimate our second fundamental parameter $\\gamma_{0}$, which together with $\\delta_{0}$ will determine all the other quantities in our model. More precise evaluations of these parameters are certainly needed and should come from an extended analysis of the Pioneer data or through a dedicated future spacecraft mission, which has already been proposed.  We argued that a direct detection of a frequency/wavelength shift in electromagnetic radiation, traveling over distances comparable to the size of our Solar System, could be explained only in terms of a conformal space-time stretching and would be the best evidence in support of our kinematical approach. If the Pioneer anomaly, or similar phenomena, will prove in the next few years to be positive indications of these effects, this might also signal a possible time-variation of dimensionless fundamental quantities, such as the fine-structure constant and our cosmological time $\\delta$. This would constitute an important step towards a deeper understanding of the role and values of all the fundamental physical constants and also impact our current understanding of the fundamental interactions."
    },
    "0812/0812.0477_arXiv.txt": {
        "abstract": "{ Umbral Dots (UDs) are thought to be manifestations of magnetoconvection in sunspot umbrae. Recent advances in their theoretical description point to the need for a thorough study of their properties and evolution based on data with the highest currently achievable resolution. } { Our UD analysis aims to provide parameters such as lifetimes, diameters, horizontal velocities, and peak intensities, as well as the evolution of selected parameters. } { We present a 106-minute TiO (705.7~nm) time series of high spatial and temporal resolution that contains thousands of UDs in the umbra of a mature sunspot in the active region NOAA~10667 at $\\mu$~=~0.95. The data were acquired with the 1-m Swedish Solar Telescope (SST) on La Palma. With the help of a multilevel tracking (MLT) algorithm the sizes, brightnesses, and trajectories of 12836 umbral dots were found and extensively analyzed. The MLT allows UDs with very low contrast to be reliably identified. } { Inside the umbra we determine a UD filling factor of 11\\,\\%. The histogram of UD lifetimes is monotonic, i.e. a UD does not have a typical lifetime. Three quarters of the UDs lived for less than 150~s and showed no or little motion. The histogram of the UD diameters exhibits a maximum at 225~km, i.e. most of the UDs are spatially resolved. UDs display a typical horizontal velocity of 420~m\\,s$^{-1}$ and a typical peak intensity of 51\\,\\% of the mean intensity of the quiet photosphere, making them on average 20\\,\\% brighter than the local umbral background. Almost all mobile UDs (large birth-death distance) were born close to the umbra-penumbra boundary, move towards the umbral center, and are brighter than average. Notably bright and mobile UDs were also observed along a prominent UD chain, both ends of which are located at the umbra-penumbra boundary. Their motion started primarily at either of the ends of the chain, continued along the chain, and ended near the chain's center. We observed the splitting and merging of UDs and the temporal succession of both. For the first time the evolution of brightness, size, and horizontal speed of a typical UD could be determined in a statistically significant way. Considerable differences between the evolution of central and peripheral UDs are found, which point to a difference in origin. } {} ",
        "introduction": "The investigation of the complex fine structure of umbrae and penumbrae is crucial to understanding the subsurface energy transport in sunspots. The energy transport from the solar interior to the solar surface outside magnetic features is mainly determined by convection, visible as granulation in images of the quiet photosphere. The strong and nearly vertical umbral magnetic field suppresses normal overturning convection inside the umbra. However, it is believed that some form of residual magnetoconvection is responsible for much of the remaining energy transport and manifests itself in the form of fine structures, such as light bridges (LBs) or umbral dots (UDs). In the present paper we consider UDs, which contribute up to 37\\,\\% of the radiative umbral flux according to \\citet{Adjabshirzadeh1983}. Different models have been proposed to explain the umbral dots. \\citet{Choudhury1986} postulated that UDs are thin columns of field-free hot gas between the cluster of small magnetic flux tubes that form the subsurface structure of a sunspot according to \\citet{Parker1979}. According to this model, a UD is formed when an upwelling brings hot material into the photosphere. An alternative model has been proposed by \\citet{Weiss1990} who consider UDs to be spatially modulated oscillations in a strong magnetic field.  A more recent, promising approach is presented by \\citet{Schuessler2006}, who used numerical simulations of three-dimensional radiative magnetoconvection to improve the physical understanding of the umbral fine structure. The simulations exhibit the emergence of small-scale upflow plumes that start off like oscillatory convection columns below the solar surface but turn into narrow overturning cells driven by the strong radiative cooling around optical depth unity. Most of those UDs show a central dark lane. The presence of dark lanes in penumbral and umbral fine structures has already been observed several times, cf. \\citet{Scharmer2002,Langhans2007,Scharmer2007}. The verification of the predicted dark lanes in large UDs by \\citet{Bharti2007} and \\citet{Rimmele2008}, as well as the verification of the predicted photospheric stratification of bright peripheral UDs by \\citet{Riethmueller2008}, support the Sch\\\"ussler \\& V\\\"ogler model of UDs. There is now a need to learn more about this phenomenon, with a statistically robust analysis of UD properties and evolution being a promising means of achieving this aim.  The most detailed analyses of UDs are more than 10 years old \\citep{Sobotka1997a,Sobotka1997b} and are based on data observed with the 50-cm SVST (Swedish Vacuum Solar Telescope), cf. the recent reviews of umbral fine structures by \\citet{Solanki2003,Thomas2004} and \\citet{Sobotka2006}. The more recent papers of \\citet{Tritschler2002,Hartkorn2003} and \\citet{Sobotka2005} have concentrated on individual properties and lack, e.g., the determination of UD trajectories. Furthermore, the possibility of improving the spatial resolution with the help of modern image reconstruction algorithms is only used by \\citet{Tritschler2002}. The present paper aims to overcome these shortcomings, by employing data from the 1-m SST (Swedish Solar Telescope) equipped with an adaptive optics system, by restoring the data employing MFBD (multi-frame blind deconvolution), and determining the evolution of UD parameters whenever possible. ",
        "conclusions": "We have analyzed a time series of images of a mature sunspot close to solar disk center. Due to the excellent image quality we were able to resolve thousands of UDs. Exhaustive UD analyses can be found in earlier papers \\citep{Sobotka1997a,Sobotka1997b,Hartkorn2003,Sobotka2005,Sobotka2006}, but the present article is the first detailed UD study of a long time series of reconstructed images with the consistent high resolution of a 1-meter telescope.  Trajectories, lifetimes, diameters, horizontal velocities, peak intensities, and distances between birth and death locations were determined by tracking single UDs over the time series. These characteristic values were used to look for reasonable separations into UD classes. In the following we summarize the obtained results and compare them with other investigations in the literature:  \\begin{enumerate} \\item There is hardly any part of the umbra which does not support UDs, but the UD brightnesses depend strongly on the location within the umbra, which confirms the previous observation of \\citet{Sobotka1997b}. \\item The histogram of lifetimes shows an exponential distribution, i.e. a UD does not have a typical lifetime. More than 3/4 of all studied UDs lived less than 150~s and their motion was negligible. The exponential distribution is in qualitative agreement with the results obtained by \\citet{Sobotka1997a,Sobotka1999}. Quantitatively, \\citet{Sobotka1997a} obtains a median lifetime of 6~min for an umbra of about 6~Mm diameter and a median of 12~min for a 4~Mm pore \\citep{Sobotka1999}, whereas we find a median value of 0.7~min for a roughly 10~Mm umbra. Note that these median values depend strongly on algorithmic constraints as well as on the cadence of the time series. For example, the method used by \\citet{Sobotka1997a} cannot lead to lifetimes shorter than 1.5~min. If we only consider UDs with lifetimes greater than 1.5~min our median increases to 4.1~min. Irrespectively of which of these two values we use, our results are consistent with the conclusion of \\citet{Sobotka1999} that UDs are more stable in a weak magnetic field if we assume a direct correlation between umbral diameter and magnetic field strength \\citep[cf.][]{Kopp1992}. Alternatively, due to the strong dependence of umbral brightness on umbral size \\citep{Mathew2007} the lifetime may be influenced mainly by the radiative flux or umbral temperature. The scatterplot of the mean umbral background intensities versus the UD lifetimes (not shown) is consistent with both possible explanations mentioned above: UDs live longer in brighter parts of the umbra. Due to the non-linear, monotonically decreasing relation between magnetic field strength and background intensity as observed by \\citet{Kopp1992} and confirmed by \\citet{Martinez1993,Solanki1993}, this implies that UDs live longer in regions of weak field. \\item The histogram of mean diameters exhibits a maximum at 225~km (0.31$^{\\prime\\prime}$) and descends from there towards the diffraction limit, so that we expect the majority of UDs to have been spatially resolved. This seems not to be the case in many of the earlier papers because there a monotonic decrease was obtained towards higher diameters \\citep[see][]{Sobotka1997a,Sobotka1999}. \\citet{Sobotka2005} also analyzed data obtained with the 1-meter SST. These data lead to a histogram that is qualitatively similar to ours. The mean diameter of 175~km (0.24$^{\\prime\\prime}$), as well as the average filling factor of 9\\,\\% is, however, noticeable smaller. This difference can be explained by the use of a different method to determine the UD boundary. An increase of our brightness threshold to determine the UD boundary leads to smaller diameters and to lower filling factors. \\citet{Hamedivafa2008} used an improved method of image segmentation and also found a mean diameter of 230~km and a similar shaped histogram. \\item The mean horizontal velocity of those of our UDs that live longer than 150~s is 420~m\\,s$^{-1}$ which is significantly higher than the 210~m\\,s$^{-1}$ reported by \\citet{Molowny1994} and higher than the 320~m\\,s$^{-1}$ found by \\citet{Sobotka1999}. In both studies the mean horizontal velocity was calculated by means of least-squares linear fits of the x and y coordinates of all trajectory points, which leads to an underestimation of velocity in case of curved trajectories. Our histogram of horizontal velocities shows a maximum at 350~m\\,s$^{-1}$, whereas some UDs can reach velocities above 1~km\\,s$^{-1}$. This is in qualitative agreement with the histograms of \\citet{Kitai1986}, \\citet{Molowny1994}, and \\citet{Hamedivafa2008} but disagrees with the results of \\citet{Sobotka1997b,Sobotka1999} whose histograms do not exhibit a maximum; they peak at zero velocity and show a monotonic decrease toward higher velocities of up to 1~km\\,s$^{-1}$. The majority of our UDs moves irregularly around the birth position. However, there are some mobile UDs that travel over long distances within their lifetime. Almost all mobile UDs emerge close to the umbral border, i.e. near the penumbra, they are brighter than the average and their horizontal motion is preferentially directed towards the center of the umbra. The mean velocity of our mobile UDs is 680~m\\,s$^{-1}$, which is in good agreement with the recent observation of \\citet{Katsukawa2007} who found a mean velocity of 700~m\\,s$^{-1}$. \\item The relation between mean UD size and lifetime is non-linear. On average, the size of UDs increases with lifetime, which was also found by \\citet{Sobotka1997a}, but in contrast to their work we find a narrow size distribution of around 290~km for long-lived UDs. \\item UDs that were born close to the penumbra show a significantly higher contrast than the UDs of the umbral interior. The mean UD intensity contrast $I_{Peak}/I_{bg}$ is 1.2 which is smaller than the value of 1.6 reported by \\citet{Sobotka2005}. This may partly be due to the longer wavelength of our observation. Additionally, our statistical ensemble contains many more UDs. In particular we took many UDs with low contrast into account, made possible by the multilevel tracking technique \\citep{Bovelet2001} we employed to identify UDs. Consequently, we believe that the lower UD contrast we find is not due to a lower resolution or higher stray light, but rather to differences in identification of UDs and in particular the difference in wavelength. We stress that the UD contrast itself depends on the background intensity; the higher the intensity the stronger the contrast. Also it cannot be ruled out that there could be systematically different contrasts between different sunspots due to intrinsicly different physical properties of the spots. \\item Whereas the temporal variation of the UD diameter is qualitatively similar for UDs formed close to the penumbra (PUDs) and those formed in the body of the umbra (CUDs), their intensity contrast and horizontal velocity display contrasting evolutions. The mean PUD shows a continuous darkening which is in agreement with the results of \\citet{Kitai2007} for a single typical PUD. The typical CUD of \\citet{Kitai2007} is found to increase in brightness linearly and then to darken linearly with time, whereas our results for the mean CUD show a non-linear increase in brightness until nearby half of the lifetime followed by an again non-linear decrease. The clear difference between PUDs and CUDs in the behavior of their contrast and mean velocity may be a result of the different origin of the two types of features \\citep{Kitai2007}. We confirm from visual inspection of a subset of PUDs that PUDs are formed when penumbral grains cross the umbral boundary. \\end{enumerate}  A comparison of the results with the simulations of \\citet{Schuessler2006} shows a better agreement with CUDs, than PUDs. For example, the simulated UDs display a gradual increase in contrast followed by a gradual decrease, just as CUDs. They also display little proper motion. This qualitative agreement further strengthens the interpretation of UDs as localized columns of overturning convection proposed by \\citet{Schuessler2006}. Hinode data had earlier suggested the presence of dark lanes in large UDs \\citep{Bharti2007} and revealed a decrease in the magnetic field strength with depth, as well as an upflow associated with a temperature enhancement \\citep{Riethmueller2008}, in good qualitative agreement with the simulations. A detailed analysis of the simulations similar to the one carried out here would allow a more quantitative comparison.  The PUDs have significantly different evolution histories than the simulated features. They start at a higher speed \\citep{Kitai1986} and in particular display their maximum brightness right after the beginning of their life \\citep[cf.][]{Kitai2007}. This, combined with the fact that they are born very close to the penumbra, or actually by breaking away from the penumbra \\citep{Thomas2004}, and move radially towards the umbral center \\citep{Kitai1986} supports that these are two distinct types of UDs based on their origin and evolution, although their physical structure is relatively similar \\citep{Riethmueller2008}."
    },
    "0812/0812.4368_arXiv.txt": {
        "abstract": "We present some results of an observational and theoretical study on unresolved stellar systems based on the Surface Brightness Fluctuations (SBF) technique. It is shown that SBF magnitudes are a valuable tracer of stellar population properties, and a reliable distance indicator. SBF magnitudes, SBF-colors, and SBF-gradients can help to constrain within relatively narrow limits the metallicity and age of the dominant stellar component in distant stellar systems, especially if coupled with other spectro-photometric indicators. ",
        "introduction": "The detailed study of the properties of distant stellar systems relies on the details to which the stellar population can be observed. The observation of single stars is, in principle, the way to extract the best information from a system. However, even with modern telescopes, single stars can be observed only in nearby galaxies, and, typically, at magnitudes significantly brighter than the main sequence. Here, we discuss a technique that has proved to be very powerful to disentangle the physical and chemical properties of distant, unresolved stellar population, the Surface Brightness Fluctuations (SBF) method. Here, we summarize the main characteristics of the technique - for a more detailed description of the SBF method see G. Raimondo, and J. P. Blakeslee's contributes to this Volume.  First introduced as a distance indicator for nearby elliptical galaxies \\citep{ts88}, the SBF method has been applied to Galactic globular clusters at a few kpc, cluster ellipticals out to $\\sim\\,$150 Mpc, and numerous lenticulars, spiral bulges, and dwarf spheroidals at intermediate distances \\citep[e.g.][]{tonry01,mieske06a,biscardi08}. By definition the SBF signal corresponds to the ratio of the 2$^{nd}$ to the 1$^{st}$ moment of the Luminosity Function (LF) of the stellar population. Therefore, the dominant contribution to SBF comes from bright stars which, acting as ``lighthouses'' on a smooth background of faint stars, generate the great part of the fluctuation signal. Further, the relative stellar luminosities depend on the observed bandpass, so the $lighthouses$ at shorter wavelengths (hot stars) are different from those at longer wavelength (cool stars), and the main contributors to the SBF change according to the filter. This makes SBF magnitudes and colors potential candidate to trace the properties of stellar population in unresolved systems.  Based on such evidences we have carried out an extensive study of SBF, and correlated photometric indicators, both on the theoretical and observational point of view. The main results of these studies and future perspectives are briefly outlined below. ",
        "conclusions": ""
    },
    "0812/0812.3853_arXiv.txt": {
        "abstract": "The existence of a massive black hole in the center of the Milky Way, coinciding with the radio source Sgr A*, is being established on more and more solid ground. In principle, this black hole, acting as a gravitational lens, is able to bend the light emitted by stars moving within its neighborhood, eventually generating secondary images. Extending a previous analysis of the gravitational lensing phenomenology to a new set of 28 stars, whose orbits have been well determined by recent observations, we have calculated all the properties of their secondary images, including time and magnitude of their luminosity peaks and their angular distances from the central black hole. The best lensing candidate is represented by the star S6, since the magnitude of its secondary image at the peak reaches $K=20.8$, with an angular separation of 0.3 mas from the central black hole, that is just at the borders of the resolution limit in the $K$ band of incoming astronomical instruments. ",
        "introduction": "One of the most fascinating arenas for testing General Relativity in the neighborhood of a real astrophysical black hole is undoubtedly represented by Sgr A*, the massive black hole (MBH) at the center of our Galaxy, whose existence has been now strongly proven by the combination of radio, sub-mm, X-ray, and Near Infrared (NIR) observations. While in the radio domain Sgr A* appears as a steady and motionless compact source \\citep{Reid07}, its X-ray, sub-mm, and NIR counterparts are variable since they exhibit outbursts of energy ({\\it flares}), which typically last for a few hours and occur several times a day \\citep{Baganoff01,Baganoff03,Genzel03, Ghez04,Clenet05,Eckart06,Eckart08,Hamaus08,Marrone08,Yusef08}. These flares are likely due to energetic events arising very close to the central MBH, on a scale of a few Schwarzschild radii, and they can be interpreted as due to emission from matter in relativistic orbits around Sgr A*. Moreover, by using large class telescopes, such as the ESO Very large Telescope (VLT) and the Keck telescopes, and thanks to the adaptive optics, it has been possible to observe a multitude of stars in the $K$ band (the so-called S-stars) moving in the gravitational potential of the Galactic MBH \\citep{Schodel03,Ghez05,Eisenhauer05}. The entire system is perfectly described by a single point mass and Newton gravity. Very recently, \\citet{Gillessen08} have been able to determine the orbits of 28 S-stars, resting upon the results of 16 years of monitoring of their motions. In particular, it has been possible to observe the whole orbit of the famous star S2, since it has completed a full revolution around Sgr A*. The value of the Galactic MBH mass, extrapolated from these improved data, is $M=4.31\\pm0.36 \\times 10^{6}$ M$_{\\sun}$.  Since the theory of General Relativity predicts that the central MBH can act as a powerful gravitational lens on S-stars, some of them, depending on the knowledge of their orbital parameters, have been the object of careful investigations. The strategy of this study is very clear. In fact, knowing the position of the S-stars in space as a function of time, it is easy to predict the time of their periapse, that is the closest approach to Sgr A*, and look for possible alignments between one of them and the lens, corresponding to the MBH, along the line of sight of an observer on the Earth. In this way, it is possible to foresee at any time where to look for a secondary image, what its apparent magnitude should be, and, thanks to our knowledge of their orbits, easily predict the best time to observe them.  S2 was the first star to be studied as a possible source for gravitational effects due to the central MBH \\citep{DePaolis03,Bozza04}. The light curve for its secondary image, as well as the first two relativistic images, were calculated in the Schwarzschild black hole hypothesis. The analysis was then enlarged by \\citet{Bozza05}, who achieved the light curves for the secondary images of five more S-stars. An important fact that emerges from the latter study is that the light curves are peaked around the periapse epoch, but two subpeaks may arise in nearly edge-on orbits, when the source is behind Sgr A* ({\\it standard lensing}) or in front of it ({\\it retrolensing}). In this respect, the star S14 represents an emblematic case.  In the present paper we want to extend the previous analysis of \\citet{Bozza05} to the new set of 28 S-stars, provided by \\citet{Gillessen08}. We will see that, within these stars, S6 turns out to be a particularly intriguing source for gravitational lensing, since its secondary image is potentially observable by incoming interferometric instruments. ",
        "conclusions": "Thanks to the great technological advances earned by the modern big-class telescopes, it has been possible to have access to very remote zones of the Milky Way, like the neighborhood of its center, where the radio source Sgr A* lies. The discovery that most of the detected stars in this region move along tight Keplerian orbits is a strong indication of the existence of a very compact object of more than $4 \\times 10^{6}$ M$_{\\sun}$, currently recognized as a massive black hole, in the geometrical center of the Milky Way. This great disclosure opens an amazing way of escape from the meanders of theoretical speculations and allows us to study gravitational lensing by a black hole on a real case and test General Relativity in its strong field regime. The idea of using the star S2 as a potential source to be lensed by the central MBH, was conceived originally by \\citet{DePaolis03}, and then acquired and rigorously developed by \\citet{Bozza04,Bozza05}, who considered other S-stars as further possible sources for this phenomenology. In this paper, we extend the previous analysis to a new set of 28 S-stars provided by \\citet{Gillessen08}. Whereas most of these stars have undetectable secondary images, there are some cases in which the lensed images have a significant probability of being detected by incoming instruments, as can be deduced from table~\\ref{tbl-2}. Among these, S6 represents the best case since its secondary image at the peak reaches $K=20.8$, with an angular separation of 0.3 mas from the central black hole. Actually, the perspectives for observing such an event are not so far from what the modern technology is developing. In fact, in 2013 the new second-generation VLTI instrument, called GRAVITY, will be scientifically operative at the ESO Paranal observational site \\citep{Eisenhauer08}. This instrument, specifically designed to observe highly relativistic motions of matter close to the event horizon of Sgr A*, will interferometrically combine the NIR light collected by the four-units telescope of the VLT, exploiting all the potential of this extremely powerful ESO observatory. The resolution that will be obtained in the NIR band is $\\sim 3$ mas for objects that can be as faint as $17<K<19$. Moreover, in its astrometric mode, GRAVITY will have an accuracy of $10$ $\\mathrm{\\mu as}$, allowing to study the motions of celestial objects to within a few times the event horizon size of the Galactic MBH. There are still many years before the secondary images of the best known S-stars will shine again (2047 for S14, 2062 for S6 and S27). Surely, new S-stars will be discovered and followed along their orbits, eventually providing even better candidates for gravitational lensing. In the meanwhile, we have no doubt that the astronomical community will fortify its stock of observational facilities with new advanced instruments, which will be able to catch this extraordinary kind of gravitational lensing events.  \\begin{table*} \\centering \\begin{scriptsize} \\begin{tabular}{lcccccccc} \\tableline % \\tableline % Star & $a['']$ & $e$ & $i[^\\circ]$ & $\\Omega[^\\circ]$ & $\\omega[^\\circ]$ & \\ $t_{\\mathrm{P}}[\\mathrm{yr}]$ & $T[\\mathrm{yr}]$ & $K$ \\\\ \\tableline % S1 & 0.508 $\\pm$ 0.028 & 0.496 $\\pm$ 0.028 & 120.82 $\\pm$ 0.46 & 341.61 $\\pm$ \\ 0.51 & 115.3 $\\pm$ 2.5 & 2000.95 $\\pm$ 0.27 & 132. $\\pm$ 11. & 14.7 \\\\ S2 & 0.123 $\\pm$ 0.001 & 0.88 $\\pm$ 0.003 & 135.25 $\\pm$ 0.47 & 225.39 $\\pm$ \\ 0.84 & 63.56 $\\pm$ 0.84 & 2002.32 $\\pm$ 0.01 & 15.8 $\\pm$ 0.11 & 14. \\\\ S4 & 0.298 $\\pm$ 0.019 & 0.406 $\\pm$ 0.022 & 77.83 $\\pm$ 0.32 & 258.11 $\\pm$ \\ 0.3 & 316.4 $\\pm$ 2.9 & 1974.4 $\\pm$ 1. & 59.5 $\\pm$ 2.6 & 14.4 \\\\ S5 & 0.25 $\\pm$ 0.042 & 0.842 $\\pm$ 0.017 & 143.7 $\\pm$ 4.7 & 109. $\\pm$ 10. \\ & 236.3 $\\pm$ 8.2 & 1983.6 $\\pm$ 2.5 & 45.7 $\\pm$ 6.9 & 15.2 \\\\ S6 & 0.436 $\\pm$ 0.153 & 0.886 $\\pm$ 0.026 & 86.44 $\\pm$ 0.59 & 83.46 $\\pm$ \\ 0.69 & 129.5 $\\pm$ 3.1 & 2063. $\\pm$ 21. & 105. $\\pm$ 34. & 15.4 \\\\ S8 & 0.411 $\\pm$ 0.004 & 0.824 $\\pm$ 0.014 & 74.01 $\\pm$ 0.73 & 315.9 $\\pm$ \\ 0.5 & 345.2 $\\pm$ 1.1 & 1983.8 $\\pm$ 0.4 & 96.1 $\\pm$ 1.6 & 14.5 \\\\ S9 & 0.293 $\\pm$ 0.052 & 0.825 $\\pm$ 0.02 & 81. $\\pm$ 0.7 & 147.58 $\\pm$ 0.44 \\ & 225.2 $\\pm$ 2.3 & 1987.8 $\\pm$ 2.1 & 58. $\\pm$ 9.5 & 15.1 \\\\ S12 & 0.308 $\\pm$ 0.008 & 0.9 $\\pm$ 0.003 & 31.61 $\\pm$ 0.76 & 240.4 $\\pm$ \\ 4.6 & 308.8 $\\pm$ 3.8 & 1995.63 $\\pm$ 0.03 & 62.5 $\\pm$ 2.3 & 15.5 \\\\ S13 & 0.297 $\\pm$ 0.012 & 0.49 $\\pm$ 0.023 & 25.5 $\\pm$ 1.6 & 73.1 $\\pm$ 4.1 \\ & 248.2 $\\pm$ 5.4 & 2004.9 $\\pm$ 0.09 & 59.2 $\\pm$ 3.8 & 15.8 \\\\ S14 & 0.256 $\\pm$ 0.01 & 0.963 $\\pm$ 0.006 & 99.4 $\\pm$ 1. & 227.74 $\\pm$ 0.7 \\ & 339. $\\pm$ 1.6 & 2000.07 $\\pm$ 0.06 & 47.3 $\\pm$ 2.9 & 15.7 \\\\ S17 & 0.311 $\\pm$ 0.004 & 0.364 $\\pm$ 0.015 & 96.44 $\\pm$ 0.18 & 188.06 $\\pm$ \\ 0.32 & 319.45 $\\pm$ 3.2 & 1992. $\\pm$ 0.3 & 63.2 $\\pm$ 2. & 15.3 \\\\ S18 & 0.265 $\\pm$ 0.08 & 0.759 $\\pm$ 0.052 & 116. $\\pm$ 2.7 & 215.2 $\\pm$ 3.6 \\ & 151.7 $\\pm$ 2.9 & 1996. $\\pm$ 0.9 & 50. $\\pm$ 16. & 16.7 \\\\ S19 & 0.798 $\\pm$ 0.064 & 0.844 $\\pm$ 0.062 & 73.58 $\\pm$ 0.61 & 342.9 $\\pm$ \\ 1.2 & 153.3 $\\pm$ 3. & 2005.1 $\\pm$ 0.22 & 260. $\\pm$ 31. & 16. \\\\ S21 & 0.213 $\\pm$ 0.041 & 0.784 $\\pm$ 0.028 & 54.8 $\\pm$ 2.7 & 252.7 $\\pm$ \\ 4.2 & 182.6 $\\pm$ 8.2 & 2028.1 $\\pm$ 5.5 & 35.8 $\\pm$ 6.9 & 16.9 \\\\ S24 & 1.06 $\\pm$ 0.178 & 0.933 $\\pm$ 0.01 & 106.3 $\\pm$ 0.93 & 4.2 $\\pm$ 1.3 \\ & 291.5 $\\pm$ 1.5 & 2024.9 $\\pm$ 5.5 & 398. $\\pm$ 73. & 15.6 \\\\ S27 & 0.454 $\\pm$ 0.078 & 0.952 $\\pm$ 0.006 & 92.91 $\\pm$ 0.73 & 191.9 $\\pm$ \\ 0.92 & 308.2 $\\pm$ 1.8 & 2059.7 $\\pm$ 9.9 & 112. $\\pm$ 18. & 15.6 \\\\ S29 & 0.397 $\\pm$ 0.335 & 0.916 $\\pm$ 0.048 & 122. $\\pm$ 11. & 157.2 $\\pm$ \\ 2.5 & 343.3 $\\pm$ 5.7 & 2021. $\\pm$ 18. & 91. $\\pm$ 79. & 16.7 \\\\ S31 & 0.298 $\\pm$ 0.044 & 0.934 $\\pm$ 0.007 & 153.8 $\\pm$ 5.8 & 103. $\\pm$ \\ 11. & 314. $\\pm$ 10. & 2013.8 $\\pm$ 2.2 & 59.4 $\\pm$ 9.2 & 15.7 \\\\ S33 & 0.41 $\\pm$ 0.088 & 0.731 $\\pm$ 0.039 & 42.9 $\\pm$ 4.5 & 82.9 $\\pm$ 5.9 \\ & 328.1 $\\pm$ 4.5 & 1967.9 $\\pm$ 6.5 & 96. $\\pm$ 21. & 16. \\\\ S38 & 0.139 $\\pm$ 0.041 & 0.802 $\\pm$ 0.041 & 166. $\\pm$ 22. & 286. $\\pm$ 68. \\ & 203. $\\pm$ 68. & 2003. $\\pm$ 0.2 & 18.9 $\\pm$ 5.8 & 17. \\\\ S66 & 1.21 $\\pm$ 0.126 & 0.178 $\\pm$ 0.039 & 135.4 $\\pm$ 2.6 & 96.8 $\\pm$ 2.9 \\ & 106. $\\pm$ 6.3 & 1782. $\\pm$ 23. & 486. $\\pm$ 41. & 14.8 \\\\ S67 & 1.1 $\\pm$ 0.102 & 0.368 $\\pm$ 0.041 & 139.9 $\\pm$ 2.3 & 106. $\\pm$ 6.1 \\ & 215.2 $\\pm$ 4.8 & 1695. $\\pm$ 16. & 419. $\\pm$ 19. & 12.1 \\\\ S71 & 1.06 $\\pm$ 0.765 & 0.844 $\\pm$ 0.075 & 76.3 $\\pm$ 3.6 & 34.6 $\\pm$ 1.5 \\ & 331.4 $\\pm$ 7.1 & 1646. $\\pm$ 251. & 399. $\\pm$ 283. & 16.1 \\\\ S83 & 2.79 $\\pm$ 0.234 & 0.657 $\\pm$ 0.096 & 123.8 $\\pm$ 1.3 & 73.6 $\\pm$ 2.1 \\ & 197.2 $\\pm$ 3.5 & 2061. $\\pm$ 25. & 1700. $\\pm$ 205. & 13.6 \\\\ S87 & 1.26 $\\pm$ 0.161 & 0.423 $\\pm$ 0.036 & 142.7 $\\pm$ 4.4 & 109.9 $\\pm$ \\ 2.9 & 41.5 $\\pm$ 3.7 & 1647. $\\pm$ 38. & 516. $\\pm$ 44. & 13.6 \\\\ S96 & 1.55 $\\pm$ 0.209 & 0.131 $\\pm$ 0.054 & 126.8 $\\pm$ 2.4 & 115.78 $\\pm$ \\ 1.93 & 231. $\\pm$ 9. & 1624. $\\pm$ 34. & 701. $\\pm$ 81. & 10. \\\\ S97 & 2.19 $\\pm$ 0.844 & 0.302 $\\pm$ 0.308 & 114.6 $\\pm$ 5. & 107.72 $\\pm$ \\ 3.15 & 38. $\\pm$ 52. & 2175. $\\pm$ 88. & 1180. $\\pm$ 688. & 10.3 \\\\ \\tableline \\end{tabular} \\end{scriptsize} \\caption{Orbital Parameters of the S-stars examined in the paper: $a$ is the semimajor axis, $e$ is the eccentricity, $i$ is the inclination of the normal of the orbit with respect to the line of sight, $\\Omega$ is the position angle of the ascending node, $\\omega$ is the periapse anomaly with respect to the ascending node, $t_{\\mathrm{P}}$ is the epoch of last/next periapse, $T$ is the orbital period, $K$ is the apparent magnitude in the $K$ band (data taken from \\citealt{Gillessen08}).} \\label{tbl-1} \\end{table*}   \\begin{table*} \\centering \\begin{scriptsize} \\begin{tabular}{lcccccc} \\tableline % \\tableline % Star & $t_0[\\mathrm{yr}]$ & $t_{1/2}[\\mathrm{yr}]$ & $K_2$ & $\\theta_2[\\mathrm{\\mu as}]$ & $\\gamma_0[^{\\circ}]$ & $D_{\\mathrm{LS},0}[AU]$ \\\\ \\tableline S1 & 2130.34 & 4.9 & 32.2 & 54 & 30.8 & 2195 \\\\ S2 & 2018.15 & 0.095 & 26.8 & 43 & 45.4 & 127 \\\\ S4 & 2048.94 & 3.9 & 28.7 & 112 & 12.2 & 2868 \\\\ S5 & 2028.8 & 2.7 & 32.9 & 32 & 83.6 & 456 \\\\ S6 & 2062.65 & 0.051 & 20.8 & 316 & 3.56 & 464 \\\\ S8 & 2082.7 & 1.6 & 28.3 & 88 & 16.2 & 1335 \\\\ S9 & 2041.39 & 1.9 & 27.1 & 144 & 9.06 & 1833 \\\\ S12 & 2058.38 & 1.4 & 32.6 & 32 & 88.5 & 319 \\\\ S13 & 2060.3 & 34 & 36.4 & 31 & 90.8 & 1571 \\\\ S14 & 2047.55 & 0.073 & 23.5 & 136 & 9.46 & 230 \\\\ S17 & 2008.24 & 2.1 & 26.9 & 196 & 6.44 & 2942 \\\\ S18 & 2045.15 & 0.72 & 30.9 & 61 & 26.2 & 670 \\\\ S19 & 2262.63 & 1.3 & 29.9 & 87 & 16.5 & 1369 \\\\ S21 & 2027.23 & 1.1 & 32.3 & 49 & 36.5 & 617 \\\\ S24 & 2047.95 & 44 & 33.1 & 81 & 17.9 & 6225 \\\\ S27 & 2062.08 & 0.55 & 22.4 & 399 & 2.92 & 1397 \\\\ S29 & 2021.73 & 0.98 & 31.6 & 51 & 34.1 & 572 \\\\ S31 & 2013.9 & 0.56 & 31.7 & 31 & 91.2 & 189 \\\\ S33 & 2067.95 & 9.1 & 34.9 & 39 & 53.7 & 1637 \\\\ S38 & 2021.8 & 0.71 & 33.5 & 32 & 88.9 & 242 \\\\ S66 & 2254.09 & 66 & 36.7 & 43 & 45.4 & 8328 \\\\ S67 & 2035.14 & 116 & 34.6 & 40 & 51.5 & 9053 \\\\ S71 & 2059.51 & 8.3 & 31.6 & 100 & 13.9 & 4073 \\\\ S83 & 3647.18 & 129 & 35.6 & 50 & 35. & 14435 \\\\ S87 & 2185.1 & 52 & 35.5 & 39 & 53.2 & 6508 \\\\ S96 & 2073.74 & 131 & 32.2 & 48 & 36.8 & 14049 \\\\ S97 & 2265.5 & 68 & 30.9 & 64 & 24.6 & 13897 \\\\ \\tableline \\end{tabular} \\end{scriptsize} \\caption{Summary of the main features of the secondary images of the S-stars: $t_{0}$ is the epoch of the peak, $t_{1/2}$ is the time spent by the secondary image at a luminosity higher than half the maximum, $K_{2}$ is the peak $K$-band magnitude for the secondary image, $\\theta_{2}$ is the maximal angular distance from the apparent shadow of the black hole, $\\gamma_{0}$ is the angle formed by the line connecting the selected star with the central black hole and the optical axis at the time of the peak, and $D_{\\mathrm{LS},0}$ is the distance of the S-star from the MBH at this time.} \\label{tbl-2} \\end{table*}  \\begin{figure} \\centering{\\includegraphics{Figure1.eps}} \\caption{Left column: light curves for the secondary images of S6 (a), S17 (c), S27(e). Right column: the alignment angle $\\gamma$ as a function of time for S6 (b), S17 (d), S27(f).} \\label{Fig 1} \\end{figure}"
    },
    "0812/0812.3312_arXiv.txt": {
        "abstract": "There is a robust upper limit on the energy of synchrotron radiation in high-energy astrophysics: $ \\sim m_{\\rm e} c^2 /\\alpha$, where $\\alpha = 1/137$ is the fine structure constant and the value refers to the comoving frame of the fluid. This is the maximal energy of synchrotron photons which can be emitted by an electron having an arbitrarily high initial energy after it turns by angle $\\sim \\pi$ in the magnetic field. This upper limit can be naturally reached if the converter mechanism contributes to the jet radiation and can be imprinted in spectra of some blazars as a cutoff or a dip in the GeV range. We use numerical simulations to probe the range of parameters of a radiating jet where the ultimate synchrotron cutoff appears. We reproduce the variety of spectra depending on the source luminosity and on the scale of the emission site. We also compare our results with the EGRET blazar spectra in order to illustrate that agreement is possible but still not statistically significant. The predicted feature, if it exists, should be observed by {\\it Fermi} in spectra of some blazars. ",
        "introduction": "The issue of maximal possible synchrotron radiation energy in astrophysics (Guilbert, Fabian \\& Rees 1983) was discussed in various frameworks: the extreme Fermi acceleration of electrons at shocks (e.g. Ghisellini 1999) and the converter mechanism (Derishev 2003; Stern 2003). In the extreme acceleration scenario the electron energy is limited by the condition that the Larmor gyration time, being the shortest available acceleration time, is shorter than the synchrotron cooling time. In the converter mechanism, a new high-energy electron appears in the relativistic jet as a result of photon-photon pair (or photomeson) production and its initial energy can be arbitrarily high in the fluid frame. However, its initial direction in the jet frame is opposite to the jet bulk motion. Therefore, its synchrotron radiation at the first moment is Doppler deboosted in the external reference frame. When the electron has turned around, the emitted synchrotron photons become Doppler boosted. When it gyrates in the magnetic field by the angle $\\pi$, it cools down below the limiting Lorentz factor (in the comoving frame) $\\gamma_{\\rm m} \\sim 10^8 B^{-1/2}$, where $B$[G] is the magnetic field strength. In the extreme acceleration scenario it is approximately the same and in  both cases the maximal comoving energy of synchrotron photons does not depend on the magnetic field: $\\varepsilon_{\\rm m} \\sim 1/\\alpha$ (i.e. $ \\sim 10^2 $MeV), hereafter $\\varepsilon = E/m_{\\rm e} c^2$ is the photon energy in electron rest mass units. In the case of blazars, this cutoff is blueshifted by factor up to $2\\Gamma$, where $\\Gamma$ is the bulk Lorentz factor of the jet. Observations of the superluminal proper motion in blazars imply that a typical Lorentz factor is about $\\Gamma \\sim 10-20$ (see Jorstad et al. 2001). In some cases $\\Gamma \\approx 40$ and its value in the emission region could be even higher, when taking into account possible deceleration of the jet (see Stern \\& Poutanen 2008, hereafter SP08). Therefore, the ultimate synchrotron cutoff in the blazar spectra should be expected in the GeV range, maybe slightly above 10 GeV in some extreme cases. This feature should be seen as a spectral dip with a continuation of the spectrum to higher energies as a result of Comptonization of the external or synchrotron radiation.  In which physical context can the observable synchrotron cutoff in GeV range appear? The main mechanism of particle acceleration, which is responsible for the gamma-ray emission of blazars is still unknown. There are at least three possibilities which are discussed:  (i) {\\it Diffusive acceleration of electrons} by various plasma perturbations or turbulence (Fermi II mechanism). The ultimate acceleration in this case would imply that the relative bulk velocities of the fluid at the gyroradius scale  are relativistic (in the conditions of examples considered here the maximal gyroradius  $r_{\\rm m} \\sim 10^{-5} - 10^{-4} R_{\\rm j}$, where $R_{\\rm j}$ is the jet radius). This is scarcely possible.    (ii) {\\it Shock acceleration of charged particles} (Fermi I mechanism). It can work at internal shocks in the jet  or at the shear layer between the jet and the external environment. In this case, acceleration rate can be much faster, especially at the shear layer, where an electron entering the jet from the external environment can gain factor $\\Gamma^2$ in energy (provided that the thickness of the transition layer is less than the gyroradius). Therefore, the particle spectra can extend up to $\\gamma_{\\rm m}$. The issue is how much power can be emitted by the particles near the cutoff. Simulations of shock acceleration (Niemiec \\& Ostrowski 2006; Ostrowski 2008) demonstrate a large variety of resulting particle spectra depending on the geometry of the magnetic field and the power spectrum of its perturbations. For some parameters the spectra are very hard. On the other hand, in these simulations the radiative energy losses are neglected and it is difficult to conclude whether they extend close to   $\\gamma_{\\rm m}$. In the shear layer scenario, an electron having the energy above $\\gamma_{\\rm m}$  cannot penetrate into the jet deeper than $r_{\\rm m}$. Therefore, the energy budget for extreme acceleration at the shear layer is restricted by the bulk energy of a thin outer layer of the jet.  (iii) {\\it Converter mechanism}, which can take a form of a runaway photon breeding in relativistic jets (Derishev et al. 2003; Stern 2003; Stern \\& Poutanen 2006). The photon breeding mechanism is probably the most efficient  way of dissipation of the relativistic jet energy into gamma-rays. The mechanism should work if the active galactic nucleus (AGN) is sufficiently powerful  (above $L_{\\rm d} \\sim 10^{42} - 10^{44}\\ergs$, depending on the scale of the emission site) and the if jet is sufficiently fast  (the Lorentz factor above $\\Gamma \\sim 10$, see SP08). When photon breeding mechanism operates, the instant energy gain at photon -- electron-positron conversion in the jet  can reach $4 \\Gamma^2$ and the ultimate energy $\\gamma_{\\rm m}$ can be easily reached. Stern \\& Poutanen (2006), studying the photon breeding regime in jets, showed that the energy of particles through the breeding cycle grows till the Lorentz factor of produced pairs reaches  $\\gamma_{\\rm m}$ in the jet frame. SP08 performing numerical simulations of the photon breeding have obtained some hard ($\\beta \\sim 2$) spectra at a moderate luminosity of the source, where the ultimate synchrotron cutoff is clearly seen (hereafter we define the photon spectral index $\\beta$ as  ${\\rm d}N/{\\rm d}\\varepsilon\\propto \\varepsilon^{-\\beta}$). As discussed above, the possibility to reproduce such cutoff with other mechanisms is uncertain at best. Therefore, the observation of such a feature would be an argument in favor of the converter mechanism as a source of the jet radiation.   Existing high-energy data supplied mainly by EGRET (30 MeV -- 20 GeV energy range),  do not reveal such a cutoff at a significant level. Nevertheless, there are some indications of a nontrivial shape of the blazar spectra in the high-energy range, more complicated than a single high-energy hump. Such an indication has been found  by Costamante, Aharonian \\& Khangulyan (2007) in the spectrum of BL Lac PKS 2155--304. This is one of five objects, which has been detected by both EGRET and Cherenkov arrays: the combination of these two data sets imply that there is a double humped structure in the ten MeV  -- TeV range with a depression around 10 GeV.  This study is motivated by the launch of {\\it Fermi} (former {\\it GLAST}). This instrument covering 200 MeV -- 200 GeV range is perfect for detection of the feature we study (see also discussion in Costamante et al. 2007).   In this work we perform a series of simulations in order to outline the variety of conditions, where the ultimate synchrotron cutoff can be observed, and to demonstrate the variety of spectral shapes in the MeV -- TeV range produced by the photon breeding mechanism. The problem setup, described in Section 2, is similar to that of SP08. In the same section we present results of numerical simulations and discuss them. In Section 3 we compare our results with the EGRET blazar data in order to illustrate that the existing data have slightly insufficient accuracy to make significant conclusions about the structure of blazar spectra in the GeV range. In section 4 we discuss the perspectives of observations. ",
        "conclusions": "Our simulations demonstrate that the ultimate synchrotron cutoff in its clear form is a relatively rare phenomenon. In the framework of the photon breeding mechanism, the cutoff in the GeV range appears when the compactness of the external radiation is just above the threshold of the supercritical photon breeding. When the compactness exceeds the threshold by an order of magnitude, the spectrum becomes almost flat or the synchrotron component moves to the softer range (see Fig. \\ref{fig:spec}). Another tendency, which can be derived from the results, is that the spectral sequence depends on the angular distribution of the external radiation. When the large-angle component of the external radiation in the emission region is relatively strong (e.g. $\\sim 0.1$ of the direct disc radiation as in model A) the spectra are smoother and the separation between the synchrotron and Compton components is weakly expressed. When the side-on radiation is weak, e.g. it is supplied only by the periphery of the accretion disc, all features are  much sharper.   The number of blazars with spectra demonstrating the clear ultimate cutoff should be relatively small. However, {\\it Fermi} will be able to reveal this kind of spectra if it exists, because it will observe simultaneously a large number of objects. The cutoff in the right place is by itself insufficient to conclude that we observe the ultimate synchrotron cutoff: the latter can also be caused by the photon-photon opacity in the same energy range (see SP08). To be sure that the observed cutoff has a synchrotron nature, one have to observe the continuation of the spectrum above the cutoff in a form of the higher energy Compton component extending to the TeV range. In principle, {\\it Fermi} can detect the  low-energy part of the Compton maximum at $E \\sim 100$ GeV and confirm the two-componentness in this way even without detection of the TeV emission by large ground-based detectors.  Although the possibility of the extreme acceleration up to the energy $\\gamma_{\\rm m}$ in internal shocks is not evident, it is not yet ruled out. Do independent signatures exist of the photon breding mechanism? A possible one is the correlation between the disc brightness and the spectral shape like those in Fig.  \\ref{fig:spec}.  If the total power of the jet is known, then the high radiative efficiency, i.e. the ratio of the gamma-ray luminosity to the total jet power, of e.g. more than 10 per cent can give an additional support to the photon breeding hypothesis and argue against the internal shock scenario. In the internal shock mechanism the luminosity budget is limited by the {\\it internal energy} of the jet, while in the case of the photon breeding a large fraction of the {\\it total bulk energy} of the jet  can be converted into radiation.  Another possible clue is the compactness threshold:  at a low compactness of the disc radiation, the photon breeding does not work (the threshold for  model C is  $L_{\\rm d}/R_{15} \\sim (1 - 10)\\times 10^{42} \\ergs$). Usually it is difficult to estimate the distance $R$ from the disc to the emission region, however, sometimes approximate constraints can be obtained using the time variability or the mass of the black hole ($R$ can not be smaller than a few $R_{\\rm g}$). In this way the photon breeding mechanism can be revealed statistically: sources with the low estimated compactness should be  less efficient and have different spectra. When the photon breeding turns on, the luminosity should increase sharply. Therefore one can expect a kind of bimodality in the efficiency and probably in the high-energy spectra of blazars."
    },
    "0812/0812.4018_arXiv.txt": {
        "abstract": "Supernovae are the dominant energy source for driving turbulence within the interstellar plasma. Until recently, their effects on magnetic field amplification in disk galaxies remained a matter of speculation. By means of self-consistent simulations of supernova-driven turbulence, we find an exponential amplification of the mean magnetic field on timescales of a few hundred million years. The robustness of the observed fast dynamo is checked at different magnetic Reynolds numbers, and we find sustained dynamo action at moderate $\\Rm$. This indicates that the mechanism might indeed be of relevance for the real ISM.  Sensing the flow via passive tracer fields, we infer that SNe produce a turbulent $\\alpha$~effect which is consistent with the predictions of quasilinear theory. To lay a foundation for global mean-field models, we aim to explore the scaling of the dynamo tensors with respect to the key parameters of our simulations. Here we give a first account on the variation with the supernova rate. ",
        "introduction": "\\label{sec:slow_fast}  In laminar dynamos, diffusion sets the relevant timescale for magnetic reconnection, thus defining an upper limit for the allowable growth rate of the mean magnetic field. Because the microscopic diffusivity is usually low, these dynamos are commonly referred to as ``slow dynamos''. Within the ISM, the diffusion time $\\tau_{\\rm d}=L^2/\\eta$ (related to the microscopic value $\\eta\\simeq10^8\\vis$) by far exceeds the Hubble time. This means that the field amplification mechanism in galaxies needs to be a ``fast dynamo'' in the sense that it works on a timescale different than $\\tau_{\\rm d}$ \\citep{1999ApJ...517..700L}. Accordingly, one assumes that the galactic dynamo is determined by some sort of effective turbulent diffusivity $\\eta_{\\rm t}$. Although the dynamo may still be limited by topological changes via reconnection, this process can be considerably faster due to the much higher value of $\\eta_{\\rm t}$ compared to $\\eta$. Turbulent reconnection has recently been studied numerically by Otmianowska-Mazur, Kowal \\& Lazarian (this volume).  According to the definition of $\\tau_{\\rm d}$, in the laminar case one expects the dynamo growth rate to increase with $\\eta$. The picture, however, changes significantly as soon as the Reynolds number is high enough to allow for developed turbulence. Structures of integral length scale $L$ are now efficiently broken down to the Kolmogorov microscale where the atomic diffusion takes over. Varying the microscopic value of $\\eta$ does only change the extent of the inertial range towards higher wavenumbers but is no longer relevant to the large-scale flow. This implies that a turbulent dynamo should be insensitive to variations in $\\eta$ as soon as a critical value $\\Rm_{\\rm c}$ is exceeded \\citep[cf.][]{1999ApJ...517..700L}.  \\begin{figure} \\center\\includegraphics[width=0.95\\columnwidth]{fig/fig1} \\caption{Evolution of the regular (a) and fluctuating (b) magnetic field strength for different $\\Rm$. For clarity, the ordinates of the different models have been offset by an order of magnitude each. In panel (c), we compare the growth rates for the turbulent (diamonds) and regular (triangles) magnetic field. The unconnected data points to the right correspond to a run with $\\eta=0$ and provide an indication for the level of numerical diffusivity.} \\label{fig:fig1} \\end{figure}  In Fig.~\\ref{fig:fig1}, we plot the evolution of the mean and turbulent magnetic field (normalised to the kinetic energy) for different magnetic Reynolds numbers $\\Rm = L^2\\,q\\Omega/\\eta$. The values for $\\Rm$ are obtained by varying $\\eta$ while keeping the rotation rate $\\Omega$ fixed. The kinematic viscosity $\\nu$ is furthermore adopted to keep the magnetic Prandtl number $\\Pm=\\nu/\\eta=2.5$ constant. Irrespective of the value of $\\Rm$, we observe a nice and steady exponential amplification of both the regular and turbulent fields. To estimate the influence of the finite grid resolution, we have performed a fiducial run (dark grey lines in panels (a) and (b) of Fig.~\\ref{fig:fig1}) at double the grid spacing for $\\Rm=10,000$. The obtained values are consistent with the higher resolved run and provide a first indication that the simulation results are reasonably converged at this level of dissipation and below.  In panel (c) of Fig.~\\ref{fig:fig1}, we compare the growth rates for the turbulent (diamonds) and regular (triangles) magnetic field as a function of the magnetic Reynolds number. The unconnected data points to the right correspond to a run with $\\eta=0$, i.e., the case of (formally) infinite $\\Rm$. As can be seen from a comparison with these points, above $\\Rm\\simeq10,000$ we are limited by the finite value of the inherent numerical diffusivity of our code, i.e., better resolved runs become mandatory to study the regime of higher $\\Rm$. With respect to the reliably converged runs, we observe growth rates that increase with $\\Rm$ -- suggesting that the effect is of genuinely turbulent nature.  In their models of the cosmic-ray-driven buoyant instability, \\citet{2007ApJ...668..110O} observe that their dynamo crucially relies on the presence of a ``microscopic'' diffusivity $\\eta$. The authors actually find the efficiency of the field amplification to scale with this parameter. Moreover, they state that no scale separation is manifest in the spectra of their simulations. In this respect, it remains disputable whether the high value for $\\eta$ might rather be interpreted as an effective turbulent dissipation $\\eta_{\\rm t}$. If so, the simulations would have to be regarded as large eddy simulations, i.e., simulations assuming turbulent effects rather than having them emerge from first principles. ",
        "conclusions": "We have presented recent simulation results on the variation of the key dynamo parameters with the supernova frequency. Yet more work is needed to understand how the vertical structure of the galactic disk is linked to the emergence of the dynamo effect and the associated transport processes. As a primary question, one is of course interested in finding a simple explanation for the observed scaling relations -- which are somewhat flatter than the square-root of the supernova rate $\\sigma$.  From quasilinear theory, it is understood that the dynamo effect depends on vertical gradients both in the density $\\rho$ and the turbulence intensity $u'$. The disk structure, on the other hand, is determined by the balance of the kinetic pressure from the forcing and the gravitation of the stellar component -- here equipartition arguments can yield a first approximation. To connect kinematic properties with dynamo parameters, we however need to specify the correlation time $\\tau_{\\rm c}$ of the turbulent flow. By comparing direct simulations with SOCA predictions, it will eventually become possible to determine how this quantity depends on the supernova frequency.  To connect the measured dynamo parameters $\\alpha$ and $\\tilde{\\eta}$ with the growth rates observed in the simulations -- and thereby check the applicability of the approach -- more mean-field modelling is necessary. We are, however, confident that the classical framework of MF-MHD remains a valuable tool in the quest of understanding galactic magnetism."
    },
    "0812/0812.3915_arXiv.txt": {
        "abstract": "Black hole and neutron star accretion flows display unusually high levels of hard coronal emission in comparison to all other optically thick, gravitationally bound, turbulent astrophysical systems.  Since these flows sit in deep relativistic gravitational potentials, their random bulk motions approach the speed of light, therefore allowing turbulent Comptonization to be an important effect.  We show that the inevitable production of hard X-ray photons results from turbulent Comptonization in the limit where the turbulence is trans-sonic and the accretion power approaches the Eddington Limit.  In this regime, the turbulent Compton {\\it y--}parameter approaches unity and the turbulent Compton temperature is a significant fraction of the electron rest mass energy, in agreement with the observed phenomena. ",
        "introduction": "Like the envelopes of late-type dwarfs, accretion disks are thought to be turbulent and optically thick. Turbulence in stellar envelopes is driven by convection, while in accretion disks turbulence is commonly thought to be driven by the magnetorotational instability (MRI; Balbus \\& Hawley 1991).  Accretion disks found in proto-stellar and proto-planetary systems, cataclysmic variables, X-ray binaries and quasars/AGN all display coronal behavior as do late-type dwarfs.  By ``corona,'' we mean the {\\it appearance}, in a spectral energy distribution, of a relatively energetic population of particles that persistently emits radiation at temperatures in excess of the given object's thermal photosphere.  In the case of black hole and neutron star accretion flows, the ratio of coronal to bolometric power $L_{_{\\rm c}}/L_{_{\\rm bol}} \\sim 0.1-1$.  In other systems that display radiatively efficient accretion, the value of $L_{_{\\rm c}}/L_{_{\\rm bol}}$ is smaller by orders of magnitude.  For example, turbulent stellar envelopes acheive $L_{_{\\rm c}}/L_{_{\\rm bol}}\\lesssim 10^{-3}$ in the most extreme (rapidly rotating) case.  Here, we ask why black hole and neutron star accretion flows display such abnormally high levels of coronal power, observed in the form of hard power-law X-ray photons.  In the directly observable case of the solar corona, mechanical energy leaks out through the photosphere in some combination of waves and relatively steady magnetic structures. The two reservoirs of stellar mechanical energy are convective turbulence and differential rotation.  Since proto-planetary and white dwarf accretion disks as well as rapidly rotating late-type dwarfs display relatively modest levels of coronal power, differential rotation cannot be the sole key ingredient in explaining the extreme coronal phenomena seen in relativistic accretion flows.  Furthermore, the difference between MRI-driven and convective turbulence must be ruled out as well because the non-relativistic accretion flows display only modest coronal phenomena.  An obvious feature that is unique to black hole and neutron star accretion flows is that close to the central sources of gravity, where most of the accretion power is generated, the bulk motions approach relativistic values.  If the electrons possess bulk motions that are trans-relativistic, then they may readily amplify soft photons up to high energies, potentially forming a hard power-law X-ray spectra though the process of second order bulk Comptonization.\\footnote{As an aside, Cosmologists refer to this radiative mechanism as the ``second-order kinetic S-Z effect.''}  In the brief discussion that follows, we focus solely on the importance of turbulent Comptonization in relativistic accretion flows.  Our approach differs from that of Socrates et al. (2004) and Thompson (2006) in that we completely ignore the relationship between the region from which soft thermal seed photons originate -- presumeably a cool thermal disk -- with the region of Comptonization.  We start with a description of an isolated eddy undergoing turbulent Comptonization. ",
        "conclusions": "At radii immediately further out from the inner-most stable circular orbit, the ability for trans-sonic turbulence to Comptonize soft photons diminishes since $v^2_{\\lambda}$ decreases, while $\\tau_{\\lambda}$ increases.  The latter implies that the input photons can only ``sample'' eddies of lower energy, on smaller scales (Socrates et al. 2004). Also at larger radii, the ability for the flow to absorb a photon via free-free and bound-free processes increases (Socrates et al. 2006). We suspect that these adjacent outer regions serve as the source of -- presumably thermal --soft input photons for the centrally located Comptonizing eddies. In the case accretion onto a neutron star, thermal emission from the stellar surface is additional source of soft seeds.  We have yet to specify the actual source of trans-sonic turbulence itself.  In the familiar language of the Shakura \\& Sunyaev (1973) thermal disk model, we have assumed throughout that $\\alpha\\lesssim 1$ i.e., $v^2_{\\lambda}\\lesssim c^2_s$.  It is now widely accepted that the MRI (Balbus \\& Hawley 1991) generically leads to a turbulent accretion stress, allowing for the transfer of angular momentum to occur in highly conducting fluids.  However, simulations of the MRI typically find low values for the Shakura \\& Sunyaev (1973) $\\alpha$-parameter such that $\\alpha \\sim 0.1-0.01$, though a clear understanding for the value of the MRI's $\\alpha$ is not currently in hand.  Our main results for Eddington-limited trans-sonic turbulence near a relativistic source of gravity are as follows:  1) Turbulence is responsible for transporting angular momentum while simultaneously converting gravitational power of protons into Comptonized photon power.  2) The dynamical, accretion, Compton cooling and eddy overturn time are all comparable to one another.  3) The turbulent Compton power of each individual eddy is comparable to the accretion power of that eddy.  4) Electrons will inevitably heat up to the a significant fraction of the electron virial temperature. In other words, turbulent Comptonization is an important electron heating mechanism.  5) The $y-$parameter due to turbulent Comptonization is close to unity.  Therefore, the outgoing hard X-ray spectrum is a power law with flat spectral index $\\Gamma\\sim 2$, in agreement with observations.  6) The results above are independent of black hole mass."
    },
    "0812/0812.2913_arXiv.txt": {
        "abstract": "Photo-heating associated with reionisation and kinetic feedback from core-collapse supernovae have previously been shown to suppress the high-redshift cosmic star formation rate.  Here we investigate the interplay between photo-heating and supernova feedback using a set of cosmological, smoothed particle hydrodynamics simulations. We show that photo-heating and supernova feedback mutually amplify each other's ability to suppress the star formation rate. Our results demonstrate the importance of the simultaneous, non-independent inclusion of these two processes in models of galaxy formation to estimate the strength of the total negative feedback they exert. They may therefore be of particular relevance to semi-analytic models in which the effects of photo-heating and supernova feedback are implicitly assumed to act independently of each other. ",
        "introduction": "\\label{Sec:Introduction} The cosmic star formation rate (SFR) is an important observable of our Universe. It is affected by a variety of physical processes, many of which are in turn regulated by the SFR, giving rise to so-called feedback loops (for an overview see, e.g., \\citealp{Ciardi:2005}).  Photo-ionisation heating due to the absorption of ionising photons from star-forming regions and the injection of kinetic energy from supernova (SN) explosions of massive stars provide two such feedback loops. Their implications for the assembly of the first generation of galaxies have been extensively discussed in studies of the epoch of reionisation (for a review of this epoch see, e.g., \\citealp{Loeb:2001}). \\par Photo-heating associated with reionisation increases the mean temperature of the intergalactic medium (IGM) to $\\sim 10^4 \\K$ (e.g., \\citealp{Hui:1997}) and reduces the rate at which hotter gas can cool (\\citealp{Efstathiou:1992,Wiersma:2008}). The increase in the gas temperature keeps the IGM smooth and prevents the assembly of low-mass galaxies, that is, galaxies with masses corresponding to a virial temperature $\\lesssim 10^4 \\K$ (e.g., \\citealp{Shapiro:1994}; \\citealp{Gnedin:1998}). Moreover, the gas in galaxies that have already collapsed is relatively quickly photo-evaporated (e.g., \\citealp{Shapiro:2004}; \\citealp{Iliev:2005}), strongly decreasing the gas fraction of low-mass halos (e.g., \\citealp{Thoul:1996}; \\citealp{Barkana:1999}; \\citealp{Dijkstra:2004}; \\citealp{Susa:2006}). Indeed, the cosmic SFR has been predicted to exhibit a distinct drop around the redshift of reionisation (\\citealp{Barkana:2000}). Photo-ionisation heating is therefore said to provide a negative feedback on reionisation.\\footnote{As pointed out by \\cite{Pawlik:2008}, photo-heating also provides a strong positive feedback on reionisation because the increase in the gas temperature smoothes out density fluctuations, reducing the recombination rate.} \\par SN explosions of massive stars typically inject a few solar masses of gas with velocity of $\\sim 10^4\\kms$, corresponding to a kinetic energy of $\\sim 10^{51} \\erg$. The ejected material sweeps up and shock-heats the surrounding gas, entraining outflows sufficiently powerful to, at least temporarily, substantially reduce the gas fractions for galaxy-scale dark matter halos (e.g., \\citealp{Yepes:1997}; \\citealp{Scannapieco:2006}). Since this leads to a suppression of the SFR, SN explosions, like photo-heating from reionisation, provide a negative feedback on reionisation. In addition to the depth of the gravitational potential, the ability of SN feedback to reduce the gas fractions generally depends on the geometry of the gas distribution (e.g., \\citealp{MacLow:1999}). \\par Studies of the effects of photo-heating and SN feedback on the SFR that considered each process in isolation have been augmented by studies that included both photo-heating and SN feedback. These studies include simulations of the formation of the first stars (e.g., \\citealp{Greif:2007}; \\citealp{Wise:2008}; \\citealp{Whalen:2008}), of the evolution of isolated galaxies (e.g., \\citealp{Fujita:2004}; \\citealp{caius:2008}, \\citealp{Tasker:2008}), and of galaxies in a cosmological volume (e.g., \\citealp{Tassis:2003}). Some of these studies also investigate the interplay between photo-heating and SN feedback. For example, \\cite{Kitayama:2005} demonstrated in a one-dimensional hydrodynamical study that a previous episode of photo-heating may increase the efficiency of the evacuation of dark matter halos by (thermal) SN feedback and enable the destruction of galaxies out to much larger masses. In this letter we report on another interaction of star formation feedbacks, using three-dimensional galaxy formation simulations. \\par We employ a set of Smoothed Particle Hydrodynamics (SPH) cosmological simulations that include star formation, and photo-ionisation heating from a uniform ultraviolet (UV) background and/or kinetic feedback from core-collapse SNe. We investigate how photo-heating affects the high-redshift ($z \\ge 6$) SFR in the presence and absence of SN feedback and, conversely, how SN feedback affects the cosmic SFR in the presence and absence of a photo-ionising background. We find that the inclusion of SN feedback amplifies the suppression of the cosmic SFR due to the inclusion of photo-heating.  On the other hand, the inclusion of photo-heating amplifies the suppression of the cosmic SFR due to the inclusion of SN feedback. \\par Photo-heating and SN feedback therefore mutually amplify each other in suppressing the SFR. Our results are relevant to current implementations of (semi-) analytic models of galaxy formation (see \\citealp{Baugh:2006} for a review), in which the effects of photo-heating and SN feedback are implicitly assumed to act independently of each other. These models may thus underestimate the strength of the combined negative feedback from photo-heating and SNe. \\par We emphasise that none of our simulations has a sufficiently high resolution to achieve convergence in the cosmic SFR. Future simulations will therefore be required to quantify our qualitative statement. We show, however, that the factor by which photo-heating and SN feedback amplify each other's ability to suppress the cosmic SFR becomes larger with increasing resolution, giving credibility to our main conclusion. \\par This letter is organised as follows. We present our simulation method in \\S\\ref{Sec:Simulations} and we illustrate our main result in \\S\\ref{Sec:Results}. Finally, we summarize our conclusions and discuss the caveats inherent to the present work in \\S\\ref{Sec:Discussion}. ",
        "conclusions": "\\label{Sec:Discussion} Photo-heating from reionisation and supernova (SN) feedback are key processes that determine the star formation rate (SFR) in the high-redshift Universe. Using a set of cosmological SPH simulations that include radiative cooling and star formation, we analysed the $z \\ge 6$ star formation history in the presence of photo-ionisation heating from a uniform UV background and/or kinetic feedback from core-collapse SNe. The inclusion of photo-heating and SN feedback both lead to a suppression of the SFR. \\par We showed that the factor by which the SFR is suppressed due to photo-heating is larger in the presence than in the absence of SN feedback. We also showed that the factor by which the SFR is suppressed due to the inclusion of SN feedback is larger in the presence than in the absence of photo-heating. This mutual amplification of SN feedback and reionisation heating is the central result of the present work. \\par We caution the reader that our simulations have not fully converged with respect to resolution and that we have ignored some potentially important physical processes. This result will therefore need to be confirmed and quantified with future simulations. In what follows we briefly discuss the most important physical effects that our analysis ignored. \\par We computed photo-heating rates from a uniform UV background in the optically thin limit. Our simulations therefore do not account for the self-shielding of gas from ionising radiation. This may lower the fraction of the gas that is photo-evaporated (e.g., \\citealp{Kitayama:2000}; \\citealp{Susa:2004}; \\citealp{Dijkstra:2004}). \\cite{Iliev:2005} (extending the work of \\citealp{Shapiro:2004}) have, however, pointed out that regions that are initially self-shielded will eventually also be photo-evaporated: as subsequent layers of gas are photo-evaporated, previously self-shielded regions become exposed to ionising radiation and are eventually stripped away. Moreover, the evaporation of the initially self-shielded gas may proceed at a speed comparable to the speed predicted by simulations that compute photo-heating rates in the optically thin limit (Fig.~3 in \\citealp{Iliev:2005}). Note also that because SN explosions decrease the gas density, the effects of self-shielding may be less prominent in simulations that include this type of feedback.  Clearly, the importance of self-shielding and the applicability of the optically thin limit to the present problem need to be critically assessed using cosmological radiation-hydrodynamical simulations. \\par Although all of our simulations employ a sufficiently high resolution to resolve all halos with virial temperatures $T_{\\rm vir} \\gtrsim 10^4 \\K$ with at least 100 particles, none of them has the resolution to properly reproduce the properties of the multi-phase medium associated with the star-forming regions these halos host. The factors by which photo-heating and SN feedback suppress the SFR are therefore not yet converged. We have, however, demonstrated that the effect of mutual amplification of photo-heating and SN feedback becomes only stronger with increasing resolution. \\par Because we have assumed the presence of a photo-dissociating background, the formation of molecular hydrogen is suppressed in our simulations. In reality, the gas may contain a significant fraction of molecular hydrogen before reionisation (but see, e.g., \\citealp{Haiman:1997}). Star formation and the associated kinetic feedback would then already be efficient in halos with virial temperatures much smaller than $10^4\\K$ (e.g., \\citealp{Tegmark:1997}). Because the gas fraction in these halos would be affected by both photo-heating and SN feedback, we expect that the inclusion of molecular hydrogen would only strengthen our main result. \\par We have also ignored the existence of atoms and ions heavier than helium. SN explosions may, however, quickly enrich the interstellar and intergalactic gas with metals (e.g., \\citealp{Bromm:2003}), which would increase its ability to cool. Metal enrichment thus provides a positive feedback that may partially offset the negative kinetic feedback from SN explosions. \\par None of the caveats we discussed seems, however, likely to invalidate our main qualitative conclusion that photo-heating and SN explosions amplify each other's effect on the cosmic SFR. Galaxy formation models that treat the effects of photo-heating and SN explosions independently of each other, like e.g. (semi-) analytic models, may therefore significantly underestimate the effect of these feedback processes on the SFR. Because photo-heating and SN feedback are important processes that affect the gas in low mass-halos at all epochs, our findings may also have implications for the understanding of the properties of low-redshift galaxies, e.g. in the context of the missing satellite problem (for a recent discussion see, e.g., \\citealp{Koposov:2009}). \\par"
    },
    "0812/0812.2252_arXiv.txt": {
        "abstract": "{} {We investigate the effects of radiative shocks on the observed X-ray emission from the Type~II supernova SN1993J. To this end, the X-ray emission is modeled as a result of the interaction between the supernova ejecta and a dense circumstellar medium at an age of 8 years.} {The circumstances under which the reverse shock is radiative are discussed and the observed X-ray emission is analyzed using the numerical code described in Nymark et al (2006). We argue that the original analysis of the X-ray observations suffered from the lack of self-consistent models for cooling shocks with high density and velocity, leading to questionable conclusions about the temperatures and elemental abundances. We reanalyze the spectra with our numerical model, and discuss the expected spectra for different explosion models for the progenitors. } {We find that the spectra of SN~1993J are compatible with a CNO-enriched composition and that  the X-ray flux is dominated by the reverse shock.} {} ",
        "introduction": "\\label{sect:intro} Many core-collapse supernovae (SNe) show strong X-ray and radio emission. This is, in particular, true for the Type~IIL, IIn and IIb SNe \\citep[see \\eg ][ for a review]{Immler03}.  We now know that this is a result of the interaction of the supernova ejecta and the circumstellar medium (CSM) of the supernova. A two-shock structure is formed consisting of the outgoing supernova shock (hereafter the forward shock) and a reverse shock, which moves backward into the ejecta \\citep[][ hereafter~\\citetalias{CF94}]{Chevalier1982a,Chevalier1982b, CF94}. The region behind the forward shock consists of shocked circumstellar material, with $T\\sim 10^8-10^9$~K, while the region behind the reverse shock, which consists of shocked ejecta, is cooler, with $T\\sim 10^7$~K.  The radio emission is caused by synchrotron emission from the forward shock \\citep[\\eg ][]{Sramek03}, while the X-ray emission originates partly from free-free and inverse Compton emission behind the forward shock, and partly from line emission in the cooler gas behind the reverse shock. Both the radio and the X-ray emission are powerful tools for understanding the physics of both the SNe and the CSM.  With the new generation of X-ray telescopes, XMM-Newton and Chandra, more detailed spectra have been obtained than ever before. This puts greater demands on the tools used for analyzing the spectra.   Traditionally, X-ray spectra have been analyzed with the XSPEC package, which includes a number of models that can be combined to get the best fit. In \\citet[][ hereafter \\citetalias{Nymark06}]{Nymark06}  we showed, however, that the SNe having the strongest X-ray emission are likely to have radiative shocks. This has not been accounted for in the analysis of the SNe observed in X-rays. Furthermore, none of the models included in the XSPEC distribution offers the possibility of fitting cooling shocks in a consistent manner.  In radiative shocks a wide range of temperatures contribute to the spectrum. The cooling also affects the hydrodynamic structure and influences the emitting volume, and thus the total luminosity from the interaction region. This is not properly accounted for  when fitting with one, or even several, single-temperature models, and the elemental abundances can be over- or under-estimated in the resulting analysis. In \\citetalias{Nymark06} we described a numerical model that computes the emission from a cooling shock self-consistently. This was applied to the coronal emission from the ring collision in SN~1987A by \\citet{Groningsson06}. Here we apply this model to the X-ray spectra of SN1993J, obtained with Chandra \\citep{Swartz03} and XMM-Newton \\citep{ZimAsch03}. ",
        "conclusions": "\\label{sect:discussion} In this paper we have investigated the X-ray spectra of SN~1993J. By applying a hydrodynamical model for a cooling shock to the X-ray observations we have discussed the fits obtained by various compositions taken from realistic ejecta models.  We find that in the spectra of SN~1993J roughly half the contribution comes from  an adiabatic shock propagating into a CNO-enriched ejecta, with the rest of the contributions from slower, radiative shocks. These could be caused by instabilities at the reverse shock front or by a clumpy CSM modifying the structure of the interaction region. Comparing ejecta models, we find that the model s15s7b2f \\citep{Woosley94} gives a somewhat better fit to the line features than the model 4H47 \\citep{NomHash88, Shigeyama94}, but the difference is hardly significant. The best fits are obtained by a composition resembling that of the circumstellar ring of SN~1987A. There is no need for a contribution from the circumstellar shock, as was also expected from the low luminosity estimated in Sect. \\ref{sect:csEmission}.   The observations discussed in the present work are close to the limit of what can be used for spectroscopic studies. If the capabilities of Chandra and XMM had been present at the early phases of this supernova, a much more detailed analysis could have been performed of the X-ray spectra. The next X-ray bright supernova within $\\sim 10$~Mpc will allow a much more detailed spectroscopic studies of the early spectra. With future missions, like ESA's, NASA's and JAXA's proposed International X-ray Observatory (IXO), these capabilities will be vastly improved. This will allow high S/N spectra to be obtained of SN at much longer distances, and will also allow detailed spectroscopic studies to be performed also at late times. This will in turn lead to a much more reliable determination of the abundances and temperature in the interaction region. We hope that the analysis in this paper will give some guidance to future X-ray observations of SNe."
    },
    "0812/0812.4043_arXiv.txt": {
        "abstract": "{We report an analysis of photometric behaviour of DI UMa, an extremely active dwarf nova. The observational campaign (completed in 2007) covers five superoutbursts and four normal outbursts.} {We examined principal parameters of the system to understand peculiarities of DI UMa, and other active cataclysmic variables.} {Based on precise photometric measurements, temporal light curve behaviour, $O-C$ analysis, and power spectrum analysis, we investigated physical parameters of the system.} {We found that the period of the supercycle now equals 31.45$\\pm$0.3 days. Observations during superoutbursts infer that the period of superhumps equals $P_{\\rm sh} = 0.055318(11)$ days ($79.66\\pm 0.02$ min). During quiescence, the light curve reveals a modulation of period $P_{\\rm orb} = 0.054579(6)$ days ($78.59 \\pm 0.01$ min), which we interpret as the orbital period of the binary system. The values obtained allowed us to determine a fractional period excess of $1.35\\% \\pm 0.02\\%$, which is surprisingly small compared to the usual value for dwarf novae (2\\%-5\\%). A detailed $O-C$ analysis was performed for two superoutbursts with the most comprehensive coverage. In both cases, we detected an increase in the superhump period with a mean rate of $\\dot P/P_{\\rm sh} = 4.4(1.0)\\times 10^{-5}$.} {Based on these measurements, we confirm that DI UMa is probably a \\emph{period bouncer}, an old system that reached its period minimum a long time ago, has a secondary that became a degenerate brown dwarf,  the entire system evolving now toward longer periods. DI UMa is an extremely interesting object because we know only one more active ER UMa star with similar characteristics (IX Dra).} ",
        "introduction": "The ER~UMa type stars (Kato and Kunjaya \\cite{kato1995})  remains one of the most intriguing systems among all cataclysmic variables. Those objects, belonging to SU UMa-type dwarf novae, exhibit various types of behaviour, such as normal outburst, superoutburst, and superhumps. Smak (\\cite{smak1984}) identified the origin of normal outbursts in dwarf novae in terms of thermal instability of the accretion disc. Whitehurst (\\cite{whitehurst1988}) and Osaki (\\cite{osaki1989}) then proposed a model that  described superoutburst phenomena in those stars (for a review see Osaki \\cite{osaki2005}). The period between two consecutive superoutbursts (also called the supercycle period)  in ER~UMa type objects is about a few tens of days. Until now, the shortest known supercycle for this class of stars has been $P_{so}=19.07$ days, the value measured for RZ~LMi (Robertson et al. \\cite{robertson1995}, Olech et al. \\cite{olech2008}).  DI UMa is a member of ER~UMa class. The history of observations of DI~UMa starts in 1959, when Hoffmeister (\\cite{hoffmeister1959}) discovered a rapidly varying star, which was later designated as DI~UMa type. Further spectroscopical studies identified this object as a cataclysmic variable (Bond \\cite{bond1978}, Szkody \\& Howell \\cite{szkody1992}). Kato, Nogami \\& Baba (\\cite{kato1996}) determined the period of the supercycle to be 25 days. Soon after, Fried et al. (\\cite{fried1999}) found that superoutbursts a have recurrence time of 30-45 days, and the superhump period and the orbital period  equal 0.05529(5) and 0.054564(1) days, respectively. Hence, they measured the superhump period excess of $1.3\\%$,  which was then one of the lowest for known CVs. Although a small deviation from regularity in the occurrence of superoutburst is admissible, this prominent change in supercycle length remains problematic from the standard thermal--tidal instability model point of view. This fact, in addition to a low period excess, that is appropriate for quiet WZ Sge stars, both with global properties and parameters characterizing this system, have encouraged us to take a closer look at the behaviour of DI~UMa. ",
        "conclusions": "Our five-month observational campaign concerning  an active SU UMa-type star DI~UMa resulted in the detection of five superoutbursts,  which repeat every 31.45 days. This is far longer than  the 25 days obtained by Kato, Nogami \\& Baba (\\cite{kato1996}) but falls within the range of 30--45 days observed by Fried et al. (\\cite{fried1999}).  Our observations indicate that DI UMa exhibits periodic, light modulations both in quiescence and during superoutbursts. The first of these modulations, with $P_{\\rm orb} = 0.054579(6)$ days ($78.59 \\pm 0.01$ min), is interpreted as the orbital period of the system and the second, with  $P_{\\rm sh} = 0.055318(11)$ days ($79.66\\pm 0.02$ min), as the superhump period.  A very orbital period and a small period excess equal to only $1.35\\% \\pm 0.02\\%$ suggest that DI UMa is a so-called \\textit{period bouncer}, i.e. an old system that reached its period minimum a long time ago, its secondary became a degenerated brown dwarf and the entire system now evolves toward longer periods (Patterson \\cite{patterson2001}). DI UMa is thus unique because we know only one more active ER UMa star with similar characteristics - namely IX Dra (Olech et al. \\cite{olech2004a}).  Inspection of Fig. \\ref{pdot} indicates that long-period systems tend to show large negative values of $\\dot P/P_{\\rm sh}$, whereas short-period systems are characterized by small and positive period derivatives. This picture has a physical interpretation proposed by Osaki and Meyer (\\cite{osaki2003}). In long-period dwarf-novae with high transfer rates, the 3:1 resonance radius in the accretion disc is close to the tidal truncation radius, and  eccentric waves in the disc may thus propagate only inwards, proding a decrease in the superhump period. In short-period systems with lower accretion rates and more infrequent outbursts, the 3:1 radius is much smaller than the tidal truncation radius and sufficient matter is stored in the disc to cause the propagation of outward eccentric waves. In these systems, the superhump period may increase.  Do the properties of DI UMa inferred from observational data agree with this scenario? It is a short-period but active dwarf nova, showing both frequent outbursts and superoutbursts, which are indicate of a high-mass transfer-rate. It may imply that the disc contains sufficient matter to reach not only the 3:1 resonance radius, but also the region of the 2:1 resonance radius, as in the case of IX Dra, another active but old dwarf nova. There is thus a possibility of the ignition of outwardly propagating eccentric waves, which may cause an increase in the superhump period.  We recall that DI UMa is the second shortest-period dwarf nova\\footnote{There are two variables, V485 Cen (Olech \\cite{olech1997}) and EI Psc (Uemura et al. \\cite{ue2002}), which show far shorter periods of superhumps but their status as classical SU UMa systems is unknown.} with a determined period derivative. Smaller values of both the superhump period and its derivative were determined for VS 0329+1250 (Shafter et al. \\cite{Sch2007})."
    },
    "0812/0812.3455_arXiv.txt": {
        "abstract": "Radio continuum emission from the supernova remnant G296.5+10.0 was observed using the Australia Telescope Compact Array. Using a 104 MHz bandwidth split into 13 $\\times$ 8 MHz spectral channels, it was possible to produce a pixel-by-pixel image of Rotation Measure (RM) across the entire remnant. A lack of correlation between RM and X-ray surface brightness reveals that the RMs originate from outside the remnant. Using this information, we will characterise the smooth component of the magnetic field within the supernova remnant and attempt to probe the magneto-ionic structure and turbulent scale sizes in the ISM and galactic halo along the line-of-sight. ",
        "introduction": "The supernova remnant G296.5+10.0 was observed using the Australia Telescope Compact Array at 1.4 GHz with an angular resolution of 30 arcseconds. Signals from the dual linear polarized feeds were combined to produce images of the remnant in linear polarisation. The Rotation Measure synthesis method (Brentjens \\& de Bruyn, 2005) was employed to obtain an RM spectrum at each pixel in the image (Figure 1). The 104 MHz bandwidth was split into 13 $\\times$ 8 MHz spectral channels, giving sufficient coverage in $\\lambda$$^2$ to remove any n$\\pi$ ambiguities in the RM values. X-ray data in the energy range 0.1--2.4 keV were taken from the ROSAT archive, via NASA's SkyView. These complimentary data allow us to estimate the electron density in the remnant. ",
        "conclusions": ""
    },
    "0812/0812.3349_arXiv.txt": {
        "abstract": "{ {We present a catalogue with the properties of all the bursts detected and localized by the IBIS instrument onboard the \\int\\ satellite from November 2002 to September 2008. The sample is composed of 56 bursts, corresponding to a rate of $\\sim 0.8$ GRB per month. Thanks to the performances of the \\int\\ Burst Alert System, 50\\% of the IBIS GRBs have detected afterglows, while 5\\% have redshift measurements. A spectral analysis of the 43 bursts in the \\int\\ public archive has been carried out using the most recent software and calibration, deriving an updated, homogeneous and accurate catalogue with the spectral features of the sample. When possible a time-resolved spectral analysis also has been carried out. The GRBs in the sample have 20-200 keV fluences in the range $5 \\times 10^{-8}$--$2.5 \\times 10^{-4}$ erg cm$^{-2}$, and peak fluxes in the range $0.11$--$56$ ph cm$^{-2}$ s$^{-1}$. While most of the spectra are well fitted by a power law with photon index $\\sim 1.6$, we found that 9 bursts are better described by a cut-off power law, resulting in E$_{p}$ values in the range 35--190 keV. Altough these results are comparable to those obtained with BAT onboard \\swi , there is marginal evidence that ISGRI detects dimmer bursts than \\swi /BAT. Using the revised spectral parameters and an updated sky exposure map that also takes into account the effects of the GRB trigger efficiency, we strengthen the evidence for a spatial correlation with the super galactic plane of the faint bursts with long spectral lag.} ",
        "introduction": "Enormous progress in the study of Gamma-Ray Bursts (GRBs) has been achieved after the results of the \\textit{Satellite per Astronomia X} (the Italian for ``Satellite for X-Ray Astronomy'') \\sax\\ , that enabled the discovery of their counterparts at lower energy \\citep{Costa1997,Vanparadijs1997,Frail1997}. Our understanding of GRBs has mainly profited from  dedicated instruments, or even satellites, specifically designed for the study of these enigmatic sources. For example, thanks to the very large GRB sample obtained with the Burst And Transient Sources Experiment (BATSE) onboard the Compton Gamma-Ray Observatory (CGRO), we learned that GRBs are isotropically distributed across the sky \\citep{briggs96}, and that their LogN-LogS is compatible with a cosmological origin \\citep{Meegan1992}. BATSE also provided accurate spectra \\citep[e.g.][]{Band1993}, and confirmed that GRBs  can be divided in two categories as a function of their duration and hardness, with about 25\\% of the GRBs being shorter than 2 s and harder compared to bursts of longer duration \\citep{Kouveliotou1993}. Quick localizations with the High Energy Transient Explorer (HETE-2) satellite allowed investigators to study the GRB emission in X-rays \\citep[e.g.][]{vanderspeck04}, and to detect the first afterglow of a short GRB \\citep{fox05}. \\swi\\ \\citep{swift}, launched in 2004, is currently localizing about 100 bursts per year, and is able to investigate the early phases of the X-ray afterglow, which were unaccessible to former instrumentation. Many key parameters are derived from the afterglow observations. Among them one of the most important is the distance, which is essential to infer the intrinsic energy and luminosity. GRBs are the most powerful explosions in the Universe, with isotropic equivalent energy $E_{iso}$ in the range from 10$^{50}$ to 10$^{54}$ erg \\citep{Amati2006}, although such values can be reduced with the hypothesis that GRBs are collimated sources \\citep{rhoads97}, as supported by the detection of achromatic breaks in some afterglow light curves. The jet opening angles inferred from the afterglow break times \\citep{Sari1999} indicate lower energy values, clustered around 10$^{51}$ erg with a reduced spread \\citep{frail2001,Bloom2003}. However, the observation of non-achromatic breaks in some afterglows, and the apparent lack of any kind of break in some others, casts doubts on this argument \\citep[see][and reference therein]{Curran2007}. In only a few years, the \\swi\\ satellite has more than doubled   the number of GRBs with known redshift, and discovered the most distant burst to date at $z \\sim$6.7 \\citep[GRB 080913,][]{Fynbo2008GCN}.   Recent progresses in GRB science concern also the identification of their progenitors. There is general consensus in considering the collapse of massive stars ($>$ 30 M$_{\\odot}$) as the origin of long GRBs \\citep{Woosley1993,Vietri1998}. Indeed, in many cases late-time bumps in the afterglow light curves of long GRBs have been reported \\citep[e.g.][]{Bloom1999, Zeh2004}, and are interpreted as an emerging supernova component. In a handful of cases the GRB/SN association has been spectroscopically confirmed \\citep[e.g.][]{Malesani2004}. However, in some cases deep optical follow-up observations failed to detect SN signatures from nearby long GRBs \\citep[e.g.][]{Dellavalle2006}, indicating the possible existence of a different class of long GRBs. Models involving the merging of two compact objects (black holes, neutron stars, or white dwarfs) are, on the other hand, preferred for short GRBs \\citep{Perna2002}.  Besides GRB-dedicated missions, also ``general purpose'' X--ray/gamma-ray satellites can significantly contribute to the GRB quest. The most remarkable example is indeed the aforementioned \\sax\\ mission, that thanks to its rapid and accurate localizations led to the discovery of afterglows at X-ray,  optical, and radio wavelengths, firmly proving the cosmological origin of these sources by allowing a measurement of their redshift.  The International Gamma-Ray Astrophysics Laboratory \\citep[\\int ,][]{integral}, launched in October 2002, carries a set of coded mask instruments dedicated to fine imaging and spectroscopy in the soft $\\gamma$-ray energy domain (15 keV--10 MeV). Even if \\int\\ is not a  mission specifically dedicated to GRBs, its imager instrument IBIS \\citep[Imager on Board the Integral Satellite,][]{ibis}, thanks to a large field of view (29$^{\\circ} \\times$29$^{\\circ}$) and a good point source location accuracy ($\\sim$ 1--2$^{\\prime}$), is well suited for GRB studies, in particular with its low energy detector plane ISGRI \\citep[\\int\\ Soft Gamma-Ray Imager,][]{isgri} (15 keV--1 MeV) . Pre-launch expectations of detecting about one GRB per month in the IBIS field of view prompted the development of the \\int\\ Burst Alert System (IBAS).  IBAS  \\citep{ibas} is a software system, running on the ground at the \\int\\ Science Data Centre \\citep[ISDC;][] {isdc}, able to detect and localize in real time the GRBs in the IBIS field of view. Thanks to IBAS, the \\int\\ satellite  has been the first to provide GRB positions with an uncertainty of only $\\sim$2-3 arcmin in near real time (less than a few tens of second after the burst trigger). To date 56 bursts have been detected in IBIS/ISGRI data\\footnote{See {\\tt http://ibas.iasf-milano.inaf.it}}, mostly automatically by IBAS, and a few in off-line searches.  In this paper we present an updated catalog of all the GRBs localized with \\int\\ , with a detailed spectral analysis of the ones for which public data are available. Some results have been published already by different authors on individual bursts, especially the ones discovered in the first months of the mission. More recently, a comprehensive analysis of a sample of 46 \\int\\ bursts has been presented by \\citet{Foley2008}, who focused mainly on their timing properties. They also extract spectra of some GRBs, but using an old version of the software and of the calibration that now we know to be not adequate when dealing with off-axis bursts. The main purpose of our work is to provide a homogeneous set of results, obtained with the latest instrument calibration and software, focusing mainly on the spectral properties. In several cases the spectral parameters we derived differ from, and supersede, those reported in previous analyses, that were based on earlier software versions and calibrations. ",
        "conclusions": ""
    },
    "0812/0812.2515_arXiv.txt": {
        "abstract": "Upper limits on the shock speeds in supernova remnants can be combined with post-shock temperatures to obtain upper limits on the ratio of cosmic ray to gas pressure ($P_{CR} / P_G$) behind the shocks.  We constrain shock speeds from proper motions and distance estimates, and we derive temperatures from X-ray spectra. The shock waves are observed as faint H$\\alpha$ filaments stretching around the Cygnus Loop supernova remnant in two epochs of the Palomar Observatory Sky Survey (POSS) separated by 39.1 years.  We measured proper motions of 18 non-radiative filaments and derived shock velocity limits based on a limit to the Cygnus Loop distance of 576 $\\pm$ 61 pc given by Blair et al. for a background star.  The PSPC instrument on-board {\\it ROSAT} observed the X-ray emission of the post-shock gas along the perimeter of the Cygnus Loop, and we measure post-shock electron temperature from spectral fits. Proper motions range from 2\\farcs 7 to 5\\farcs 4 over the POSS epochs and post-shock temperatures range from $kT \\sim 100-200 \\rm~eV$.  Our analysis suggests a cosmic ray to post-shock gas pressure consistent with zero, and in some positions $P_{CR}$ is formally smaller than zero.  We conclude that the distance to the Cygnus Loop is close to the upper limit given by the distance to the background star and that either the electron temperatures are lower than those measured from {\\it ROSAT}~PSPC X-ray spectral fits or an additional heat input for the electrons, possibly due to thermal conduction, is required. ",
        "introduction": "Supernova explosions are responsible for showering the surrounding interstellar medium (ISM) with heavy elements, influencing the distribution and composition of gas in the host galaxy.  A typical supernova releases $\\sim$10$^{51}$ ergs, sending a highly supersonic shock wave into its surroundings at speeds of up to 30,000 km~s$^{-1}$ (Ghavamian et al. 2001).  This expanding surge of energy heats and ionizes the neighboring gas out to hundreds of parsecs, initiating cloud collapse which leads to star formation and galactic evolution. Although a supernova is short-lived, the ejected remains propagate outward and interact with the ISM.  These supernova remnants (SNRs) radiate over a large spectral range, transporting information about the blast wave morphology and local ISM.  The Cygnus Loop is a well-known SNR whose distance is less than 576 $\\pm$ 61 pc based on the distance to a star whose spectrum shows absorption features from shocked gas in the remnant (Blair et al., 2008). From a global perspective, its structure is governed by interactions with the surrounding inhomogeneous ISM.  Despite its apparent spherical symmetry, both optical and X-ray observations suggest that the Cygnus Loop does not follow the simplified Sedov-Taylor morphology of an adiabatic blast wave expanding into a homogeneous, low density medium.  Instead, the blast wave deceleration is the result of encounters with dense, extended clouds distributed inhomogeneously in the pre-shocked gas (Levenson et al., 1997; Levenson et al., 1998).  We focus on the non-radiative H$\\alpha$ shock filaments, primarily in the northeastern Cygnus Loop, several of which were previously observed by Hester et al. (1994), Blair et al. (1999, 2008), Ghavamian et al. (2001) and Raymond et al. (2003).  In a non-radiative shock, the cooling timescale of the hot shocked gas is much greater than the age of the shock.  As the shocked gas has not had sufficient time to cool, post-shock radiation does not affect the shock evolution (McKee \\& Hollenbach, 1980).  Therefore, non-radiative shock filaments provide a laboratory for probing the conditions at the shock front before this information is lost by radiative or collisional processes.  When a non-radiative shock encounters neutral gas, Balmer line emission is produced by the excitation of neutral hydrogen atoms prior to ionization.  The resulting H$\\alpha$ filaments trace the outer edge of the Cygnus Loop non-radiative blast wave as it expands through low density, partially neutral gas.  These H$\\alpha$ filaments are useful for proper motion measurements of shock fronts and investigations of shocked gas parameters.  SNR shock waves are believed to accelerate cosmic rays up to energies of 10$^{15}$ eV and to be responsible for a large fraction of interstellar cosmic rays, implying an efficient process for particle acceleration.  The theory of diffusive shock acceleration describes how charged particles (dominantly protons and electrons) passing back and forth across a shock front reach high energies, resulting in a power-law spectrum $\\propto$ E$^{-2}$ for cosmic rays above $\\sim$ 10$^{9}$ GeV (Blandford \\& Eichler, 1987).  This particle acceleration process requires a collisionless medium, as frequent collisions would return the velocity profile to Maxwellian.  The efficiency of energy conversion in SNR shocks into the acceleration of cosmic rays is highly dependent on shock structure and is a topic of active research.  Theory indicates that the fraction of the energy dissipated in a shock going into cosmic rays is bistable, being either very small or $\\sim$ 80\\% (Malkov et al. 2000).  Recent calculations of the evolution of SNR shocks show that the efficiency starts off low and evolves gradually to somewhat smaller values than that predicted by steady shock models (Caprioli et al. 2008). Synchrotron emission in the radio and X-ray bands provides direct observation of particle acceleration in SNRs; however, the majority of the SNR energy going into particle acceleration is claimed by protons. Our current observational ability to detect $\\gamma$-rays (and neutrinos) is limited, so we must settle for inferences of the total cosmic ray acceleration efficiency through indirect methods.  Warren et al. (2005) inferred shock compression ratios greater than 4 from the distance between the blast wave and the contact discontinuity in Tycho's SNR.  Their results suggest that the cosmic ray to gas pressure ratio, $P_{CR}/P_G > 1$, though projection effects might allow a smaller compression ratio (Cassam-Chenai et al. 2008).  High cosmic ray acceleration efficiencies have also been inferred from the low electron temperature in 1E102-72.6 (Hughes et al. 2000) and from shock wave precursors to Balmer line filaments (Smith et al. 1994; Hester et al. 1994).  In this paper we find upper limits to $P_{CR}/P_{G}$ by combining measurements of post-shock electron temperature with the shock speed given by proper motions and distance upper limits.  Our sample includes 18 non-radiative H$\\alpha$ shock segments dominantly in the northeastern rim of the Cygnus Loop.  This sample, although limited by the uncertainties in proper motion, temperature, and distance, provides constraints on $P_{CR}/P_{G}$.  Several positions show apparently negative values of $P_{CR}$, and we discuss the uncertainties in temperature and distance and the possibility of additional electron heating that could account for that non-physical result. ",
        "conclusions": "Inspection of the derived cosmic ray to gas pressure equation emphasizes the importance of constraining proper motion, distance, and temperature to infer tight upper limits on $P_{CR}/P_{G}$. $P_{CR}/P_{G}$ scales with the square of the shock speed, given by the product of proper motion and distance.  In general, the upper limits indicate a small value for the ratio, and in some cases it is formally negative.  In the context of the discussions of the uncertainties given above, it is clear that the proper motion measurements cannot account for the unphysical values of $P_{CR}/P_{G}$.  Most distance estimates are significantly smaller than the 637 pc we have adopted as the nominal upper limit, and it seems like a strain to make the distance large enough ($\\sim 1$ kpc) to make all the upper limits to $P_{CR}/P_{G}$ positive.  However, it may be possible given the general uncertainties in SdO star distances.  Therefore, we conclude that the heart of uncertainty likely lies in post-shock temperature measurements and electron-ion thermal equilibration assumptions.  Assuming ionization equilibrium and that shock velocity measurements are reasonably accurate, we require lower post-shock electron temperatures than measured from our {\\it ROSAT}~PSPC X-ray spectral fits.  The fits based on the Raymond \\& Smith (1977) code give temperatures about 10\\% lower than those using APEC, but that is several times too small to explain the discepancy.  \\subsection{Heating by Thermal Conduction} There is another possible way out of the discrepancy.  We have assumed that $T_e = T_i$ behind the shock, and it is possible that $T_e > T_i$.  However, such behavior is not seen in shocks in the interplanetary medium, and there is no obvious physical reason for such electron heating in the shock.  Ghavamian et al. (2006) describe a wave heating mechanism that would produce particle temperatures somewhat higher than observed here, but $T_p$ from Ghavamian et al. (2001) and a higher $T_e$ would then require a large value of $V_s$.  An interesting alternative is electron heating by thermal conduction from gas farther behind the shock.  A number of papers have explored the global structure of SNRs with thermal conduction (e.g., Slavin \\& Cox 1992; Cox et al. 1999; Shelton et al. 1999).  We can obtain an observational estimate from X-ray observations.  Nemes et al. (2008) show the variation of $T_e$ with radius in the northeastern Cygnus Loop based on XMM spectra, with a gradient of 0.05 keV over a distance of 0.07 times the shock radius, or  $\\nabla T \\sim 1.8 \\times 10^{-13}~\\rm K~cm^{-1}$. This implies a volumetric heating rate of about $3.4 \\times 10^{-22}~ \\rm erg~cm^{-3}~s^{-1}$, which could heat the gas by about $3 \\times 10^5$ K as it travels over the 100\\arcsec\\ region we analyze.  Thus the true post-shock temperature could be somewhat lower than the X-ray temperatures we measure by approximately the amount required for $P_{CR}$ greater than zero.  The above estimate requires a temperature gradient parallel to the magnetic field, so it could not be correct everywhere.  However, only 7 of the 18 positions show $P_{CR}$ nominally smaller than zero."
    },
    "0812/0812.0383_arXiv.txt": {
        "abstract": "We use a deep \\chandra\\ observation to examine the structure of the hot intra--group medium of the compact group of galaxies Stephan's Quintet. The group is thought to be undergoing a strong dynamical interaction as an interloper, NGC~7318b, passes through the group core at $\\sim850$\\kmps. Previous studies have interpreted a bright ridge of X--ray and radio continuum emission as the result of shock heating, with support from observations at other wavelengths. We find that gas in this ridge has a similar temperature ($\\sim0.6$~keV) and abundance ($\\sim0.3$\\Zsol) to the surrounding diffuse emission, and that a hard emission component is consistent with that expected from high--mass X--ray binaries associated with star--formation in the ridge. The cooling rate of gas in the ridge is consistent with the current star formation rate, suggesting that radiative cooling is driving the observed star formation.  The lack of a high--temperature gas component is used to place constraints on the nature of the interaction and shock, and we find that an oblique shock heating a pre--existing filament of \\Hi\\ may be the most likely explanation of the X--ray gas in the ridge. The mass of hot gas in the group is roughly equal to the deficit in observed \\Hi\\ mass compared to predictions, but only $\\sim2\\%$ of the gas is contained in the ridge. The hot gas component is too extended to have been heated by the current interaction, strongly suggesting that it must have been heated during previous dynamical encounters. ",
        "introduction": "Stephan's Quintet, also known as \\object[HCG92]{HCG~92} \\citep[][hereafter SQ]{Hickson82}, was discovered roughly 130 years ago \\citep{Stephan1877} and is perhaps the most extensively studied compact galaxy group. Several factors contribute to the interest shown in the system. Of the original five galaxies identified as part of the group, one (\\object[NGC7320]{NGC~7320}) has a discordant redshift and is now recognised as a superimposed foreground object; attempts to determine the true distances of the various members motivated much early work. Among the remaining galaxies, three (\\object[NGC7317]{NGC~7317}, \\object[NGC7318a]{NGC~7318a} and \\object[NGC7319]{NGC~7319}) form a central kernel with similar recession velocities, while the fifth (\\object[NGC7318b]{NGC~7318b}) appears to be passing through the group core with a velocity of $\\sim$850\\kmps. The group shows signs of complex past tidal interactions, including stellar and \\Hi\\ tidal tails \\citep{Arp73,Shostaketal84,Sulenticetal01,Williamsetal02} hosting ongoing star formation \\citep{Xuetal99,Gallagheretal01,MendesdeOliveiraetal04,Xuetal05}, and possible tidal dwarf galaxies \\citep{Xuetal03}. These features were probably formed during one or more passages through the group by the nearby \\object[NGC7320c]{NGC~7320c} \\citep{Shostaketal84,Molesetal97}.  However, perhaps the most interesting aspect is the evidence of ongoing interactions as NGC7318b passes through the group. 1.4~GHz radio continuum observations reveal a ridge of emission lying along the eastern edge of NGC~7318b \\citep{AllenHartsuiker72,VanderhulstRots81} which has been shown to correspond to similar features in \\Ha\\ \\citep{Sulenticetal01}, UV \\citep{Xuetal05}, H$_2$ \\citep{Appletonetal06}, 15$\\mu$m IR \\citep{Xuetal99} and X-ray \\citep{Pietschetal97} emission. This structure has been widely interpreted as a shock caused by the collision of NGC~7318b with a pre-existing IGM or tidally stripped material \\citep[e.g.,][]{Sulenticetal01,Xuetal03}. As such, it may serve as an example of the process by which the X-ray emitting IGM seen in more evolved groups is formed.  SQ has been observed previously by both \\chandra\\ and \\xmm\\ \\citep{Trinchierietal03,Trinchierietal05}. Unfortunately the IGM is not particularly bright in X-rays ($\\sim$2$\\times$10$^{41}$\\ergps\\ for the shock/ridge region), the \\chandra\\ pointing was relatively short ($<$20~ks) and the group is quite compact compared to the \\xmms\\ field of view and imaging resolution. Spectral fitting of the ridge suggested a multi-temperature model was needed but did not strongly constrain the shock properties. The data did show the presence of substructures within the ridge, and indicated the possibility of temperature differences between these structures.  In this paper we combine a new, deeper \\chandra\\ exposure with the previous short observation to provide stronger constraints on the properties of the ridge and associated shock, and the surrounding IGM. SQ is an exceptionally complex system, with X-ray emission arising from many other sources, including the individual galaxies, tidal features and star forming regions.  We will discuss these in detail and compare them with results from our own observations in other wavebands, which will in turn be described in detail in future papers. ",
        "conclusions": "Stephan's Quintet provides a rare opportunity to observe a galaxy group in the process of development from an X--ray faint spiral--rich system to a more evolved, elliptical--dominated, X--ray bright state. The possibility that a significant fraction of the cold gas in the group is being heated by the ongoing dynamical interactions has an obvious importance for our understanding of the origins of the hot gas halos of X-ray luminous groups. We have analysed a deep, $\\sim$95~ks \\chandra\\ observation of SQ with the goal of placing strong constraints on the nature of the interactions in the system and improving our understanding of the origin of the emission ridge. Our results can be summarized as follows:  \\begin{enumerate}  \\item The X--ray emission in the ridge is well described by a plasma of temperature 0.6~keV and abundance 0.3\\Zsol\\ plasma, with a powerlaw component consistent with that expected from HMXBs associated with the star formation in this region. The fitted temperature and abundance are very similar to the values found for the diffuse emission surrounding the ridge. This strongly suggests that the gas cannot be the product of a normal shock driven by the interaction with NGC~7318b. An oblique shock with an angle of $\\sim30$\\degree\\ could produce the temperatures we observe if the pre--shock material were \\Hi\\ or $\\sim0.4$~keV gas. This angle matches that of the disk of NGC~7318b, suggesting that the shock heating of an \\Hi\\ filament is a viable mechanism to produce much of the X--ray gas observed in the emission ridge. However, for 0.4~keV or hotter gas, the opening angle of the shock would be wide, and the interactions with the surrounding, pre--existing X--ray halo must at some point produce a normal shock. It is therefore unclear why no high--temperature plasma component, associated with shocking of the pre--existing hot gas, is observed.  \\item The total mass of X--ray emitting gas in SQ is similar to the deficit in \\Hi\\ compared to expected values. It is therefore possible that the hot gas component of the group halo arises entirely from shock heating of \\Hi. However, the extent of X-ray emission is large compared to the expected rate of expansion of a shock heated filament of \\Hi. This strongly suggests that a hot intra--group medium was already present in the group before the current interaction. This could have arisen during previous dynamical interactions, for which there is plentiful evidence, in particular multiple tidal tails pointing roughly toward the nearby galaxy NGC~7320c. Alternatively, the hot gas component could be primordial material shock heated during its accretion into the group. The expected Virial temperature of SQ, based on the group velocity dispersion excluding NGC~7318b, is similar to the observed gas temperature outside the group core, though with large uncertainties owing to the small number of galaxies involved and the dynamical state of the group.  \\item Some of the brightest X-ray and radio emission regions in the ridge appear to be spatially correlated with star forming regions, identified by UV emission. The southern knot and the point at which the eastern spiral arm of NGC~7318b turns south are both bright in X--ray and radio. However, outside these regions there are significant disagreements between radio and X--ray distributions. This supports the idea that a combination of emission mechanisms is required to produced the structures observed.  \\item The cooling rate of hot gas in the ridge is consistent with the current star formation rate, suggesting that star formation in this region may be largely driven by radiative cooling. The mass of metals expected from the predicted number of supernovae also provides a good match to the observed gas abundance. However the soft X--ray luminosity expected from supernovae and star formation is smaller than that observed by a factor $\\sim10$ and the energy available from supernovae associated with star formation is also too small to support the emission observed at various wavelengths. The energy available from shock heating is probably sufficient to produce the observed luminosities.  \\item Efficient energy loss from the hot gas via collisional heating of dust appears unlikely to be the dominant source of cooling at present, since it would be expected to introduce strong temperature gradients which are not observed in the temperature map. Assuming that an \\Hi\\ filament has been shocked, the dust associated with this filament would now be in the emission ridge, and would most effectively cool the gas in the densest parts of the ridge. As the ridge has a similar (or slightly warmer) temperature to its surroundings, a coincidence would be required for us to observe the gas at just this point in its cooling. Dust may have been a more important factor in cooling the gas immediately after the passage of the shock, but must have been rapidly destroyed by sputtering.  \\item It seems unlikely that our measured temperatures are being biased by the presence of non-equilibrium ionisation gas. Testing suggests that a non-equilibrium plasma capable of producing the temperature and abundance we observe would be very poorly described by an APEC model and would have to have been shocked within the last $\\sim10^5$~yr. By $10^6$~yr after shocking we would expect spectral fitting to accurately reproduce the plasma temperature. We therefore conclude that our measured temperatures are correct and that the great majority of the gas is in ionisation equilibrium.   \\end{enumerate}  In general, we conclude that the source of much of the X--ray gas in the ridge, and the majority of the energy powering the multi--wavelength emission from this region, arises from shocks driven by the infall of NGC~7318b. However, star formation plays an important role in determining the exact morphology of the ridge and is responsible for the majority of the hard emission via HMXBs. Considering the ridge in isolation, it appears possible that \\Hi\\ in a pre--existing filament and in the disk of NGC~7318b has been shock heated by the interaction, with additional heating and enrichment by supernovae and stellar winds. However, this scenario cannot explain the larger X--ray halo of SQ, which may be evidence of shock--heating during previous dynamical interactions."
    },
    "0812/0812.0660_arXiv.txt": {
        "abstract": "$f(R)$ gravity, capable of driving the late-time acceleration of the universe, is emerging as a promising alternative to dark energy. Various $f(R)$ gravity models have been intensively tested against probes of the expansion history, including type Ia supernovae (SNIa), the cosmic microwave background (CMB) and baryon acoustic oscillations (BAO). In this paper we propose to use the statistical lens sample from Sloan Digital Sky Survey Quasar Lens Search Data Release 3 (SQLS DR3) to constrain $f(R)$ gravity models. This sample can probe the expansion history up to $z\\sim2.2$, higher than what probed by current SNIa and BAO data. We adopt a typical parameterization of the form $f(R)=R-\\alpha H^2_0(-\\frac{R}{H^2_0})^\\beta$ with $\\alpha$ and $\\beta$ constants. For $\\beta=0$ ($\\Lambda$CDM), we obtain the best-fit value of the parameter $\\alpha=-4.193$, for which the 95\\% confidence interval that is [-4.633, -3.754]. This best-fit value of $\\alpha$ corresponds to the matter density parameter $\\Omega_{m0}=0.301$, consistent with constraints from other probes. Allowing $\\beta$ to be free, the best-fit parameters are $(\\alpha, \\beta)=(-3.777, 0.06195)$. Consequently, we give $\\Omega_{m0}=0.285$ and the deceleration parameter $q_0=-0.544$. At the 95\\% confidence level, $\\alpha$ and $\\beta$ are constrained to [-4.67, -2.89] and [-0.078, 0.202] respectively. Clearly, given the currently limited sample size, we can only constrain $\\beta$ within the accuracy of $\\Delta\\beta\\sim 0.1$ and thus can not distinguish between $\\Lambda$CDM and $f(R)$ gravity with high significance, and actually, the former lies in the 68\\% confidence contour. We expect that the extension of the SQLS DR3 lens sample to the SDSS DR5 and SDSS-II will make constraints on the model more stringent.  \\end {abstract} \\begin{keywords} cosmology: theory -- csomological parameters -- gravitational lensing -- dark matter \\end{keywords} ",
        "introduction": "In modern cosmology, one of the most striking discoveries is that our expanding universe is undergoing a phase of acceleration. The key observational results that support this discovery are: the luminosity-redshift relationship from SNIa surveys \\citep{Astier2006, Perlmutter1999, Riess1998, Riess2004, Riess2007}, the CMB anisotropy spectrum \\citep{de Bernardis2000, Spergel2003, Spergel2007, Hinshaw2007}, the large scale structure from galaxy redshift surveys \\citep{Cole2005, Seljak2005, Tegmark2004a} and BAO \\citep{Eisenstein2005} . With the combination of all these observational results, one gets the standard $\\Lambda$CDM cosmology as: $\\Omega_K=0,\\Omega_M=0.27,\\Omega_\\Lambda=0.73$.  Amongst the matter content baryonic matter amounts to only $4\\%$, other $23\\%$ is the so-called dark matter (DM), and the rest component of $73\\%$ dominate the universe that is often referred to as dark energy (DE), which is a negative-pressure ideal fluid smoothly permeated the universe. It is worth noting that, notwithstanding the standard cosmology is very well in agreement with the astrophysical data, the nature of the dark matter and dark energy remains mystery in the modern Cosmology and Physics, so that there are so many theoretical proposals to account for them on the ground. Possible models to explain the acceleration include a classical cosmology constant \\citep{Weinberg1989, Carroll1992}, Chaplygin gas \\citep{kamenshchik2001,zhu2004,wuyu2007a,wuyu2007b}, a wide variety of scalar-field models such as Quintessence \\cite[e.g.,][]{Caldwell1998,wu2008}, K-essence \\citep{Armendariz-Picon2000, Armendariz-Picon2001}, Phantom field \\citep{Caldwell2002, Dabrowski2003} and so on.  Except above models, there are several popular alternative ideas for the accelerating universe by modifying general relativity rather than resorting to some kinds of exotic fluid. It is a natural consequence to modify the theory of gravity, after the current observations mightily suggesting the failure of General Relativity (GR) as a large cosmological scale gravity theory. Actually, on the right side of Einstein's field equation, the energy momentum tensor is related to the matter content, and on the left side,the Einstein tensor consists of pure geometrical terms, which can be modified to account for DM and DE phenomena \\citep{Copeland2006}. Accordingly, both of accelerating behaviour (DE) and dynamical phenomena (DM) can be interpreted as curvature effects. Recently, a lot of people attempt to modify the General Relativity in different ways and propose various modified gravity theories to address the DM or DE phenomenon or both. Amongst them, $f(R)$ gravity theories, where $R$ is curvature scalar, have attracted much attention, which are put forward to account for DE by adopting an arbitrary function of R in Einstein-Hilbert Lagrangian, the so-called $f(R)$ term, instead of $R$ in traditional General Relativity. Obviously, an important reason for this interest is that $f(R)$ term in types of negative as well as positive powers of R can produce the inflation at early times and now observed acceleration phase at late times, following the well-known sequence of universal evolution \\citep{Nojiri2003, Nojiri2007, Sotiriou2006-2, Capozziello2006-1, Vollick2003, Chiba2003, Erickcek2006, Navarro2007, Olmo2005, Olmo2007, Amendola2007, Li B.2007-1, Li B.2006}. Usually, two possibilities referring to as the metric approach and the Palatini approach can derive the Einstein field equation from the Hilbert action. By comparison, the metric approach assumes that the equation has only one variable with respect to metric where the affine connection is the function of metric, while in the Palatini approach the metric and the connection are treated to be independent of each other as two variables. When $f(R)$ is in form of linear function, both approaches lead to the same results.  Once adding any non-linear term to the Hilbert action, however, the two methods would produce enormous differences. The reasons for the Palatini approach seeming appealing compared with another one are that, on one hand, the metric approach leads to the fourth order field equations, which is difficult to solve analytically, whereas the Palatini approach results in the two order field equations. On the other hand, although  $f(R)$ theories via metric approach can give some interesting and successful results, some of them suffer from certain fatal defects. For example, they cannot pass the solar system test, have instabilities, have incorrect Newtonian limit and cannot totally describe all epoches of the universe.  Here we adopt the $f(R)$ gravity theory within the Palatini approach which can avoid the above mentioned problems. In this framework, the form $f(R)=R+\\alpha(-R)^\\beta$ is chosen so that it can pass the solar system test and has the correct Newtonian limit \\citep{Sotiriou2006-3}, and can explain the late accelerating phase in the universe \\citep{Fay2007, Capozziello2005, Sotiriou2006-1}. Furthermore, it is important to go beyond the quantitative studying and test these theories using the observations. Recently constraint from data on above type of $f(R)$ theory is intensively discussed by many authors. Among them, in these papers \\citep{Santos2008, Fay2007, Borowiec2006, Amarzguioui2006, Sotiriou2006-1}, the CMB shift parameter, supernovae Ia surveys data and baryon acoutic oscillations were combined to constrain the parameters and, in particular, they give $\\beta\\sim 10^{-1}$. In a previous work, \\citet{Koivisto2006} used the matter power spectrum from the SDSS to get further restriction $\\beta\\sim 10^{-5}$, which reduced allowed parameter space to a tiny around the $\\Lambda$CDM cosmology. Recently, \\citet{Li B.2007-2} jointed the WMAP, supernovae Legacy Survey (SNLS) data and Sloan Digital Sky Survey (SDSS) data to tighten the parameter up to $\\beta\\sim 10^{-6}$, which made this model hard to distinguish from the standard one, where the parameter $\\beta$ is zero. So far, there has been no attempt to use the strong lensing observation to test the $f(R)$ gravity theories in the Palatini formalism.  Nowadays, with more and more lens surveys available to enlarge the statistical lensing samples, gravitational lensing has developed into a powerful tool to study a host of important subjects on different scales in astrophysics, from stars to galaxies and clusters, further, to the large structure of the universe. Since the gravitational lensing phenomena involve the source information, two dimensional mass distribution of the lens and the geometry of the universe, it is useful to not only infer the distant source properties far below the resolution limit or sensitivity limit of current observations and offer an ideal way to probe the mass distribution of the universe, but also constrain the parameters of cosmological models \\citep{Copeland2006, Wu1996}. So far the large lens surveys have built both radio (the Cosmic-Lens All Sky Survey, CLASS) and optical (SQLS DR3) lens samples. The CLASS forms a well-defined lens sample containing 13 lenses from 8958 radio sources with image separations of $0.3^{''}<\\theta<3^{''}$ and the i-band flux ratio $q\\leq10$ (bright to faint) \\citep{Browne2002}. Meanwhile, the SQLS constructs a statistical sample of 11 lensed quasars from 22683 optical quasars, of which the range of redshift is $0.6<z<2.2$ and the apparent magnitude is brighter than $i=19.1$ \\citep{Inada2008}. As compared with CLASS, the latter sample has larger image separations ($1^{''}<\\theta<20^{''}$) and smaller flux ratio limit ($q_c=10^{-0.5}$, faint to bright). Many authors use the strong gravitational lensing statistics to study DE, including the equation of state of DE \\citep{Chae2002, Chen2004, Oguri2008} and the test for the modified gravity theories as alternatives to DE \\citep{Zhu2008}.  The main goal of this paper is to investigate the constraints of strong lensing observation on the parameters of $f(R)$ theory in the Palatini approach using the SQLS statistical sample. But there has a big question that how the non-linear terms in $f(R)$ gravity theories contribute to the light-bending and correct the Einstein deflection angle. If this influence is large enough, it is inevitable that all the models built on GR to compute the lensing statistical probability must be changed. But this question is starting to be solved, here, we still use the same way as that based on GR to calculate the lensing probability because of the following considerations. Firstly, the $f(R)$ gravity theories are applied to explain the nature of DE, but as far as we know it, DE smoothly fills the space to power the acceleration expansion as the same way in everywhere of the universe and make the universe to be flat. \\citet{Dave2002} showed that DE does not cluster on scales less than 100Mpc. Naturally, as an alternative to DE, any modified gravity theory, like $f(R)$ non-linear term, should not affect the gravitational potential distribution of dark halos as the virialized systems such as galaxies and clusters, and in particular, should have not detectable contribution to the light-bending if they are some successful models. Moreover, it is worth stressing that how the modifying the lagrangian of the gravitational field affects the standard theory of lensing (built on GR) is not well investigated. Recently, theories with $f(R)\\propto R^n$ were been investigated within the metric approach for a point-like lens and results were given that the $R^n$ modified gravity signatures possibly to be detected through a careful examination of galactic microlensing \\citep{Capozziello2006-2}. And \\citet{Zhang2007} derived the complete set of the linearized field equations of the two Newtonian potentials $\\phi$ and $\\psi$ in the metric formalism, and predicted that, for some $f(R)$ gravity models, the gravitational lensing was virtually identical to that based on GR, under the environments of galaxies and clusters. While, under the Palatini approach, the paper \\citep{Ruggiero2008} is a first step to evaluate some effects of the non linearity of the gravity Lagrangian on lensing phenomenology. As expected, for a spherically symmetrical point-like model, the estimated results suggest that the effects of $f(R)$ are confined around a cosmological scale, and hence, they have no effects on smaller scales such as our Galaxy. For galaxies, the weak-field approximation is valid. The standard lensing theory (based on GR) is on the basis of this assumption. In GR, the deflection angle for a point mass can be easily obtained. For the weak-field assumption, the field equations of GR can be linearized. Therefore, we can first divide the whole galaxy into small elements, and each mass element can be treated as point mass, then the deflection angle of the general galactic mass distribution models can be obtained by the sum (or integration) of the deflections due to the individual mass components \\citep{Kochanek2004}. In $f(R)$ gravity, if the gravitational field is weak, we are able to perturb the General Einstein field equations (based on $f(R)$ gravity theories) and simplify them to linear equations. In this case, the superimposition principle works, so we can get that the field equations which describe the point-like objects can also be used to describe the ensemble of mass points. That is to say,  both of the mass point and the ensemble of mass points have field equations in the same formalism. So we can safely extrapolate the conclusion for point-like lenses, that standard lensing theory is a good approximation in the case of $f(R)$ gravity theories, to galaxy-scale lenses. However, there are no works to prove it through mathematical process, and further theoretical studies are needed to test the $f(R)$ gravity theories in the Palatini formalism. Under these situations, we consider that the standard theory of lensing is still correct in the caseof $f(R)$ theory we used.  According to astronomical observations, dynamical analysis and numerical simulations, there are three kinds of popular mass density profile of dark halos as a lens model in the standard lensing theory, namely, the singular isothermal sphere (SIS), the Navarro-Frenk-White (NFW) and the generalized NFW (GNFW) profile. They can reasonably reproduce the results of strong lensing survey \\citep{Li2002, Li2003, Sarbu2001}. As been well known, the lensing probability is very sensitive to the density profile of lenses. Moreover, the lensing is determined almost entirely by the fraction of the halo mass that is contained within a fiducial radius that is $\\sim\\%4$ of the virial radius. Since in the inner regions the slope of SIS density is steeper than that of NFW density, for the multiple image separations of $\\Delta\\theta<5{''}$, the lensing probability for NFW halos is lower than the corresponding probability for SIS halos by about 3 orders of magnitude \\citep{Li2002, Li2003, Keeton2001-2, Wyithe2001}. The cooling mass dividing galaxies and clusters for transition from SIS to NFW is $M_c\\sim10^{13}M_\\odot$. The SQLS DR3 has 11 multi-imaged lens systems, including ten galaxy-scale lenses and one cluster-scale lens with large image separation ($14.62{''}$) \\citep{Inada2008}. It is the first lens sample that contains both galaxy-scale lenses and cluster-scale lens. In lensing statistics, if we consider lensing halos with all mass scales, from galaxies to clusters,  then the effects of substructures \\citep{Oguri2006} and baryon infall \\citep{kochanek2001,keeton2001-1} cannot be neglected, in particular to galaxy groups, for which neither SIS nor NFW can be applied. For early-type galaxies, which dominate the galactic strong lensing, SIS is a good model of lenses, in particular, it results from the original NFW-like halos through the baryon infall effects \\citep{rusin2005,koopmans2006}. Moreover, the non-spherical lens profiles do not significantly affect the lensing cross section \\citep{Oguri2005, Huterer2005, Kochanek1996}. Therefore, in this paper, we use the SIS model to compute the strong lensing probability for galaxies. To avoid the complexity from large mass scales mentioned above, we consider only galaxy-scale lenses in the sample, and thus compute the differential lensing probability rather than the usual integrated lensing probability to match the SQLS DR3 sample (containing 10 galaxy-scale lenses).  A more realistic lensing model including SIS and NFW to fit the complete SQLS sample and comparing to the results of this paper is our future work. As comparison, \\citet{Oguri2008} adds two additional cuts to select a 7 lensed quasar sub-sample from the DR3 statistical lens sample for constraining the parameter $\\omega$ of state of dark energy.  The rest of the paper is organized as follows: In section 2 we outline the $f(R)$ theories in the Palatini approach, including the cosmological dynamical equation and parameters with the FRW universe setting.  In section 3 we present the differential lensing probability based on the $f(R)$ gravity theories with SIS model. In section 4, we discuss the specific analytic function $f(R)=R-\\alpha H^2_0(-\\frac{R}{H^2_0})^\\beta$ used in this paper, give the corresponding cosmological evolution behaviour and investigate the observational constraint on our $f(R)$ theory arising from SQLS DR3 statistical lens sample. Finally, we give discussion and conclusions in section 5.  \\section[]{the $f(R)$ gravity theories and corresponding cosmology parameters}  \\subsection{The generalized Einstein Equations}  The starting point of the $f(R)$ theories in the Palatini approach is Einstein-Hilbert action, whose formalism is: \\begin{equation} S=-\\frac{1}{2\\kappa}\\int{d^4x\\sqrt{-g}f(R)}+S_M, \\label{action} \\end{equation} where $\\kappa=8\\pi G$, light velocity $c=1$ and $S_M$ is the matter action which is a functional of metric $g_{\\mu\\nu}$ and the matter fields $\\psi_M$. As mentioned previously, the metric and the affine connection are two independent quantities in the Palatini approach, so the Ricci scalar is defined by two variables, the metric and the connection \\cite[see][for details]{Vollick2003}: \\begin{equation} R=g^{\\mu\\nu}\\hat{R}_{\\mu\\nu}, \\end{equation} in which \\begin{equation} \\hat{R}_{\\mu\\nu}={\\hat{\\Gamma}^\\alpha}_{\\mu\\nu,\\alpha}-{\\hat{\\Gamma}^\\alpha}_{ \\mu\\alpha,\\nu}+{\\hat{\\Gamma}^\\alpha}_{\\alpha\\lambda}{\\hat{\\Gamma}^\\lambda}_{ \\mu\\nu}-{\\hat{\\Gamma}^\\alpha}_{\\mu\\lambda}{\\hat{\\Gamma}^\\lambda}_{\\alpha\\nu}, \\end{equation} where R is negative. From Eq.(\\ref{action}), we derive the generalized Einstein equations. Varying with respect to metric $g_\\mu\\nu$, we get one equation: \\begin{equation} f'(R)\\hat{R}_{\\mu\\nu}-\\frac{1}{2}g_{\\mu\\nu}f(R)=-\\kappa T_{\\mu\\nu}, \\label{equation} \\end{equation} where $f'(R)=df/dR$ and the energy momentum tensor $T_{\\mu\\nu}$ is given as \\begin{equation} T_{\\mu\\nu}=-\\frac{2}{\\sqrt{-g}}\\frac{\\delta S_M}{\\delta g^{\\mu\\nu}}. \\end{equation} Taking the trace of Eq.(\\ref{equation}) gives the so-called structural equation and it controls the solutions of Eq.(\\ref{equation}). \\begin{equation} Rf'(R)-2f(R)=-\\kappa T. \\label{structure} \\end{equation} Varying with respect to the connection ${\\hat{\\Gamma^\\lambda}_{\\mu\\nu}}$ gives another equation \\begin{equation} \\hat{\\nabla}_\\alpha[f'(R)\\sqrt{-g}g^{\\mu\\nu}]=0, \\end{equation} where $\\nabla$ denotes the covariant derivative with respect to the affine connection. From this equation it is found that the new metric $h_{\\mu\\nu}$, which can describe the affine connection as the Levi-Civita connections, is conformal to $g_{\\mu\\nu}$, \\begin{equation} h_{\\mu\\nu}=f'(R)g_{\\mu\\nu}. \\end{equation} Finally combined above equations, we could easily get the relation between $\\hat{R}_{\\mu\\nu}$ and ${R}_{\\mu\\nu}$ which is the Ricci tensor associated with the metric $g_{\\mu\\nu}$. \\begin{equation} \\hat{R}_{\\mu\\nu}=R_{\\mu\\nu}-\\frac{3}{2}\\frac{\\nabla_\\mu f'\\nabla_\\nu f'}{f'^2}+\\frac{\\nabla_\\mu\\nabla_\\nu f'}{f'}+\\frac{1}{2}g_{\\mu\\nu}\\frac{\\nabla^\\mu\\nabla_\\mu f'}{f'}. \\label{ricci} \\end{equation} Note that the covariant derivative $\\nabla$ we refer to is associated with the Levi-Civita connection of metric $g_{\\mu\\nu}$.  \\subsection{The Background Cosmology} In this subsection, we shall make a detailed study of the cosmological viability of the model based on $f(R)$ theories in Palatini formalism. It is well-known that WMAP's data is a strong evidence to support a flat universe, so we choose a spatially flat FRW (Friedmann-Robertson-Walker) metric to describe our background universe. The metric takes the standard form: \\begin{equation} ds^2=-dt^2+a(t)^2\\delta_{ij}dx^idx^j             (i,j=1,2,3), \\end{equation} where $a(t)$ is the scale factor. As usual, we assume a perfect fluid energy-momentum tensor ${T_\\mu}^\\nu=diag(\\rho,p,p,p)$. By contracting Eq.(\\ref{ricci}) we get the generalized Friedmann equation: \\begin{equation} (H+\\frac{1}{2}\\frac{\\dot{f'}}{f'})^2=\\frac{1}{6}\\frac{\\kappa(\\rho+3p)}{f'}-\\frac {1}{6}\\frac{f}{f'}. \\label{friedmann} \\end{equation} In this paper we consider a matter dominated universe, so the constant equation of state is $p=\\omega\\rho(\\omega=0)$ and the relation between the matter density and scalar factor is $\\rho=\\rho_0R^{-3(\\omega+1)}(\\omega=0)$. Using Eq.(\\ref{structure}),  we can obtain: \\begin{equation} a(R)=(\\kappa\\rho_{m0})^\\frac{1}{3}(Rf'-2f)^{-\\frac{1}{3}}, \\label{ar} \\end{equation} where we have chosen $a_0=1$. On the other hand, from Eq.(\\ref{structure}) and the conservation of energy $T_{\\mu\\nu;\\lambda}=0$, we have as: \\begin{equation} \\dot{R}=-\\frac{3H\\rho_M}{Rf''(R)-f'(R)}. \\label{dotr} \\end{equation} Using Eqs.(\\ref{structure}), (\\ref{friedmann}) and (\\ref{dotr}) we get the Friedmann equation $H(R)$ in the form of R, \\begin{equation} H^2(R)=\\frac{1}{6f'}\\frac{Rf'-3f}{(1-\\frac{3}{2}\\frac{f''(Rf'-2f)}{f'(Rf''-f')} )^2}. \\label{h2r} \\end{equation} Combing Eqs.(\\ref{ar}) and (\\ref{h2r}) we can know the whole expansion history that is determined by $H(a)$, for any specific expression of $f(R)$.  Let us now consider the cosmological distance based on the given model. From the explicit expression for the Hubble parameter $H(R)$ and the relation between z and $a(R)$: $1+z=a^{-1}(R)$, it is convenient to rewrite the proper distance, luminosity distance, angular diameter distance and the deceleration parameter in terms of R. The proper distance and luminosity distance to the object at redshift $z$ are \\begin{eqnarray} D^P(z) &=&\\int^z_0{\\frac{dz}{(1+z)H(z)}} \\\\ \\nonumber &=&\\frac{1}{3}\\int^{R_z}_{R_0}{\\frac{Rf''-f'}{Rf'-2f}\\frac{dR}{H(R)}}=D^P(R), \\end{eqnarray} \\begin{eqnarray} D^L(z)&=&(1+z)\\int^z_0{\\frac{dz}{H(z)}} \\\\ \\nonumber &=&\\frac{1}{3}(k\\rho_0)^{-\\frac{2}{3}}(Rf'-2f)^{\\frac{1}{3}}\\int^{R_z}_{R_0}{ \\frac{Rf''-f'}{(Rf'-2f)^\\frac{2}{3}}\\frac{dR}{H(R)}} \\\\ \\nonumber &=&D^L(R). \\end{eqnarray} The angular diameter distance from an object at red-shift $z_1$ to an object at red-shift $z_2$ is \\begin{eqnarray} D^A(z_1,z_2)&=&\\frac{1}{1+z_2}\\int^{z_2}_{z_1}{\\frac{dz}{H(z)}} \\\\ \\nonumber &=&\\frac{1}{3}(Rf'-2f)^{-\\frac{1}{3}}\\int^{R_{z_2}}_{R_{z_1}}{\\frac{Rf''-f'}{ (Rf'-2f)^\\frac{2}{3}}\\frac{dR}{H(R)}} \\\\ \\nonumber &=&D^A(R_1,R_2). \\end{eqnarray} The deceleration parameter is \\begin{equation} q(t)=-\\frac{a(t)\\ddot{a}(t)}{\\dot{a}^2(t)}=-\\left(1+\\frac{H'(R)a(R)}{H(R)a'(R)} \\right)=q(R), \\end{equation} when $q>0$ the universe is in early-time matter-dominated phase and  $q<0$ the universe is in dark energy dominated late-time acceleration phase. ",
        "conclusions": "To summarize, in this work we have derived new constraints on $f(R)=R-\\alpha(-R)^\\beta$ in the Palatini formalism, using the statistics of strong gravitational lensing and the lens sample come from SQLS DR3. We use  $f(R)$ within Palatini approach because of it being free from the instabilities, passing the solar system and having the correct Newtonian limit. This typical type of $f(R)$ theories in form of $f(R)=R-\\alpha(-R)^\\beta$ have been recently put forward to explain the late-time expansion phase dominated by dark energy and have been well studied by using the cosmology measurements including background universe evolution parameters, matter power spectrum and CMB. In our paper, we assume that the $f(R)$ gravity theories have hardly detectable effect on the gravitational lensing phenomena of galaxy-scale halos, and then the standard lensing is a realistic approximation of lensing in Palatini $f(R)$ gravity theories. It is worth emphasizing that in the SQlS lens sample, there is a cluster-scale lens, we use SIS model to calculate the differential lensing probability to fit the other 10 lens data to avoid destroying the well-defined statistical lens sample.  Considering the FRW setting, we have shown the background evolution for the late expansion era of matter-dominated phase based on our $f(R)$ theory. The expected lens number density arising from the SQLS DR3 sample is not as sensitive as the matter power spectrum to the parameter $\\beta$.  The number density distributions for different values of $\\beta$ are indistinguishable to $10^{-3}$ in magnitude, while the matter power spectrum changes a lot even if $\\beta$ varies only by a tiny amount, such as an order of  magnitude $10^{-5}$ \\citep{Li B.2007-2, Koivisto2006}. Therefore, the matter spectrum can constrain $\\beta$ to a small region and make the model hardly deviation from the $\\Lambda$CDM model. But the values of our best fit model, $\\alpha=-3.777$ and $\\beta=0.06195$, are not much more different from the standard one with $\\alpha=-4.38$ and $\\beta=0$, implying that the model studied is not significantly preferred over the standard model. The results are consistent with the works of \\citet{Amarzguioui2006}, \\citet{Fay2007} and \\citet{Carvalho2008}, in which the best-fitting models are $(\\alpha, \\beta)=(-3.6, 0.09)$, $(\\alpha, \\beta)=(-4.63, 0.027)$ and $(\\alpha, \\beta)=(-4.7, 0.03)$, respectively, and we all found that $\\Lambda$CDM lies in the 68.3\\% confidence level. Finally, we confine the allowed values of $\\alpha$ and $\\beta$ to the ranges [-4.67,-2.89] and [-0.078, 0.202] at the 95\\% confidence level, and conclude that the SQLS DR3 statistical lens sample can only constrain the parameter $\\beta$ to $10^{-1}$ in magnitude.  In our expected number density distribution (Figure \\ref{dn1}), for any value of  $\\beta$ with the fixed $\\Omega_{m0}=0.27$, the peak is at about $\\Delta\\theta=0.6{''}$. But the lens sample we used, with the image-separation ranging from $1{''}$ to $6.17{''}$, provides no data in $\\Delta\\theta<1{''}$, this significantly restrict the constraint precision of the parameters. The CLASS radio lens sample has the lens image separations of $0.3{''}<\\Delta\\theta<3{''}$. Naturally, jointing the CLASS and SQLS samples to constrain the model  will give better results. Of course, a more larger lens sample from future surveys, \\cite[e.g., Square Kilometer Array,][]{Koopmans2004}, will be able to make the constraints on the model more stringent. Furthermore, using the SIS model and NFW model together to constrain the $f(R)$ theories models arising from the lensing observation data containing both galaxy lens and cluster lens is future work.  Finally, we emphasize that our paper is the first try to test $f(R)$ gravity theories using strong lensing statistics. The reason may be that, the possible deviation of gravitational lensing based on $f(R)$  gravity theories from that built on GR is not well studied practically. Our results hold on the valid hypotheses that this possible deviation is hardly detectable, and they are not contrary to other works. As mentioned in the introduction, if the $f(R)$ theories are successful models to explain the nature of DE,  they should have no important effects which can be detected on the galaxy-scale and cluster-scale dark halos. The reason is that DE is introduced to explain the acceleration of the universe, which is, of course, in the cosmological scale. Namely, DE smoothly permeates the regions of smaller scales, such as galaxies and clusters, then it has no contributions to the inhomogeneity of gravitational potential for such halos. For the dark matter as the dominator of the gravity field, the SIS and NFW lens models are well consistent with the observations. It is therefore necessary to test the $f(R)$ theories through verifying that the phenomenology is not contradict those observational results that do agree with the predictions of the standard lensing theory. This critical testing will be presented in future work."
    },
    "0812/0812.0984_arXiv.txt": {
        "abstract": "Active galactic nuclei (AGN) produce vast amounts of high energy radiation deep in their central engines. X-rays either escape the AGN or are absorbed and re-emitted mostly as IR. By studying the dispersion in the ratio of observed mid-IR luminosity to observed 2-10keV X-ray luminosity ($R_{ir/x}$) in AGN we can investigate the reprocessing material (possibly a torus or donut of dust) in the AGN central engine, independent of model assumptions.  We studied the ratio of observed mid-IR and 2-10keV X-ray luminosities in a heterogeneous sample of 245 AGN from the literature. We found that when we removed AGN with prominent jets, $\\sim 90\\%$ of Type I AGN lay within a very tight dispersion in luminosity ratio ($1<R_{ir/x}<30$). This implies that the AGN central engine is extremely uniform and models of the physical AGN environment (e.g. cloud cover, turbulent disk, opening angle of absorbing structures such as dusty tori) must span a very narrow range of parameters. We also found that the far-IR(100$\\micron$) to mid-IR (12$\\micron$) observed luminosity ratio is an effective descriminator between heavily obscured AGN and relatively unobscured AGN. ",
        "introduction": "\\label{sec:intro} Active galactic nuclei (AGN) are powered by the accretion of matter onto a supermassive ($\\sim 10^{[6,9]}M_{\\odot}$) black hole. X-ray continuum emission from AGN (without jets) is believed to originate in the innermost regions of the AGN whereas IR continuum emission is believed to originate in dusty, cooler outer regions. The IR emission from AGN is due to a combination of reprocessed higher energy emission as well as thermal emission from star formation. Much if not most of the IR due to reprocessing comes from energy dumped in cool dusty material (mostly out of the observers' X-ray sightline) by photons with the most energy (X-rays). A comparison of the dispersion in the IR to X-ray luminosity ratios ($R_{ir/x}$) of a wide variety of AGN should therefore allow us to constrain the range of physical properties and geometry of the X-ray reprocessing material in AGN, independent of model assumptions.  The standard model of AGN explains the wide range of properties of observed AGN in terms of viewing angle, since the equatorial flared disk expected at the heart of AGN breaks spherical symmetry \\citep{b2}. However, partially or heavily obscured AGN need not be edge on, since obscuring material only needs to block the observer's sightline (see e.g. McKernan \\& Yaqoob 1998; Risaliti, Elvis \\& Nicastro 2002; Miller, Turner \\& Reeves 2008). Radiation reprocessing occurs in cool (possibly obscuring) material distributed around the AGN. Independent of assumptions about the distribution of cool material, the \\emph{dispersion} in $R_{ir/x}$ can constrain the range of possible geometries and physical conditions of the obscuring, reprocessing material around the central engine. The dispersion of $R_{ir/x}$ among AGN also provides information about depopulation or selection effects (bias), if there are values of $R_{ir/x}$ that are not favoured by AGN. Genuine depopulation in $R_{ir/x}$ space would yield model-independent information on allowed configurations of the AGN central engine.  The IR and X-ray emission in AGN have been compared in many previous studies \\citep{b69,b62,b3,b4,b5,b71}. Most recently \\citet{b5} reveal a strong correlation between modelled intrinsic 2-10keV X-ray luminosity and near-IR nuclear luminosity, suggestive of a 'leaky' torus of dust. \\citet{b71} find a strong correlation between mid-IR luminosity and the very hard X-ray (14-195keV) observed luminosity, which is unaffected by absorption and indicates that the same phenomenon underlies all AGN. However, as yet there have been no comparisons of the observed hard 2-10keV X-ray luminosity with the IRAS band (12-100$\\micron$) IR luminosities. The 2-10keV X-rays are energetic enough to emerge with relatively mild attentuation through Compton-thin obscurers with absorption becoming more substantial above $N_{H}\\sim 10^{22} \\rm{cm}^{-2}$. Attenuated X-rays will mostly re-emerge as IR emission. Therefore, in tandem with IRAS band IR emission, 2-10keV X-rays are a good probe of absorption and subsequent reprocessing in the obscuring material around AGN. In this paper, for the first time, we studied the ratio of observed 2-10keV X-ray emission to IRAS band (12-100$\\micron$) IR emission for a large, heterogeneous sample of AGN. By studying dispersions in the luminosity ratio $R_{ir/x}$ of our sample, we hope to obtain information about the \\emph{dispersion} of physical properties of the AGN environment, independent of model assumptions.  In section~\\ref{sec:sed} we briefly discuss X-ray and IR luminosity of AGN, including previous studies which compared emission in the two bands, and in section~\\ref{sec:obs} we discuss our sample of AGN and associated caveats. In section~\\ref{sec:xir}, we explore the distribution of our sample of AGN in $R_{ir/x}$ parameter space and we use AGN classification and well-known AGN to constrain the dispersion of AGN in $R_{ir/x}$. In section~\\ref{sec:blank} we attempt to rule out selection bias as an explanation of the dispersion of AGN in $R_{ir/x}$ space. In section~\\ref{sec:discuss} we discuss our results and establish model-independent constraints on the reprocessing material around AGN. Section~\\ref{sec:conclusions} outlines our conclusions. ",
        "conclusions": "\\label{sec:conclusions} Our aim was to compare common observed AGN properties that encompass the continuum of AGN activity in an effort to, if possible, derive model-independent constraints on AGN structure and the accretion neighbourhood. We studied the observed X-ray to IR luminosity ratios of a heterogeneous sample of 245 galaxies from the literature and, in spite of our simple approach, AGN variability and non-simultaneity of most of the data, we found strong trends.  Fig.~\\ref{fig:irx12} reveals the emergence of observed 'activity' at $L_{IR}, L_{X} >10^{42} \\rm{erg} \\rm{s}^{-1}$. Below these luminosities, AGN may compete with X-ray and IR emitters in the host galaxy (such as ULXs, see e.g. \\citet{b30}) or may be obscured by large columns of absorbing material. Since emission from jets considerably complicates our interpretation of the observed luminosity ratios, we only considered those AGN without prominent observed jets. We find that AGN considered 'face-on' or unobscured in the standard model (our group 1) are tightly clustered by observed luminosity ratio ($\\sim 90\\%$ have $1<R_{ir/x}<30$), indicating that there is very little variation between individual group 1 AGN in: (a) intrinsic SED (i.e. accretion flow, whether ADAF, tubulent disk or both), (b) solid angle of the reprocessing ($X \\rightarrow IR$) absorber (whether clouds, torus, flared-disk or some combination) and (c) details of the absorbing material (chemistry \\& dust grain sizes). Thus, Fig.~\\ref{fig:irx12} demonstrates that underlying Group 1 AGN is a remarkably uniform phenomenon.  Fig.~\\ref{fig:irx12} shows a clear difference in $R_{ir/x}$ between unobscured AGN (group 1) and obscured \\& low luminosity AGN (our group 2). Despite the fact that our sample is \\emph{not} complete, our findings for the dispersion of AGN are confirmed in a comparison with other, independent AGN samples (e.g. \\citet{b37,b71}, which are not limited by X-ray absorption, or \\citet{b32,b27,b28} from $z\\sim 1$). Furthermore, subsets of our sample of 245 AGN are derived from complete IRAS samples \\citep{b60,b59} and these subsets at least do not span values of $R_{ir/x}$ different from the rest of our sample. Bias is of course a potential problem with any approach such as ours, but it is comforting that we can find some independent verification for our simple, model-independent approach. Estimates of the intrinsic 2-10 keV luminosity for Compton-thick AGN from very hard X-ray (absorption independent) observations (e.g. \\citet{b37,b38,b36}) reveal that the intrinsic luminosity ratio is consistent with the group 1 AGN (in agreement with the findings of e.g. \\citet{b3,b4,b5,b71}). However, unlike \\citet{b3,b4,b5} we have endeavoured to avoid the highly model-dependent business of estimating intrinsic 2-10keV X-ray luminosity of group 2 AGN based on the 2-10keV observed luminosity. We also studied the far-IR to mid-IR luminosity ratios of the nuclei in our sample. We found that the $100\\micron$ to $12\\micron$ luminosity ratio ($L_{100\\micron}/L_{12\\micron} \\sim 1$) is an extremely effective separator of the group 1(unobscured) and group 2 (obscured \\& low luminosity) populations (see Fig.\\ref{fig:ir12v100}). Using archetypal AGN as a guide, this plot ranks group 2 AGN in order of decreasing luminosity ratio for increasing obscuration.  In principle our approach could be applied to intermediate mass black holes (IMBHs) and Galactic black hole candidates (GBHCs) to investigate the accretion neighbourhood around black holes on all mass scales. The outstanding question is whether reprocessing material around IMBHs and GBHCs is substantially different in covering factor, structure, chemistry or reprocessing details from that around AGN. Are a flared disk or dusty torus associated with a minimum mass scale, or are we looking at a basically identical phenomenon across 7-8 orders of magnitude in mass?"
    },
    "0812/0812.3043_arXiv.txt": {
        "abstract": "The evolution  of white  dwarfs can be  described as a  simple cooling process.   Recently,  it  has  been  possible  to  determine  with  an unprecedented precision their luminosity function, that is, the number of stars per unit volume  and luminosity interval.  Since the shape of the bright  branch of this function  is only sensitive  to the average cooling rate, we use this  property to check the possible existence of axions, a  proposed but not yet detected  weakly interacting particle. We  show  here that  the  inclusion of  the  axion  emissivity in  the evolutionary models of white  dwarfs noticeably improves the agreement between the theoretical calculations and the observational white dwarf luminosity function, thus providing the first positive indication that axions  could  exist.  Our  results  indicate  that  the best  fit  is obtained for $m_{\\rm  a} cos^2 \\beta \\simeq 2 -  6$ meV, where $m_{\\rm a}$ is the mass of the axion and $cos^2\\beta$ is a free parameter, and that values larger than 10 meV are clearly excluded. ",
        "introduction": "White   dwarfs  are  the   final  evolutionary   stage  of   low-  and intermediate-mass  stars ($M\\leq  10\\pm 2  M_\\odot$).  Since  they are degenerate  objects,  they  cannot  obtain energy  from  thermonuclear reactions  and  their evolution  is  just  a  gravothermal process  of cooling.   They  have a  relatively  simple  structure  composed by  a degenerate core that contains the bulk of mass and acts as a reservoir of  energy, and  a  partially degenerate  envelope  that controls  the energy outflow.   White dwarfs  masses are in  the range of  $0.3 \\leq M/M_\\odot \\leq 1.4$ (the  Chandrasekhar's mass). Those with $M\\leq 0.4 M_\\odot$  have a core  made of  He, while  those with  a $M  \\geq 1.05 M_\\odot$ have  a core made of O  and Ne. The remaining  ones, the vast majority, have a  core made of a mixture  of C and O. All  of them are surrounded by a  thin helium layer with a  mass ranging from $10^{-2}$ to  $10^{-4}\\, M_\\odot$,  which, in  turn,  is surrounded  by an  even thinner layer of hydrogen with a mass between $10^{-4}$ and $10^{-15}$ $M_\\odot$,  although about  25\\%  of  white dwarfs  do  not have  such hydrogen envelopes.  White dwarfs displaying hydrogen in their spectra are known  as DA  and the  remainning ones as  non-DA. Because  of the different  opacities involved,  DA white  dwarfs cool  down  much more slowly than  non-DA white dwarfs.  Because of the simplicity of their structures, it has been proposed to use white dwarfs as laboratories  to test new physics. The reasons for such  a proposal  are the  following ones.  Firstly, the  evolution of white dwarfs is just a  simple process of cooling. Secondly, the basic physical  ingredients necessary  to predict  their evolution  are well identified,  although not necessarily  well understood,  and, finally, there  is a  solid and  continuously growing  observational background that  allows  to  test  the  different  theories.   Therefore,  if  an additional  source or  sink  of energy  is  added, the  characteristic cooling time is modified and these changes can be detected through the anomalies  introduced in  the luminosity  function or  in  the secular drift of the pulsation period of degenerate variables --- see Isern \\& Garc\\'{\\i}a--Berro  (2008) for details.  In this  paper we  will limit ourselves  to study  the  changes introduced  by  the introduction  of axions in the white dwarf luminosity function. ",
        "conclusions": "Figure \\ref{fig5} displays the luminosity function for different axion masses, a constant star formation rate and an age  of the  Galactic disk  of 11  Gyr.  As  already mentioned,  it is important  to  remember  that  the  bright branch  of  the  luminosity function  is  not  sensitive  to  these  last  assumptions.   All  the luminosity  functions have been  normalized to  the luminosity  bin at $\\log  (L/L_{\\odot})\\simeq-3$ or,  equivalently,  $M_{\\rm bol}  \\simeq 12.2$.   The  best fit  model  ---  namely  that which  minimizes  the $\\chi^2$ test in the region $-1>\\log (L/L_{\\odot})>-3$ (that is, $ 7.2 <  M_{\\rm  bol}  <  12.2$),   which  is  the  region  where  both  the observational  data and  the theoretical  models are  reliable  --- is obtained  for $m_{\\rm a}  \\cos^2 \\beta\\approx  5.5$ meV  and solutions with $m_{\\rm a} \\cos^2 \\beta >  10$ meV are clearly excluded (Isern et al 2008).   Figure \\ref{fig6} displays  the behavior of $\\chi^2$  as a function of the mass of  the axion for different normalization points. In all cases, $\\chi^2$ displays a pronounced minimum for masses in the range of  roughly 2 to 6 meV  except for $l=2.51$ and  $2.67$ that are compatible  with $m_{\\rm  a} =  0$. However,  as it  can be  seen from figure  \\ref{fig6},  both points  deviate  from  the  behavior of  the neighbouring values.  \\begin{figure}[t] \\vskip 9cm \\special{psfile=fig6.ps hscale=80 vscale=80 hoffset=-20 voffset=-300} \\caption{Value  of $\\chi^2$ as  a function  of the  mass of  the axion adopting different normalization values.} \\label{fig6} \\end{figure}  Taken  at  face  value,  these  results  not  only  provide  a  strong constraint to the allowed mass of axions, but also a first evidence of their  existence and a  rough estimation  of their  mass.  This  is of course  a strong  statement and  in  order to  be accepted  it has  to fulfill   the   following  conditions:   it   has   to  be   confirmed independently, it  has to resist the introduction  of any conventional effect not previously included, it has to provide testable predictions and it can not enter  in contradiction with well established facts.  A detailed discussion of the existing  uncertainties is out of the scope of the present  paper but it is worthwhile to  enumerate some of them: influence of  the metallicity in the  age and core  composition of the progenitor,  transparency of the  envelope and  conversion of  DA into non-DA white dwarfs, IMF, intial-final mass relationship, pathological SFRs  or  uncertainties  in  the  observational  luminosity  function, especially on  the role of He  white dwarfs and the  detailed shape of the brightest part of this function and others.  The results  found here are  completely compatible with  the presently known constraints (Raffelt 2007), but since the predicted mass is near the allowed upper  bound, it is natural to  expect they will introduce some subtle changes  in the late stages of stellar  evolution as it is indeed the case  of white dwarfs. A detailed  discussion of this point is also out of  the scope of this paper and we  will just mention that these values  are compatible with the  bounds imposed by  the drift of the  pulsational period  of the  ZZ  Ceti star  G117$-$B15A (Isern  et al. 1992;  C\\'orsico et al. 2001;  Bischoff-Kim et al.   2008).  It is also worthwhile  to mention  here that axions  with $m_{\\rm  a} \\cos^2 \\beta  \\approx  5$ meV  would  change  the  expected period  drift  of variable DB  white dwarfs --- which  have values between  $\\dot P \\sim 10^{-13}$ and $10^{-14}$ s~s$^{-1}$ (C\\'orsico \\& Althaus 2004) --- by a factor of 2, the exact value depending on the adopted temperature of the stellar core.   The structure of the Sun would  not be modified by the  emission of  axions of  this  masses since  they would  represent $L_{\\rm  a} \\sim  10^{-6}L_\\odot$. However,  an instrument  like CAST, able to operate in the region  of masses $m_{\\rm a}\\sim {\\rm meV}$ and coupling constant  $g_{{\\rm a}\\gamma} \\sim  10^{-12}\\, {\\rm GeV}^{-1}$ could  be able  to detect  them  and confirm  or rule  out the  masses predicted  here.  Cosmology could  also provide  some insight  to this problem. Axions have been proposed as dark matter candidates but if we assume  $\\cos^2\\beta=1$,  their contribution  would be of  the order  of $\\Omega_{\\rm a}  h^2 \\sim 10^{-4}$,  where $h$ is the  Hubble constant normalized to 100 km/s/Mpc (Raffelt  2007), which rules out them as the main  dark matter  candidate.   On  the contrary,  if  dark energy  is interpreted  as the  energy density  of vacuum,  axions could  have an important role. The vacuum energy is obtained from the sum of the zero point  energy  of all  the  quantum  fields  (where bosons  contribute positively and  fermions negatively)  and if we  try to  reproduce the observed value, that is positive,  with just one quantum field we need a boson  with a mass  of the order  or smaller than 10  meV (Friemann, Turner \\& Huterer,  2008), that is just what we  have found here. This is indeed a remarkable coincidence.  To  conclude, it  is evident  that our  claim about  the  existence of axions has  to be taken  with caution, but  it clearly shows  that the white dwarf luminosity function  can provide strong arguments to solve the long-lasting problem  of the CP violation.  Future  work should be devoted  to  improving  the  observational data,  especially  for  the brightest  part   of  the   white  dwarf  luminosity   function.   The theoretical models could also  be improved and also additional efforts should  be   devoted  to  find  new  independent   methods  to  detect axions. Finally,  it is worthwhile to say  that even in the  case of a negative result,  an accurate and precise luminosity  function could be used to provide  insight to many other problems  like the hypothetical drift  of  the  gravitation  constant  or  the  magnetic  momentum  of neutrinos just to cite two cases.  \\ack Part of this  work was supported by the  MEC grants AYA05-08013-C03-01 and 02,  and ESP2007-61593, by the  European Union FEDER  funds and by the AGAUR."
    },
    "0812/0812.3569_arXiv.txt": {
        "abstract": "Using time evolutions of the relevant linearised equations we study non-axisymmetric oscillations of rapidly rotating and superfluid neutron stars. We consider perturbations of Newtonian axisymmetric background configurations and account for the presence of superfluid components via the standard two-fluid model.  Within the Cowling approximation, we are able to carry out evolutions for uniformly rotating stars up to the mass-shedding limit. This leads to the first detailed analysis of superfluid neutron star oscillations in the fast rotation regime, where the star is significantly deformed by the centrifugal force. For simplicity, we focus on background models where the two fluids (superfluid neutrons and protons) co-rotate, are in $\\beta$-equilibrium and coexist throughout the volume of the star. We construct sequences of rotating stars for two analytical model equations of state. These models represent relatively simple generalisations of single fluid, polytropic stars. We study the effects of entrainment, rotation and symmetry energy on non-radial oscillations of these models. Our results show that entrainment and symmetry energy can have a significant effect on the rotational splitting of non-axisymmetric modes. In particular, the symmetry energy modifies the inertial mode frequencies considerably in the regime of fast rotation. ",
        "introduction": "\\label{sec:Intro}  According to the standard paradigm, millisecond pulsars are accelerated to their fast rotation rates by accreting matter from a close companion. This means that they tend to be relatively old. Moreover, the fastest spinning pulsars should have weak (exterior) magnetic fields.  In the standard accretion model, neutron stars with canonical $10^{12}$~G dipole fields will reach equilibrium already at a modest spin. A weak surface field is also expected since accretion leads to magnetic field burial. This picture agrees well with observational data.  Rapidly rotating neutron stars are most commonly found in binary systems. It is well established that accreting neutron stars in low-mass X-ray binaries (where the angular momentum transfer is more efficient due to the long evolution time of the low mass partner) can reach a millisecond rotation period. Furthermore, the fastest known millisecond pulsar J1748-2446ad, with a period of 1.39~ms~\\citep{2006Sci...311.1901H}, is in a binary system. In fact, its companion, with mass $M \\geq 0.14 M_{\\odot}$, could still fill the Roche lobe powering the spin-up phase further. The mass and radius of J1748-2446ad are unknown, but combining reasonable ranges for these parameters, $M = 1.4-2 M_{\\odot}$ and $R=10-14~\\rm{km}$, with an empirical formula for the maximum rotation of the star~\\citep{2004Sci...304..536L} \\begin{equation} \\Omega_K \\approx 6566 \\left( \\frac{M}{M_{\\odot}} \\right) ^{1/2} \\left( \\frac{ 10 \\rm{km} }{ R } \\right) ^{3/2}  \\rm{Hz} \\, , \\end{equation} one finds that the spin of this object lies in the range $0.48 \\lesssim \\Omega / \\Omega_K \\lesssim 0.96 $.  In other words, it could be close to the mass-shedding limit.  The temperature of a mature neutron star is likely below the critical temperature, $T_c \\simeq 5\\times 10 ^9$~K, where neutrons and protons are superfluid and superconducting, respectively. Depending on the cooling mechanism, neutrino emission can cool a hot proto-neutron star below this temperature shortly after its formation in a core collapse supernova~\\citep*[see e.g,][]{2006NuPhA.777..497P}.  In an accreting system, the neutron star core temperature is not expected to increase beyond $T_c$ (nuclear burning in the accreted surface layers is thought to the heat the core to $\\sim 10^8$~K).  Hence, all mature neutron stars should contain degenerate superfluid neutrons in the outer core and the inner crust and degenerate superconducting protons in the outer core. The deep core may contain more exotic phases of matter, like superfluid hyperons and/or colour superconducting quarks.  Superfluidity influences the thermal evolution and the dynamical properties of a neutron star. In particular, the dynamics is strongly affected by entrainment, the formation of quantized neutron vortices, and the presence of new dissipative mechanisms like mutual friction.  An understanding of superfluid dynamics is crucial for modelling many aspects of the neutron star physics, ranging from pulsar glitches and free precession to the mutual friction damping of stellar oscillations and associated instabilities.   Tidal forces, accretion and glitches may trigger oscillations and/or instabilities in rapidly rotating neutron stars. Observations of such oscillations, either via electromagnetic or gravitational radiation, would help us explore the exotic physics of these compact objects \\citep{1998MNRAS.299.1059A,benhar-2004-70,2007MNRAS.374..256S}.  In this context it is interesting to note the differences in dynamics between neutron stars above the superfluid transition temperature and colder systems.  The core of a hot young neutron star is relatively well described by the Navier-Stokes equations.  In contrast, a star with a superfluid core requires a multi-fluid description. The standard model for such systems is inspired by the two-fluid model for superfluid Helium~\\citep{Landau:1959, Khalatnikov:1965, Tilley:1990}. The oscillation spectrum of superfluid neutron stars reflects the presence of the additional degree of freedom~\\citep{1988ApJ...333..880E}. Basically, the perturbed fluid elements in a two-fluid system can oscillate either in phase or in a counter-phase motion. Previous studies, see for example, \\citet{1995A&A...303..515L, 2001MNRAS.328.1129A, 2002A&A...393..949P, 2003MNRAS.344..207Y}, have established that co-moving pulsations have spectral properties similar to single fluid stars. Hence, it is natural to refer to such modes as ``ordinary'' modes. The counter-moving degree of freedom leads to new oscillation modes that are specific to the two-fluid systems. These are often referred to as ``superfluid'' modes. Later, when we write down the perturbation equations for a superfluid neutron star core, we choose to work with variables that are directly associated with the two types of motion. This is natural if we want to distinguish spectral properties associated with the ``superfluid\" degree of freedom. In addition, we use the standard classification of neutron star oscillation modes, based on the main restoring force that acts on a perturbed fluid element. A rotating single fluid star can sustain acoustic and inertial modes restored by pressure and the Coriolis force, respectively.  When thermal or composition gradients are present in the star, buoyancy acts as restoring force for the so-called gravity modes~\\citep{1989nos..book.....U, 1992ApJ...395..240R}. Previous work has shown that superfluid neutron stars have two families of acoustic and inertial modes, more or less clearly (depending on the stellar model) associated with the co- and counter-moving fluid motion. It is, however, not the case that all single fluid modes have a \"double\" in the superfluid problem. The gravity modes are not present at all in a superfluid core~\\citep{1995A&A...303..515L, 2001MNRAS.328.1129A, 2002A&A...393..949P}.  Their absence provides a potentially important signature for neutron star seismology.  Rapidly rotating neutron stars have been studied in detail with a variety of methods~\\citep[see e.g.,][]{lrr-2003-3}. Yet, there have not been any previous studies of multi-fluid dynamics in the rapid rotation regime near the break-up limit.  The oscillation modes of superfluid neutron stars have only been calculated in the frequency domain using the slow rotation approximation~\\citep{2000PhRvD..61j4003L, 2003MNRAS.344..207Y, 2003PhRvD..67l4019Y}.  In that framework, the effects of the stellar rotation are determined perturbatively as corrections to the non-rotating results.  The only previous attempt (that we are aware of) to study superfluid oscillations for truly fast spinning stars is the work by \\citet{2000PhRvD..61j4003L}.  They extended the two-potential formalism of \\citet{1990ApJ...355..226I} to the superfluid case.  However, due to numerical difficulties they could not study rotating models near the break-up limit.  Neither did they manage to determine the superfluid modes.   In this work, we study the time evolution of perturbed fast rotating, Newtonian superfluid neutron stars within the Cowling approximation. As far as we are aware, this is the first study that evolves in time the oscillations of superfluid neutron stars. Moreover, it is the first detailed analysis of the rapid rotation regime.  Within the framework of the two-fluid formalism, we carry out a linear perturbation analysis for stationary and axisymmetric equilibrium configurations. As preparation for more detailed studies, we consider relatively simple models where the two fluids coexist throughout the star, and where the unperturbed configuration is in $\\beta$-equilibrium.  These assumptions imply that the two fluids are co-rotating and share the same external stellar surface in our background configurations. We use these models to investigate the effects of entrainment, symmetry energy and rotation on the superfluid oscillation spectrum. In order to establish the reliability of our numerical evolution code, we compare our results to previous work for non-rotating and slowly rotating models. We then consider, for the first time, the dynamics of superfluid models rotating up to the mass shedding limit. ",
        "conclusions": "} \\label{sec:concl}  In this paper we have considered, for the first time, the oscillations of a superfluid neutron star as an initial-value problem. Using time evolutions of the relevant linearised equations we studied non-axisymmetric oscillations of rapidly rotating superfluid neutron stars. We considered perturbations of axisymmetric background configurations in Newtonian gravity and accounted for the presence of superfluid components via the standard two-fluid model.  Within the Cowling approximation, we were able to carry out evolutions for uniformly rotating stars up to the mass-shedding limit.  Our results represent the first detailed analysis of superfluid neutron star oscillations in the fast rotation regime, where the star is significantly deformed by the centrifugal force.  For simplicity, we focused on background models such that the two fluids (superfluid neutrons and protons) co-rotate, are in $\\beta$-equilibrium and coexist in all the volume of the star. Two different analytical model equations of state were considered. The models were chosen to represent relatively simple generalizations of single fluid, polytropic stars. We investigated the effects of entrainment, rotation and symmetry energy on various non-radial oscillation modes of these models. Our results show that entrainment and symmetry energy can have a significant effect on the rotational splitting of non-axisymmetric modes. In particular, the symmetry energy modifies the inertial mode frequencies considerably in the regime of fast rotation.  The perturbative time-evolution framework provides a useful tool that should allow us to consider more realistic (and by necessity complicated) neutron star models in the future.  The clear advantage over frequency-domain studies is that it is straightforward to study oscillations corresponding to eigenfunctions with a complex set of rotational couplings. This is particularly useful in the rapid rotation regime.  The obvious downside is that time evolutions can never provide the ``complete\" mode spectrum of the star.  Initial data has to be chosen in such a way that the oscillations of interest are excited at a significant level. It is difficult to, without prior knowledge, find initial data that excites only a few modes. In many cases this is, however, less relevant. The main question is if one can extract accurate information regarding the nature of the star's oscillations from the numerical data. The results we have presented demonstrate that this is, undoubtedly, the case. Hence, it is relevant to develop the perturbative evolution framework further. We are currently considering more general background models, with a relative velocity between the two fluid components. We are also adding the dissipative coupling associated with mutual friction to the code. Once these features are incorporated we will be able to consider (obviously still at a basic level) the dynamics associated with the superfluid two-stream instability \\citep{2stream1,2stream2} and the possible relation with pulsar glitches~\\citep[see][for a recent discussion]{2008arXiv0806.3664G}.  This is a very exciting prospect. Looking further ahead, we would like to add layering to the stellar model by introducing both an elastic crust and distinct superfluid/normal regions. There are challenges associated with these aspects, but there is no reason why these developments should be prohibitively difficult."
    },
    "0812/0812.3872_arXiv.txt": {
        "abstract": "Based on a multiwavelength study, the interstellar medium around the \\hii\\ region Sh2-173 has been analyzed. The ionized region is clearly detected in the optical and in the radio continuum images. The analysis of the \\hi\\ data shows a region of low emissivity that has an excellent morphological correlation with the radio continuum emission. The \\hii\\ region is partially  bordered by a photodissociation region, which, in turn, is encircled  by a molecular structure. The \\hi\\ and CO structures  related to Sh2-173 are observed in the velocity ranges from --25 to --31\\kms, and from --27 to --39 \\kms, respectively. Taking into account the presence of noncircular motions in the Perseus spiral arm, together with previous distance estimates for the region, we adopt a distance of 2.5 $\\pm$ 0.5 kpc for Sh2-173. Seven hot stars were identified in the field of Sh2-173, being only one an O-type star. The amount of energetic photons emitted by this star is enough to keep the region ionized and heat the dust. Given that an expanding \\hii\\ region may trigger star formation, a search for YSO candidates was made using different infrared point source catalogues. A population of 46 YSO candidates was identified projected onto the molecular clouds.\\\\ On the other hand, Sh2-173 is located in a dense edge of a large ($\\sim5^{\\circ}$) \\hi\\ shell, \\gsh. The possibility for Sh2-173 of being part of a hierarchical system of three generations is suggested. In this scenario, the large \\hi\\ shell, which was probably originated due to  the action of Cas\\,OB5, would have triggered the formation of Sh2-173, which, in turn, is triggering new stars in its surrounding molecular cloud. To test this hypothesis, the ages of both, the \\hii\\ region and the large shell, were estimated and compared. We concluded that Sh2-173 is a young \\hii\\ region of about  0.6 - 1.0  Myr old. As for the large shell we obtained a dynamical age of 5 $\\pm$ 1 Myr. These age estimates, together with the relative location of the different structures, support the hypothesis that Sh2-173 is part of a hierarchical system. ",
        "introduction": "Massive stars affect their environment by ionizing radiation, stellar winds, and, finally, supernova explosions. These stars produce intense extreme ultraviolet ($h\\nu > 13.6$ eV) radiation, creating  \\hii\\ regions around them, and intense far ultraviolet ($4.5 < h\\nu < 13.6$ eV) radiation that can penetrate into regions of molecular gas, dissociating it. They also have  powerful winds which sweep up the surrounding gas creating what is known  as  an interstellar bubble, i.e. a minimum in the gas distribution characterized by a low volume density and a high temperature, surrounded by an expanding envelope \\citep{wea77}. In the case of OB associations, their collective effect has a huge impact in the surrounding interstellar medium (ISM), creating large structures of about hundred parsecs. Winds and supernova explosions transfer to the ISM vast amounts of mechanical energy, generating the compression of nearby molecular clouds. The effects of the massive stars are mostly destructive in their immediate neighborhood, since they tend to disrupt the parental molecular cloud. However, a little further away, massive stars stimulate star formation. A number of papers have been devoted to investigate star formation in the dense molecular clouds in the  surroundings of \\hii\\ regions, both from a theoretical and an observational point of view.  Different processes  can trigger stellar formation at the periphery of \\hii\\ regions \\citep{elm77, kle80, san82, lef94, whi94, lef95, deh05, dal07}. Star formation can be induced in the environs of massive stars due to the action of expanding shocks which modify the surrounding material, either favouring the formation of cometary globules in preexistent clumps (\"Radiative Driven Implosion\" model, RDI) or sweeping up the surrounding material into a collected layer, which latter becomes gravitationally unstable generating massive fragments (\"collect and collapse\" model). Both processes, which may act simultaneously in an \\hii\\ region, lead to the formation of a new generation of stars. Observational evidence of both (the RDI and \"collect and collapse\") processes  has been found in a number of \\hii\\ regions \\citep[e.g.][]{tho04, deh08, cap08, kir08}.    The goal of this paper is to achieve a better understanding of the phenomena associated with the interaction between young massive stars and their surrounding interstellar medium and, specifically, to attempt to shed some light on the process of triggered star formation. To carry out this study a multiwavelength approach is essential in order to obtain a complete picture of the different components of the ISM, the young stellar population, and the interactions among them.  In this work we investigate the environs of the \\hii\\ region Sh2-173 by analyzing the distribution of the ionized and neutral material, and that of the interstellar dust. Sh2-173 is an optically visible \\hii\\ region \\citep{sha59} located at ({\\it l, b}) = (119\\fdg4, --0\\fdg94) in the Perseus spiral arm. The DSS2-R image of Sh2-173 is displayed in Fig. \\ref{optico}. The region is approximately circular in shape with a diameter of about 30\\am. An arc-like structure of enhanced optical emission is evident on the north-west side of the region, while fainter emission  is present at higher  galactic longitudes and lower galactic latitudes. Inside the ionized region, two voids lacking optical emission are evident  at ({\\it l, b}) = (119\\fdg4, --0\\fdg85) and  ({\\it l, b}) = (119\\fdg55, --0\\fdg95). \\citet*{rus07} identified seven hot stars in the field of Sh2-173, which are listed in Table\\ref{tabest}, and indicated with small triangles in Fig. \\ref{optico}. All of the stars but one are projected close to the center of the region. ALS\\,6155, on the contrary, is seen projected on its southeast border.  The distance to Sh2-173 has been estimated by several authors using independent methods. A spectrophotometric distance of $ 2.7 \\pm 0.9$ kpc was derived  by \\citet{geo75}, while \\citet{rus07} estimated a distance of $ 3.12 \\pm 0.34$ kpc based on UBV photometric and spectroscopic data of the exciting stars  (see Table \\ref{tabest}). Using  a new method for determination of distances, which is  based on \\hi\\ column densities and is independent of Galactic kinematics, \\citet{fos03} estimated for this \\hii\\ region a distance of $ 2.2 \\pm 0.4$ kpc. \\citet*{bli82} in their CO survey towards \\hii\\ regions found  molecular gas related to Sh2-173 at --34.5 \\kms\\ (all the velocities are referred to the LSR). This velocity coincides with  the radial velocity of the ionized gas derived from Fabry-Perot H$\\alpha$ observations, $-34.3 \\pm 0.3$ \\kms\\, \\citep*{fic90}. It is well known that  noncircular motions on a large scale are present in the Perseus spiral arm \\citep{bra93}. An inspection of  Fig. 2b of \\citet{bra93}, which takes into account noncircular motions, suggests  for a velocity of about --34 \\kms\\,, a kinematical distance of 1.8-2 kpc. Thus, the distance of Sh2-173 is between 1.8 and 3.1 kpc. As a working  hypothesis, we will adopt a distance of 2.5 $\\pm$ 0.5 kpc for Sh2-173.   Sh2-173 is seen located  in the dense border of a large \\hi\\ shell ($\\sim 5^{\\circ}$ diameter) that has been found in the Perseus spiral arm by \\citet{fic86} using \\citet{wea73} \\hi\\ data. \\citet{fic86} suggested that the presence of SNR remnants and several \\hii\\ regions in the field indicates a great deal of star-formation activity going on nearby. \\citet*{kis04} found a far infrared ring close to the \\hi\\ shell defined by \\citet{fic86}. \\citet{hey98} found a large void of CO emission centered at ({\\it l, b}) = (117\\fdg6, 0\\fdg0), coincident with the \\hi\\ shell. Based on a multiwavelength study, \\citet{moo03} determined the main parameters of the shell and associated it with the OB association Cas\\,OB5 ($\\equiv$ Rupr\\,33). According to \\citet {gar92}, Cas\\,OB5 has a zero-age main sequence (ZAMS) fairly good defined with a distance modulus DM = 11.5 mag (D $\\sim$ 2 kpc). The diameter of the association is about 4$^{\\circ}$, or  140 pc at 2 kpc. Its earliest star is HD\\,108 (O6f). Among their members there are six O-type stars, being the others B stars \\citep{gar92}. A WN4 star (WR\\,159) was found in the area of Cas\\,OB5,  located at  ({\\it l, b}) = (115\\fdg78, +1\\fdg24) \\citep{neg03}.  \\begin{figure} \\includegraphics[]{fig1.ps} \\caption{DSS2-R optical image of the \\hii\\ region Sh2-173. The small triangles indicate the position of the seven OB stars listed in Table \\ref{tabest}. } \\label{optico} \\end{figure}    \\begin{table} \\centering \\begin{minipage}{100mm} \\caption{Stellar parameters of the OB stars associated with Sh2-173.\\label{tabest}} \\begin{tabular}{l c c c c} \\hline Star & Spectral  & Galactic  &  Distance$^1$& Comments \\\\ & Type$^1$& coordinates &  [kpc] &\\\\ \\hline ALS6145 & B0.5V & 119\\fdg39, --0\\fdg93 &  $ 3.17 \\pm 1.27 $  &\\\\ ALS6150 & B2IV & 119\\fdg44, --0\\fdg91 &  $ 3.78 \\pm 1.89$ &\\\\ ALS6151 & O9V & 119\\fdg44, --0\\fdg92 &  $2.57 \\pm 0.5$ & BD+60\\,39\\\\ ALS6155 & B0.5V & 119\\fdg44, --1\\fdg11 &  $3.36 \\pm 1.34$ & \\\\ ALS6156 & B2V & 119\\fdg48, --0\\fdg93 &  $4.41 \\pm 1.77 $ &\\\\ ALS6157 & B2IV & 119\\fdg48, --1\\fdg02 &  $ 2.27 \\pm 1.14 $ & BD+60\\,42 \\\\ ALS6158 & B1V & 119\\fdg52, --0\\fdg85 &  $ 4.02 \\pm 0.66 $ &\\\\ \\hline  \\end{tabular}  \\medskip 1- From \\citet{rus07}.  \\end{minipage} \\end{table} ",
        "conclusions": "The \\hii\\ region Sh2-173 has been analyzed in order to study the phenomena associated with the interaction of massive stars and the interstellar medium, and specifically, the process of triggering star formation. The analysis of the available radio continuum data have enabled us to estimate some parameters that characterize the \\hii\\ region, like the ionized mass, the emission measure, and the electron density. From the radio flux densities estimated at different wavelengths, the thermal nature of the source was confirmed by the estimation of the spectral index $\\alpha = 0.0 \\pm 0.1, S_{\\nu} \\sim \\nu^{\\alpha}$. We noted that the massive stars present in the field can provide the necessary UV photon flux to keep the \\hii\\ region ionized and heat the dust. As expected for \\hii\\, regions, Sh2-173 displays a strong correlation between the 1420 MHz brightness temperature and infrared emission. An inspection of the \\hi\\ images has shown the existence of an \\hi\\ cavity which presents a good morphological agreement with the ionized gas. The \\hi\\ feature is observed within the velocity range --24.6 and --31.6 \\kms\\, and has a systemic velocity of $-27.0 \\pm 1.3$ \\kms\\,, compatible with the velocity of the ionized gas as obtained from H$\\alpha$ observations. The distribution of the  MSX emission at 8.3 $\\mu$m clearly shows the presence of a PDR bordering the brightest part of the \\hii\\ region. Molecular gas probably related to Sh2-173 is observed in the velocity range  --27.4 to --40.6 \\kms. Particularly, the CO feature named cloud A is clearly observed encircling both the \\hii\\ region and the PDR. Applying  different colour criteria, 46 YSO candidates have been identified  in the vicinity of Sh2-173. Their positions in the CC and CM diagrams, as well as their correlation with molecular emission highly suggest the presence of an active star formation process in the molecular envelope of Sh2-173.  An analysis of the \\hi\\ distribution in a larger area shows that Sh2-173 is located at the edge of a large shell, together with several other \\hii\\ regions. This morphology suggests that Sh2-173 may be part of a hierarchical system of three generations. In this context, the age of the different structures were estimated. Based on evolutionary models of \\hii\\ regions, we estimated for Sh2-173 a dynamical age of  0.6 - 1.0 Myr. As for the large \\hi\\ shell, we have obtained a dynamical age $t_{\\rm dyn} = 5 \\pm 1 $ Myr. Based on energetics considerations, we have concluded that the massive stars belonging to Cas\\,OB5 could not alone create the shell, being highly probable that at least three supernova explosions have contributed in its formation. The obtained age difference between Sh2-173 and the \\hi\\ shell, together with  their relative location, lead us to the conclusion that Sh2-173 may have been created as a consequence of the  action of a strong shock produced by the expansion of \\gsh\\, into the surrounding molecular gas. The analysis of the ages and distribution of the other \\hii\\ regions observed in the area would be useful to confirm this scenario."
    },
    "0812/0812.4349_arXiv.txt": {
        "abstract": "We report a discovery of extended counterrotating gaseous disks in early-type disk galaxies NGC~2551 and NGC~5631. To find them, we have undertaken complex spectral observations including integral-field spectroscopy for the central parts of the galaxies and long-slit deep spectroscopy to probe the external parts. The line-of-sight velocity fields have been constructed and compared to the photometric structure of the galaxies. As a result, we have revealed full-size counterrotating gaseous disks, the one coplanar to the stellar disk in NGC~2551 and the other inclined to the main stellar disk in NGC~5631. We suggest that we observe the early stages of minor-merger events which may be two different stages of the process of lenticular galaxy formation in rather sparse environments. ",
        "introduction": "In principle, a significant role of mergers implied by the hierarchical paradigm for the galaxy evolution must result in frequent visible misalignments of rotation momentum between various stellar and gaseous galactic subsystems. Especially it must be true for non-cluster lenticular galaxies whose origin should be probably due to minor merger events.  However, findings of extended counterrotating gaseous disks are still rare. In the Sa-galaxy NGC~3626 \\citep{ciri,n3626mol}, in the S0 NGC~4546 \\citep{n4546}, and in the Sb-galaxy NGC~7742 \\citep{sau2,n7742we} all the gas counterrotates the stars. In the Sa-galaxies NGC~3593 \\citep{n3593_1,n3593_2,n3593mol}, NGC~7217 \\citep{mk7217,we7217}, NGC~5719 \\citep{n5719}, and NGC~4138 \\citep{jore4138,n4138thak} the counterrotating gas is already partly processed into stars, so these galaxies have two stellar counterrotating disks one of which corotates the gas. In NGC~4550 one can see already the full-size counterrotating stellar disk whereas the counterrotating gas is mainly exhausted \\citep{n4550_1,n4550_2,n4550we,sau3,sau5}; the similar situation may be suspected in NGC~7331 \\citep{prada,me7331}. In the Sab-galaxy NGC~4826 \\citep{n4826rubin,n4826rix,n4826b92,n4826b94} and in the S0 NGC~1596 \\citep{n1596} the outer gas counterrotates the inner parts of the galaxies, certainly being accreted quite recently. These few examples include almost all known extended counterrotating subsystems. Statistical estimates by \\citet{stat_kan} and \\citet{i_stat} put an upper limit of 8\\% --12\\%\\ of all spiral, S0/a-Scd, galaxies to possess such structures. For S0 galaxies the appearance of counterrotating gas may be more frequent, \\citet{km96} gives the estimate of $24\\pm 8$\\%; it can be consistent with the idea of S0 galaxy (trans-)formation from a spiral by minor merger: in such event some external gas with decoupled momentum must be accreted. But in S0 galaxies the counterrotating gas is observed to be mostly confined to the very central part of the galaxies as it can be seen in the sample by \\citet{b92}; the extended counterrotating gaseous disks are rare in S0s as extended gaseous disks in general.  Another related phenomenon is inner gaseous polar disks in disk galaxies \\citep{inpoldisk}. We found them as well in S0 galaxies with generally small amount of gas \\citep{n7280we,polars0} as in spiral galaxies with normal extended HI disks -- NGC~2841 \\citep{silvb97}, NGC~7217 \\citep{we7217}, NGC~7468 \\citep{n7468}. For the inner polar disk origin the most popular hypothesis is also external gas accretion; however in some cases dynamical simulations predict strongly inclined circumnuclear gaseous disks produced by secular evolution processes. The simulations by \\citet{sec1} of the isolated stellar-gaseous disk evolution gave such `polar disks' as a result of gas redistribution in the global disk of a galaxy, if initially all the gas in the disk counterrotated the stars. Interestingly, the old question, what is the primary, a hen or an egg, is still actual concerning this problem. \\citet{vaks82} obtained the similar configuration, the inner polar disk plus the outer counterrotating gas, starting from a single inner polar disk: in a tumbling triaxial potential the outer parts of the gaseous polar disk warped in their model so that the outer gas counterrotated the stars almost in the main symmetry plane. By paying attention to the outer extension of the inner gaseous polar disks found by us, we have revealed indeed some cases of the configuration required, the inner polar disk plus the more outer counterrotation: these are the lenticular galaxies NGC~7280 and NGC~7332 \\citep{mesau5}, IC~1548 \\citep{n80gr}, and the late-type spiral galaxy NGC~7625 \\citep{arp212}.  An exceptionality of the extended counterrotating disks means that there must exist some additional conditions for a disk galaxy to retain large masses of accreted counterrotating gas; perhaps, it may be a very low rate of the acquisition process \\citep{thakryd}, or an absence of large amount of initial, `own' galactic gas in the recipient galaxy  \\citep{i_stat}. Every new disk galaxy with a globally counterrotating gas component may in principle help to determine these conditions. In this paper we report a discovery of two more extended counterrotating gaseous disks in early-type disk galaxies. We present the results of the kinematical study of the nearby S0 galaxies NGC~2551 and NGC~5631. NGC~2551 and NGC~5631 considered in this paper are early-type disk galaxies of intermediate luminosity. Both belong to spiral-dominated groups \\citep{nog}, and both have a substantial amount of rotating neutral hydrogen with unknown sense of rotation, according to single-dish radioobservations at 21 cm \\citep{bal_h1,h1fr}. Their main parameters retrieved in databases and from literature are presented in the Table~\\ref{tab1}. The layout of the paper is the following. In Section~2 we describe our observations, our data reduction and some additional information. In Section~3 we present the counterrotating gaseous disk in NGC~2551, and in Section~4 -- the complex kinematics including inclined counterrotating stellar-gaseous disk in NGC~5631. Section~5 contains some discussion of the origin and possible future fate of the counterrotating gas in these two galaxies.  \\begin{table} \\caption[ ] {Global parameters of the galaxies} \\label{tab1} \\begin{flushleft} \\begin{tabular}{lcc} \\hline\\noalign{\\smallskip} NGC & 2551 & 5631  \\\\ \\hline Type (NED$^1$) & SA(s)0/a & SA(s)$0^0$\\\\ $R_{25}$, kpc (LEDA$^2$) & 8.5 & 9.3\\\\ $B_T^0$ (RC3$^3$) & 12.78 & 12.35 \\\\ $M_B$ (LEDA)  & --20.0 & --20.2  \\\\ $(B-V)_T^0$ (RC3) & 0.92 & 0.90  \\\\ $V_r $ (NED), $\\mbox{km} \\cdot \\mbox{s}^{-1}$ & 2344  & 1979  \\\\ Distance, Mpc (LEDA) &  37.1 & 32.1 \\\\ Inclination (LEDA) & $50^{\\circ}$ & $21^{\\circ}$  \\\\ {\\it PA}$_{phot}$ (LEDA) & $52^{\\circ}$ & --  \\\\ $V_{rot} \\sin i$, $\\mbox{km} \\cdot \\mbox{s}^{-1}$ (LEDA, HI) & $83.6 \\pm 5.1$ & $165.4 \\pm 14.3$ \\\\ $M_{HI} ^4$, $10^9\\,M_{\\odot}$ & 1.4 & 1.4 \\\\ \\hline \\multicolumn{3}{l}{$^1$\\rule{0pt}{11pt}\\footnotesize NASA/IPAC Extragalactic Database}\\\\ \\multicolumn{3}{l}{$^2$\\rule{0pt}{11pt}\\footnotesize Lyon-Meudon Extragalactic Database}\\\\ \\multicolumn{3}{l}{$^3$\\rule{0pt}{11pt}\\footnotesize Third Reference Catalogue of Bright Galaxies}\\\\ \\multicolumn{3}{l}{$^4$\\rule{0pt}{11pt}\\footnotesize Bettoni et al.(2003)} \\end{tabular} \\end{flushleft} \\end{table} ",
        "conclusions": "Both NGC~2551 and NGC~5631 possess the extended (up to 0.7--1.0$R_{25}$, or up to 5--7 kpc from the center) counterrotating gaseous disks. Both galaxies belong to loose groups dominating by spiral galaxies \\citep{garcia,groups_gh}. However either NGC~2551 nor NGC 5631 have a close neighbor within the circle of 100 kpc radius, to provide interaction and smooth gas accretion that is suggested by \\citet{thakryd} to be the most probable mechanism of massive counterrotating disk formation. The only alternative which is available for NGC~2551 and NGC~5631 is a minor merger with a gas-rich satellite. We do not know what morphological type the galaxies NGC~2551 and NGC~5631 had before their merging, and if they have had their own (corotating) gas. But in any case the accreted gas had to suffer instaneous star formation triggered by shock compression during the merging. Perhaps, in NGC~2551 and NGC~5631 we see different stages of the same process. Ultraviolet imaging with the UIT and later on with the GALEX has revealed an extended, up to $R=40^{\\prime \\prime}$, star-forming disk in NGC~2551 \\citep{uit,galex}. An intensity ratio H$\\alpha$/[NII] observed by us with the SCORPIO implies an excitation by young stars up to $40^{\\prime \\prime}$; however between  $R=40^{\\prime \\prime}$ and $R=50^{\\prime \\prime}$ we see only one emission line, [NII]$\\lambda$6583, so in this ring the gas excitation may be of shock origin. In NGC~5631 there is no current star formation in the counterrotating gaseous disk, and the gas is excited by shock over the full extension of the gaseous disk. Perhaps, it is the preceding evolutionary stage with respect to NGC~2551, the accreted gas disk has not yet settled into the symmetry plane of the main galaxy, and star formation is only going to start in the compressed inner dusty ring at $R=10^{\\prime \\prime} - 15^{\\prime \\prime}$. And what may be the subsequent stage? Perhaps, it is NGC~4138 where a counterrotating extended gas is supplemented by the substantial counterrotating young stellar component \\citep{n4138thak} -- the evident consequence of the star formation in the counterrotating gaseous disk.  To give a conclusion, we summarize that by applying complex spectral methods including integral-field spectroscopy to the central parts of the galaxies and long-slit deep spectroscopy to probe the external parts, we have found two more global gas counterrotating systems in non-interacting early-type disk galaxies. In NGC~2551 two counterrotating disks, gaseous and stellar ones, may be coplanar: the orientation parameters of the optical-band image and of the UV-band, star-formation related image are similar. In NGC~5631 the gaseous disk is inclined by some $35^{\\circ}$ to the main stellar disk; it may contain also significant coupled stellar component. The totality of the spectral and photometric data give evidence for the minor merging as the most probable origin of the counterrotating gas in these galaxies. Perhaps, we observe two different stages of the process of lenticular galaxy formation in rather sparse group environments.   \\begin{figure} \\epsscale{1} \\plotone{f8a.eps}\\\\ \\plotone{f8b.eps} \\caption{The color map (g$'$-r$'$) of NGC 5631, from the SDSS data. The $r'$-band isophotes (top) and schematic view of the both disks (bottom) are overlapped onto the color map. The derived orientation of the inner inclined gaseous/dust disk is $i=35^\\circ$, $PA=122^\\circ$.} \\end{figure}"
    },
    "0812/0812.3934_arXiv.txt": {
        "abstract": "We present new results of our program to systematically search for strongly lensed galaxies in the Sloan Digital Sky Survey (SDSS) imaging data.  In this study six strong lens systems are presented which we have \\emph{confirmed} with follow-up spectroscopy and imaging using the 3.5m telescope at the Apache Point Observatory.  Preliminary mass models indicate that the lenses are group-scale systems with velocity dispersions ranging from $466-878$ km $\\rm{s^{-1}}$ at $z=0.17-0.45$ which are strongly lensing source galaxies at $z=0.4-1.4$.  Galaxy groups are a relatively new mass scale just beginning to be probed with strong lensing.  Our sample of lenses roughly doubles the confirmed number of group-scale lenses in the SDSS and complements ongoing strong lens searches in other imaging surveys such as the CFHTLS \\citep{cabanac07}.  As our arcs were discovered in the SDSS imaging data they are all bright ($r\\lesssim22$), making them ideally suited for detailed follow-up studies. ",
        "introduction": "Gravitational lens systems provide the opportunity to study in detail the properties of distant galaxies through the magnification provided by lensing.  Models of these systems also provide insight into the underlying mass distribution in the foreground lens.  New samples of strong lenses have been enabled by the Sloan Digital Sky Survey \\citep{york00} with systems initially discovered in the SDSS \\emph{spectroscopic} data \\citep{bolton06}.  Because of the nature of the selection these are smaller separation systems ($<3\\arcsec$, the width of the SDSS fiber radius) consisting of a single Luminous Red Galaxy (LRG) lensing a background source galaxy.  An additional parameter space to discover strong lenses is also provided by the SDSS \\emph{imaging} data, and a number of complementary programs are now underway in the SDSS to search for these systems \\citep{belokurov07,estrada07,hennawi08,shin08,wen08}.  Because of the moderate imaging quality (1.4\\arcsec FWHM in r) in the SDSS these systems typically have a larger lens-arc separation ($>3.0\\arcsec$) which is outside of the regime probed by the SDSS spectroscopic lens survey.  Motivated by the discovery of the 8 o'clock arc \\citep{allam07}, the brightest lensed Lyman Break galaxy (LBG) currently known $(z=2.73)$, we have also initiated a program to systematically search for strong lens systems in the SDSS.  Our program is focused on searching for lensed high redshift galaxies in particular $z>2$ lensed LBG's similar to the 8 o'clock arc.  Since Fall 2006 we have been following up candidate systems with imaging and spectroscopy using the 3.5m telescope at Apache Point Observatory (APO) in New Mexico.  The first lensed system discovered in our systematic search, a lensed LBG at $z=2.0$ which we named the `Clone', was recently reported in \\citet{lin08}.  In addition to these lensed high redshift galaxies we are also discovering a number of interesting lensed galaxies at lower redshift.  In this paper we present our first results on six of these lower-redshift systems which we have spectroscopically confirmed to be lensed galaxies, ranging in redshift from $z=0.4-1.4$.  This letter is organized as follows: details of initial candidate selection and follow-up imaging/spectroscopy are described in $\\S \\ref{sec:data}$.  Preliminary mass models of each system are presented in $\\S \\ref{sec:sample}$.  In $\\S \\ref{sec:summary}$ we summarize our work and describe future directions of our survey.  Throughout we assume a standard cosmology of $\\Omega_{m}=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "\\label{sec:summary} We have presented six confirmed group-scale strong lenses from our search program in the SDSS. Although the focus of our survey is to confirm high redshift $(z=2\\sim3)$ lensed galaxies similar to the 8 o'clock arc \\citep{allam07} and the Clone \\citep{lin08} we are discovering a number of lensed galaxies lying at lower redshift.  These systems were initially discovered via two systematic searches of the SDSS DR5 database and from a separate search of a catalog of candidate merging galaxies.  The arcs in our confirmed systems are relatively bright $r<22$, making them ideal candidates for detailed follow-up studies.  Our initial follow-up APO SPICAM imaging has relatively poor resolution but preliminary modeling indicates that these systems are group-scale lenses, an intermediate regime between isolated galaxies and galaxy clusters.  This is a relatively new regime that imaging surveys are just beginning to probe with strong lensing \\citep{cabanac07}.  Additional high resolution imaging will allow for more detailed models of these systems which will yield precise measurements of the mass distribution in the lens and detailed studies of the source galaxy in each system.  Future papers in this series will describe our sample of confirmed lensed $z=2-3$ galaxies as well as the results of our ongoing follow-up programs using HST, Spitzer, and other telescopes."
    },
    "0812/0812.1931_arXiv.txt": {
        "abstract": "{The DAMA collaboration have claimed to detect particle dark matter (DM) via an annual modulation in their observed recoil event rate.  This appears to be in strong disagreement with the null results of other experiments if interpreted in terms of elastic DM scattering, while agreement for a small region of parameter space is possible for inelastic DM (iDM) due to the altered kinematics of the collision.  To date most analyses assume a simple galactic halo DM velocity distribution, the Standard Halo Model, but direct experimental support for the SHM is severely lacking and theoretical studies indicate possible significant differences.   We investigate the dependence of DAMA and the other direct detection experiments on the local DM velocity distribution, utilizing the results of the Via Lactea and Dark Disc numerical simulations.  We also investigate effects of varying the solar circular velocity, the DM escape velocity, and the DAMA quenching factor within experimental limits. Our data set includes the latest ZEPLIN-III results, as well as full publicly available data sets.   Due to the more sensitive dependence of the inelastic cross section on the velocity distribution, we find that with Via Lactea the DAMA results can be consistent with all other experiments over an enlarged region of iDM parameter space, with higher mass particles being preferred, while Dark Disc does not lead to an improvement.  A definitive test of DAMA for iDM requires heavy element detectors. } ",
        "introduction": "First-hand information on the nature of the dark matter (DM) in our universe may be revealed by one or more of the current generation of direct detection experiments.  The DAMA collaboration have measured an annual modulation in their scintillation event rate, initially with the DAMA/NaI set-up \\cite{Bernabei:2005hj}, and more recently with the DAMA/LIBRA configuration \\cite{Bernabei:2008yi}, which they interpret as being due to the modulation in the DM-nuclear scattering event rate following from the annual variation of the Earth's velocity through the DM in our galaxy.  At face value, this measurement appears to be in disagreement with other direct detection experiments if we make the usual assumptions about the nature of the DM; namely that the DM is a weakly interacting massive particle (WIMP) that elastically scatters off the target material with a spin independent interaction.  Inelastic dark matter (iDM) \\cite{TuckerSmith:2001hy} was proposed as a way to reconcile the positive result from DAMA/NaI with the null result from germanium based detectors.  In the iDM scenario, WIMP-nucleon elastic scattering is suppressed, while inelastic scattering from a ground-state WIMP to a slightly higher mass excited WIMP is allowed and dominates the recoil event rate.  As we summarize in Section \\ref{basics}, the kinematics of the recoil scattering are changed by the inelastic nature of the collision, and this can bring the DAMA results closer to agreement with the other experiments for three principal reasons: \\begin{itemize} \\item Heavier nuclei are favoured. \\item The recoil spectrum is changed at low energies. \\item The ratio of modulated to unmodulated signal is higher. \\end{itemize} Nevertheless the other direct detection experiments still strongly constrain the DM interpretation of DAMA/LIBRA, with, apparently, only a relatively small region of parameter space being allowed \\cite{Chang:2008gd}.  One unchecked assumption that goes into the analysis is that the velocity distribution of the DM in the galactic rest frame is well described by so-called Standard Halo Model (SHM).   The SHM assumes that in the galactic frame the WIMP distribution is an isotropic isothermal sphere, which leads to an essentially structureless isotropic Maxwell-Boltzmann velocity distribution with dispersion set by the local circular velocity. This is a questionable assumption: Little is known about the phase-space distribution of DM on the scales relevant for direct detection experiments, and as we outline in Section \\ref{haloes} numerical simulations of DM distributions for Milky-Way-like galaxies lead to results differing from the SHM in potentially significant ways, especially for iDM which is more sensitive to the DM velocity distribution.  In this paper, we compare the results obtained using the SHM to those obtained using two recent computer simulations of the DM distribution in a Milky-Way-like galaxy: \\begin{itemize} \\item Via Lactea \\cite{Diemand:2006ik} - a Milky Way size DM halo distribution containing 234 million particles. \\item Dark Disc \\cite{Read:2008fh} - a simulation which contains in addition to the SHM, a slowly rotating disc of DM \\end{itemize} With these less idealised halo models, we find that the iDM scenario allows the DAMA region to be consistent with all other experiments, and that in the Via Lactea distribution, more parameter space can be opened up at high WIMP masses, compared to the SHM.   After reviewing the details of each experiment and describing how we calculate the allowed DAMA region and exclusion curves for the null experiments in Section \\ref{exper}, we present our results together with their physical interpretation in Section \\ref{Results}. ",
        "conclusions": "\\label{Conclusions}  We have shown, in the context of iDM, that the region of agreement between the DAMA data and results from other experiments is sensitive to the uncertainties present in the galactic WIMP velocity distribution.  In particular we have found that the other direct detection experiments, including the most recent ZEPLIN-II and III, CRESST-II, XENON10, KIMS, and CDMS-II data sets, do not exclude the region of parameter space preferred by the DAMA results up to WIMP masses $M_\\chi \\sim \\text{1 TeV}$ when the $\\text{VL}_{220}$ velocity distribution is used (see Figures \\ref{deltaplots2} and \\ref{massplots}), while for the $\\text{VL}_{270}$ velocity distribution, WIMP masses $M_\\chi \\leq \\text{200 GeV}$ are allowed (see Figures \\ref{VL270deltaplots} and \\ref{VL270massplots}).  Furthermore, we have also argued that the region of agreement between experiments is very sensitive to the quenching factor used to interpret the DAMA data (see Figure \\ref{quenchfig}), and also to the local galactic escape velocity (see Figure \\ref{escape}), both of which presently have $\\sim$10\\% uncertainties, and somewhat to the Sun's circular velocity (see Figure \\ref{circ}).  Independent of experimental set-ups we have also shown that, in the iDM scenario, detectors would be insensitive to WIMPs traveling with a small velocity relative to the sun due to the increased minimum velocity for scattering.  This would lead to dark matter clumps or streams being undetectable if rotating in the galactic plane, as in the Dark Disc halo model, or streaming with a small relative velocity (see Figure \\ref{darkdisc}).  New data from heavy element experiments such as CRESST-II and XENON100 could have the potential to completely exclude the DAMA preferred region for iDM, especially if taken during the modulation maximum.  In particular, if the planned EURECA experiment were to use a tungsten detector it would provide significant limits on both the elastic and iDM scattering cross-sections.  Furthermore because of the low energy cut-off in the recoil energy spectrum for iDM, it is important that the experiments increase their sensitivity at higher recoil energies, where the iDM recoil spectrum peaks.  Without a better understanding of the details of the velocity distribution in our galaxy, including the effects of the baryons, precise statements about the consistency of various direct detection experiments are not possible. \\\\  \\noindent{{\\bf Note added:} Subsequent to the submission of this work two recent DM simulations were brought to our attention:  \\begin{itemize} \\item  Recently the results of a second simulation, Via Lactea II \\cite{Kuhlen:2008qj}, were made public.  These results showed further detail in the substructure of the DM halo, with the clumpiness extending over six orders of magnitude, down to the resolution limit.  \\item  Another Milky Way-sized DM halo simulation, the Aquarius Project, also recently released their results \\cite{Springel:2008cc}, and the phase-space structure of these simulated halos was the subject of further study \\cite{Vogelsberger:2008qb}.  The velocity distributions found varied from simulation to simulation and showed interesting bumpy features.  Further to this, in \\cite{Vogelsberger:2008qb} the influence of these velocity distributions on direct detection signals for elastically scattering dark matter was investigated and it was found that the recoil spectra could deviate by up to 10\\% from that expected from the best-fit multivariate Gaussian.  \\end{itemize}  These simulations provide further strong support for deviations from the SHM, but as data from these simulations are not (to the authors knowledge) publicly available, we have not been able to update our study to include them, although we hope to return to this subject in a future publication."
    },
    "0812/0812.2932_arXiv.txt": {
        "abstract": "An interesting strategy for indirect detection of Dark Matter comes through the amounts of electrons and positrons usually emitted by DM pair annihilation. The $e^+e^-$ gyrating in the galactic magnetic field then produce secondary synchrotron radiation. The radio emission from the galactic halo as well as from its expected substructures if compared with the measured diffuse radio background can provide constraints on the physics of WIMPs. In particular one gets the bound of $\\sigva= 10^{-24}$\\,cm$^3$s$^{-1}$ for a DM mass $m_{\\chi}=100$\\,GeV even though sensibly depending on the astrophysical uncertainties. ",
        "introduction": "Among the indirect DM detection channels, the radio emission due to secondary electrons or positrons can represent a chance to look for DM annihilation. During the process of thermalization in the galactic medium the high energy $e^+$ and $e^-$ release secondary low energy radiation, in particular in the radio and X-ray band, which in principle could be detected. Furthermore, while the astrophysical uncertainties affecting this signal are similar to the case of direct $e^+$, $e^-$ detection, the sensitivities are quite different, and, in particular, the radio band allows for a the discrimination of tiny signals even in a background many order of magnitudes more intense.  Indirect detection of DM annihilation through secondary photons has received recently an increasing attention, exploring the expected signature in X-rays \\cite{Bergstrom:2006ny,Regis:2008ij,Jeltema:2008ax}, at radio wavelengths \\cite{Blasi:2002ct,Aloisio:2004hy,Tasitsiomi:2003vw,Zhang:2008rs} , or both \\cite{Colafrancesco:2005ji,Colafrancesco:2006he,Baltz:2004bb}. In this paper we will focus our analysis on the radio signal expected from the Milky Way (MW) halo and its substructures. The results and the details of the following approach can be found in Ref. \\cite{Borriello:2008gy}. ",
        "conclusions": "Using conservative assumptions for the DM distribution in our galaxy we  derive the expected secondary radiation due to synchrotron emission from high energy electrons produced in DM annihilation. The signal from single bright clumps offers only poor sensitivities because of diffusion effects which spread the electrons over large areas diluting the radio signal. The diffuse signal from the halo and the unresolved clumps is instead relevant and can be compared to the radio astrophysical background to derive constraints on the DM mass and annihilation cross section.  Constraints in the radio band, in particular, are complementary to similar (less stringent but less model dependent) constraints in the X-ray/gamma band \\cite{Mack:2008wu,Kachelriess:2007aj} and from neutrinos \\cite{Yuksel:2007ac}. Radio data, in particular, are more sensitive in the GeV-TeV region while neutrinos provide more stringent bounds for very high DM masses ($\\agt 10$ TeV). Gammas, instead, are more constraining for $m_\\chi \\alt 1$ GeV. The combination of the various observations provides thus interesting constraints over a wide range of masses  pushing the allowed window significantly near the thermal relic possibility.    More into details, we obtain conservative constraints at the level of $\\sigva \\sim 10^{-23}$\\,cm$^3$s$^{-1}$ for a DM mass $m_{\\chi}=100$\\,GeV from the WMAP Haze at 23 GHz. However, depending on the astrophysical uncertainties, in particular on the assumption on the galactic magnetic field model, constraints as strong as $\\sigva \\sim 10^{-25}$\\,cm$^3$s$^{-1}$ can be achieved. Complementary to other works which employ the WMAP Haze at 23 GHz, we also use  the information in a wide frequency band in the range 100 MHz-100 GHz. Adding this information the constraints become of the order of $\\sigva \\sim 10^{-24}$\\,cm$^3$s$^{-1}$ for a DM mass $m_{\\chi}=100$\\,GeV. The multi-frequency approach thus gives comparable constraints with respect to the WMAP Haze only, or generally better for $m_\\chi\\alt100$ GeV where the best sensitivity is achieved at $\\sim$ GHz frequencies.   The derived constraints are quite conservative because no attempt to model the astrophysical background is made differently from the case of the WMAP Haze. Indeed, the Haze residual map itself should be interpreted with some caution, given that the significance of the feature is at the moment still debated and complementary analyses from different groups (as the WMAP one) miss in finding a clear evidence of the feature. Definitely the multi-frequency approach will be necessary to test in a convincing way a possible DM signal like the claim related to the WMAP Haze. Progresses are expected with the forthcoming data at high frequencies from Planck and at low frequencies from LOFAR and, in a more distant future, from SKA. These surveys will help in disentangling the various astrophysical contributions thus assessing the real significance of the Haze feature. Further, the low frequency data in particular, will help to improve our knowledge of the galactic magnetic field. Progresses in these fields will provide a major improvement for the interpretation of the DM-radio connection.   \\vspace{1pc}"
    },
    "0812/0812.0795_arXiv.txt": {
        "abstract": "This is a report on the findings of the dark matter science working group for the white paper on the status and future of TeV gamma-ray astronomy. The white paper was commissioned by the American Physical Society, and the full white paper can be found on astro-ph (arXiv:0810.0444). This detailed section discusses the prospects for dark matter detection with future gamma-ray experiments, and the complementarity of gamma-ray measurements with other indirect, direct or accelerator-based searches. We conclude that any comprehensive search for dark matter should include gamma-ray observations, both to identify the dark matter particle (through the characteristics of the gamma-ray spectrum) and to measure the distribution of dark matter in galactic halos. \\vspace*{1.5cm} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.0040_arXiv.txt": {
        "abstract": "Magnetic field modifies the properties of waves in a complex way. Significant advances has been made recently in our understanding of the physics of waves in solar active regions with the help of analytical theories, numerical simulations, as well as hi-resolution observations. In this contribution we review the current ideas in the field, with the emphasis on theoretical models of waves in sunspots. ",
        "introduction": "%   Photospheres and chromospheres of sunspots are the regions where different physical agents come into play with almost equal weight. The restoring forces in the wave equation, such as magnetic Lorentz force, gas pressure gradient and buoyancy, are of the same order of magnitude, making arbitrary the division into pure wave modes and complicating theoretical models. A recent summary on the observational properties of sunspot oscillations together with some theoretical issues is presented by \\citet{Bogdan+Judge2006}. The most recent observational results suggest that wave phenomena at different layers of the sunspots umbrae and penumbrae are related with each other \\citep[see \\eg][]{Marco+etal1996, Maltby+etal1999, Maltby+etal2001, Brynildsen+etal2000, Brynildsen+etal2002, Christopoulou+etal2000, Christopoulou+etal2001, Rouppe+etal2003, Tziotziou+etal2006, Bloomfield+etal2007}. The interpretation of the observational material defines several groups of questions concerning the sunspot wave physics that can be summarized in the following way: \\begin{itemize} \\item What drives the waves observed in sunspots? Are they externally driven by the quiet Sun $p$-modes? Are there sources of oscillations inside the umbra due to weak convection? \\item How can the oscillations observed at different sunspot layers be interpreted in terms of MHD waves? What are the relationships between photospheric and chromospheric oscillations? What causes the complex spatial pattern of oscillations such as chromospheric umbral flashes, penumbral waves, spatial coherency of waves over the umbra, etc.? \\item Why is the wave power suppressed in the umbral photosphere compared to the quiet Sun? Why is it enhanced in the chromosphere? \\item What mechanisms produce the change of the dominating frequency of waves in the umbra from 3 mHz in the photosphere to 5--6 mHz in the chromosphere? What are the reasons for the spatial power distribution at different frequency intervals observed in the umbra and penumbra? \\item Are there observational evidences of the mode transformation in sunspots? \\item What is the source of helioseismological velocity signal detected in active regions? Are they due to fast, slow MHD waves? What are the consequences of the strong magnetic field of sunspots onto helioseismology measurements and inversions of sub-photospheric structure of sunspots? \\end{itemize} Some of these questions, like the interpretation of the wave modes or change of the frequency with height, received larger attention in the literature in the past and can be considered answered to a smaller or larger extent. Some others, like those concerning local helioseismology in active regions, only received attention recently and no unique answer has been found yet. In the rest of paper we will describe the existing models and the theoretical progress made in understanding sunspot oscillations. ",
        "conclusions": "Theoretical interpretation of waves observed in sunspots has made big advances during the last decades. The influence of the $\\beta=1$ level and the magnetic field inclination on the observed wave properties is being clarified, allowing the interpretation of the observed oscillations in term of MHD wave modes. The observed wave pattern in the upper photosphere and the chromosphere is found to be compatible with the low-$\\beta$ acoustic-like slow mode propagation along the inclined magnetic field lines. Several physical effects are proposed to be responsible for the frequency change with height in the umbra from the 3 mHz in the photosphere to 5--6 mHz in the chromosphere. The mode transformation at the $\\beta=1$ level is proposed to be responsible (at least in part) for the wave power distribution over the active regions, including the wave suppression in the umbra and, possibly, enhanced power of acoustic halos. Theoretical interpretation of the local helioseismology measurements including magnetic field has been initiated. The first simulations of helioseismic waves in sunspots show a general agreement of the wave travel times of the high-$\\beta$ fast modes with observations. Still, the final issue in this field is far from clear and the major developments are expected in the future. The 3-dimensional wave propagation and mode transformation in sunspots are important to analyze from simulations, including the chromospheric layers."
    },
    "0812/0812.0865_arXiv.txt": {
        "abstract": "In the $J$, $H$, and $K_{S}$ bands survey of the the Galactic Center region over an area of \\timeform{2D}  $\\times$ \\timeform{5D}, we have found many dark clouds, among which a distinguished chain of dark clouds can be identified with a quiescent CO cloud. The distances of the clouds is estimated to be 3.2-4.2 kpc, corresponding to the Norma arm by our new method to determine distance to dark clouds using the cumulative number of stars against $J-K_{S}$ colors. Adopting these estimated distances, the size is about $\\sim$70 pc in length and the total mass of the cloud is 6$\\times$10$^{4}$M$\\solar$. Three compact HII regions harbor in the cloud, indicating that star forming activities are going on at the cores of the quiescent CO cloud on the spiral arm. ",
        "introduction": "The central region of our Galaxy suffers large interstellar extinction, and images taken toward the Galactic Center (GC) in the optical wavelengths are generally featureless star fields with few stars. Since the discovery of an infrared source corresponding to the GC by Becklin and Neugebauer (1968), the GC region has been observed extensively in the infrared, where interstellar extinction is one order of magnitude smaller than in the optical. Succeeding to the larger-scale near-infrared (NIR) surveys of 2MASS (Skrutskie et al. 2006) and DENIS \\citep{Epchtein94}, which have given a general view of the Galactic structure, we have carried out a $J$, $H$, and $K_{S}$ bands survey of the $\\timeform{2D}$  $\\times$ $\\timeform{5D}$ region centered on the GC with 2 mag deeper sensitivity and 2 times finer resolution than those of 2MASS. In the course of the survey, we found numerous dark clouds silhouetted against the dense star fields.      Toward the GC, existence of many molecular clouds is well known from radio observations (Morris and Serabyn 1996 and references therein). Their location in the line of sight is uncertain, but part of them, located on the near side of the GC, should appear as dark clouds in the near infrared images.  However, the correspondence of near infrared dark clouds to radio molecular clouds has not been studied well.    The distance to dark/molecular clouds is crucial for discussing their physical parameters associated with them. In radio observations, kinematic distances to molecular clouds are often derived from their radial velocity on the assumption that each cloud follows the circular Galactic rotation. However, since the Galactic rotation is perpendicular to our line of sight to the GC, we cannot determine the distance to the molecular clouds. Moreover, non-circular components such as expanding rings make situations complicated.  Sofue (2006) proposed a method to determine the distance of CO clouds in the direction of GC from the gradient in the $l-v$ diagram, and determined the distance of CO clouds toward the GC. This method is powerful for determining the distance to arm-like structures, which extend long in the $l$ direction, but unsuited for small CO clouds because the velocity gradient is not determined well.  In this paper, we report a chain of dark clouds found in our near infrared survey, which shows good positional and morphological agreements with a velocity component of $^{12}$CO ($J$=1-0) observed with the Nobeyama radio telescope (Oka et al. 1998). A new method to determine the distance to dark clouds from a diagram of $J-K_S$ color and cumulative star number, and the resultant distance to the chain of dark clouds are presented. ",
        "conclusions": "We estimated a projected length of the chain of dark clouds to be $\\sim$70 pc and the typical thickness of the filaments is as thin as 0.1 degree, or about 7 pc, located at $\\sim$30 pc off from the Galactic plane. The gaseous filaments are themselves divided into finer structures, and several high-intensity knots are seen on the ridge of the $^{12}$CO contours. By integrating the $^{12}$CO excess in the Fig. 2, the masses of each ridges are estimated to be an order of $\\sim$10$^{4}$M\\solar, and the total mass for the entire CO cloud results in 6$\\times$10$^{4}$M\\solar. If we assume the same thickness in the line of sight (7 pc) as the apparent thickness and an extent of 70 pc, this total mass yields molecular gas density on the order of 18 M\\solar pc$^{-3}$, or 4$\\times$10$^{2}$ H$_{2}$ cm$^{-3}$. This is a typical density for normal molecular clouds.  In the three high-intensity knots, we found four NIR sources, SF1-A, SF1-B, SF2, and SF3, with extents of $\\sim10\\arcsec$ to $\\sim40\\arcsec$ or 0.2-0.8 pc at 4 kpc. Pseudo color images and position of them are shown in Fig. 7 and Table 2. Mid infrared sources (MSX-PSC, \\cite{Egan96}) are associated with all of them. OH, H$_{2}$O, and/or CH$_{3}$OH masers (SIMBAD, or \\cite{Caswell83a}, \\cite{Caswell83b} and \\cite{Menten91}) are also associated with SF1-A, SF1-B, and SF3 with $V_{LSR}$=$+15$ to $+23$ km s$^{-1}$, quite consistent to $V_{LSR}$ of $^{12}$CO. Both their sizes and  $K_{S}$ magnitudes in 2MASS Point Source Catalog (9.15, 7.67, and 7.80 for SF1-A, SF2, and SF3) are consistent with those of typical compact HII regions at a distance $\\sim$4 kpc.% We infer that the dark/molecular cloud discovered in this survey are molecular clouds in the Norma spiral arm at 4 kpc from us.  \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){fig5.eps} \\end{center} \\caption{Positions of RF1-2 and DC1-8 in the galactic coordinate, together with $^{12}$CO contour. } \\label{fig:dc_rf_pos} \\end{figure}   \\begin{figure*} \\begin{center} \\FigureFile(160mm,120mm){fig6.eps} \\end{center} \\caption{Cumulative star-number histogram against $J-Ks$ toward two reference fields(RF1 and RF2) and eight DCs (DC1-8) are drawn. The dashed line indicates the predicted number from the Galaxy model. For RF1 and RF2, they are calculated with no  screen (interstellar extinction only). For DC1-8, the best-fit combination of the $D$ and $A_{K_S}$ are drawn. The best-fit parameters are written at bottom in each panel. } \\label{fig:model_comp_dc1-8_J-K} \\end{figure*}"
    },
    "0812/0812.0748_arXiv.txt": {
        "abstract": "We construct a mass model for the spiral lens galaxy 2237+0305, at redshift $z_l$=0.04, based on gravitational-lensing constraints, HI rotation, and new stellar-kinematic information, based on data taken with the ESI spectrograph on the 10m Keck-II Telescope.  High resolution rotation curves and velocity dispersion profiles along two perpendicular directions, close to the major and minor axes of the lens galaxy, were obtained by fitting the Mgb-Fe absorption line region. The stellar rotation curve rises slowly and flattens at $r\\sim{1.5}\\arcsec$ ($\\sim$1.1 kpc). The velocity dispersion profile is approximately flat. A combination of photometric, kinematic and lensing information is used to construct a mass model for the four major mass components of the system --- the dark matter halo, disc, bulge, and bar. The best-fitting solution has a dark matter halo with a logarithmic inner density slope of $\\gamma$=0.9$\\pm$0.3 for $\\rho_{\\rm DM} \\propto r^{-\\gamma}$, a bulge with M/L$_B$=6.6$\\pm$0.3$\\Upsilon_\\odot$, and a disc with M/L$_B$ =1.2$\\pm$0.3$\\Upsilon_\\odot$, in agreement with measurements of late-type spirals.  The bulge dominates support in the inner regions where the multiple images are located and is therefore tightly constrained by the observations.  The disc is sub-maximal and contributes $45\\pm11$ per cent of the rotational support of the galaxy at 2.2r$_d$. The halo mass is $(2.0\\pm0.6)\\times10^{12} M_{\\odot}$, and the stellar to virial mass ratio is $7.0\\pm2.3$ per cent, consistent with typical galaxies of the same mass. ",
        "introduction": "Strong gravitational lensing has been used extensively to model galactic- and cluster-mass systems \\citep[e.g.][]{maller00,trott02,winn03,sand04,bradac06,halkola06,yoo06,jiang07,limousin07,bolton08b,grillo08}. The combination of lensing information with photometric and kinematic data can be used to break degeneracies inherent to each source of information (e.g. the mass-sheet degeneracy and the mass-profile anisotropy degeneracy) to produce a well-constrained mass model for the system \\citep[e.g.][]{koopmans98,winn03,koopmans03,treu04,czoske08,koopmans06}. This is particularly important in spiral galaxies which have multiple mass components that have to be modelled simultaneously, and the limitations of individual methods are particularly stringent. Specifically, in the absence of kinematic information, lens image positions can accurately constrain the enclosed mass, and to some extent also constrain the mass distribution up to a constant mass-sheet \\citep{gorenstein88,wambsganss94,wucknitz02}.  When assuming a stellar mass-to-light ratio for the different optical components (e.g. disc and bulge), based, for example, on stellar-population synthesis models, the dark matter contribution to the galaxy can then be inferred. In the absence of lensing information, but with kinematic data in hand, the overall galactic mass distribution is known, but again a stellar mass-to-light ratio must be assumed for the optical components in order to infer the contribution of the dark matter \\citep[the disc-halo degeneracy,][]{perez04,verheijen07}. Because of the rising contribution of the stellar mass components, this becomes increasingly more difficult towards the inner regions of the galaxy. With the combination of lensing and high-resolution kinematic information, however, the stellar mass-to-light ratios of the optical components can be further constrained, breaking some of the degeneracies between mass, mass-slope and anisotropy \\citep[e.g.][and references therein]{trott02,treu04}.  An accurate model for the mass distribution in spiral galaxies can provide interesting information about both the baryonic and the non-baryonic components. Disentangling these components provides information about the mass-to-light ratios as well as the influence, shape and density profile of the dark matter.  The main motivations of this work are to: (1) break the classical degeneracy between the stellar component and the dark-matter halo in spiral galaxies \\citep{vanalbada86}, using the additional information that lensing provides; (2) determine independent mass-to-light ratios for the bulge and disc components; (3) determine whether the disc component is ``maximal'' \\cite[e.g.][]{sackett97,courteau99}. In more general terms, understanding the relative mass of the different components can help constrain models for disc galaxy formation \\citep[e.g.][and references therein, see also \\citet{shin07}]{dutton05,dutton07}. In addition, the presence of massive discs in lens galaxies gives rise to interesting perturbations to the standard elliptical models adopted for early-type galaxies with a number of interesting consequences for the topology of the time delay surface \\citep{moller03}.  In this paper, we address these issues by studying the gravitational lens galaxy 2237+0305.  Newly measured rotation curves and line-of-sight velocity dispersion profiles along two directions close to the major and minor axes -- obtained from Keck-ESI spectroscopy -- are combined with photometric and lensing data from HST to construct a detailed mass model for the system. This is used to separate luminous and dark matter over scales ranging from the central parts dominated by the bulge out to several exponential scale lengths of the disc component, with the aid of HI data.  The lens system 2237+0305 -- chosen for this study due to the low redshift of the lens galaxy $z_l=0.04$ -- has been observed extensively since its discovery in 1985 \\citep{huchra85}. In addition to providing information about the global mass structure of the galaxy \\citep{schneider88,kent88,mihov01,trott02}, it has also been used to study the properties of stars, or compact objects, in the lensing galaxy through microlensing \\citep[e.g.][]{gilmerino05}, and the structure of the broad-line region of the lensed quasar \\citep{wayth05,vakulik07,eigenbrod08a}. Previously published spectra of 2237+0305 have either concentrated on studying the quasar spectrum, or only presented a small wavelength range for the galaxy spectrum. The discovery paper, \\citet{huchra85}, presented a quasar spectrum taken with the MMT and featured both the CIV and CIII] broad lines. In addition, they identified the H+K break at $\\sim$4100$\\ang$ and Mgb absorption (5381$\\ang$). This spectrum was not of sufficient quality to identify any further features, but was successful in its primary aim: to confirm the system as a lens. \\citet{foltz92} observed the galaxy to determine a velocity dispersion for the bulge. They identified Mgb and FeI absorption features, but also only presented a spectrum over a small wavelength range (5100--5550$\\ang$). \\citet{lewis98} concentrated their observations around the quasar emission lines to investigate variations between the four images, thereby inferring microlensing effects (line and continuum regions are expected to behave differently under microlensing conditions because the flux emanates from regions of different source size). The system 2237+0305 has also been the subject of several microlensing monitoring experiments to study the size of the quasar emitting regions \\citep{wozniak00,udalski06,eigenbrod08b}. \\citet{rauch02} identified absorption lines in the spectrum from intervening (between lens and source) objects using the HIRES spectrograph on Keck. These spectra display high-resolution absorption line profiles of the intervening systems and as such do not concentrate on the spectrum of the galaxy as a whole.  \\citet{vandenven08} have recently obtained GMOS IFU data of the CaII near-infrared triplet from the central region of 2237+0305 to determine the contribution of dark matter in the central regions of the galaxy. They demonstrated that within the central projected 4$\\arcsec$, dark matter can contribute no more than 20\\% of dynamical mass, using single stellar population models to constrain the stellar mass contribution, and both lensing and inferred kinematics to constrain the total dynamical mass.  Our results agree with those of \\citet{vandenven08} in the region of overlap. However, our methodology and aims are significantly different, as our goal is to construct a simply parametrized model to describe the data over a tenfold radial extent.  In a previous analysis of the system, \\citet{trott02} combined lensing and photometric data to constrain the inner mass distribution of the galaxy. In addition, neutral hydrogen rotation points at large radius were used to study the influence of the outer dark matter halo. They found that additional kinematic information was required in the image region to break the remaining model degeneracies. Specifically, the inner logarithmic slope of the dark matter halo, a contentious issue \\citep[see for example][]{hayashi04,deblok05}, could not be constrained without additional information about the inner regions of the galaxy.  In contrast, this paper presents mass-to-light ratios for the bulge and disc components that do not require stellar population-synthesis models to constrain the stellar M/L ratios, and are therefore \\textit{independent} measurements. The only assumption is that the mass-to-light ratio is constant across the component. This also allows us to quantitatively assess the contribution and density profile of the inner dark-matter halo in this high-mass spiral galaxy, complementing work done in dwarf and low surface brightness spiral galaxies \\citep{deblok01,borriello01,deblok05,zackrisson06}. Furthermore, whereas in \\citet{trott02} the halo was constrained to have a softened isothermal shape due to the absence of kinematic data, in this paper we have allowed for a more flexible model.  The paper is organized as follows. Section \\ref{kinematics} describes spectroscopic observations of 2237+0305 with ESI on Keck, presents the resulting optical spectrum and uses absorption features to derive both an optical rotation curve and line-of-sight velocity dispersion profile along perpendicular axes. Section \\ref{galaxy_DF} introduces the modelling: Section \\ref{mass_models} describes the mass density models used for each component; a two-integral axisymmetric model of the galaxy that is used to generate the observed kinematics is introduced in Section \\ref{kine_models}; then Section \\ref{algorithm} describes the algorithm used to find the solution.  Sections \\ref{results} and \\ref{discussion} present the results of the modelling and some discussion, and Section \\ref{conclusion} concludes. We assume $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\Lambda}=0.7$ and H$_0=70$ kms$^{-1}$Mpc$^{-1}$ throughout this paper. ",
        "conclusions": "\\label{discussion}  The constraints used to model the galaxy in \\citet{trott02} were image positions and the two neutral hydrogen rotation points at large radius (and well away from the image positions). Given that no kinematic information was available in the central regions, the inferred velocity curve due to the four mass components was assumed to be circular. This means that an acceptable fit to the data was attainable with the four smoothly modelled mass components because the constraints effectively lie at only two radii --- 0.9$\\arcsec$ (images) and $\\sim$30$\\arcsec$ (HI data). There were considerable degeneracies in the model at other radii \\citep[e.g. see Figure 2 from][]{trott04}.  In the present analysis, the addition of kinematic data provides additional constraints out to a radius of $\\sim$5$\\arcsec$ ($\\sim$4 kpc). As displayed in Figure \\ref{best_fit_rc_no_prior}, constraining information is available in the central 5$\\arcsec$ and at two outer points in the galaxy, leaving the models to fit the remainder of radii according to their input density profiles. The bulge is well-constrained because it contributes considerably to the mass within the inner 5$\\arcsec$. The mass-to-light ratios of this study are more tightly constrained than those of \\citet{trott02} --- their 1$\\sigma$ uncertainties were calculated while holding the values of other parameters at their minimum, thereby underestimating the true uncertainties. The bulge mass-to-light ratio is consistent with previous measurements from early-type galaxies M/L$_B$=7.9$\\pm$2.3$\\Upsilon_{\\odot,h=0.70}$. \\citep{gerhard01,treu04}. Similarly, \\citet{jiang07} report M/L$_B$ = 7.2$\\pm$0.5$\\Upsilon_{\\odot,h=0.70}$ from an analysis of 22 lens galaxies.  As far as the disc mass-to-light ratio is concerned, our best fitting value is also consistent with typical measurements reported in the literature \\citep[e.g.][and references therein]{courteau99}.  The dark matter halo, however, contributes mostly in the outer regions of the galaxy where few constraints are available. In addition, unlike the bulge and disc, the absence of a light profile to guide the density profile of the dark matter yields additional freedom in the model. This allows a more flexible model, but also increases degeneracy. It is thus remarkable to see how well the dark matter halo is constrained by the available data.  However, more detailed information about the structure of the dark matter halo would require information at radii greater than 6$\\arcsec$, where its influence increases. Alternatively, to gain more information about the halo from these data, one could place a prior on the bulge and disc mass-to-light ratios based on studies of other galaxies (or on stellar synthesis models) and determine whether such constraints narrowed the allowed value of the inner slope. This is left for future work.  The best-fitting solution described in \\citet{trott02} had a more massive disc and less massive halo compared with the best-fitting solution described here. The bulge mass was consistent between the two studies, within quoted errors. The errors quoted in \\citet{trott02} were calculated assuming no parameter covariances, and therefore underestimate the true errors as calculated by marginalising over the remaining parameters (as performed in this work). Therefore it is difficult to determine whether the halo and disc masses (and therefore the disc mass-to-light ratio) are consistent between the two studies. Note, however, that the dark matter halo mass profile used in this study was different to that in \\citet{trott02}, allowing freedom in the inner slope.  The fit to the image positions in this study is not as good as the fit in \\citet{trott02}. This is due to the increased complexity of the model in this work, and the need for the fit to balance the image positions and the kinematics. In \\citet{trott02}, the image positions formed the most precise portion of the data (the HI points, being far from the galactic centre and with relatively large error bars are easier to fit), and therefore were able to be well-fitted. In this work, the first and second order kinematics formed a relatively larger proportion of the $\\chi^2$ calculation, and therefore were fitted at the expense of the image positions.  In future work the modelling could be improved to include any effects of external shear (G2237+0305 has other less massive systems identified along the line-of-sight) and to include the kinematic effects of the bar more consistently. The former could improve the lensing component of the $\\chi^2$ without affecting the dynamical fit. In addition, the dark matter halo could be modelled to be non-spherical with flattening aligned with that of the disc (as suggested by studies of other lensing systems).  Despite the comparatively good model presented in this present work, it is important to keep in mind the following caveats. The lens galaxy 2237+0305 has additional structure that is not accounted for in the four component model, e.g. prominent spiral arms that are attached to the bar component. In addition, the bar is not treated in a self-consistent manner.  These may not be important for the lensing but still produce anisotropy in the kinematics. The bar is difficult to model kinematically and an attempt has not been made in this work. However, in the central few arcseconds its presence is small and we opted to neglect it for the inner kinematics, including its effect only in the lensing and for the overall rotation curve.  The image magnifications and total magnification are dependent on the balance between the steepness and the anisotropy of the overall mass model. As demonstrated by \\citet{wambsganss94}, 2237+0305 allows a range of models with degeneracy between the density slope and ellipticity. Inclusion of kinematics in our model has reduced that degeneracy. The magnification ratios are consistent with the mid-IR measurements of \\citet{agol00}. They argued that mid-IR measurements are the least likely to be contaminated by stellar microlensing, electron scattering and extinction \\citep[see also][]{chiba05,agol06}.  In terms of global properties, the central galaxy is bulge-dominated, and modelling the mass distribution of the bulge accurately is most important for an accurate overall model. Beyond $r\\sim$20$\\arcsec$, the halo appears to continue to increase in rotational support in the outer regions, and to absorb the rotation beyond the optical radius.  Our results can be compared with those of \\citet{vandenven08}. Their measured central velocity dispersion ($166\\pm2$ kms$^{-1}$) is consistent with that presented in this work ($172\\pm9$ kms$^{-1}$) within the errors. Although they had the advantage of two-dimensional data and the additional information that accompanies that, our data are based on stellar kinematics from the Mgb-Fe region, observed in a cleaner region of the sky than the CaII triplet targeted by their observations. This allowed us to reach further out in projected distance from the centre, albeit along preferred directions.  The results of \\citet{vandenven08} are complementary to ours: they fit a fourier-based lens model, fitting the image positions and radio flux ratios, and combine the information with the implied mass distribution from the kinematics and light distribution, to obtain an estimate of the luminous mass-to-light ratio. A comparison of this value with that obtained using a single stellar population model yields a maximum dark matter contribution within the Einstein radius of 20\\%, which is higher than our best fit value of 7\\% and our upper limit of 15\\% (3$\\sigma$). In addition, their comparison of the lens-based mass model and the observed luminosity distribution indicates that mass follows light in the inner regions, and any halo contribution is relatively constant. Our analysis focuses on determining a simple yet comprehensive description of the galaxy over a broader range of scales. Although the results for the dark matter halo can be improved by more extended and precise kinematic data, we achieved our aim of determining the bulge and disc mass-to-light ratios with good precision, without the need to use stellar population synthesis models. Future studies with the addition of stellar population synthesis models will be very valuable.  The good agreement between our results and those of \\citet{vandenven08} -- based on largely independent datasets (except for the strong lensing) and methodologies -- is encouraging: the combination of lensing and dynamics is a powerful tool to study spiral galaxies, not only elliptical galaxies. As larger samples of spiral lens galaxies are being discovered by the SLACS \\citep{bolton08a} and other surveys -- including many where the Einstein radius will be equal or larger than the scale length of the disc -- the combination of lensing, kinematics, and stellar population synthesis models should enable substantial progress in our understanding of the structure of disc galaxies, and therefore their formation and evolution."
    },
    "0812/0812.2267_arXiv.txt": {
        "abstract": "We used near-infrared 2MASS data to construct visual extinction maps of a sample of Southern Bok globules utilizing the NICE method. We derived radial extinction profiles of dense cores identified in the globules and analyzed their stability against gravitational collapse with isothermal Bonnor-Ebert spheres. The frequency distribution of the stability parameter ($\\xi_{max}$) of these cores shows that a large number of them are located in stable states, followed by an abrupt decrease of cores in unstable states. This decrease is steeper for globules with associated IRAS point sources than for starless globules. Moreover, globules in stable states have a Bonnor-Ebert temperature of $T = 15 \\pm 6$~K, while the group of critical plus unstable globules has a different temperature of $T = 10 \\pm 3$~K. Distances were estimated to all the globules studied in this work and the spectral class of the IRAS sources was calculated. No variations were found in the stability parameters of the cores and the spectral class of their associated IRAS sources. On the basis of \\tco\\ J = 1--0 molecular line observations, we identified and modeled a blue-assymetric line profile toward a globule of the sample, obtaining an upper limit infall speed of 0.25 km s$^{-1}$. ",
        "introduction": "\\label{sec:intro}  \\citet{bok47} were the first to call attention to small dark regions, dense and rounded, which called \\emph{globules}, and suggested the possibility that the star formation could occur within them. These objects were called later Bok globules and in general they are isolated, cold ($T \\sim$ 10~K), dense ($n \\sim 10^{4-5}$~cm$^{-3}$), and seen as potential sites for the formation of low-mass stars ($<$ 10~M$_{\\odot}$). Following \\citet{beichman86}, the globules that do not contain protostars or young stellar objects (YSOs) associated, are called starless cores. When they show some evidence of star formation, such as Class 0 or I protostars, are called protostellar-cores \\citep{andre93}.  As the majority of the stars form in molecular cloud complexes and in clusters embedded in giant molecular clouds and not in isolation, it is very important to know in detail the physical processes that give rise to the formation of a single star in an isolated environment. For this purpose, Bok globules constitute an ideal laboratory to study the isolated low-mass star formation \\citep[][hereafter BHR]{bourke95a}.  Visual extinction was used by \\citet{clemens88} to study 248 Bok globules in the Northern Hemisphere, and by BHR to study 169 globules in the Southern Hemisphere. Sub-millimeter and millimeter wavelength studies of starless cores \\citep{ward94}, indicate that the 1D profiles (mean of azimuthal distributions) are flat for radii $\\sim$ 3000-7000~AU, and steeper for larger radii. These profiles deviate from the distribution $n(r) \\propto r^{-2}$ of the singular isothermal sphere, proposed by \\citet{shu77} as the initial state of isolated dense cores before gravitational collapse. Similar profiles are obtained by mapping the dust distribution in a dense core or Bok globule. These configurations are consistent with ``Bonnor-Ebert spheres'' \\citep{ebert55,bonnor56}, which are non-singular solutions to the equations of hydrostatic equilibrium confined by an external pressure. In the 1D isothermal approximation, there is a family of Bonnor-Ebert solutions parametrized by the central density $n_{c}$, where the density profile is characterized by two regimes: a slow decrease in the density for small radii, and a faster decrease, of power-law type ($ \\propto r^{-2}$), for larger radii.  Detailed modelling of profiles in the infrared \\citep{kandori05,teixeira05} and continuum emission \\citep[e.g.,][]{ward99} confirm that Bonnor-Ebert spheres are a good approach to the internal structure of starless cores. However, it is not currently possibly to distinguish between a static Bonnor-Ebert sphere and a collapsing one for profiles based in dust absorption or emission, since the structure of the profile, for a given $n_c$, does not change significantly until the last stages of collapse \\citep{myers05}.  Using near-infrared observations, \\citet{alves01} showed that a Bonnor-Ebert sphere with stability parameter $\\xi_{max}$ = 6.9 adjust perfectly the observed column density profile of the starless globule Barnard 68. \\citet{harvey01} modeled the profile of the protostellar globule B335 with $\\xi_{max}$ = 12.5, and \\citet{harvey03} estimated $\\xi_{max}$ = 25 for the starless core L694-2. The Coalsack globule 2 was modeled by \\citet{racca02} that derived $\\xi_{max}$ = 7.3, while \\citet{lada04} obtained $\\xi_{max}$ = 5.8 using a more sensitive telescope. In recent works, \\citet{huard06} found $\\xi_{max}$ = 35.8 for the protostellar core L1014, and \\citet{kainu07} obtained $\\xi_{max}$ = 23 for the protostellar globule DCld303.8-14.2 and $\\xi_{max} \\gtrsim$ 8 for the starless globule Thumbprint Nebula.  The largest sample of Bok globules studied in the near-infrared was done by \\citet{kandori05} where 14 globules were selected. They found that more than half of the starless globules (7 out of 11) are located close to the critical state, with $\\xi _{max}$ = 6.5 $\\pm$ 2. This led \\citet{kandori05} to suggest that a Bonnor-Ebert sphere in the critical state characterizes the structure of typical starless cores. The remaining starless globules and those that present evidence of star formation show clearly unstable states, with $\\xi _{max}>$ 10.  In this work, we present a study of 21 Southern Bok globules, where 11 are starless and 10 have associated IRAS\\footnote{Infrared Astronomical Satellite \\citep{neugebauer84}} point sources. In \\S \\ref{sec:observ} we present the source list and the observations. Using 2MASS data, we constructed visual extinction maps in \\S \\ref{sec:extmaps}, which enabled us to detect dense cores embedded in the globules. In \\S \\ref{sec:distance} we estimated new distances to the globules, and by means of the Bonnor-Ebert model, in \\S \\ref{sec:structure} we studied their internal structure and stability against gravitational collapse. In \\S \\ref{sec:starfrom} we analized the IRAS point sources associated to the globules and in \\S \\ref{sec:collapsing} we found evidence of infall motions in a globule using millimetric observations of the \\tco\\ and \\ceio\\ molecules. Our main conclusions are summarized in \\S \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In this work, we studied a sample of 21 isolated Southern Bok globules selected among 169 globules catalogued by BHR. In this sample, 11 are starless and 10 have associated IRAS point sources. Employing the NICE method of \\citet{lada94}, we used the 2MASS PSC to construct visual extinction maps for the Bok globules with 20$''$ resolution. This enabled us to detect dense cores and substructures embedded in the globules, which were modelled with Bonnor-Ebert spheres to derive their physical parameters and investigate their stability. Two methods were used to determine the spectral class of the IRAS sources associated to the globules. In addition, we searched for spectral lines observations data obtained toward these globules to explore whether they have or not evidence of collapse. We found observations of \\tco\\ and \\ceio\\ J = 1--0 lines toward two globules of the sample showing evidences of collapse in one of them. Distances were estimated to all these globules from plots of color excess versus distance. The main results of this paper are summarized as follows:  \\begin{enumerate}  \\item Dense cores were identified embedded in the globules, with mean volumetric densities greater than 10$^{4}$~cm$^{-3}$ and masses between $\\sim $ 1-4~M$_{\\odot }$.  \\item We inferred distances to several globules different from that ones obtained by BHR in spite of using the same technique. Due to the uncertain nature of this method, in order to improve the distance determination, we used smaller areas and a catalog containing information for approximately 10 times more stars than the catalog utilized by BHR.  \\item We found that there is a large number of cores in the stable regime, followed by an abrupt decrease of cores in the unstable regime, but this decrease is steeper for IRAS globules than for starless. The globules that are in the stable state have a temperature of $T = 15 \\pm 6$ K, while the globules that are in the critical and unstable states have a remarkably different temperature of $T = 10 \\pm 3$ K.  \\item We found that, according to the bolometric temperature, 6 IRAS sources are classified as Class 0, 4 as Class I, and 3 as Class II. The spectral index also produced similar results. No correlation was identified between the Bonnor-Ebert stability parameter of the condensations and the spectral class of the associated IRAS sources.  \\item By means of the \\tco\\ J = 1--0 molecular line, we identified a blue-assymetric line profile toward BHR 138. Using the \\emph{two-layer} collapse model, it was possible to estimate an infall speed of 0.25 km s$^{-1}$. However, taking into account possible rotational effects, this value could be as small as $\\sim$ 0.1 km s$^{-1}$.  \\end{enumerate}"
    },
    "0812/0812.3327_arXiv.txt": {
        "abstract": "Using new and archival radio data, we have measured the proper motion of the black hole X-ray binary V404 Cyg to be $9.2\\pm0.3$\\,mas\\,yr$^{-1}$.  Combined with the systemic radial velocity from the literature, we derive the full three-dimensional heliocentric space velocity of the system, which we use to calculate a peculiar velocity in the range 47--102\\,km\\,s$^{-1}$, with a best fitting value of 64\\,km\\,s$^{-1}$.  We consider possible explanations for the observed peculiar velocity, and find that the black hole cannot have formed via direct collapse.  A natal supernova is required, in which either significant mass ($\\sim 11M_{\\odot}$) was lost, giving rise to a symmetric Blaauw kick of up to $\\sim65$\\,km\\,s$^{-1}$, or, more probably, asymmetries in the supernova led to an additional kick out of the orbital plane of the binary system.  In the case of a purely symmetric kick, the black hole must have been formed with a mass $\\sim 9M_{\\odot}$, since when it has accreted 0.5--1.5\\,$M_{\\odot}$ from its companion. ",
        "introduction": "\\label{sec:intro}  The proper motions of X-ray binary systems can be used to derive important information on the birthplaces and formation mechanisms of their compact objects.  Since typical proper motions are of order a few milliarcseconds per year, high-resolution observations and long time baselines are required to measure the transverse motions of such systems across the sky.  Proper motions have only been measured for a handful of X-ray binary systems to date \\citep{Mir01,Mir02,Rib02,Mir03,Mir03b,Dha06,Dha07}, and only one X-ray binary, Sco X-1, has a measured proper motion, parallactic distance, and radial velocity \\citep{Bra99,Cow75}.  With the position, proper motion, radial velocity, and distance to the system, the full three-dimensional space velocity of the system can be derived. Along with system parameters such as the component masses, orbital period, donor temperature and luminosity, these parameters may be used to reconstruct the full evolutionary history of the binary system back to the time of compact object formation, as done for the systems GRO J\\,1655-40 \\citep{Wil05} and XTE J\\,1118+480 \\citep{Fra08}.  By studying the distribution of black hole X-ray binary velocities with compact object masses, we can derive constraints on theoretical models of black hole formation \\citep[e.g.][]{Fry01}.  The two most common theoretical scenarios for creating a black hole involve either a massive star collapsing directly into a black hole with very little or no mass ejection, or delayed formation in a supernova, as fallback onto the neutron star of material ejected during the explosion creates the black hole.  In the latter case, constraints on the magnitude and symmetry of any natal kick can also be derived.  When the primary star reaches the end of its life and explodes as a supernova, the centre of mass of the ejected material continues to move with the velocity the progenitor had immediately before the explosion.  The centre of mass of the binary system then recoils in the opposite direction.  Such a kick \\citep{Bla61} is thus constrained to lie in the orbital plane. In addition, in the presence of asymmetries in the supernova explosion, a further, asymmetric, kick, which need not be in the orbital plane, may be imparted to the binary \\citep[see, e.g.,][for more detailed overviews]{Bra95a,Por98,Lai01}.  While pulsar proper motions \\citep{Lyn94} and the high eccentricities of Be/X-ray binaries provide strong evidence for asymmetric kicks in neutron star formation \\citep{van97,van00}, less attention has been paid to black hole binary systems.  In two of the black hole X-ray binaries with some of the best observational constraints on the system parameters, XTE J\\,1118+480 and GRO J1655-40, there is evidence for an asymmetric supernova kick \\citep{Gua05,Wil05}.  There is also some evidence suggesting an asymmetric kick in GRS\\,1915+105 \\citep{Dha07}, although in that case the uncertainty in the distance to the system precludes the derivation of strong constraints on the kick velocity.  However, the black hole in Cygnus X-1 is inferred to have formed via direct collapse \\citep{Mir03}.  More black hole sources need to be studied to measure black hole kicks and investigate consequent formation mechanisms in order to constrain theoretical models of black hole formation.  \\subsection{V404 Cyg} V404 Cyg is a dynamically-confirmed black hole X-ray binary system, with a mass function of $6.08\\pm0.06M_{\\odot}$ \\citep{Cas94}.  The system comprises a black hole accretor of mass $12^{+3}_{-2}M_{\\odot}$ \\citep{Sha94} in a 6.5-d orbit with a $0.7^{+0.3}_{-0.2}M_{\\odot}$ K0 IV stripped giant star \\citep{Cas93,Kin93}.  The orbit is highly circular, with an eccentricity of $e<3\\times10^{-4}$, in agreement with the fact that tidal forces just before the onset of mass transfer should have recircularised the orbit following the supernova.  Its low radial velocity \\citep[$-0.4\\pm2.2$\\,km\\,s$^{-1}$;][]{Cas94} has been taken as evidence for a small natal kick \\citep{Bra95,Nel99}.  In this paper we present measurements of the proper motion of V404 Cyg, and go on to derive its full three-dimensional space velocity, infer its Galactocentric orbit, and discuss the implications for the formation of the black hole in the light of the inferred natal kick from the supernova. ",
        "conclusions": "We have measured the proper motion of V404 Cyg using 20 years' worth of VLA and VLBI radio observations.  Together with the radial velocity and constraints on the distance of the system, this translates to a peculiar motion of $64.1\\pm3.7^{+37.8}_{-16.6}$\\,km\\,s$^{-1}$.  Given the measured proper motion, the black hole cannot have been formed via direct collapse.  A supernova is required to achieve the observed peculiar velocity, with either a large amount of mass ($\\sim11M_{\\odot}$) being lost in the explosion, or, more probably, the system being subject to an asymmetric kick.  In the case of a pure Blaauw kick, $\\sim 1M_{\\odot}$ must have been transferred from the donor to the black hole since the onset of mass transfer, implying an initial black hole mass of $\\sim 9M_{\\odot}$ with a donor mass of $\\sim 2M_{\\odot}$ prior to the onset of mass transfer."
    },
    "0812/0812.1608_arXiv.txt": {
        "abstract": "{This paper reports on the detection of two new multiple planet systems around solar-like stars HD\\,47186 and HD\\,181433. The first system includes a hot Neptune of 22.78 M$_{\\oplus}$ at 4.08-days period and a Saturn of 0.35 M$_{\\rm JUP}$ at 3.7-years period. The second system includes a Super-Earth of 7.5 M$_{\\oplus}$ at 9.4-days period, a 0.64 M$_{\\rm JUP}$ at 2.6-years period as well as a third companion of 0.54 M$_{\\rm JUP}$ with a period of about 6 years. These detections increase to 20 the number of close-in low-mass exoplanets (below 0.1 M$_{\\rm JUP}$) and strengthen the fact that 80\\% of these planets are in a multiple planetary systems.} ",
        "introduction": "The HARPS spectrograph based on the 3.6-m ESO telescope at La Silla Observatory is now in operation since 2003. The HARPS consortium started 5 years ago an ambitious and comprehensive Guaranteed Time Observations (GTO) program of high-precision systematic search for exoplanet in the Southern sky (Mayor et al. \\cite{mayor03}, \\cite{mayor08}). Relevant efforts were made for the search for very low-mass planets. Our consortium dedicates ~50\\% of the GTO on HARPS since 2003 to monitoring about 200 nearby non-active stars. Thank to radial velocity accuracy better \\ than 1 {\\ms} (Pepe et al. \\cite{pepe04}, Lovis et al. \\cite{lovis06b}) HARPS succeeds in the discovery of several Neptune-mass and Super-Earth exoplanets around solar-type stars (Santos et al. \\cite{santos04}, Udry et al. \\cite{udry06}, Lovis et al. \\cite{lovis06}, Melo et al. \\cite{melo08}, Mayor et al. \\cite{mayor08}) and M dwarfs (Bonfils et al \\cite{bonfils05}, Udry et al. \\cite{udry07}, Bonfils et al. \\cite{bonfils07}, Forveille et al. \\cite{forveille08}). Such efficiency mainly comes from the following factors : 1) a careful and frequent monitoring of instrumental performances, 2) a continuous improvement of the data reduction software (e.g. Lovis \\& Pepe \\cite{lovis07}), 3) a dedicated and careful observing strategy to deal with stellar seismic noise (e.g. Bouchy et al. \\cite{bouchy05}), 4) a long-duration monitoring of non-active stars, 5) an accumulation of measurements in order to identify multi-planetary systems. Our large program is now leading to an increasing list of close-in low-mass exoplanets which will improve our knowledge of the planet distribution in the mass-period diagram and will allow comparisons with theoretical predictions. Furthermore this increasing number of close-in low-mass planets stimulates dedicated photometric follow-up in order to detect transiting Neptunes like GJ436b (Gillon et al. \\cite{gillon07}). Indeed we expect statistically that 5-10\\% of this close-in low-mass exoplanets offer the appropriate configuration to transit their parent stars. In that case a direct measurement of the planetary radius as well as the exact mass will be provided. In this paper, we present the discovery of two new multiple planet systems including one hot Neptune and one Super-Earth orbiting the stars HD\\,47186 and HD\\,181433 respectively. ",
        "conclusions": "Figure~\\ref{ma_diag} represents the $\\sim$300 known exoplanets\\footnote{from The Extrasolar Planets Encyclopedia (http://exoplanet.eu) June 2008.} in the mass-separation diagram. The triangles refer to exoplanets found by radial velocities. The dark triangles refer to transiting exoplanets. The circles refer to exoplanets found by microlensing. The bold triangles correspond to HARPS discovered exoplanets including the 5 ones described in this paper and the 3 ones orbiting HD\\,40307 (Mayor et al. 2008). HARPS is presently filling the area of close-in low-mass exoplanets with minimum mass below 0.1 M$_{\\rm JUP}$. One expects that few \\% of these exoplanets offer the appropriate configuration to transit their parent stars. In that case a direct measurement of the planetary radius as well as the exact mass will be done providing crutial informations and constraints about their composition. This was the case of the up-to-day unique transiting Neptune GJ436b (Gillon et al. \\cite{gillon07}) which was first discovered by radial velocity (Butler et al. \\cite{butler04}). Such close-in low-mass planets are expected to have radius of few R$_\\oplus$ which means transit depth in the sub-millimagnitude regime. Although impossible to be detected in ground-based photometry surveys, such small transiting exoplanets may be detected in specific high-precision ground-based photometric follow-up, especially in case of small stellar radius (K and M dwarves), or using spaced-based facilities. One advantage comes from the fact that these planets orbit around nearby bright stars which makes possible deep and further characterization.  \\begin{figure} \\centering \\includegraphics[width=8.5cm]{bouchy_f4.pdf} \\caption{Mass-separation diagram of the 300 known exoplanets. The triangles refer to exoplanets found by radial velocities. The dark triangles refer to transiting exoplanets. The circles refer to exoplanets found by microlensing. The bold triangles correspond to HARPS discovered exoplanets. Lines of radial-velocity semi-amplitude of 1 and 3 {\\ms} are added assuming a 1 solar-mass star.} \\label{ma_diag} \\end{figure}   The two new planetary systems described here strengthen the fact low-mass exoplanets are found in multiple planetary systems. Indeed, from the 20 known exoplanets with masses lower than 0.1 M$_{\\rm JUP}$, 16 are in a multiple planetary systems, hence 80\\%. This fraction should be compared to the 23\\% of the $\\sim$300 known exoplanets which are in a multiple planetary systems. If we restrict the analysis to planetary period shorter than 50 days, which correspond to the detectability period cutoff of low mass explanets, the conclusion remains the same, 19\\% of exoplanets and 72\\% of low mass explanets are in a multiple systems. We note that from the 4 low-mass planets which are not up-to-day identified as part of a multiple planetary systems, 3 are orbiting a M-dwarf (GJ674b, GJ436b and HD\\,285968b).  It is interesting to discuss the discoveries presented in this paper from a theoretical point of view. In the planet population synthesis calculations of Mordasini et al. (\\cite{mordasini08a}) based on the core accretion scenario, a large number of hot-Neptunian and Super-Earth planets as the ones presented here are predicted to exist at small distances from the host star. These low mass planets form a distinct sub-population of planets which did not undergo gas runaway accretion, in contrast to the sub-population of Hot Jupiters. In the predicted IMF of the close-in planets, a local minimum is found between the two groups, occurring  for solar type stars at a mass of about 30 M$_{\\oplus}$. Indeed, Fig.~\\ref{ma_diag} indicates that the high precision program of HARPS has started to explore this new sub-population of close-in low mass planets. One might even tentatively see in Fig.~\\ref{ma_diag} a bimodal mass distribution of the ``Hot'' planets. This emerging bimodal shape of the mass distribution from gaseous giant planets to the super-Earth regime was already pointed out by Mayor \\& Udry (\\cite{mayor08b}) and is also visible on the mass distribution of detected planet with period shorter than 50 days (see Fig.~\\ref{bimodal}). Planetary formation simulations also based on the core accretion scenario and done by Ida \\& Lin (\\cite{ida04a}, \\cite{ida08}) also predict a bimodal distribution from gaseous giants to Super-Earth and furthermore a paucity of extrasolar planets with mass in the range 10-100  M$_{\\oplus}$.   \\begin{figure} \\centering \\includegraphics[width=8cm]{bouchy_f5.pdf} \\caption{Distribution of planetary masses from giant gaseous planets to the super-Earth regime for close-in planets (P$\\le$50 days). The hatched histogram represents HARPS detections.} \\label{bimodal} \\end{figure}   Discoveries of Hot Neptunes and close-in Super-Earth planets raise questions about how they were exactly formed. In-situ formation seems very unlikely for HD\\,47186b. This is due to the fact that in-situ formation allows the assembly of planets of only a few times the Earth mass even at the supersolar metallicity of about 0.2-0.3 dex of the stars considered here. HD\\,47186b has in contrast a clearly larger, Neptunian like mass and is located at a very small semi-major axis, where only tiny amounts of planetary building blocks are available, if any (due to evaporation of solids very close to the star). For HD\\,181433b, the planetary mass is only about a third, while the semi-major axis is larger, so that its mass is no more so clearly out of reach of what might can form in-situ. Still, for reasonable disk masses and profiles, it is unlikely that a such a massive Super-Earth planet can form at its current location. Therefore, it seem necessary that some migration process was at work. Various processes can bring planets close to the star: Planet-disc interaction in the form of type I and II migration (e.g. Terquem \\& Papaloizou \\cite{terquem07}), planet-planet interaction in the form of shepherding (e.g. Mandell et al. \\cite{mandell07}), scattering with subsequent circularization or the Kozai mechanism, among some others (e.g. Raymond et al. \\cite{raymond08}; Zhou \\& Lin \\cite{zhou08}).  It is interesting to discuss the discoveries in the context of the two competing giant planet formation mechanisms. In the core accretion mechanism, the formation of a system with both giant and low mass (Neptunian or Super-Earth) planets can be regarded as a natural outcome, which is not necessary the case in the gravitational instability model. Especially a planetary system architecture with low mass planets at smaller semi-major axes and one or several giant planets at larger distances is expected in the baseline core accretion model without long distance migration, as the increase of the available planetary building blocks with distance facilitates giant planet growth at larger distances, while the smaller amounts available at shorter distances should lead to the formation of low mass planets. This simple picture is reflected in both HD\\,47186 and HD\\,181433, even if the masses of the inner planets are large compared to those of the Solar System. We therefore may regard these systems as cases where some, but not extreme migration of both the low mass planets and the giants occurred. This would indicate an initial disk mass larger than that of the Solar System, but less massive than one leading to the formation of Hot Jupiters.  Population synthesis calculations based on the core accretion paradigm reproduce the ''metallicity effect'' (Ida \\& Lin \\cite{ida04b}) i.e. the strong positive correlation between the stellar metallicity and the planetary detection probability. Recent discoveries mainly by HARPS have shown that this ``metallicity effect'' might not exists for close-in low mass planets, even though the data set is currently too small for definitive conclusions (Udry \\&  Santos \\cite{udry07b}). However the multiple planetary systems like HD\\,47186 and HD\\,181433 with both a giant and a low mass planet (55 Cnc, Gl876, HD\\,160691, HD\\,190360, HD\\,219828) have all supersolar metallicities. In the core accretion scenario this is understood in the following way: High [Fe/H] stars (disks) are able to produce both giant and low mass planets, while low [Fe/H] systems are able to produce low mass planets only.  More discoveries of low mass planets will help clarify such correlations."
    },
    "0812/0812.4677_arXiv.txt": {
        "abstract": "Most gamma-ray bursts (GRBs) observed by the {\\it Swift} satellite show an early steep decay phase (SDP) in their X-ray lightcurve, which is usually a smooth continuation of the prompt gamma-ray emission, strongly suggesting that it is its tail. However, the mechanism behind it is still not clear. The most popular model for this SDP is High Latitude Emission (HLE), in which after the prompt emission from a (quasi-) spherical shell stops photons from increasingly large angles relative to the line of sight still reach the observer, with a smaller Doppler factor. This results in a simple relation between the temporal and spectral indexes, $\\alpha = 2 + \\beta$ where $F_\\nu \\propto t^{-\\alpha}\\nu^{-\\beta}$. While HLE is expected in many models for the prompt GRB emission, such as the popular internal shocks model, there are models in which it is not expected, such as sporadic magnetic reconnection events. Therefore, testing whether the SDP is consistent with HLE can help distinguish between different prompt emission models. In order to adequately address this question in a careful quantitative manner we develop a realistic self-consistent model for the prompt emission and its HLE tail, which can be used for combined temporal and spectral fits to GRB data that would provide strict tests for the HLE model. We model the prompt emission as the sum of its individual pulses with their HLE tails, where each pulse arises from an ultra-relativistic uniform thin spherical shell that emits isotropically in its own rest frame over a finite range of radii. Analytic expressions for the observed flux density are obtained for the internal shock case with a Band function emission spectrum. We find that the observed instantaneous spectrum is also a Band function. Our model naturally produces, at least qualitatively, the observed spectral softening and steepening of the flux decay as the peak photon energy sweeps across the observed energy range. The observed flux during the SDP is initially dominated by the tail of the last pulse, but the tails of one or more earlier pulses can become dominant later on. A simple criterion is given for the dominant pulse at late times. The relation $\\alpha = 2+\\beta$ holds also as $\\beta$ and $\\alpha$ change in time. Modeling several overlapping pulses as a single wider pulse would over-predict the emission tail. ",
        "introduction": "Before the launch of the {\\it Swift} satellite (Gehrels et al. 2004), Gamma Ray burst (GRB) X-ray afterglows were detected at least several hours after the burst (Soffitta et al. 2004 and references therein). They typically displayed a power law decay $\\sim t^{-1} - t^{-1.5}$ around their detection time (De Pasquale et al. 2006). {\\it Swift}'s ability to rapidly and autonomously slew when the Burst Alert Telescope (BAT, observing in the energy range $15-350\\;$keV; Barthelmy et al. 2005) instrument detects a GRB enables it to point its other instruments - the X-Ray Telescope (XRT, observing in the energy range $0.2-10\\;$keV; Burrows et al. 2005a) and UV/Optical Telescope (UVOT, observing at wavelengths $170-650\\;$nm, i.e. from the optical to the near UV; Roming et al. 2005) - toward the GRB within tens of seconds from the GRB trigger time.The XRT thus filled the observational gap between the end of the prompt emission and the beginning of the pre-{\\it Swift} afterglow observations several hours later. It revealed a complex behaviour usually consisting of three phases, followed by most GRBs, and referred to as a canonical light curve (Nousek et al. 2006), consisting of three distinct power-law segments where $F_\\nu \\propto t^{-\\alpha}$: an initial (at $t < t_{\\rm break,1}$, with $300\\;{\\rm s} \\lesssim t_{\\rm break,1} \\lesssim 500\\;{\\rm s}$) very steep decay with time $t$ (with a power-law index $3 \\lesssim \\alpha_1 \\lesssim 5$; see also Bartherlmy et al. 2005; Tagliaferri et al. 2005); a subsequent (at $t_{\\rm break,1} < t < t_{\\rm break,2}$, with $10^3\\;{\\rm s} \\lesssim t_{\\rm break,2} \\lesssim 10^4\\;{\\rm s}$) very shallow decay ($0.2 \\lesssim \\alpha_2 \\lesssim 0.8$); and a final steepening (at $t > t_{\\rm break,2}$) to the familiar pre-{\\it Swift} power-law behaviour ($1 \\lesssim \\alpha_3 \\lesssim 1.5$). In many cases there are X-ray flares superimposed on this underlying smooth component (typically during the first two phases, at $t < t_{\\rm break,2}$; Burrows et al. 2005b; Falcone et al. 2006; Krimm et al. 2007), and in some cases there is a later (at $t_j > t_{\\rm break,2}$) further steepening due to a jet.  The third decay phase ($F_\\nu \\propto t^{-\\alpha_3}$) is the afterglow emission that was observed before {\\it Swift} and is well explained by synchrotron radiation from the forward shock that is driven into the external medium as the GRB ejecta are decelerated, where the energy in the afterglow shock is constant in time (no significant energy gains or losses). The plateau (or shallow decay) phase can be explained either by pre-{\\it Swift} models or by later models that have been developed especially for this purpose (Nousek et al. 2006; Panaitescu et al. 2006; Granot 2007).  It could be due to energy injection, either by a tail of decreasing Lorentz factors at the end of the ejection phase (Rees \\& M\\'eszaros 1998; Sari \\& M\\'eszaros 2000; Ramirez-Ruiz, Merloni \\& Rees 2001; Granot \\& Kumar 2006) or by a relativistic wind produced by a long lasting source activity (Rees \\& M\\'eszaros 2000; McFadyen et al. 2001; Lee \\& Ramirez-Ruiz 2002; Dai 2004; Ramirez-Ruiz 2004), by an increasing efficiency of X-ray afterglow emission due to time dependence of the shock microphysics parameters (Granot, K\\\"onigl \\& Piran 2006), by a viewing angle slightly outside the region of prominent afterglow emission (Eichler \\& Granot 2006), by a contribution from the reverse-shock (Genet, Daigne \\& Mochkovitch 2007) or by a two component jet model (Peng, K\\\"onigl \\& Granot 2005; Granot, K\\\"onigl \\& Piran 2006).  The steep decay phase is observed in most bursts, and is in the great majority of cases a smooth continuation of the prompt emission, both temporally and spectrally (O'Brien et al. 2006). This strongly suggests that it is the tail of the prompt emission. Several explanations for this phase have been suggested in the context of previously existing models (Tagliaferri et al. 2005; Nousek et al. 2006), such as emission from the hot cocoon in the collapsar model (M\\'eszaros \\& Rees 2001, Ramirez-Ruiz et al. 2002). The most popular model, by far, is High Latitude Emission (HLE) originally referred to as emission from a ``naked'' GRB (Kumar and Panaitescu 2000a). In this model the prompt GRB emission is from a (quasi-) spherical shell, and after it turns off at some radius then photons keep reaching the observer from increasingly larger angles relative to the line of sight, due to the the added path length caused by the curvature of the emitting region. Such late arriving photons experience a smaller Doppler factor. This leads to a simple relation between the temporal and spectral indexes, $\\alpha = 2+\\beta$ where $F_{\\nu}(t) \\propto t^{-\\alpha} \\nu^{-\\beta}$, that holds at late times when $t-t_0 \\gg \\Delta t$, where $t_0$ and $\\Delta t$ are the start time and width of the pulse, respectively. The steep decay phase also shows a softening of the spectrum with time (see Zhang et al. 2007 and references therein).  The consistency of the steep decay phase with HLE has been studied by several authors (Nousek et al. 2006; Liang et al. 2006; Butler \\& Kocevski 2007; Zhang et al. 2007; Qin 2008). However, some simplifying assumptions were usually made, which may affect the comparison between this model and the observations. One such assumption is the choice of the reference time $t_0$ for the steep decay, especially when the prompt emission consists of several pulses. Liang et al. (2006) find that when assuming the HLE relation $\\alpha = 2+\\beta$ and fitting for $t_0$ its derived value is consistent with the onset of the last pulse of the prompt emission (or of the individual spike or flare whose tail is being fit). Zhang et al. (2007) find that the HLE cannot explain the steep decays accompanied by a spectral softening, but can explain the cases with no observed spectral evolution. Barniol Duran and Kumar (2008) find that only $20\\%$ of their sample is consistent with HLE. Butler \\& Kocevski (2007) find that for a (physically motivated) time independent soft X-ray absorption (fixed N$_{\\rm H}$) the spectrum during the steep decay phase, is much better fit by an intrinsic Band function spectrum (Band et al. 1993) rather than by a power-law, and that the peak photon energy shifts to lower energies with time. Qin (2008) finds that such a behavior can, at least qualitatively, be produced for a delta function emission in radius with a Band function spectrum. It is therefore still a largely open question whether the temporal and spectral properties of the steep decay are consistent with HLE.  Moreover, it appear that a physically motivated model for the prompt emission with realistic assumptions about the emission (e.g. over a finite range of radii with a Band function emissivity) is needed in order to address this question in a more quantitative and fully self consistent manner.  The nature of the prompt GRB emission is what ultimately determines the properties of its tail. HLE is expected only in models where the prompt emission is from a quasi-spherical shell and turns off rather abruptly at some finite radius (or lab frame time). The best example for this type of model is internal shocks (Rees \\& M\\'eszaros 1994; Sari \\& Piran 1997) where variability in the Lorentz factor of the relativistic GRB outflow causes faster shells of ejecta to collide with slower sells resulting in shocks going into the shell over a finite range of radii (typically $\\Delta R \\sim R$).  On the other hands, there are models in which HLE is not expected, such as in the case of isolated sporadic magnetic reconnection events within a Poynting flux dominated outflow (e.g. Lyutikov \\& Blandford 2003) in which each spike in the GRB light curve is from a distinct small and well localized region. Therefore, testing whether the steep decay phase is consistent with HLE would help distinguish between these two types (or classes) of prompt GRB models. This can be an important step toward identifying the basic underlying mechanism for the prompt emission, which is still one the the most striking open questions in GRB research more than four decades after the discovery of GRBs.  In order to address this question, we develop a model for the prompt and its HLE tail that is physically motivated, realistic, and easy to use (fully analytic in its simplest version) in global joint fits (to all of the available data at all times and photon energies) of the prompt GRB and its SDP tail. Such global fits can provide a stringent and fully self-consistent test of HLE model for the SDP in GRBs.  The prompt emission is modeled as the sum of a finite number of pulses. Each pulse corresponds to a spike in the prompt GRB light curve and has its own HLE tail. An individual spike is modeled as arising from a thin uniform spherical relativistic shell that emits isotropically in its own rest frame over a finite range of radii, while the observed flux is calculated by integrating over the equal arrival time surface (Granot, Piran \\& Sari 1999; Granot 2005; Granot, Cohen-Tanugi \\& do Couto e Silva 2008) of photons to a distant observer. Our model is particularly suitable for internal shocks, which we focus on in this paper. For the emitted spectrum we consider the phenomenological Band function, which provides a good fit the the prompt emission spectrum of the vast majority of GRBs. We point out that our model can also be used for X-ray flares, which appear to have temporal and spectral properties similar to the spikes of the prompt GRB emission. The main text provides the most useful results in an easy to use form, while the full derivations of these results are provided in appendixes in order to help understand their origin and make it easier to extend or generalize our model. We stress here that our main aim is not necessarily to uniquely determine all of the model parameters, which may be subject to various degeneracies and may prove hard when fitting to real data, but instead to test whether our model can provide a good fit to the data for any set of physical parameters. While such a good fit would still not prove that the HLE must be at work, it would definitely support HLE as a viable and arguably most plausible model.  Our model for an individual pulse is described in \\S~\\ref{sec_dyn}, and results for the flux in the case for internal shocks with a Band function spectrum are given in \\S~\\ref{sec_res_flux}. The dependence of a single pulse on the model parameters is then investigated in \\S~\\ref{sec_emission_single}, while \\S~\\ref{sec_prompt} discusses how to combine several pulses in order obtain to the total prompt emission and its tail. Both are intended to help the reader when using our model to fit data, which is one of the main aims of our paper. Our conclusions are discussed in \\S~\\ref{sec_concl}. This paper describes in detail our theoretical model and its main properties, and stresses some important caveats that one should keep in mind when using it to fit data in order to test the HLE model. In subsequent work we intend to confront it with {\\it Swift} BAT+XRT data. ",
        "conclusions": "\\label{sec_concl}   We have presented and explored a model for the prompt GRB emission and its high latitude emission (HLE) tail. This model is physically motivated and realistic: it consists of a finite number of emission episodes, each of which corresponds to a single spike in the prompt light curve, and is modeled by a relativistically expanding thin spherical uniform shell emitting isotropically in its own rest frame within a finite range of radii. Our model thus describes the prompt emission and the steep decay pahse as a whole from its very start to its late tail. Yet this model is easy to use (fully analytic in its simplest form described here), making it particularly suitable for detailed combined temporal and spectral global fits to the prompt GRB emission and the following steep decay phase (SDP). Such fits can provide a stricter test of the HLE model for the SDP compared to most previous models, since we use a single self-consistent model to fit both the prompt emission and the SDP, while most previous models fit only the SDP and are largely decoupled from the details of the prompt emission. Moreover, our model is also physically motivated, and more realistic than previous models. We have derived analytic expressions for the flux in the realistic case of a Band function spectrum (eqs.~[\\ref{eq_bandDR_finalcondensed}] and [\\ref{F_expanded_latetime}]), which consists of two power laws that smoothly join at some typical photon energy.  The temporal evolution of the instantaneous values of the temporal ($\\alpha$) and spectral ($\\beta$) indexes for a single emission episode was studied, corresponding to a single observed pulse in the light curve. The definition of $\\alpha$ is not unique as it depends on the choice of reference time. We explored two options for the reference time, either the ejection time ($\\alpha_{ej}$) or the onset time of the spike ($\\alpha_{on}$), and found that for the former the HLE relation ($\\alpha_{ej} = 2 + \\beta$) is satisfied from immediately after the peak of the spike ($\\bar{T} > \\bar{T}_f$), while for the former it is only approached at late times ($\\alpha_{on} \\approx 2 + \\beta$ at for $\\bar{T} > \\bar{T}_f$ and $\\bar{T} \\gg 1$).  We have intentionally chosen a simple model to describe the pulses, in order to reduce the number of free parameters. For a single emission episode (or pulse), in the most generic case there are ten free parameters: the power $m = -2d\\log\\Gamma/d\\log R$, $d = d\\log \\nu'_p/d\\log R$, $a = d\\log L'_{\\nu'_p}/d\\log R$, the normalization factor $F_0$ (or $L_0$), three additional parameters for the Band function (the two spectral slopes, $b_1$ and $b_2$, as well as the peak energy at the onset of the pulse $E_0$), the two timescales $T_0$ and $T_f$, and the ejection time $T_{\\rm ej}$. We have the general constraint $\\Delta R > 0$, which implies $\\tilde{T}_f= 1 + \\Delta R/R_0 > 1$. Focusing on the internal shocks model fixes some of these parameters: as the outflow is typically in the coasting phase, $m=0$, while for synchrotron emission from fast cooling electrons $d = -1$ and $a = 1$.  Since we expect $\\Delta R/R_0 = \\bar{T}_f \\sim 1$ we can fix $\\bar{T}_f \\sim 1$ (although a wider range, such as $0.2 \\Delta R/R_0 \\lesssim 5$, may still be considered as plausible). Fixing $m$, $d$, $a$, and $T_f/T_0$ in this manner would leave only six free parameters. For a prompt emission with several pulses, one may be able in some cases to neglect the spectral evolution and use the same values of $b_1$, $b_2$, and $E_0$ for all the different pulses (or at least two of them, e.g. $b_1$ and $b_2$), which leads to a total number of free parameter of $3(N_{\\rm pulses}+1)$ (or $4N_{\\rm pulses}+2$ if $E_0$ cannot be fixed for all the pulses) for a burst with $N_{\\rm pulses}$ pulses.  The shape of the pulses in our model can vary considerably, from very spiky peaks to rounder ones, from a very sharp rise to shallower rise, and so on (see Figs.~\\ref{fig_LC_onepulse_selon_freq} -- \\ref{fig_var_a}). This can help reproduce some of the observed diversity in the shape of spike in the prompt light curve. This appears to be a promising feature of our model. However, we have an abrupt change in the temporal index at $\\bar{T_f}$, that usually corresponds to a sharp peak of the spike. This is caused by our model assumption that the emission abruptly shuts off at the outer emission radius $R_f$. Therefore, we also consider an alternative and more realistic assumption, which leads to a rounder peak for the spikes, where the emission more gradually turns off at $R > R_f$. This is done by introducing and exponential turn-off with radius of the comoving spectral luminosity, $L'_{\\nu'}(R)$, and is examined in Appendix \\ref{ap_expcutoff}. The more gradual the turn-off of the emission with radius the rounder the peak of the pulse in the light curve. This can help fit the observed variety of pulse shapes even better (at the cost of adding an additional free parameter).  In the particular case of synchrotron emission from internal shocks, we find that the observed spectrum has the same shape as the emitted one, which in our case is modeled as a Band function. The observed peak photon energy of the Band function decreases with time, $E_p(\\tilde{T}) = E_0/\\tilde{T}$, naturally leading to a softening of the spectrum with time, similar to what is observed by {\\it Swift}. Thus, our model can at least qualitatively reproduce the main temporal and spectral features observed by {\\it Swift}. The spectral index $\\beta$ evolves from its value below $E_p$ ($\\beta = -b_1$) to its value above $E_p$ ($\\beta = -b_2$), where the transition that corresponds to the passage of $E_p$ through the observed energy band occurs at earlier times for higher observed photon energies (or frequencies).  When modeling the prompt emission by combining several pulses, the SDP is initially dominated by the last pulse (just after its peak, if it is above the flux fro the other pulses), but can later be dominated by the tail of other pulses. The relative contribution of a pulse to the late time flux scales as $\\sim F_{peak}T_f^{2+\\beta}$, and therefore wider pulses (with a larger $T_f$), and to a lesser extent pulses with a larger peak flux ($F_{\\rm peak}$), tend to dominate the late time flux, deep into the SDP. Moreover, often the contribution to the total flux from the tails of several pulses can be comparable, so it cannot be adequately modeled using a single pulse model. Therefore, we caution here that modeling the steep decay phase using the HLE of a single pulse, $F_\\nu \\propto (T-T_{\\rm ref})^{-(2+\\beta)}$, may lead to wrong conclusions, and all the more so if the reference time $T_{\\rm ref}$ is arbitrarily set to the GRB trigger time. Even if $T_{\\rm ref}$ is set to the onset time of the last spike, this may still be a bad approximation in many cases since (i) we find that $\\alpha_{ej} = 2 + \\beta$ (with $T_{\\rm ref} = T_{\\rm ej}$) rather than $\\alpha_{on}$ (with $T_{\\rm ref} = T_{\\rm ej}+T_0$, corresponding to the onset of the spike) while $\\alpha_{on}$ approaches $\\alpha_{ej} = 2+\\beta$ only at late times well after the peak of the last pulse, and (ii) at such late times the flux often becomes dominated by the tails of earlier pulses.  Our model can produce different shapes for the tail of the prompt emission, from close to a power law (which can have a different temporal index than its asymptotic late time value) to a curved shape with decreasing temporal index $\\alpha$. This is qualitatively consistent with observations, where these type of behaviour are observed. We have demonstrated that just after the peak of the last pulse, the decay index of the prompt emission tail can reach very large values, far greater than the typical average value observed during the SDP by {\\it Swift}, of $3 \\lesssim \\alpha \\lesssim 5$ (Nousek et al. 2006). Larger values for the temporal index, however, are sometimes observed close to the end of the prompt emission (for example in GRB050422, GRB050803 or GRB050916; see figure 2 from O'Brien et al. 2006), in accord with our model.  Because of the large number of free parameters, the fitting of actual data should be handled with care, and there may be various degeneracies involved. The results of such fits to data should also be taken cautiously because of the difficulty in properly resolving distinct pulses in the prompt emission. For different reasons (such as noisy data, coarse time bins, pulse overlap, etc.), a group of distinct pulses may be fitted by a single broader pulse, resulting in an over-prediction of the flux during the SDP, as well different spectral and temporal evolution, which might lead to a misinterpretation of the SDP. Nevertheless, when handled with care, a fit of our model to a good combined data set of the prompt GRB emission and its SDP tail can serve as a powerful test of the HLE model for the SDP, and thus help distinguish between different models for the prompt GRB emission.\\\\   \\vspace{0.5cm} \\noindent J.~G. gratefully acknowledges a Royal Society Wolfson Research Merit Award.             \\newpage"
    },
    "0812/0812.1052_arXiv.txt": {
        "abstract": "Near-future cosmological observations targeted at investigations of dark energy pose stringent requirements on the accuracy of theoretical predictions for the nonlinear clustering of matter. Currently, $N$-body simulations comprise the only viable approach to this problem. In this paper we study various sources of computational error and methods to control them. By applying our methodology to a large suite of cosmological simulations we show that results for the (gravity-only) nonlinear matter power spectrum can be obtained at 1\\% accuracy out to $k\\sim 1\\,h\\,{\\rm Mpc}^{-1}$. The key components of these high accuracy simulations are: precise initial conditions, very large simulation volumes, sufficient mass resolution, and accurate time stepping. This paper is the first in a series of three; the final aim is a high-accuracy prediction scheme for the nonlinear matter power spectrum that improves current fitting formulae by an order of magnitude. ",
        "introduction": "The nature of the dark energy believed to be causing the current accelerated expansion of the Universe is one of the greatest puzzles in the physical sciences, with deep implications for our understanding of the Universe and fundamental physics. The twin aims of better characterizing and further understanding the nature of dark energy are widely recognized as key science goals for the next decade. Although dark energy remains very poorly understood, theory nevertheless plays an essential role in furthering this enterprise.  The phenomenology of cosmological models is theory-driven not only in terms of providing explanations for the diverse phenomena that are observed, as well as promoting alternative explanations of existing measurements, but also due to the increasing reliance on theorists to produce sophisticated numerical models of the Universe which can be used to refine and calibrate experimental probes. Without a dedicated effort to develop the tools and skill-sets necessary for the interpretation of the next generation of experiments, we risk being ``theory limited'' in essentially all areas of dark energy studies.  As a concrete example of this general trend, forecasts for determination of the dark energy equation of state and other cosmological parameters from next-generation observations of cosmological structure typically assume calibration against simulations accurate to the level of 1\\% or better. This target has rarely been met for simulations of complex nonlinear phenomena such as the formation of large-scale structure in the Universe. However it is precisely these probes, which provide information on both the geometry of space-time and the growth of large-scale structure, which will be key to unraveling the mystery of dark energy.  For upcoming measurements to be exploited to the full, theory must reach not only the levels of accuracy justified by the measurements but also cover a sufficiently wide range of cosmologies.  The problem breaks down to two questions: (i) What is a reasonable coverage of cosmological parameters, given the expected set of observations?  (ii) What is the required accuracy for theoretical predictions -- over this range of parameters -- for the given set of observations? It is crucial to realize that the ultimate requirement is on controlling the {\\em absolute\\/} error -- taking into account all of the relevant physics: gravity, hydrodynamics, and feedback mechanisms. This is much more difficult to achieve than {\\it relative\\/} error control -- e.g.~asking what the relative importance of baryonic physics is versus a baseline gravity-only simulation. Most recent papers discuss the latter, implicitly assuming the existence of a reference spectrum.  One aim of our work is to provide just such a reference spectrum within the boundaries outlined. We fully expect that the answers to both (i) and (ii) will evolve, requiring more accurate modeling of a smaller range of models, so we are most interested here in the near-term needs. Associated with the first problem is the fact that, given the impossibility of running complex simulations over the many thousands of cosmologies necessary for grid based or Markov Chain Monte Carlo (MCMC) estimation of cosmological parameters, one must develop efficient interpolation methods for theoretical predictions.  These methods must of course also satisfy the accuracy requirements of question (ii).  The control of errors in the underlying theory for the CMB is adequate to analyze results from Planck \\citep{CMBcompare,Won08}. This is, however, not the case for predictions of gravitational clustering in the nonlinear regime, as is required for cluster counts, redshift space distortions, baryon acoustic oscillations (BAO) and weak lensing (WL) observations. In the case of BAO, the galaxy power spectrum in the quasi-linear regime should be known to sub-percent accuracy, and for WL the same is true for the mass power spectrum to significantly smaller scales. Perturbation theory has errors on the mass power spectrum currently estimated to be at the percent level in the weakly nonlinear regime (see, e.g., \\citealt{JeoKom06} and \\citealt{Car09} for recent treatments, or \\citealt{Ber02} for an earlier review). To reduce these errors, test the approximations, and model galaxy bias, numerical simulations are unavoidable. Theoretical templates, in terms of current power spectrum fits based on simulations (with errors at the 5\\% level), are already a limiting factor for WL observations at wavenumbers $k\\sim 1\\,h\\,{\\rm Mpc}^{-1}$. \\citet{HutTak05} show that in order to avoid errors from imprecise theoretical templates mimicking the effect of cosmological parameter variations, the power spectrum has to be calibrated at about 0.5-1\\% for $0.1\\,h\\,{\\rm Mpc}^{-1}\\le k \\le 10\\,h\\,{\\rm Mpc}^{-1}$. The scale most sensitive for WL measurements is around $k\\sim 1\\,h\\,$Mpc$^{-1}$ and $z\\sim 0.5$ and the power spectrum therefore needs to be calibrated the most accurately at that point (see, e.g, \\citealt{HutTak05}, Figure~1). In a very recent paper, \\citet{Hil08} re-emphasize the need for very accurate predictions for the theoretical power spectrum, pointing out that currently used fitting functions such as the \\citet{PD96} formula or the fit derived by \\citet{Smi03} underestimate the cosmic shear-power spectra by $>30\\%$ for $k>10\\,h\\,{\\rm Mpc}^{-1}$.  \\begin{figure}[t] \\centerline{ \\includegraphics[width=3in]{xi_rat.eps}} \\caption{\\label{theta} Ratio of the E-mode correlation function with and without an assumed suppression of the power spectrum mimicking a possible systematic error in the matter power spectrum. This figure demonstrates that a gradual decrease in the accuracy of the matter power spectrum on small scales will not lead to a catastrophic error in the weak lensing prediction. The green line with $k_F=10\\,h\\,{\\rm Mpc}^{-1}$ corresponds to error properties which will be close to the degradation we expect for the matter power spectrum presented in this paper (see text).} \\end{figure}  We have independently assessed the impact a mis-modeled power spectrum would have on the predictions of weak lensing observables, including the fact that a wide range of spatial scales can be mapped into a given angular scale. Assuming a distribution of sources with $\\langle z\\rangle=1$, and using the Limber approximation, we compute the observable shear-shear correlation function, $\\xi(\\theta)=\\langle\\gamma(0)\\cdot\\gamma(\\theta)\\rangle$, given an estimate of the $z$-dependent mass power spectrum, $\\Delta^2(k,z)$. To mimic the inaccuracy of $\\Delta^2 (k,z)$ on scales smaller than $1\\,h\\,{\\rm Mpc}^{-1}$, we multiply it by a $z$-independent filter of the form $(1+k^2/k_F^2)^{-1}$ for a variety of $k_F$. At $\\ell=1000$ the suppression is 2-3\\% for $k_F=1\\,0\\,h\\,{\\rm Mpc}$ ($k_F$ being the assumed suppression of the power spectrum) and it drops to 1\\% at $\\ell~500$. Assuming that $k_F\\simeq 10\\,h\\,{\\rm Mpc}^{-1}$ reflects the error properties we are aiming at in this paper (i.e.~$\\Delta^2\\sim 1\\%$ low at $k\\simeq\\,1\\,h\\,$Mpc$^{-1}$ and smoothly but increasingly low for smaller scales) we expect our results could be used to predict the shear correlation function at the percent level for separations larger than $2^\\prime$. Figure~\\ref{theta} shows the expected error for different filter scales. Assuming sources at higher $z$ shifts all of the curves to larger scales, while a lower source redshift shifts the curves to smaller scales.  In order to extract precise cosmological information from WL measurements, additional physics beyond the gravitational contribution must be taken into account. At length scales smaller than $k\\sim 1\\,h\\,{\\rm Mpc}^{-1}$, baryonic effects are expected to be larger than 1\\% \\citep{Whi04,ZhaKno04,Jin06,Rud08,GuiTeyCol09} and will have to be treated separately, either directly via hydrodynamic simulations or, as is more likely, by a combination of simulations and self-calibration techniques (e.g.~constraining cluster profiles by cluster-galaxy lensing at the same time as constraining the shear). In any case, gravitational $N$-body simulations must remain the bedrock on which all of these techniques are based.  Taking all of these considerations into account, the purpose of this paper is to establish that gravitational $N$-body simulations can produce power spectra accurate to 1\\% out to $k\\sim 1\\,h\\,{\\rm Mpc}^{-1}$ between $z=0-1$ for a range of cosmological models. Given the success of the CDM paradigm in explaining current observational data we shall consider cosmologies within that framework. All of our models will assume a spatially flat Universe with purely adiabatic fluctuations and a power-law power spectrum. Since it is unlikely that near-term observations can place meaningful constraints on the temporal variation of the equation of state of the dark energy, we will restrict attention to cosmologies with a constant equation of state parameter $w=-p/\\rho$ (where $p$ is the pressure and $\\rho$ the density of the dark energy with $w=-1$ in a $\\Lambda$CDM cosmology). Since $\\Lambda$CDM is a good fit to the data, the accuracy of simulations can be established primarily around this point.  In this paper we will establish that gravitational $N$-body simulations can meet the above demands and derive a set of simulation criteria which balance the need for accuracy against computational costs. The target regime covers the most important range for current and near-future WL surveys and additional physics is controllable at the required level of accuracy. Showing that the required accuracy can be obtained from $N$-body simulations is only the first step in setting up a power spectrum determination scheme useful for weak lensing surveys. In order to analyze observational data and infer cosmological parameters, precise predictions for the power spectrum over a large range of cosmologies are required. This paper -- establishing that achieving the base accuracy is possible -- is the first in a series of three communications. In the second, we will demonstrate that a relatively small number of numerically obtained power spectra are sufficient to derive an accurate prediction scheme -- or emulator -- for the power spectrum covering the full range of desired cosmologies. The third paper of the series will present results from the complete simulation suite, named the ``Coyote Universe'' after the computing cluster on which it has been carried out. The third paper will also contain a public release of a precision power spectrum emulator.  In order to establish the accuracy over the required spatial dynamic range, as well as over the redshifts probed, a variety of tests need to be conducted. These include studies of the initial conditions, convergence to linear theory at very large length scales, the mass resolution requirement, and other evolution-specific requirements such as force resolution and time-stepping errors. To establish robustness of the final results, codes based on different $N$-body algorithms should independently converge to the same results (within error bounds). While some of these studies have been conducted separately and within the confines of the cosmic code verification project \\cite[]{CodeCompare}, this is the first time that the more or less complete set of possible problems has been investigated in realistic simulations.  We find that it is indeed possible to control the accuracy of $N$-body simulations at 1\\% out to $k\\sim 1\\,h\\,{\\rm Mpc}^{-1}$. Even though these scales are not very small, the simulation requirements are rather demanding. First, the simulation volume needs to be large enough to capture the linear regime accurately. Due to mode-mode coupling, nonlinear effects influence scales as large as 500~$h^{-1}$Mpc. Therefore, the simulation volume needs to cover at least 1~($h^{-1}$Gpc)$^3$. Second, with this requirement imposed, the number of particles necessary to avoid errors from discreteness effects at the smallest length scales of interest, also becomes substantial. As we discuss later, because we are measuring the mass power spectrum (which is sensitive to near-mean-density regions) numerical results aiming for accuracy at the sub-percent level can only be trusted at scales below the particle Nyquist wavenumber \\citep[see also][]{Joy08}. A $1\\,(h^{-1}{\\rm Gpc})^3$ simulation volume requires a minimum particle loading of a billion particles. Third, it is important to start the simulation at a high enough redshift to allow enough dynamic range (in time) for structures to evolve correctly and for the initial perturbations to be captured accurately by the Zel'dovich approximation (ZA). Lastly, the force resolution and time stepping has to be accurate enough to ensure convergence of the simulation results.  The paper is organized as follows. In Section~\\ref{chall} we use a simple example to demonstrate the need for precision predictions from theory. Section~\\ref{nbody} contains a description of the $N$-body codes used in this paper and some basic information about the simulations. In Section~\\ref{power} we briefly describe the power spectrum estimator. In Sections~\\ref{ics} and \\ref{tests} investigations of initial conditions and time evolution are reported, demonstrating that the simulations can achieve the required accuracy levels. Finally, we compare the numerical results to the commonly used semi-analytic {\\sc HaloFit\\/} approach \\citep{Smi03} in Section~\\ref{sec:halofit}, finding a discrepancy of $\\sim 5-10\\%$ between the fit and the simulations. We provide a summary discussion of our results in Section~\\ref{conclusion}. Appendix~\\ref{app:2lpt} discusses errors in setting up the initial conditions, comparing the Zel'dovich and 2LPT approximations. Appendix~\\ref{app:rich} provides details of the Richardson extrapolation procedure used for some of the convergence tests. ",
        "conclusions": "\\label{conclusion}  The advent of precision cosmological observations poses a major challenge to computational cosmology. With observational results accurate to the percent level a significant uncertainty in extracting cosmological information from the data is due to inaccuracies in theoretical templates. At the required level of accuracy large scale simulations are unavoidable, since the nonlinear nature of the problem makes it impossible to derive analytic or semi-analytic expressions for statistics such as the matter power spectrum, at an accuracy better than $\\sim 10\\%$. While simulations in principle should yield results at sub-percent accuracy, in practice this is a non-trivial task due to uncertainties in the numerical implementation and modeling of relevant physical processes.  Motivated by this realization, we decided to carry out an end-to-end calculation of one of the simplest non-trivial problems we could imagine: a percent level computation of the nonlinear mass power spectrum to $k\\sim 1\\,h\\,{\\rm Mpc}^{-1}$ over the range $0<z<1$. This was a problem which appeared useful and timely as well as tractable (if not straightforward) while still providing a meaningful learning environment -- by actually going through all of the steps we would map out the necessary infrastructure which would be required, find the most difficult pieces of the problem and present a proof-of-principle demonstration that meaningful, precision theoretical predictions could be used in support of future cosmological measurements.  We have broken the problem into three steps, to be presented in three publications. In this, first, paper we showed that it is possible to obtain a calibration of the nonlinear matter power spectrum at sub-percent/percent accuracy out to $k\\sim 1\\,h\\,{\\rm Mpc}^{-1}$ between $z=1$ and $z=0$. This wavelength regime is important for ongoing and near-future weak-lensing surveys. The restriction to these (large) length scales has two major advantages: baryonic effects are subdominant on these scales \\citep[e.g.][]{Whi04,ZhaKno04,Jin06,Rud08,GuiTeyCol09} and the numerical requirements in this regime remain rather modest. Each simulation can be carried out in a matter of days on parallel computers with several hundred processors and the data volume is manageable with arrays of inexpensive disk. Pushing beyond $k\\sim 1\\,h\\,{\\rm Mpc}^{-1}$ will require advances in our understanding of the implementation of baryonic physics, or self-calibration techniques, as well as advances in algorithms and computational power.  We derived a set of numerical requirements to obtain an accurate power spectrum by performing a large suite of convergence and comparison tests. The goal was a set of code settings which balance the need for precision and the limitation of computational resources. As shown here, the simulation volume and, especially, the particle loading are two major concerns in obtaining an accurate matter power spectrum. The simulation volume has to be in the $\\sim$Gpc$^3$ range, leading to a minimum requirement of $\\sim 1$ billion particles. Further increase in volume would be helpful, but would require a concomitant increase in the number of particles, greatly adding to the computational burden. The 1~Gpc$^3$/1~billion particle simulation is a good compromise between sufficient accuracy and computational cost.  Besides a large simulation volume and good particle sampling, initialization of the simulation also plays an important role. To guarantee converged results, the simulation must be started at a high enough redshift. We found that a starting redshift of $z_{\\rm in}\\simeq 200$ is sufficient to get accurate results between $z=1$ and $z=0$.  The results for the power spectrum are rather stable to changes in the number of time steps. This is clearly related to the fact that our resolution demands are relatively modest. For the PM runs, a few hundred time steps are sufficient, while for the tree-PM runs the overall number of time steps is a factor of ten larger. We emphasize that the simulation settings discussed here will lead to the required accuracy only up to $k\\sim1\\,h\\,{\\rm Mpc}^{-1}$. While these settings can be used as a guideline for other simulation aims, they do not replace convergence tests that must be performed for each new problem, if one desires high precision results.  While weak lensing was a primary motivation for this study, our efforts are of wider interest as an exercise in precision ``theoretical'' cosmology.  We demonstrated that it is possible to achieve 1\\% accuracy in the mass power spectrum in gravity only simulations on relatively large scales for a limited range of cosmological models.  Had this not been the case the field would have needed to rethink its demands on theory.  The non-trivial computational and human cost of even this ``first step'' argues for increased efforts in these directions in order to satisfy the increasingly stringent demands of future observations.  Having established the ability to generate power spectra with sufficient accuracy from $N$-body simulations, the next major question that arises is how to use these costly simulations for parameter estimation, e.g., via Markov Chain Monte Carlo. To address this problem, we have recently introduced the cosmic calibration framework \\citep{Hei06,Hab07,Sch08} which is based on an interpolation scheme for the power spectrum (or any other statistic of interest) derived from a relatively small number of training runs.  The next step in generating precise predictions for the matter power spectrum is to determine the minimum number of cosmological models needed to build an accurate emulator and then to construct the emulator from a set of high-precision simulations.  In the second paper of this series we establish that 30-40 cosmological models are sufficient to explore the parameter space for $w$CDM cosmologies (constant $w$) given the current constraints on parameter values.  The third and final paper will present results from the simulation suite designed and discussed in the second paper, and will include a power spectrum emulator that will be publicly released."
    },
    "0812/0812.1578_arXiv.txt": {
        "abstract": "We constrain the iron abundance in a sample of 33 low-ionization Galactic planetary nebulae (PNe) using [\\ion{Fe}{3}] lines and correcting for the contribution of higher ionization states with ionization correction factors that take into account uncertainties in the atomic data. We find very low iron abundances in all the objects, suggesting that more than 90\\% of their iron atoms are condensed onto dust grains. This number is based on the solar iron abundance and implies a lower limit on the dust-to-gas mass ratio, solely due to iron, of $M_{\\rm dust}/M_{\\rm gas}\\geq1.3\\times10^{-3}$ for our sample. The depletion factors of different PNe cover about two orders of magnitude, probably reflecting differences in the formation, growth, or destruction of their dust grains. However, we do not find any systematic difference between the gaseous iron abundances calculated for C-rich and O-rich PNe, suggesting similar iron depletion efficiencies in both environments. The iron abundances of our sample PNe are similar to those derived following the same procedure for a group of 10 Galactic \\ion{H}{2} regions. These high depletion factors argue for high depletion efficiencies of refractory elements onto dust grains both in molecular clouds and asymptotic giant brach stars, and low dust destruction efficiencies both in interstellar and circumstellar ionized gas. ",
        "introduction": "\\label{intro}  The progenitors of planetary nebulae (PNe), asymptotic giant branch (AGB) stars, have atmospheres particularly favorable for grain formation, and are considered the most efficient source of circumstellar dust \\citep[and references therein]{Whittet_03,Ferrarotti_06}. However, it is not clear yet how much dust do PNe have and whether this dust is destroyed or modified during their lifetime. \\citet{Pottasch_84} and \\citet{Lenzuni_89} studied PNe with {\\it IRAS} data and found that their derived dust-to-gas mass ratios decreased with nebular radius (which they used as a proxy for nebular age). However, these results strongly depend on the poorly known distances to the studied PNe, and were called into question by \\citet{Stasinska_99}, who derived dust-to-gas ratios using distance-independent quantities and found no correlation with the surface brightness in H$\\beta$, their proxy for nebular age. \\citeauthor{Stasinska_99} concluded that there is no evidence for a decrease in the dust-to-gas mass ratio as PNe evolve, but since there are many uncertainties involved, the issue is far from being settled.  As an alternative to dust-to-gas mass ratios derived from infrared emission, one might consider studying dust through element depletions. Elements such as Al, Ca, Si, Ni and Fe have abundances in the interstellar medium (ISM) much lower than solar \\citep{Morton_73, Morton_74}, and this is generally interpreted as due to their depletion in dust grains. The differences in depletion factors found in different environments give important clues on the nature of the formation and destruction mechanisms for dust grains  in the ISM \\citep[see, e.g.,][]{Whittet_03}, and the depletion factors in PNe can provide clues on the mechanisms that operate in ionized gas. To study the sensitivity of depletions to environment one must choose an element that is mostly condensed into dust grains, like those mentioned above, since in that case the destruction of a small quantity of dust will translate into a measurable increase of the element abundance in the gas. However, the abundances of these elements are usually difficult to measure in ionized gas due to the lack of suitable emission lines or atomic data, and due to the highly uncertain corrections for unobserved ions. These problems, combined with the wide spread in the degrees of ionization, and hence of ionization states, found in PNe, imply that the depletion factors derived so far for PNe use different ions and ionization correction factors (ICFs) and are not only uncertain, but also difficult to compare between them. The published values for the abundances of refractory elements in PNe cover the ranges: 1/6--1/300 the solar abundance for Ca, 1/2--1/350 for Al, 1/3--1/300 for Fe, near solar to 1/10 solar for Mg, and near solar to 1/20 solar for Si \\citep{Shields_75, Garstang_78, Shields_78,Pequignot_80,Aller_81, Shields_81,Aller_83,Beckwith_84,Pwa_86,Clegg_87a, Clegg_87b, Middlemass_90, Keyes_90, Kingdon_95,Pottasch_99,Perinotto_99,Casassus_00,Pottasch_01, Pottasch_02,Pottasch_03,Liu_04a,Pottasch_05,Sterling_05,Georgiev_06, Likkel_06, Pottasch_07a,Pottasch_07b,Pottasch_08}.  The problem with the determination of depletion factors in ionized gas is somewhat alleviated in the case of \\ion{H}{2} regions, where the range of degrees of ionization is much smaller. However, dust grains in PNe and \\ion{H}{2} regions are likely to have very different characteristics. The dust grains present in PNe formed in the cool atmospheres of the progenitor stars, whereas those grains now present in \\ion{H}{2} regions were located before in the associated molecular clouds and can be considered processed interstellar dust grains. Therefore, it is important to perform a homogeneous study of depletion factors in a sample of PNe, and especially so if the results can also be compared with those found in \\ion{H}{2} regions. Differences in depletion factors can provide much information on the efficiency of dust formation and destruction processes.  Of all the refractory elements we mentioned above, Fe has the strongest lines in the visible range of the spectrum. Furthermore, since most of the Fe atoms are condensed into dust grains and since the cosmic abundance of Fe is relatively high, this element is an important contributor to the mass of refractory dust grains \\citep{Sofia_94}, and the Fe gaseous abundance will probably reflect the abundance of refractory elements in dust. These reasons make Fe a good choice to study depletion factors in ionized gas.  The comparison between the results derived in PNe and \\ion{H}{2} regions will be more meaningful if the sample of PNe is restricted to objects with a low degree of ionization, because in that case the same ions need to be considered in the abundance determination for both types of objects. In \\ion{H}{2} regions, Fe will be mostly found in three ionization states: Fe$^{+}$, Fe$^{++}$, and Fe$^{+3}$. Fe$^{+}$ has a low ionization potential and its contribution to the total abundance is often negligible \\citep{Rodriguez_02}. On the other hand, [\\ion{Fe}{4}] lines are weak and more difficult to measure than [\\ion{Fe}{3}] lines. Hence, Fe abundances are usually calculated from Fe$^{++}$ abundances and an ICF derived from photoionization models. However, for the handful of objects in which [\\ion{Fe}{4}] lines have been measured, this ICF can be compared with that implied by the derived Fe$^{+3}$ abundances, and a discrepancy has been found between them \\citep[and references therein]{Rodriguez_03}. \\citet{Rodriguez_05} determined what changes in all the atomic data involved in the calculations would explain this discrepancy: (1) a decrease in the collision strengths for Fe$^{+3}$ by factors of 2--3, (2) an increase in the collision strengths for Fe$^{++}$ by factors of 2--3, or (3) an increase in the total recombination  coefficient or the rate of the charge-exchange reaction with  H$^0$ for Fe$^{+3}$ by a factor of $\\sim10$. \\citet{Rodriguez_05} argued that the three explanations are equally plausible, and derived two different ICFs: \\begin{equation}\\label{eq1} \\frac{\\mbox{Fe}}{\\mbox{O}}=0.9\\left(\\frac{\\mbox{O}^{+}}{\\mbox{O}^{++}}\\right)^{0.08}\\frac{\\mbox{Fe}^{++}}{\\mbox{O}^{+}} \\end{equation} \\begin{equation}\\label{eq2} \\frac{\\mbox{Fe}}{\\mbox{O}}=1.1\\left(\\frac{\\mbox{O}^{+}}{\\mbox{O}^{++}}\\right)^{0.58}\\frac{\\mbox{Fe}^{++}}{\\mbox{O}^{+}} \\end{equation} Equation~(\\ref{eq1}) is based on photoionization models that use the state-of-the-art values for the atomic data relevant to the problem, and Equation~(\\ref{eq2}) is derived from those objects with measurements of [\\ion{Fe}{4}] lines \\citep[see][]{Rodriguez_05}. Equation~(\\ref{eq2}) should be replaced by \\begin{equation}\\label{eq3} \\frac{\\rm {Fe}}{\\rm {O}}=\\frac{\\mbox{Fe}^+ + \\mbox{Fe}^{++}}{\\mbox{O}^+} \\end{equation} for those objects with $\\log(\\mbox{O}^+/\\mbox{O}^{++})\\geq-0.1$, since $\\mbox{Fe}^{++}$ and $\\mbox{O}^+$ will then dominate the total abundances of Fe and O. If the discrepancy were completely due to errors in the collision strengths  for Fe$^{+3}$, good values of the Fe abundance could be obtained from the  ICF of Equation~(\\ref{eq1}). If the collision strengths for Fe$^{++}$ were the  ones to blame, the correct abundance would be the previous value lowered  by $\\sim 0.3$~dex. Finally, if the models predictions were wrong  -- due to errors in the total recombination coefficient or the rate of the charge-exchange reaction -- the ICF of equation~(\\ref{eq2}) would give the best values of the Fe abundance. The discrepancy is probably due to some combination of the aforementioned causes, therefore the errors required in any of the atomic data are likely to be lower than those considered above, and the values of the Fe abundance will consequently be intermediate between the extreme values obtained with the ICF scheme described above.  The three ionization correction schemes give values of the Fe abundance that can be similar or differ by more than a factor of 10, but since these schemes involve drastic changes in the atomic data involved in the abundance calculation, the extreme values of the Fe abundance implied by them can be used to constrain the true values of the Fe abundances in the gas. \\citet{Rodriguez_05} followed this procedure and found that the two Galactic \\ion{H}{2} regions and the four Galactic PNe of their sample have less than 5\\% of their Fe atoms in the gas phase. In this paper, we apply this procedure to constrain the Fe abundances in a sample of 33 low-ionization Galactic PNe, and compare the results  with the values obtained for 10 Galactic \\ion{H}{2} regions. ",
        "conclusions": "We have constrained the iron abundances in a sample of 33 low-ionization PNe using the ICFs scheme developed by \\citet{Rodriguez_05}. This is the largest sample of PNe where the iron abundance has been calculated, including 18 nebulae with first determinations. The fact that we considered quite drastic changes in the atomic data involved, and the fact that we analyzed the whole sample with the same procedure, allow us to constrain the real value of the Fe abundances and compare the results between objects. The depletion factors ($[\\mbox{Fe/H}] = \\log(\\mbox{Fe/H}) - \\log(\\mbox{Fe/H})_{\\odot}$) of the PNe are below $-0.9$ dex, which implies that more than $\\sim$90\\% of the total Fe abundance is condensed onto dust grains. This result suggests that the efficiency with which Fe atoms attach to dust grains is higher that the one predicted by the models of AGB dust production of \\citet{Ferrarotti_06}. We derived a lower limit to the dust-to-gas mass ratio for the sample PNe just by considering that 90\\% of their iron atoms are condensed onto dust grains, $M_{\\rm dust}/M_{\\rm gas}\\geq1.3\\times10^{-3}$. Other elements, such as Si, Mg, C, and O, will have similar or higher contributions to the mass of dust grains.  The derived depletion factors span about two orders of magnitude, but we have not found any significant correlation between the derived Fe abundances and the evolutionary parameters of our sample PNe (surface brightness or electron density), in agreement with the results obtained by \\citet{Stasinska_99}. This result suggests that no significant destruction of dust grains is taking place in these objects. The different depletion factors shown by the PNe could be due either to different dust condensation efficiencies or to the destruction of a minor dust component in some objects.  We do not find any systematic difference in the Fe abundances that can be related to the morphology, the type of the central star (i.e., H-rich or H-poor central star), or the dust chemistry. This result suggests that C-rich and O-rich progenitor stars have similar iron depletion efficiencies. However, the lack of identifications of dust features in most of the sample PNe, and the uncertainties associated with the value of the $\\mbox{C/O}$ ratio, imply that this issue cannot be considered settled.  We have compared our results with the values derived for a group of 10 \\ion{H}{2} regions using the same procedure. The high depletion factors found for both kinds of objects imply that the atmospheres of AGB stars deplete refractory elements onto dust grains as efficiently as molecular clouds, whereas dust destruction processes are not very efficient in either \\ion{H}{2} regions or PNe."
    },
    "0812/0812.4441_arXiv.txt": {
        "abstract": "We explore  the dynamical and thermal evolution of the ejected neutron star crust in a Quark-Nova explosion. Typical explosion energies and ejected crust masses result in relativistic ejection with Lorentz factors of a few to a few hundred.  The ejecta undergoes a rapid cooling and stretching resulting in  break up into many small pieces (clumps) when the ejecta is only $\\sim 100$ km from the explosion site. The number and size of the clumps depends on whether the breakup occurs in the liquid or solid phase. For these two cases, the clump number is $\\sim 10^3$ (liquid phase) or $\\sim 10^7$ (solid phase) and, at break up, are spherical  (size $\\sim 10^4$ cm; liquid phase) or needle shaped ($\\sim 10^4\\times 10^2$ cm; solid phase). ",
        "introduction": "In the quark-nova (QN) picture, (Ouyed et al. 2002;  Ker\\\"anen et al. 2005; hereafter ODD and KOJ respectively) the core of a neutron star, that undergoes the phase transition to the quark phase, shrinks in a spherically symmetric fashion to a stable, more compact strange matter configuration faster than the overlaying material (the neutron-rich hadronic envelope) can respond, leading to an effective core collapse. The core of the neutron star is a few kilometers in radius initially, but shrinks to 1-2 km in a collapse time of about 0.1 ms \\cite{LBV}. The gravitational potential energy released (plus latent heat of phase transition) during this event is converted partly into internal energy  and partly into outward propagating shock waves which impart kinetic energy to the material that eventually forms the ejecta.   There are three previously proposed mechanisms for  ejection of the outer layers of the neutron star (i.e. crust): (i) Unstable baryon to quark combustion leading to a shock-driven  ejection (Horvath \\& Benvenuto 1988). More recent work, assuming realistic quark matter equations of state, argues for strong deflagration (Drago et al. 2007)  that can expel surface material.  In these models up to $10^{-2} M_{\\odot}$ can be ejected. These calculations focus on the microphysics and not on the effect of the global state of the resulting quark core which collapses prior to complete combustion (KOJ), leading to conversion only of the inner core ($\\sim$ 1-2 km) of the neutron star; (ii) Neutrino-driven explosion where the energy is deposited in a thin (the densest) layer at the bottom of the crust above a gap separating it from the  collapsing core (KOJ). For the neutrino-driven mechanism, the core bounce was neglected, and neutrinos emitted from the conversion to strange matter  transported the energy into the outer regions of the star, leading to heating and subsequent mass ejection. Consequently, mass ejection is limited to about $10^{-5}M_{\\odot}$ (corresponding to the crust mass below neutron drip density)  for compact quark cores of size (1-2) km; (iii) Thermal fireball driven ejection which we consider for the present study. The fireball is inherent to the properties of the quark star at birth. The birth temperature  was found to be of the order of 10-20 MeV since the collapse is adiabatic rather than isothermal (ODD; KOJ).  In this temperature regime the quark matter is in the superconducting Color-Flavor Locked (CFL) phase (Rajagopal\\&Wilczek 2001)  where the photon emissivity dwarfs the neutrino emissivity (Vogt, Rapp, \\& Ouyed 2004; Ouyed, Rapp, \\& Vogt 2005). The average photon energy is $\\sim$3$T$ for a CFL temperature $T$: since the plasma frequency of CFL matter is $\\hbar\\omega\\sim 23$ MeV (e.g. Usov 2001), the photon emissivity is highly attenuated as soon as the surface temperature of the star cools below $T_{\\rm a} = \\hbar\\omega_{\\rm p}/3\\simeq 7.7$~MeV. To summarize, the thermal fireball is generated as the star cools from its birth temperature down to $\\sim 7.7$ MeV.   Here we focus on the case where a QN goes off in isolation, neglecting any interaction with the surroundings, and the crust is ejected by a thermal fireball (i.e. case (iii) above). The remainder of this  paper is presented as follows: In Sect. 2 we discuss the energetics of the QN and the resulting fireball and ensuing crust ejection. Sect. 3   deals with the hydrodynamical and thermal evolution of the ejecta. We   conclude in Sect. 4. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we investigate the thermal and dynamic evolution of a relativistically expanding iron-rich shell from a QN explosion. The QN produces a photon fireball which acts as piston to eject and accelerate the crust of the parent neutron star to mildly relativistic speeds.  We find that the shell rapidly cools and breaks up into numerous ($10^3$ to $10^7$) chunks because of the rapid lateral expansion. Breakup may occur while the material is in liquid or solid phase: if liquid the chunks are nearly spherical; if solid they are needle shaped moving parallel to their long axis.  Although the presented model is based on physical arguments, most of these are in reality more complicated and so would require more detailed studies and the help of numerical simulations. For example,  the process of clumping, crystallization, and breakup of the ejecta,  would require better knowledge of the ambient conditions surrounding the ejecta. The astrophysical implications will be discussed elsewhere."
    },
    "0812/0812.2995_arXiv.txt": {
        "abstract": "We present cosmological solutions from the dimensionally reduced Chern-Simons term and obtain the smooth transition solution from the decelerated phase (AdS) to the accelerated phase (dS). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.1028_arXiv.txt": {
        "abstract": "Tidal dissipation may be important for the internal evolution as well as the orbits of short-period massive planets---hot Jupiters. We revisit a mechanism proposed by Ogilvie and Lin for tidal forcing of inertial waves, which are short-wavelength, low-frequency disturbances restored primarily by Coriolis rather than buoyancy forces.  This mechanism is of particular interest for hot Jupiters because it relies upon a rocky core, and because these bodies are otherwise largely convective.  Compared to waves excited at the base of the stratified, externally heated atmosphere, waves excited at the core are more likely to deposit heat in the convective region and thereby affect the planetary radius.  However, Ogilvie and Lin's results were numerical, and the manner of the wave excitation was not clear.  Using WKB methods, we demonstrate the production of short waves by scattering of the equilibrium tide off the core at critical latitudes.  The tidal dissipation rate associated with these waves scales as the fifth power of the core radius, and the implied tidal $Q$ is of order ten million for nominal values of the planet's mass, radius, orbital period, and core size. We comment upon an alternative proposal by Wu for exciting inertial waves in an unstratified fluid body by means of compressibility rather than a core.  We also find that even a core of rock is unlikely to be rigid.  But Ogilvie and Lin's mechanism should still operate if the core is substantially denser than its immediate surroundings. ",
        "introduction": "\\label{sec:intro}  The discovery of extrasolar planets has sharpened the need for a predictive theory of tidal circularization and synchronization.  Some $2\\%$ of nearby single FGK stars harbor roughly Jupiter-mass planets with orbital periods below ten days \\citep[and references therein]{Cumming_etal08}.  The orbital eccentricities of the shorter-period planets, especially those below four days, are markedly less than those of the longer period systems, presumably as a result of tidal dissipation \\citep{Rasio_etal96}.  Synchronization is likely to occur more easily and at longer periods than circularization because of the smaller moment of inertia associated with the spin as compared to the orbit \\citep{Lubow_etal97}.  For close stellar binaries, the lack of an adequate tidal theory causes little uncertainty concerning the internal evolution of the components since tidal heating contributes negligibly to the luminosity, and since it is assumed that tides are always sufficient to synchronize the spins of stars that come into contact with their Roche lobes.  The intrinsic luminosities of jovian planets due to their gradual contraction and loss of primordial heat are so small, however, that tidal heating may be competitive.  Indeed tides have been invoked to account for the anomalously large radii and low densities, compared to baseline models, of some planets that are observed to transit their host stars \\citep{Bodenheimer_etal01}.  Most studies of orbital evolution reduce the tidal uncertainties to a single parameter, the tidal quality factor $Q$, which is inversely proportional to the tidal dissipation rate, and which one hopes to calibrate by reference to stellar binaries of known age, such as those in star clusters, or to the inferred tidal interactions between Jupiter and its Galilean satellites \\citep[and references therein]{GS66,Mazeh08}.  A difficulty with such pure empiricism is that $Q$ may depend upon structural details such as composition, equation of state, rate of rotation relative to the body's density or to the tidal period, and so on, which differ between the object of interest and the rather limited set of calibrators.  Indeed there is some indirect evidence that $Q$ does vary.  As has been pointed out \\citep{Rasio_etal96,Sasselov03,Jackson_etal08}, differences in $Q$ between planets and stars, as well as possible dependence of $Q$ on the ratio of the orbital to spin frequency \\citep{Ogilvie_Lin07}, may be crucial to the survival of hot Jupiters: since the host stars rotate subsynchronously \\citep{Fabrycky_etal07} and the mass ratio is large, tidal dissipation within the star would tend to drag the planet inward.  In order to assess the importance of tidal inputs for hot Jupiters, furthermore, it is necessary to predict not only the overall rate of tidal dissipation, but also where within the planet mechanical energy is converted to heat.  This requires additional assumptions or knowledge beyond an empirically calibrated $Q$ alone.  For example, perhaps the best understood and most predictive tidal mechanism, at least for nonrotating bodies, is the excitation of shortwavelength g-modes at an interface between convective (isentropic) and radiative (stratified) regions, originally proposed by Zahn for application to early-type stars \\citep{Zahn70,Zahn75} and first applied to hot Jupiters by \\cite{Lubow_etal97}.  The g-modes would dissipate within the radiative region since they do not propagate in isentropic regions. For irradiated planets, as for early-type stars, this means that dissipation would occur in the outer parts.  On the other hand, if turbulent viscosity due to convection were effective \\citep{Zahn66}, then the heat would be deposited locally at depth. Because the thermal timescales and density scale heights differ greatly between the convective and radiative regions, the g-mode and turbulent mechanisms would have different consequences for the planetary radius even if they produced the same $Q$.  As argued by \\cite{Goldreich_Nicholson77}, the effective viscosity of convection is probably very strongly suppressed when the turnover time of the larger convective eddies exceeds the tidal period \\citep[see also][]{Zahn89,Goodman_Oh97,Penev_etal07}.  Since jovian planets are very deeply into this regime ($\\tau_{\\rm conv}\\sim 10^3\\textrm{yr}$), it is unlikely that convection can be responsible for $Q$ values as low as are inferred for Jupiter, $Q_p\\approx 10^5-10^6$ \\citep{GS66,Peale_Greenberg80}.  While bearing in mind that phase transitions or other nontrivial microphysics may contribute to tidal dissipation \\citep[e.g.][]{Stevenson80}, we take the traditional view that the likely alternative is a dynamical tide: that is, resonant tidal forcing of a low-frequency, short-wavelength mode or traveling wave at special locations within the planet, such as the convective-radiative interface already discussed.  Short waves have two basic physical advantages as candidates for a tidal theory.  First, they are more easily damped than the large-scale ellipsoidal distortion (``equilibrium tide''), by radiative diffusion \\citep{Zahn75} or by nonlinear breaking \\citep[e.g.][]{Goodman_Dickson98}.  Second, the number of short-wavelength modes is potentially large, scaling as $(\\Rp/\\lambda)^2$ for modes with a fixed azimuthal dependence $\\exp(im\\phi)$.  For direct forcing, the frequency of the mode as well as the azimuthal order $m$ must match those of the tidal potential. Viewed in the corotating frame of the planet (we assume a uniformly rotating background state), this frequency is normally much lower than the fundamental dynamical freqency $\\omega_{\\rm dyn}\\equiv (3G\\Mp/\\Rp^3)^{1/2}$, so the modes of interest must be approximately noncompressive and restored by buoyancy and/or rotation rather than pressure.  The nomenclature for such low-frequency motions is rich and bewildering: internal waves, g modes, r modes, toroidal modes, hybrid modes, Hough modes.  These names distinguish the relative importance of buoyancy versus rotation and other special properties. In this paper, we concentrate on modes restored by rotation rather than buoyancy, which we refer to collectively as ``inertial'' modes or waves.  All of these nearly incompressible motions have frequencies in limited ranges controlled by the Coriolis and \\BV parameters, regardless of wavelength.  Local dispersion relations for the wave frequency depend upon the ratios of components of the three-dimensional wavevector, but hardly depend upon the wavelength itself, except via viscous or radiative damping terms.  This behavior is completely different from that of modes depending upon compressibility (acoustic or p modes) or selfgravity.  In the ideal-fluid limit, the spectrum of modes that can be resonant with a tide in the allowed frequency range is therefore dense \\citep{Papaloizou_Pringle81}. But, the tidal potential must have a finite projection onto a resonant mode in order to excite it. So, the main challenge for a theory of dynamical tides is to estimate these projections, also called overlap integrals.  Two recent and independent studies have proposed that inertial modes may be tidally forced even in the absence of a stably stratified surface layer, and at rates that approach what is needed to explain the observationally inferred $Q_p$, provided that the tidal frequency in the corotating frame is less than twice the rotation frequency; this is the relevant regime for circularization of a planet that is already synchronized.  By direct numerical methods, \\citet[hereafter OL04]{Ogilvie_Lin04} calculated the linear excitation of inertial waves/modes in a compressible spherical annulus surrounding a solid core.  The current belief is that Jupiter itself contains $15-30 M_\\oplus$ of heavy elements (atomic number $Z>2$), or $5-10\\%$ of the planet's total mass, and also that many or most of the hot Jupiters may be even more enriched; whether these ``metals'' are concentrated in a distinct core is uncertain \\citep{Guillot05,Guillot_etal06,Burrows_etal07}.  OL04 calculate the tidal dissipation in steady state by balancing the excitation against an artificial viscous term, but they conclude that the tidal dissipation due to the excitation of inertial modes, though varying erratically with frequency in the inertial range (i.e., $|\\omega_{\\rm tide}|<|2\\Omega|$), tends towards finite values corresponding to $Q\\sim 10^6$ in the inviscid limit.  This they explain by the presence in their low-viscosity models of very short-wavelength disturbances concentrated on a ``web of rays'' in the poloidal plane.  They find that the intensity of this short-wavelength response correlates with the size of the core, and they suggest that it has something to do with closed ray paths (``wave attractors'') reflected alternately by the inner and outer boundaries of the spherical annulus.  OL04's interpretation of their own numerical results seems to have been informed by previous work by \\cite{Rieutord_Valdettaro97} and by \\cite{Rieutord_etal01}, who studied the linear modes of an incompressible, slightly viscous liquid in a spherical shell, by a combination of numerical and analytic methods.  The latter authors demonstrated the existence of very fine features in the velocity, and they explored the relationship of these features to critical latitudes and wave attractors, but they did not consider tidal excitation. Subsequently, \\citet{Ogilvie05} used a simplified analytical model to argue that wave attractors can absorb energy at nonzero rates that vary continuously with forcing frequency in the inviscid limit.  The models of OL04 include a stably stratified atmosphere, with an apparently independent set of ``Hough modes'' excited at the interface between this zone and the convective region, as previously demonstrated for rotating massive stars by Savonije and Papaloizou \\citep{Savonije_Papaloizou97, Papaloizou_Savonije97}.  The Hough modes appear to be much less dependent on the core, and the dissipation associated with them varies much more smoothly with tidal frequency; also, these modes extend outside the inertial range, i.e. to $|\\omega_{\\rm tide}|>|2\\Omega|$.  Shortly after the work of OL04, Y.~Wu claimed to demonstrate the tidal excitation of inertial modes in entirely unstratified and coreless bodies.  In \\cite{Wu05a}, exploiting methods developed by \\cite{Bryan} for incompressible rotating bodies, she analyzes the properties of free (unforced) modes of oscillation in compressible models with special radial density profiles.  In \\cite{Wu05b}, she calculates spatial overlap integrals between these modes and a quadrupolar perturbing potential to find $Q\\sim 10^9$; she argues that $Q$ may fall to $\\lesssim 10^7$ in more realistic models with a radial density jump due, for example, to a first-order phase transition.  Using a different mathematical formalism, but similar underlying low-frequency approximation, \\cite{Ivanov_Papaloizou07} have calculated tidal excitation of inertial modes in planets on highly eccentric orbits, with realistic but isentropic equations of state.  This is an extension of earlier calculations for rotating $n=3/2$ polytropes by \\cite{Papaloizou_Ivanov05}.  Their models are coreless, like those of Wu, but unlike hers, their tidal excitation is dominated by two large-scale inertial modes, perhaps because they treat the perturbing potential as uncorrelated from one periastron encounter to the next, so that the excitation is non-resonant.  It is clear from the above that considerable progress has been made in recent years on the dynamical tides of rotating planets, but the importance of short-wavelength inertial waves remains obscure: in particular, the circumstances under which they can be tidally excited. OL04's results indicate that a solid core is somehow important, while Wu's work suggests that it may not be.  It is desirable to achieve a semianalytic understanding of the excitation before making detailed calculations for realistic planetary interiors. For one thing, since the structure of hot Jupiters is still much less well understood than that of stars, it will be useful to know what features of the structure are most important for the tidal problem. For another, the apparently chaotic variation of tidal torque with tidal frequency seen in the results of OL04, and to some extent in the earlier ones of \\cite{Savonije_Papaloizou97}, raise questions about what is required of a numerical calculation to obtain convergence in the inviscid limit, or indeed what it means to converge.  \\goodbreak ",
        "conclusions": "\\label{sec:discussion}  Up to this point, we have had little to say about the work of Yanquin Wu on the excitation of inertial waves in \\emph{coreless} but compressible isentropic fluid bodies \\citep{Wu05a,Wu05b}.  Here we explain why we believe that her calculations, though technically superb, are not applicable to tides in real planets or stars.  Some of these criticisms may apply also to the work of \\cite{Ivanov_Papaloizou07}, who adopted a ``low-frequency'' approximation that appears to be physically equivalent to Wu's.  But we focus on Wu's work because of the admirable clarity of her exposition, and because she is concerned with resonant excitation rather than the quasi-impulsive excitation analyzed by \\cite{Ivanov_Papaloizou07}.  We then go on to discuss whether the rock-and-ice core of a jovian planet can be regarded as rigid, or even solid, and the consequences for the production of inertial waves if it is not.  \\subsection{Previous tidal calculations for coreless isentropic bodies} \\label{subsec:Wu}  The salient claims by Wu that we address are the following: \\begin{enumerate} \\item The overlap integrals between the short-wavelength inertial modes and the perturbing tidal potential diminish as a negative power of the number of radial or latitudinal nodes, rather than exponentially, even for unperturbed radial density profiles that are smooth apart from a power-law convergence to zero at the surface. \\item For many radial nodes, the overlap is concentrated toward the surface; this is where the excitation mainly occurs. \\item The overlap integrals vanish for a constant-density, incompressible body. \\end{enumerate} Wu suggested that the tidal coupling could be enhanced by a density discontinuity associated with a first-order phase transition within the convection zone, perhaps at the interface between molecular and metallic hydrogen.  This is probably true in principle, but we address here only the claims made for strictly isentropic bodies, thereby excluding first-order phase transitions because of the entropy jump associated with latent heat.  There does not appear to be a consensus as to the first-order nature of the molecular-to-metallic transition in hydrogen \\citep[][and references therein]{Militzer_etal08,Sumi_Sekino08}.  We agree in part with the third of the ennumerated claims above but take issue with the first two.  Our counter-arguments are mainly these: \\begin{enumerate} \\item By selectively neglecting compressibility in some places while retaining it in others, and especially by oversimplifying the low-frequency limit of the outer boundary condition, Wu has created a singularity at the boundary, though the dynamics there should actually be smooth if the boundary is free.  We suspect that the tidal coupling she calculates is due mainly or entirely to this singularity. \\item For barytropic bodies [$P=P(\\rho)$], albeit with an inconsistent treatment of their gravitational potentials, there exist tidal responses that completely lack short-wavelength components and that can be exhibited in closed form.  The fluid displacement and enthalpy are independent of adiabatic index, and so are the same for a compressible as for an incompressible body. \\end{enumerate} We now expand upon these last two points.  After deriving the equivalent of equation (\\ref{eq:wavey}), Wu argues that the lefthand side, which contains the sound speed $\\cs$ in the denominator, can be neglected on the grounds that compression of the fluid is very slight for modes that are both short in wavelength ($kR\\gg 1$) and low in frequency ($\\omega^2\\ll\\cs^2k^2$).  This term comes directly from the time derivative in the continuity equation, so dropping it is equivalent to replacing the continuity equation by $\\diver(\\rho\\vv_1)=0$.  But according to her formalism, the tidal forcing is explicitly $\\propto\\cs^{-2}$, as shown by eq.~(7) of \\cite{Wu05b}, so the tidal coupling itself should vanish in the limit $\\cs^{-2}\\to0$.  Indeed, her Appendix~B makes this explicit.  However the neglected term is singular at the free surface, because $\\cs^2\\to0$ there. In the full equation (\\ref{eq:wavey}), the singularity involving $\\cs^{-2}$ is balanced by another singularity involving $\\propto\\rho^{-1}\\grad\\rho$, where $\\rho$ is the unperturbed density, Using our equations (\\ref{eq:vsol}) \\& (\\ref{eq:wavey}), it can be seen that the condition under which the two singularities cancel one another is precisely the free boundary condition (\\ref{eq:freebc}), whose physical interpretation is the vanishing of the lagrangian enthalpy perturbation.  If, following \\cite[\\S2.1 \\& \\S4.1]{Wu05a}, one neglects the righthand side of eq.~(\\ref{eq:wavey}) but retains the gradient of the unperturbed density profile on the righthand side, then the only possible well-behaved modes of the system are those for which the normal components of the fluid velocity and displacement [$\\vxi=\\vv/(-i\\omega)$] vanish at the boundary.  Evidently, Wu believed that taking $\\boldsymbol{\\hat n\\cdot\\xi}=0$ at the boundary is an acceptable approximation because (i) the inertial modes don't move the boundary very far, and (ii) the density vanishes at the boundary anyway.  The modification to the boundary condition affects the structure of the mode not just at the surface, however, but down to depths comparable to that of the first radial node of $\\psi$ below the surface. Let $h$ be half the nodal depth, so that $\\psi(h)\\approx\\psi(0)$, where $\\psi(0)$ is the surface value of the eigenfunction.  Denote the unperturbed density at this depth by $\\rho(h)$, and the sound speed by $\\cs(h)= \\sqrt{gh(N+1)/N}$, assuming a polytropic equation of state $P\\propto\\rho^{(N+1)/N}$.  The eulerian density perturbation at this depth is \\begin{equation} \\label{eq:eulerianpert} \\rho_1(h)=\\rho(h)\\frac{\\psi(h)}{\\cs^2(h)}\\approx\\frac{N\\rho(h)}{N+1} \\frac{\\psi(0)}{gh}\\,. \\end{equation} (With Wu's definition of $\\psi$, there would be a factor of $\\omega^2$ here.) On the other hand, since $\\rho(R)\\propto(\\Rp-R)^N$ near the surface, the contribution to $\\rho_1=-\\diver(\\rho\\vxi)$ from the surface displacement, which Wu neglects (as do Ivanov \\& Papaloizou), is \\begin{equation} \\label{eq:disprho} \\left.-\\vxi\\bdot\\grad\\rho\\right|_{R=\\Rp-h}\\approx \\frac{N\\rho(h)}{h}\\xi_R(0) \\approx \\frac{N\\rho(h)}{h}\\frac{\\psi(0)}{g}\\,, \\end{equation} where we have evaluated $\\xi_R(0)$ from the free boundary condition (\\ref{eq:freebc}) in the last step, taking $\\Phi_1=0$ as appropriate for a free mode of oscillation with negligible self-gravity. Evidently, the neglected contribution (\\ref{eq:disprho}) is comparable to the total (\\ref{eq:eulerianpert}), and therefore is not negligible within the first node.  In Wu's powerlaw-sphere model, where \\begin{equation} \\label{eq:wusphere} \\rho\\propto (R_{\\max}^2-R^2)^\\beta \\end{equation} the nodal depth is $\\propto R/n$, except near the critical colatitudes $\\theta=\\cos^{-1}(\\pm\\omega/2\\Omega)$ (the ``singularity belt'' in Wu's parlance) where it scales $\\propto R/n^2$. Here $n=n_1+n_2$ in terms of Wu's modal indices $n_i\\ge0$, which are roughly proportional to the WKB wavenumbers of the inertial modes: that is, $n\\gg 1$ for shortwavelength modes. Wu finds that the tidal forcing, which is proportional to the overlap integral between the perturbing potential $\\Phi_1$ and the modal eigenfunction, scales with this index approximately as $n^{-2\\beta-1}$.  This is consistent with the idea that the excitation occurs within the first node from the surface (and probably also within the singularity belt) since the fraction of the planetary mass within depth $2h$ of the surface for the density profile \\eqref{eq:wusphere} is $\\Delta M/M \\propto(2h)^{\\beta+1}$.  It is true that the horizontal components of the velocity and displacement near the boundary are larger than their radial components by a factor $\\sim\\omega_{\\rm dyn}^2/\\Omega^2\\gg1$, where $\\omega_{\\rm dyn}\\equiv(3g/R)^{1/2}$ is the dynamical frequency of the planet. So it is likely that the errors in modal energies and eigenfrequencies caused by the approximate boundary condition $\\xi_R=0$ are slight.  But the overlap integrals are of a higher order of smallness in $(\\Omega/\\omega_{\\rm dyn})^2$, so that their relative errors could be large.  The discussion so far does not make clear the sign of the error (if there is one): perhaps the tidal coupling would be \\emph{larger} with the exact boundary condition.  The following model system, which is borrowed from \\cite{GNG87}, suggests that the error is in the direction of overestimating the coupling.  Let the equation of state again be polytropic, with $\\rho\\propto w^N$, $P\\propto w^{N+1}$, and $\\cs^2=w/N$, where $w\\equiv(N+1)P/\\rho$ is the enthalpy, which remains finite in the incompressible limit $N\\to 0^+$. To match (\\ref{eq:wusphere}), the unperturbed enthalpy should be \\begin{equation} \\label{eq:w0} w(x,y,z)= w(0)\\left(1-\\frac{x^2+y^2+z^2}{R_{\\max}^2}\\right), \\end{equation} and $N=\\beta$.  Hydrostatic equilibrium in the corotating frame (where $\\vv=0$) requires \\begin{equation} \\label{eq:hydrostatic} 0=  \\grad\\left[w+\\Phi-\\frac{1}{2}\\Omega^2(x^2+y^2)\\right], \\end{equation} where $\\Phi$ is the unperturbed potential.  Since $w$ and the centrifugal term are quadratic functions, $\\Phi$ must also be such a function. In this regard, the model differs from the one considered by Wu.  In effect, she calculates the potential due to the density profile (\\ref{eq:wusphere}) from Poisson's equation and uses this to determine the enthalpy.  Consequently her pressure profile is not simply a power of the density profile, and $\\Gamma_1\\equiv (\\pd\\ln P/\\pd\\ln\\rho)_S$ is not constant in her models, at least not for $\\beta>0$.  However, she argues at several points\\footnote{For example, in \\citet[after eq.~(C7)]{Wu05b}: ``The results only depend on the boundary behavior of $f(\\Theta)$ as long as it is sufficiently smooth.  This explains why models with different polytrope representations ($\\rho\\propto[1-(r/R)]^\\beta$ or $p\\propto\\rho^{1+1/\\beta}$) give rise to essentially the same overlap integrals.''} that the tidal forcings she calculates depend, at least for a smooth density profile, only on the behavior near the boundary, where it is indeed approximately true that $P\\propto\\rho^{(\\beta+1)/\\beta}$.  As a matter of fact, because of her neglect of the $\\cs^{-2}$ term in eq.~(\\ref{eq:wavey}), the equation of state doesn't enter her calculations of the normal modes: only the density profile does, which could result from many isentropic equations of state paired with an appropriate background potential. The overlap integrals do depend upon the equation of state via the sound speed, but as long as $\\cs^2$ approaches zero linearly near the boundary, it is hard to see how the overlap integrals for short-wavelength modes could be sensitive to the full functional forms of $\\cs^2(R)$, and therefore of $\\Phi_0(R)$, if they are excited near the boundary.  For these reasons, it does not seem crucial that the unperturbed potential be fully consistent with the mass distribution.  After elimination of the density in favor of the enthalpy, the linearized equations become \\begin{subequations}\\label{eq:polytide} \\begin{align} \\label{eq:pteuler} -i\\omega\\vv+\\boldsymbol{2\\Omega\\times v}+\\grad w_1  &= -\\grad\\Phi_1,\\\\ \\label{eq:ptcont} -i\\omega w_1 +\\vv\\bdot\\grad w+N^{-1}w\\diver\\vv &=0. \\end{align} \\end{subequations} Now suppose that the tidal potential is quadrupolar: that is, a homogeneous and harmonic quadratic polynomial in $(x,y,z)$; for definiteness, \\begin{equation} \\label{eq:htide} \\Phi_1 = \\frac{A}{2}(x+iy)^2 e^{-i\\omega t} \\end{equation} where $A$ is a constant, and the exponential factor will be taken as read hereafter.  Then since $w$ is also a second-degree polynomial, eqs.~(\\ref{eq:polytide}) can be satisfied by taking the components of $\\vv$ to be polynomials of the first degree, and $w_1$ of the second degree.  After some algebra, \\begin{align}\\label{eq:ptsol} v_x &= -i v_y = a(x+iy), & v_z&=0, & w_1= \\frac{B}{2}(x+iy)^2,\\nonumber\\\\ \\mbox{where}~~a&= \\frac{-i\\omega A}{\\omega(\\omega+2\\Omega)-4w(0)/R_{\\max}^2}\\,, &\\mbox{and}~B &= \\frac{4iw(0)}{\\omega R_{\\max}^2}a\\,. \\end{align} Since $\\diver\\vv=0$, the equation of state (i.e. $N$) doesn't enter the solution (\\ref{eq:ptsol}).  Also, the relative vorticity $\\boldsymbol{\\nabla\\times}\\vv=0$, so this solution has the same total vorticity $2\\vOmega$ as it would have in the absence of the tide and therefore might be the solution of an initial value problem in which the tide was ``turned on'' slowly.  The denominator in the expression for $a$ vanishes at $\\omega=-\\Omega\\pm(\\omega_0^2-\\Omega^2)^{1/2}$, where $\\omega_0\\equiv 2\\sqrt{w(0)}/R_{\\max}$ is the dynamical frequency of the model and therefore presumably $\\gg\\Omega$.  At the unperturbed surface where $w=0$, eq.~(\\ref{eq:ptcont}) reduces to the free boundary condition that the lagrangian enthalpy perturbation vanishes.  These details aside, the important point is that {\\it the tidal response is entirely long-wavelength for any polytropic index when a free rather than rigid outer boundary condition is used,} at least in this idealized coreless model, which uses a nonselfconsistent but smooth unperturbed potential, and at least for a nonsynchronous body in a circular orbit.  Short-wavelength inertial modes are not tidally forced even though the fluid is compressible.  \\subsection{Rigidity of the core}\\label{subsec:rigidity}  OL04 assumed the core to be solid, and therefore impenetrable by low-frequency but short-wavelength disturbances such as inertial waves, but sufficiently plastic as to comply with the large-scale equilibrium tide as if it were fluid. Here we re-examine the strength and solidity of the core.  In agreement with OL04, we find that the elastic strength of even a solid core would be negligible as regards the equilibrium tide.  However, we evaluate the equilibrium tide differently than OL04, and we allow for the density contrast between the core and its immediate surroundings. Furthermore, we estimate that the core is most probably fluid rather than solid.  It is presumed that the cores of Jovian planets consist of elements heavier than hydrogen and helium, more specifically of some combination ``rock'' (refractory minerals such as silicates and iron) and ``ice'' (molecular species such as ${\\rm H_2O}$, ${\\rm CH_4}$, and ${\\rm NH_3}$) \\citep{Guillot05}.  To support shear stress, these materials would have to be in a solid phase.  The pressure at the surface of the core is comparable to the central pressure of an $N=1$ polytrope: \\begin{equation} \\label{eq:Pcore} P_{\\rm c}\\approx \\frac{\\pi}{8}\\frac{G\\Mp^2}{\\Rp^4}\\approx 4\\times 10^{13} \\left(\\frac{\\Mp}{M_{\\rm J}}\\right)^2\\left(\\frac{\\Rp}{R_{\\rm J}}\\right)^{-4}\\ {\\rm dyn\\,cm^{-2}}\\,, \\end{equation} or $\\sim 40\\,{\\rm Mbar}$.  The temperature is more difficult to predict.  Jovian planets are supported mainly by degeneracy pressure, so the central temperature has only a modest effect on the planetary radius.   The temperature depends upon the planet's age and rate of cooling, which in turn depend upon on uncertain opacities in the envelope, not to mention the possibility of internal heating by tides. Present estimates are in the range \\begin{equation} \\label{eq:Tcore} T_{\\rm c}\\sim 2\\mbox{-}4\\times 10^4\\,{\\rm K} \\end{equation} based on standard models fit to radii of transiting planets and the estimated ages of their host stars \\citep[e.g.][]{Arras_Bildsten06}.  The pressure (\\ref{eq:Pcore}) is large compared to bulk moduli of common refractory materials at room temperature---e.g. $K=1.7$ and $1.0$~Mbar for iron and silicon, respectively---so rocky cores will be compressed to densities $\\gtrsim 20\\,{\\rm g\\,cm^{-3}}$ [we derive this number from a generic equation of state for ``rock'' by \\cite{Hubbard_Marley89}], i.e. a factor 3-10 times larger than their densities at atmospheric pressure.  Under standard conditions, the shear moduli of such materials are comparable to their bulk moduli ($\\mu_{\\rm Fe}\\approx \\mu_{\\rm Si}\\approx 0.8\\,{\\rm Mbar}$). Under compression, the bulk moduli rise more quickly than the shear moduli because the former is associated with the degeneracy pressure of the electrons, whereas the latter is a Coulomb effect having to do with the ion lattice; for very large compression factors, one therefore expects $\\mu/K\\propto (\\rho/\\rho_0)^{1/3}$.  With these scalings, we can compare the elastic energy of the core under its distortion by an equilibrium tide with the corresponding gravitational energy.  Approximating the core by a constant density $\\rho_{\\rm c}$, we have \\begin{eqnarray} \\label{eq:distortions} \\Delta E_{\\rm grav}&=&\\frac{k_p}{(1+2k_p)^2}\\frac{G\\Mc^2}{\\Rc}\\sigma^2, \\quad \\Delta E_{\\rm elas}= \\frac{3}{4}\\frac{\\mu\\Mc}{\\rho_{\\rm c}}\\sigma^2,\\nonumber\\\\ \\frac{\\Delta E_{\\rm elas}}{\\Delta E_{\\rm grav}} &\\approx& \\frac{25\\mu\\Rc}{4G\\Mc\\rho_{\\rm c}} \\,, \\end{eqnarray} where $\\sigma$ is the radial strain $\\xi_{r,\\max}/\\Rc$ caused by the quadrupolar tide.  The elastic modulus in the core's compressed state should be $\\mu\\approx(\\rho/\\rho_0)^{1/3} \\mu_0\\approx 2\\mu_0$, where the subscript ``0'' refers to atmospheric conditions, hence no more than a few Mbar, whereas $G\\Mc\\rho_{\\rm c}/\\Rc\\approx 200\\,{\\rm Mbar}$ if $\\Mc=30\\,{\\rm M}_{\\oplus}$ and $\\rho_{\\rm c}=20\\,{\\rm g\\,cm^{-3}}$, which is of course comparable to the pressure (\\ref{eq:Pcore}).  So the elastic energy of the equilibrium tide in the core is only $5\\%$ of the gravitational energy, and therefore the large-scale tidal distortion of the core should be nearly the same as for a fluid.  However, elastic strength could still be enough to prevent the propagation of inertial waves inside the core, because the tidal period is long compared to the crossing time of an elastic wave in the core, which is of order half an hour for the numbers above.  The discussion so far has been based on a solid core, but a liquid one may be more likely.  Melting of an ionic lattice tends to occur at $\\Gamma\\equiv (Z_{\\rm eff}e)^2/r_i kT\\sim 10^2$, a dimensionless measure of the relative importance of Coulomb to thermal energies; here $r_i\\approx(3\\rho/4\\pi m_i)^{1/3}$ is the mean distance between ions in terms of their mass, $m_i$ \\citep[e.g.][and references therein]{Shapiro_Teukolsky83}.  The question is what to use for the effective charge $Z_{\\rm eff}$ governing the ionic interactions.  For silicon at its reference density of $2.33\\,{\\rm g\\,cm^{-3}}$, for example, this formula would predict a melting temperature $\\approx 2\\times 10^5(10^{-2}\\Gamma)^{-1}\\,{\\rm K}$ if one were to take $Z_{\\rm eff}=14$, the full charge on the nucleus.  In fact, most of that charge is shielded by electrons whose orbits are much smaller than $r_i$, so that $Z_{\\rm eff}\\approx 1$ is a more reasonable choice; this leads to $990\\Gamma_2^{-1}\\,{\\rm K}$, which is comparable to the actual value, $T_{\\rm m}=1693\\,{\\rm K}$.  As far as we know, there is no experimental measurement of the melt temperature near $44\\,{\\rm Mbar}$, but if one simply scales it from 1~bar by the the reciprocal of the inter-ionic distance, assuming that the Coulomb interaction is characterized by a constant $Z_{\\rm eff}$ throughout this range, then $T_{\\rm m}\\,@\\,20\\,{\\rm g\\,cm^{-3}}\\approx 3500\\,{\\rm K}$.  Despite the crudeness of this argument, it therefore seems likely that ``rock'' should be molten at the much higher temperatures in eq.~(\\ref{eq:Tcore}).  To recap, the core is probably fluid, and even if it is solid, its elastic strength will likely be unimportant for the equilibrium tide. However, the density contrast between the high-$Z$ core and its surroundings will affect its tidal distortion.  For simplicity, consider the equilibrium tide in a nonrotating ``planet'' composed of two incompressible fluids having different densities: $\\rho=\\rho_2$ in the core, $0\\le R<\\Rc$, and $\\rho=\\rho_1<\\rho_2$ outside it, $\\Rc<R<\\Rp$.  As usual, the perturbing tidal potential is quadrupolar, $\\Phi_{1,\\rm ext}= \\epsilon R^2P_2(\\cos\\theta)$. Since the vorticity vanishes except at the interface between the two fluids, the displacements can be assumed to be proportional to the gradient of a scalar, $\\vxi=\\grad\\chi$, where $\\chi$ satisfies Laplace's equation and is $\\propto P_2(\\cos\\theta)$; $\\chi$ is discontinuous at the interface but its radial derivative must be continuous. This idealized problem can then be worked out analytically by matching the radial parts of $\\chi$ and of $\\Phi_{1,\\rm self}$ across the interface. Included among these conditions is that the perturbed interface should remain an equipotential, \\begin{equation*} \\frac{\\partial\\chi}{\\partial R}= -\\frac{\\Phi_1}{G\\Mc/\\Rc^2}\\quad\\mbox{at}~R=\\Rc, \\end{equation*} which is analogous to the free boundary condition at the surface.  If $\\delta\\equiv\\rho_2/\\rho_1\\ge 1$ is the ratio of densities and $\\eta\\equiv\\Rc/\\Rp\\le 1$, it can then be shown that the radial strain in the core is related to the radial strain at the interface by \\begin{equation} \\label{eq:strainrat} \\left.\\frac{\\xi_R}{R}\\right|_{R<\\Rc} = \\frac{5+(\\delta-1)\\eta^3}{5\\delta-3(\\delta-1)(1-\\eta^5)} \\,\\left.\\frac{\\xi_R}{R}\\right|_{R=\\Rp}\\,. \\end{equation} In the limit $\\eta\\ll 1$ (recall that we expect $\\Rc\\lesssim 0.2\\Rp$), the factor on the righthand side reduces approximately to $5/(2\\delta+3)$.  Since $\\delta\\approx2\\mbox{-}4$, the core distorts substantially less than the main body of the planet, though it does not remain perfectly spherical.  We would therefore still expect the tide to generate short-wavelength inertial waves in a rotating planet with a core, but compared to the estimate (\\ref{eq:Qval}) made for a rigid core, the tidal $Q$ would increase by a factor $\\approx[1-5/(2\\delta+3)]^{-2}=[(2\\delta+3)/2(\\delta-1)]^2$, i.e. by one half to one order of magnitude."
    },
    "0812/0812.1672_arXiv.txt": {
        "abstract": "In this paper, we study the cosmological implications of the 100 square degree Weak Lensing survey (the CFHTLS-Wide, RCS, VIRMOS-DESCART and GaBoDS surveys). We combine these weak lensing data with the cosmic microwave background (CMB) measurements from the WMAP5, BOOMERanG, CBI, VSA, ACBAR, the SDSS LRG matter power spectrum and the Type Ia Supernoave (SNIa) data with the ``Union\" compilation (307 sample), using the Markov Chain Monte Carlo method to determine the cosmological parameters, such as the equation-of-state (EoS) of dark energy $w$, the density fluctuation amplitude $\\sigma_8$, the total neutrino mass $\\sum m_{\\nu}$ and the parameters associated with the power spectrum of the primordial fluctuations. Our results show that the $\\Lambda$CDM model remains a good fit to all of these data. In a flat universe, we obtain a tight limit on the constant EoS of dark energy, $w=-0.97\\pm 0.041$ ($1\\sigma$). For the dynamical dark energy model with time evolving EoS parameterized as $w_{\\DE}(a) = w_0 + w_a (1-a)$, we find that the best-fit values are $w_0=-1.064$ and $w_a=0.375$, implying the mildly preference of Quintom model whose EoS gets across the cosmological constant boundary during evolution. Regarding the total neutrino mass limit, we obtain the upper limit, $\\sum m_{\\nu}< 0.471$ eV ($95\\%$ C.L.) within the framework of the flat $\\Lambda$CDM model. Due to the obvious degeneracies between the neutrino mass and the EoS of dark energy model, this upper limit will be relaxed by a factor of $2$ in the framework of dynamical dark energy models. Assuming that the primordial fluctuations are adiabatic with a power law spectrum, within the $\\Lambda$CDM model, we find that the upper limit on the ratio of the tensor to scalar is $r<0.35$ ($95\\%$ C.L.) and the inflationary models with the slope $n_s\\geq1$ are excluded at more than $2~\\sigma$ confidence level. In this paper we pay particular attention to the contribution from the weak lensing data and find that the current weak lensing data do improve the constraints on matter density $\\Omega_m$, $\\sigma_8$, $\\sum{m_{\\nu}}$, and the EoS of dark energy. ",
        "introduction": "\\label{Int}  Cosmological observations play a crucial role in understanding our universe. With the accumulation of observational data from CMB measurements, large scale structure (LSS) surveys and SNIa observations as well as the improvements of the data quality, the accurate constraints on cosmological parameters can be expected from the data analysis. Recently WMAP group has released the five-year data of temperature and polarization power spectra \\cite{WMAP5GF1,WMAP5GF2,WMAP5Other}. And the Arcminute Cosmology Bolometer Array (ACBAR) experiment has also published its new CMB temperature power spectrum \\cite{ACBAR}. These new CMB data can strengthen the constraints on the cosmological parameters, especially for those associated with the inflationary models \\cite{KinneyWMAP5,WMAP5GF1,WMAP5GF2}. Furthermore, the Supernova Cosmology Project has made an unified analysis of the world's supernovae datasets and presented a new compilation ``Union\" (307 sample) \\cite{Union} which includes the recent samples of SNIa from SNLS and ESSENCE Survey, as well as some older datasets. In the literature \\cite{Union,LinderUnion,xia_union}, these data has been widely used to constrain various cosmological models. However, one should keep in mind that the degeneracies of cosmological parameters exist in almost all cosmological observations, \\emph{i.e.}, they are not sensitive to single parameters but to some specific combinations of them. It is therefore highly necessary to combine different probes to break parameter degeneracies so as to achieve tight constraints. Furthermore, different observations are affected by different systematic errors, and it is thus helpful to reduce potential biases by combining different probes.  Weak gravitational lensing which is directly related to the dark matter distribution, the geometry and the dynamics of the universe provides an useful way to break the degeneracies mentioned above. Cosmic shear, i.e. the distortion of images of the high-redshift source galaxies through the tidal gravitational field of the large-scale matter distribution in the universe,  provides a powerful probe of the mass fluctuations in the universe and gives a direct measurement of the matter power spectrum down to the non-linear regime. Since the mass distribution and its late time evolution are sensitive to the dark energy and the neutrino mass, we expect to get interesting information about both through the study of cosmic shear.  Recently, the cosmic shear analysis of the 100 square degree weak lensing survey which includes the Canada-France-Hawaii Telescope Legacy Survey (CFHTLSWide)\\cite{Hoekstra2006}, the Garching-Bonn Deep Survey (GaBoDs)\\cite{Hetterscheidt2006}, the Red-Sequence Cluster Survey (RCS)\\cite{Hoekstra2002a}, and the VIRMOS-DESCART survey (VIRMOS)\\cite{Waerbeke2005} has been presented by Ref.\\cite{Benjamin2007}. These combined surveys cover $113~{\\rm degree}^{2}$ which is the largest sky coverage so far. With these data, there are some studies on the cosmological implications recently in the literature \\cite{ichiki2008,Gongyan2008,Lesgourgues2007,Shaun2008,FuLiping2008,Fu2008,Kilbinger2008}. In this paper we study the constraints on cosmological parameters from the weak lensing cosmic shear as well as the current cosmological observational data, such as CMB, LSS and SN Ia. Specifically, we focus on the effects from the weak lensing.    Our paper is organized as follows: In Section II we describe the method and the latest observational datasets we use in this paper; Section III contains our main global fitting results on the cosmological parameters and the last section is the summary. ",
        "conclusions": "\\label{Sum}  In this paper we have studied the constraints on the cosmological parameters from the latest observational data including CMB, LSS, SN Ia and weak lensing effects. We have paid particular attention to the additional contributions from the 100 square degree cosmic shear data. Our results show that the current weak lensing data do help to improve the constraints on $\\Omega_m$, $\\sigma_8$, the total neutrino mass $\\Sigma{m_{\\nu}}$, and the constant EoS of dark energy. For other parameters, they do not add too much value due to their large statistical errors."
    },
    "0812/0812.1996_arXiv.txt": {
        "abstract": "Over the past one hundred years Brownian motion theory has contributed substantially to our understanding of various microscopic phenomena. Originally proposed as a phenomenological  paradigm for atomistic matter interactions, the theory has since evolved into a broad and vivid research area, with an ever increasing number of applications in biology, chemistry, finance, and physics. The mathematical description of stochastic processes has led to new approaches in other fields, culminating in the path integral formulation of modern quantum theory. Stimulated by experimental progress in high energy physics and astrophysics,  the  unification of relativistic and stochastic concepts has re-attracted considerable interest during the past decade. Focusing on the framework of special relativity, we review here recent progress in the phenomenological description of relativistic diffusion processes. After a brief historical overview, we will summarize basic concepts from the Langevin theory of nonrelativistic Brownian motions and discuss relevant aspects of relativistic equilibrium thermostatistics. The introductory parts are followed by a detailed discussion of  relativistic Langevin equations in phase space.  We address the choice of time parameters, discretization rules, relativistic fluctuation-dissipation theorems, and Lorentz transformations of stochastic differential equations. The general theory is illustrated through analytical and numerical results for the diffusion of free relativistic Brownian particles. Subsequently, we discuss how Langevin-type equations can be obtained as approximations to  microscopic models.  The final part of the article is dedicated to relativistic diffusion processes in Minkowski spacetime. Since the velocities of relativistic particles are bounded by the speed of light,  nontrivial relativistic Markov processes in spacetime do not exist; i.e., relativistic generalizations of the nonrelativistic diffusion equation and its Gaussian solutions must necessarily be non-Markovian. We compare different proposals that were made in the literature and discuss their respective benefits and drawbacks. The review concludes with a summary of open questions, which may serve as a starting point for future investigations and extensions of the theory. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.3993_arXiv.txt": {
        "abstract": "Stationary solutions to the equations of non-linear diffusive shock acceleration play a fundamental role in the theory of cosmic-ray acceleration. Their existence usually requires that a fraction of the accelerated particles be allowed to escape from the system. Because the scattering mean-free-path is thought to be an increasing function of energy, this condition is conventionally implemented as an upper cut-off in energy space --- particles are then permitted to escape from any part of the system, once their energy exceeds this limit. However, because accelerated particles are responsible for substantial amplification of the ambient magnetic field in a region upstream of the shock front, we examine an alternative approach in which particles escape over a spatial boundary. We use a simple iterative scheme that constructs stationary numerical solutions to the coupled kinetic and hydrodynamic equations.  For parameters appropriate for supernova remnants, we find stationary solutions with efficient acceleration when the escape boundary is placed at the point where growth and advection of strongly driven non-resonant waves are in balance. We also present the energy dependence of the distribution function close to the energy where it cuts off - a diagnostic that is in principle accessible to observation. ",
        "introduction": "\\label{Introduction}  It is widely thought that the diffusive acceleration of charged particles at shock fronts can be very efficient \\citep[for a review see][]{MalkovDrury}. The precise value of the efficiency is controlled by the microphysics of the injection of low-energy particles into the acceleration process, i.e., by processes operating on small length-scales. However, in the case of non-relativistic shocks, such as those encountered in supernova remnants, the bulk of the energy in nonthermal particles is carried off by those of the highest attainable energy, i.e., by particles that interact with relatively large length-scale structures.  Conventionally, the physics of this system is captured by combining a hydrodynamic description of the background plasma with the diffusion-advection equation obeyed by the distribution function of the accelerated particles. Analysis of the {\\em stationary} solutions of these equations is the foundation on which the study of the overall efficiency of the process and the maximum energy to which particles can be accelerated rests.  The importance of stationary solutions was realized early on by \\citet{DruryVolk81} who used a reduced, two-fluid description to analyze acceleration by plane shock fronts in a one-dimensional flow. Provided accelerated particles are not permitted to leave the system, the two-fluid description can be derived from the full description, including the diffusion-advection equation, using only two plausible assumptions about the particle distribution function. The more important of these is that the so-called {\\em effective diffusion coefficient} is always positive. \\citet{DruryVolk81} proved that this restricts the possible stationary flow patterns to those containing a {\\em precursor} in which the accelerated particles decelerate the incoming upstream flow, followed by a hydrodynamic shock front. This important result, which implies that only a finite number of stationary solutions exist for shocks of a given Mach number, was subsequently generalized to the relativistic case by \\citet{baringkirk91}.  However, the relevance of these studies is questionable if the diffusion coefficient is an increasing function of particle energy, as is the case, for example, if the transport is dominated by scattering off Alfv\\'en waves via the cyclotron resonance. One reason for this is that the timescale on which a stationary solution is approached becomes large at high particle energy. In this case, quasi-stationary solutions with a constant distribution function below a slowly evolving upper cut-off in energy can be expected to establish themselves, and have been found numerically by \\citet{Bell87} and by \\citet{FalleGiddings}. Another reason is that, as particles are accelerated to higher and higher energy, their mean-free path increases, and at some point becomes comparable to the size of any realistic system. Thus, even if one considers only strictly stationary solutions, the escape of high-energy particles appears to be an important property that will limit the maximum energy to which particles can be accelerated in any given system.  One way of accounting for particle escape is to truncate the distribution function above some finite value of the energy. But even with this simplification, finding stationary solutions of the combined kinetic and hydrodynamic equations is much more difficult than solving the two-fluid system. An important advance was the discovery of an approximate analytic solution of the diffusion-advection equation by \\cite{Malkov97}. This solution has a particle distribution function that vanishes above an upper cut-off, $\\pmax$, in the magnitude of the particle momentum. Wherever the hydrodynamic flow is compressed (usually throughout the precursor and at the shock front) there exists a flux of particles across this boundary in momentum space. These particles cease to contribute to the stress-energy tensor, and, therefore, effectively escape from the system. A similar distribution was adopted by \\citet{achterberg87} when investigating numerical solutions to this equation. Under this assumption it is possible to find approximate analytic solutions, and relatively straightforward to construct numerical solutions to the full set of equations \\citep{achterberg87,MDV00,Blasi02,AB05}.  Two problems intrinsic to this approach are immediately obvious. Firstly, the momentum dependence of the distribution close to $\\pmax$ --- a diagnostic that is, in principle, accessible to observation --- is not well-approximated. Secondly, the exclusion of a post-cursor, although made plausible by the two-fluid approach, is not always justified. The conditions that must be fulfilled by the distribution function and the momentum-dependent diffusion coefficient for this assumption to be valid \\citep{kirk90} can in principle be checked {\\em a posteriori}, but this is not straightforward for discontinuous distributions.  These problems arise because an upper cut-off in momentum is used to describe the physics of particle escape. Recently, however, observational evidence has been accumulating suggesting that cosmic rays are responsible for substantial amplification of the ambient magnetic field in the precursors of shock fronts in supernova remnants \\citep{Hwang,VinkLam,Bamba,Uchiyama}. This implies that the scattering mean-free-path is not only a function of energy, but also depends strongly on position with respect to the shock front. It, therefore, highlights a more serious short-coming of the approach that uses a cut-off in momentum space: The amplification of the field is likely to be connected with the spatial flux of energetic particles, which is artificially distorted if particles are assumed to vanish across a momentum boundary. In this paper we examine stationary solutions with escape through a spatial boundary instead of through a boundary in energy space.  Spatial boundaries have previously been implemented in Monte-Carlo simulations of the non-linear acceleration problem. In particular, \\citet{VEB06} used a model equation to describe Alfv\\'enic turbulence and coupled it to a Monte-Carlo simulation that permitted escape over a spatial boundary. In this way they were able to investigate the effects of an enhanced resonant interaction between the turbulent waves and the accelerated particles, although the position of the boundary itself remained arbitrary.  Recently, \\cite{ZirPtusk08} used MHD simulations to describe the excited turbulence and find the diffusion coefficient of the highest energy accelerated particles. They allowed these particles to escape over a spatial boundary, but used the test-particle approximation in which the flow speed is unaffected by the particles. In discussing shock fronts in supernova remnants, they suggested that the boundary leads the shock front by a distance given roughly by the radius of the remnant.  Our approach differs from these, not only in the numerical method used to solve the acceleration problem, but also in the input physics. Resonant interactions between energetic particles and Alfv\\'en waves were long thought to be responsible for coupling these particles to the background plasma \\citep{bell78,voelkmckenzie,achterberg83,belllucek00}. However, well ahead of the shock front, non-resonant processes are more strongly driven and can be expected to dominate under the conditions present in supernova remnants \\citep{bell04,Pelletier,bykovtoptyghin,revilleetal07,ZPV08}. In the linear phase, these instabilities inject short-wavelength turbulence into the plasma, resulting in a relatively large diffusion coefficient that is proportional to the square of the particle momentum \\citep{ZirPtusk08}. However, the non-linear evolution includes not only a cascade of energy to smaller length scales, but also the appearance of large-scale structures such as cavities \\citep{bell05}. In this regime, the mean-free-path is reduced, and its $p^2$ dependence eliminated, as can be seen from large-scale numerical simulations of the transport properties \\citep{revilleetal08}. These simulations have not yet advanced to the stage where they can provide a model diffusion coefficient over a wide dynamic range of momentum. Consequently, we model this effect by assuming the diffusion to be of Bohm type in a background field that is amplified to the value at which the non-resonant instabilities are expected to saturate. Although this is a relatively crude approach, it enables us to solve the non-linear problem of finding stationary solutions. We are then able to check the location of the spatial boundary, which should be located where the growth rate of the non-resonant modes is approximately equal to the speed of advance of the shock divided by the distance from the shock front. Since the escaping flux is dominated by the highest energy particles, we do not expect that changing the form of the diffusion coefficient will alter significantly our conclusions, provided it remains an increasing function of momentum.  The paper is set out as follows: In Section~\\ref{basic} we set up the advection-diffusion equation and the hydrodynamic equations governing the system of accelerated particles and background plasma. The two ways of allowing for particle escape are discussed in Section~\\ref{escape}. In Section~\\ref{codedescription} we describe the iteration and finite difference methods used to find stationary solutions of the combined advection-diffusion and hydrodynamic equations using the two different boundary conditions, and in Section~\\ref{comparison} we compare and discuss the results. A discussion of the self-consistency of the position of the spatial boundary and a summary of our conclusions is presented in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  The assumption underlying most previous investigations of non-linear diffusive shock acceleration is the non-existence of waves able to scatter particles with energy above an upper cut-off. In this paper, we examine an alternative picture, in which the current generated by the streaming cosmic rays falls below some critical value at a distance $L_{\\rm esc}$ upstream of the subshock, and that beyond this the turbulence is insufficient to scatter the particles.  In terms of the non-linear response of the system to changes in the injection parameter, we find the two pictures are quite similar (see Fig.~\\ref{fig1}). One difference is that our spatial boundary method can be used to model the shape of the distribution close to the cut-off (see Fig.~\\ref{fig3} and Fig.~\\ref{fig4}). This is an advantage because it potentially enables one to model the radiative signatures of the acceleration process.  The main difference, however, is that we are able to address the physics that determines the location of the boundary. Previous work that implemented such a boundary \\cite{VEB06} did not constrain its location. In time-dependent models of acceleration in supernova remnants \\cite{berezhkoetal97}, an effective spatial boundary is imposed by the spherical geometry. But they assume a value for the amplified magnetic field in the entire computational box, without considering whether or not this is consistent with the location of the boundary.  We argue that the location of the boundary is determined by the growth rate of the instability responsible for field amplification. It has recently been shown that, in the case of efficient shock acceleration, a short-wavelength non-resonant mode \\cite{bell04} plays a crucial role. In a series of papers, \\cite{Pelletier,Marcowith} argue that this non-resonant mode dominates far from the shock, while the resonant streaming instability takes over closer to the shock, driving the diffusion towards Bohm-type in the precursor. According to these arguments, the position of the free-escape boundary should be fixed by the properties of the non-resonant mode, once the shock is modified by the cosmic-ray pressure. We show explicitly that the non-resonant modes are strongly driven at the escape boundary, and compare the local growth rate with the rate at which the waves are advected towards the shock front (see Fig.~\\ref{fig6}). In this picture, particles stream freely ahead (upstream) of the boundary, whereas behind (downstream of) it, the turbulence is assumed to be fully developed and particles undergo Bohm diffusion in an amplified magnetic field.   \\cite{ZirPtusk08} also suggest that the geometry of a supernova shock front sets the length scale for the precursor, and thus determines the maximum particle energy. However, we find that when the non-linear dynamics of the acceleration process are included, the length scale is set by the strength of the amplified magnetic field and the efficiency of the injection process. Ultimately, it is the microphysics of injection --- represented in Figs.~\\ref{fig1} and \\ref{fig6} by the parameter $\\nu$ --- that determines not only the efficiency of the acceleration process, but also the maximum attainable particle energy."
    },
    "0812/0812.3598_arXiv.txt": {
        "abstract": "We present high resolution 240 and 607 MHz GMRT radio observations, complemented with 74 MHz archival VLA radio observations of the Ophiuchus cluster of galaxies, whose radio mini-halo has been recently detected at 1400 MHz. We also present archival \\emph{Chandra} and \\emph{XMM-Newton} data of the Ophiuchus cluster.  Our observations do not show significant radio emission from the mini-halo, hence we present upper limits to the integrated, diffuse non-thermal radio emission of the core of the Ophiuchus cluster.  The \\emph{XMM-Newton} observations can be well explained by a two-temperature thermal model with temperatures of $\\simeq$1.8 keV and $\\simeq$9.0 keV, respectively, which confirms previous results that suggest that the innermost central region of the Ophiuchus cluster is a cooling core.  This result is consistent with the occurrence of a mini-halo, as expected to be found in hot clusters with cool cores.  We also used the \\emph{XMM-Newton} data to set up an upper limit to the (non-thermal) X-ray emission from the cluster.  We also emphasize that the non-thermal X-ray emission obtained with \\emph{XMM-Newton} and \\emph{INTEGRAL} cannot be produced by the putative AGN of the galaxy at the cluster center.  The combination of available radio and X-ray data has strong implications for the currently proposed models of the spectral energy distribution (SED) from the Ophiuchus cluster. In particular, a synchrotron+IC model is in agreement with the currently available data, if the average magnetic field is in the range (0.02$-$0.3)$\\mu$G. A pure WIMP annihilation scenario can in principle reproduce both radio and X-ray emission, but at the expense of postulating very large boost factors from dark matter substructures, jointly with extremely low values of the average magnetic field. Finally, a scenario where synchrotron and inverse Compton emission arise from PeV electron-positron pairs (via interactions with the CMB), can be ruled out, as it predicts a non-thermal soft X-ray emission that largely exceeds the thermal Bremsstrahlung measured by \\emph{INTEGRAL}. ",
        "introduction": "\\label{intro}  Clusters of galaxies are the largest bound structures of the Universe and, according to hierarchical scenarios of structure formation, the latest ones to form. They are filled by a hot ($10^7-10^8$ K) plasma, called intra-cluster medium (ICM), and thus radiate in soft X-ray bands through thermal Bremsstrahlung.  In hierarchical cosmological scenarios, the most massive clusters form as a result of merging of smaller clusters. The current scenario predict that a fraction of the energy dissipated during the merging is channeled into particles acceleration via shocks and turbulence (e.g. Sarazin 1999a; Blasi, Gabici \\& Brunetti 2007).  This lead to a complex population of primary protons and electrons in the ICM.  Models predict that a large population of \\emph{secondary electrons} could be injected by collisions between relativistic and thermal protons in the ICM. Alternatively, the ICM relativistic electrons population could be \\emph{re-accelerated in situ} by various mechanisms associated with turbulence in massive merger events. The recent observation of a radio halo with a very steep synchrotron spectrum in Abell~521 by Brunetti et al. (2008) supports the turbulence acceleration mechanism.  This electron population is expected to be responsible for both the synchrotron radio emission and the hard X-ray emission. The origin of the latter emission is disputed. It has been proposed that it is due to inverse-Compton (IC) scattering of relativistic electrons with the cosmic microwave background (CMB) (Atoyan \\& V\\\"olk 2000, Ensslin \\& Biermann 1998, Sarazin 1999b), or to a population of PeV electrons that would radiate in hard X-rays through synchrotron emission (Timokhin, Aharonian \\& Neronov 2004, Inoue, Aharonian \\& Sugiyama 2005). Those latter models have the disadvantage that do not explain the well-known synchrotron radio emission.  Extended radio emission in clusters of galaxies has been known to exist for a long time (see e.g. Feretti \\& Giovannini 2007). However, the detection of the putative hard X-ray emission has been rather elusive, apart from some rather weak and controversial detections of a hard tail in the X-ray spectrum of Coma (Fusco-Femiano et al. 2004) and Abell 2256 (Fusco-Femiano et al. 2005) with \\emph{BeppoSAX}, and the \\emph{Chandra} observations of the Perseus cluster (Sanders et al. 2004, Sanders, Fabian \\& Dunn 2005). Nevertheless, \\emph{INTEGRAL} failed to find an X-ray non-thermal component in the Coma cluster (Renaud et al. 2006), and recent \\emph{XMM-Newton} observations of Perseus (Molendi \\& Gastaldello 2008) did not confirm Sanders et al. (2004) and Sanders, Fabian \\& Dunn (2005) results.  The Ophiuchus cluster is a nearby (z=0.028, Johnston et al. 1981) rich cluster located above the direction of the Galactic Center ($l=0.5^\\circ, b=9.4^\\circ$), with a very high plasma temperature ($kT\\sim10$ keV).  Johnston et al. (1981) claimed, based on radio data available at the time, that the cluster was associated with the steep-spectrum radio source MSH 17-203 (also dubbed Cul 1709-231). The identification of this radio source as a radio halo from the Ophiuchus cluster would thus imply the presence of relativistic electrons, and hence the presence of a non-thermal, high-energy tail in the X-ray band would be expected. And, in fact, while this paper was being refereed, Govoni et al. (2009) and Murgia et al. (2009) have reported the detection of a radio mini-halo  in the Ophiuchus cluster (see Fig. 6 in Govoni et al. 2009), using deep VLA 1400 MHz observations.   \\begin{figure} \\begin{center} \\resizebox{8.6 cm}{!}{\\includegraphics[width=8.5cm]{ophiuchus-nvss-vlss.eps}} \\caption[]{VLSS radio image at 74~MHz (top) and NVSS radio image at 1400~MHz (bottom) of the Ophiuchus Cluster. In both panels the positions of the \\emph{INTEGRAL} X-ray detection and of the nearest X-ray source 1RXS J171209.5-231005 are shown.  The synthesized beams (FWHM) at 74 MHz and 1400 MHz were (80''$\\times$80'') and (45''$\\times$45''), respectively, and are drawn at the bottom left of each panel.  The big circles in the panels correspond to the approximate size of the synthesized beam of the FLEURS and SYDNEY (big circle in top panel) and of the OSU (big circle in the bottom panel) radio observations (see Table \\ref{all_data}), carried out in the 70's. Note that the radio emission of a large amount of different sources enters the beam, and therefore heavily overestimates the flux at those frequencies. The off-source rms noise in the VLSS and NVSS images is of 220 mJy/b and 0.4 mJy/b, respectively, and contours are drawn starting at three times the quoted rms, increasing in steps of $\\sqrt{3}$.} \\label{fig,vla} \\end{center} \\end{figure}  Watanabe et al. (2001), based on \\emph{ASCA} data, found that the core of the cluster has an angular diameter of 6.4'.  Measurements of the ICM temperature vary from 8.5$\\pm$0.5~keV (INTEGRAL; Eckert et al. 2008) up to 9.5$^{+1.4}_{-1.1}$~keV (Swift/BAT; Ajello et al. 2008).  Watanabe et al. also found a large (20'$\\times$30'), hot ($kT>13$ keV) region, 20' west of the cluster center, from which they concluded that the cluster is not dynamically relaxed, and suggested it experienced a major merging event in the recent past ($t \\lsim 1$ Gyr).  Eckert et al. (2008) have recently reported a tentatively resolved ($\\sim$5') X-ray source at the cluster center, and claimed the presence of a non-thermal tail. Eckert et al. interpret the non-thermal hard X-ray emission as due to IC radiation from relativistic electrons scattered off the CMB in the cluster medium. The evidence of a non-thermal hard X-ray tail in the Ophiuchus cluster from \\emph{INTEGRAL} could be the first significant detection of such non-thermal X-ray emission from a galaxy cluster.  More recent Suzaku observations of the Ophiuchus cluster by Fujita et al. (2008) have, however, failed to detect the non-thermal component claimed by Eckert et al. (2008), although their quoted upper limit of $2.8\\times10^{-11}$~erg~cm$^{-2}$~s$^{-1}$ in the $20-60$~keV energy band is still compatible with the \\emph{INTEGRAL} detection.  Fujita et al. (2008) found that the cluster was hot, except for the innermost $\\sim$50 kpc, and that the intra-cluster medium (ICM) was in ionization equilibrium state. From those results, Fujita et al. conclude that Ophiuchus cluster is not a major merger, but one of the hottest clusters with a cool core.   Even more recently, Ajello et al. (2008) have found, combining Swift/BAT and \\emph{Chandra} spectra, an upper limit on the Ophiuchus non-thermal X-ray emission in the $20-60$~keV band, of $7.2\\times10^{-12}$~erg~cm$^{-2}$~s$^{-1}$, which is below the \\emph{INTEGRAL} detection.  Nevertheless, the claimed \\emph{INTEGRAL} detection, the Suzaku and the combined Swift/BAT and \\emph{Chandra} upper limits are consistent.  In this Paper, we present and discuss archival and new radio data for the Ophiuchus cluster, along with archival X-ray data.  Together with the published X-ray data of the cluster, we model all the data to shed light on the origin of the diffuse non-thermal emission from this cluster of galaxies.  In Section 2, we present the VLA and GMRT radio data from the Ophiuchus cluster, along with archival \\emph{Chandra} and \\emph{XMM-Newton} data. In Section 3, we model the SED of the Ophiuchus cluster using several different scenarios.  Finally, we summarize our main results and their implications in Section 4.  All sky positions in the paper are for the J2000.0 equinox. We assume cosmological parameters of $\\Omega_{0}=0.3$, $\\lambda_{0}=0.7$ and $H_{0}=70$~km~s$^{-1}$~Mpc$^{-1}$. For these parameters, our assumed redshift of 0.028 to the Ophiuchus cluster (Johnston et al. 1981) corresponds to a luminosity distance of 122.6~Mpc and a comoving distance of 126.03~Mpc. At this distance, 1'' corresponds to 560 pc. ",
        "conclusions": "\\label{summary}  We present high resolution 240 and 607 MHz GMRT radio observations, complemented with 74 MHz archival VLA radio observations of the Ophiuchus cluster of galaxies, whose mini-radio halo has been recently detected at 1400 MHz. We also present archival \\emph{Chandra} and \\emph{XMM-Newton} data of the Ophiuchus cluster.  Our observations do not show significant radio emission from the mini-halo, hence we present upper limits to the integrated, diffuse non-thermal radio emission of the core of the Ophiuchus cluster.   The \\emph{XMM-Newton} observations can be well explained by a two-temperature thermal model with temperatures of $\\simeq$1.8 keV and $\\simeq$9.0 keV, respectively. This result is in agreement with previous ones by Fujita et al. (2008), which suggested that the central region of the Ophiuchus cluster is a cooling core.  In turn, this result is consistent with the occurrence of a mini-halo, as expected to be found in hot clusters with cool cores, and therefore support the conclusion by Fujita et al. (2008) that Ophiuchus is not the result of a major merger.  We also used the \\emph{XMM-Newton} data to set up an upper limit to the (non-thermal) X-ray emission from the cluster and therefore constrain the SED of the Ophiuchus cluster.  The high spatial resolution of the radio and X-ray data have allowed us to study separately the diffuse emission of the cluster core from that due to the population of point sources. In particular, we have found a relatively bright radio and X-ray (\\emph{Chandra}) point source at the center of the cluster (source A in our images), which is spatially coincident with the \\emph{INTEGRAL} hard X-ray source. The flat radio spectrum of this point source and the lack of non-thermal emission has allowed us to associate it to a low-level AGN activity, most likely associated to the cD central galaxy 2MASX J17122774-2322108. This central source is very unlikely to be responsible for the hard X-ray emission detected by \\emph{INTEGRAL}, in spite of its apparent spatial coincidence.  Neither our newly obtained radio data at 240 and 607 MHz, nor the archival VLA data at 74 MHz for the Ophiuchus cluster of galaxies detect significant diffuse radio emission from the mini-halo in Ophiuchus, which has been recently at 1400 MHz.  The radio and X-ray non-thermal data have dramatic implications for currently proposed scenarios to model the spectral energy distribution (SED) of the Ophiuchus cluster of galaxies.  In particular, a synchrotron+IC model is in agreement with the currently available data, if the average magnetic field is between $\\sim$0.02$\\mu$G and $\\sim$0.3$\\mu$G.  A pure dark matter scenario can also, in principle, reproduce both radio and X-ray emission, but very large boost factors from dark matter substructures and extremely low values of the average magnetic field are needed.  We notice that for a magnetic field value comparable with the expectations from Faraday rotation measures ($\\sim 1 \\mu$G), GLAST-Fermi and high-resolution radio searches for a diffuse signal from galaxy clusters will have a comparable and complementary reach in the particle dark matter parameter space.  Finally, we also explored the possibility of synchrotron and inverse Compton emission arising from PeV electron-positron pairs (via interactions with the CMB), as described in Inoue, Aharonian \\& Sugiyama (2005). Those models can be ruled out, since they predict a non-thermal soft X-ray emission that largely exceeds the thermal Bremsstrahlung measured by \\emph{INTEGRAL} and the non-thermal X-ray emission inferred from the \\emph{XMM}-Newton data.  Given that Ophiuchus is one of the closest galaxy clusters, it can be studied at high resolution and with high sensitivity. Therefore, its study at all wavelengths is encouraged, since it can result in much better constraints of its SED, and will allow for tests on acceleration scenarios, as we have discussed in this paper. In particular, optical spectroscopy is advisable, as it would allow to confirm, or rule out, the membership of some of the sources shown in the radio images presented in this work. Gamma-ray observations would also be very useful to constrain the SED and the particle acceleration mechanisms."
    },
    "0812/0812.3621_arXiv.txt": {
        "abstract": "We consider a sample of type-I active galactic nuclei (AGN) that were observed by {\\it Chandra/HETG} and resulted in high signal-to-noise grating spectra, which we study in detail. All objects show signatures for very high ionization outflows. Using a novel scheme to model the physics and spectral signatures of gaseous winds from these objects, we are able to estimate the mass loss rates and kinetic luminosities associated with the highly ionized gas and investigate its physical properties. Our conclusions are as follows: 1) There is a strong indication that the outflowing gas in those objects is multi-phase with similar kinematics for the different phases. 2) The X-ray spectrum is consistent with such flows being thermally driven from $\\sim$\\,pc scales, and are therefore unlikely to be associated with the inner accretion disk. 3) The underlying X-ray spectrum consists of a hard X-ray powerlaw which is similar for all objects shining below their Eddington rate and a soft excess whose contribution becomes more prominent for objects shining close to their Eddington limit. 4) The physical properties of the outflow are similar in all cases and a coherent picture emerges concerning its physical properties. 5) The deduced mass loss rates are, roughly, of the order of the mass accretion rate in those objects so that the kinetic luminosity carried by such winds is only a tiny fraction ($\\ll$1\\%) of the bolometric luminosity. We discuss the implications of our results for AGN structure and AGN interaction with the environment. ",
        "introduction": "Recent {\\it Chandra} and {\\it XMM} grating observations of type-I (broad emission lines) active galactic nuclei (AGN) show the presence of highly ionized gas (HIG) outflowing from the centers of many objects (e.g., Kaastra et al. 2000, Kaspi et al. 2001, Pounds et al. 2003, McKernan et al. 2007). This suggests that HIG may play an important role in AGN physics and so must be included in future unification schemes (e.g., Elvis 2000). Currently, however, little is known with confidence about the physical properties of this component and a clear physical picture is yet to emerge. In addition to being important for AGN study, HIG flows carry energy and momentum out from their inner engine to potentially large scales and, as such, may have a profound impact on the environment of quasars, be it their host galaxies or even the inter-galactic medium (e.g., Di Matteo et al. 2005, Scannapieco \\& Oh 2003). Thus, constraining the physical properties of AGN outflows is a fundamental question in AGN physics with great implications for a wide range of astrophysical problems.  Detailed X-ray spectroscopic observations of a few nearby low-luminosity AGN (Seyfert 1 galaxies) indicate that HIG outflows have a stratified ionization structure and that they are located within a parsec or so from the central continuum source (e.g., Blustin et al. 2007,  Kaspi \\& Behar 2006, Kraemer et al. 2005, Krongold et al. 2003, 2009, Netzer et al. 2003, Schulz et al. 2008, Smith et al. 2007, Steenbrugge et al. 2005) . The outflowing gas appears to be multiphase with various phases being in a rough pressure equilibrium and lying on the thermally stable parts of the heating-cooling curve (e.g., Chakravorty et al. 2008, Holczer et al. 2007, Netzer et al. 2003). Despite the considerable progress in our phenomenological understanding of HIG flows, their underlying physics remains elusive. For example, a detailed investigation of the best studied HIG in AGN (the case of NGC\\,3783) yielded  mass loss rates in the range of a few percent to nearly a hundred solar masses per year (e.g., Behar et al. 2003, Blustin et al. 2005, Netzer et al. 2003). Some studies pre-suppose that the mass loss rate is less than or equal to the mass accretion rate (e.g., Steenbrugge et al. 2005) yet it is not clear that this must be the case. Clearly, the current situation is unsatisfactory if we wish to understand the role of HIG flows in the global picture of AGN and assess their effect on the environment of such objects.  Perhaps the main limitation that prevents us from fully understanding the HIG phenomenon in AGN results from model incompleteness: for example, some works dealing with HIG absorption measure the column densities from individual transitions and attempt to estimate the mass loss rate by requiring full volume filling gas, which does not seem to be supported by recent observational and theoretical work; e.g., Arav, Li, \\& Begelman 1994, Kraemer et al. 2005, Proga et al. 2008). Other works make use of detailed photoionization calculations and spectral modeling to determine the temperature and column density range of the HIG (e.g., Kaspi et al. 2002,  Krongold et al. 2003, Netzer et al. 2003, R{\\'o}{\\.z}a{\\'n}ska et al. 2006). Nevertheless, pure photoionization modelling does not account for the kinematics of the outflow and cannot be used to reveal the location of the gas. This is due to an inherent degeneracy between the density of the photoionized gas and the distance from the ionizing source for a wide range of densities relevant to HIG outflows (see however R{\\'o}{\\.z}a{\\'n}ska et al. 2008). Limits on the location of the gas may be obtained by studying the reaction of the photo-ionized gas to the varying ionizing flux level of the source, as has been done for the case of NGC\\,3783 (Netzer et al. 2003, Krongold et al. 2005; see also Chevallier et al. 2007). New numerical schemes that self-consistently model the gas dynamics as well as its thermal state were recently calculated by Dorodnitsyn et al. (2008a,b) and their spectral predictions are yet to be compared to observations. Thus, the physics of HIG outflows in AGN and, in particular, the estimates for the mass carried by them are subject to considerable uncertainties.  Recently, Chelouche \\& Netzer (2005; hereafter CN05) have demonstrated that, by employing a detailed and self-consistent modelling of the dynamics and photoionization properties of the outflowing gas, it may be possible to constrain the mass loss rate to much better precision than before. By applying their model to NGC\\,3783, they found that the mass loss rate is considerably smaller compared to previous estimates. While their model seems to be consistent with most observational constraints for the case of NGC\\,3783, it is yet to be confirmed for other AGN. Here we wish to adopt a more general approach and study a sample of X-ray bright, nearby type-I AGN and investigate their properties using a consistent analysis method. Specifically, we wish to test the thermally driven wind model as a possible new ingredient of the AGN structure and do that by considering several objects where such advanced analysis is warranted. Our aim is to gain better physical understanding of the AGN phenomenon as is emerging from high resolution and S/N observations and understand the dynamics of highly ionized gas in those systems. In addition, we wish to quantify the mass loss rate from these objects which have been claimed to have a paramounting effect on the environment of such objects.  In this paper we consider a sample of X-ray bright, nearby type-I AGN (Seyfert 1-1.5 galaxies) for which high quality grating spectra are available and the mass of the central black-hole has been measured. We improve upon the model of CN05 and apply it to study HIG outflows in those objects. The paper is organized as follows: In \\S 2 we present the sample of objects used in this work and discuss their properties. Section 3 outlines the basics of our model and the improvements and generalizations made to the CN05 scheme. Results of our spectral modelling pertaining to individual objects are presented in \\S 4. We discuss the implications of our results for the unification scheme of AGN and AGN interaction with their environment in \\S 5. Summary follows in \\S 6. ",
        "conclusions": "We present the first systematic study of  type-I AGN with high resolution spectroscopic X-ray observations. Our sample consists of five objects for which good signal-to-noise  {\\it Chandra/HETG} data are available. We report on our findings concerning the physical properties of highly ionized, low velocity gas seen to be outflowing from those systems.  Using a novel scheme for the calculation of the spectral features from thermally and radiatively driven winds we are able to constrain, for the first time, physical properties of the ejected gas pertaining to its location, the mass loss rate, and the associated kinetic luminosity. In addition, the number of free parameters in our model is typically much smaller than that used by pure photoionization models. We find that all AGN in our sample possess highly ionized outflows whose mass loss rate is, on average, of order of the mass accretion (but with a scatter of order a few around this value) while the kinetic luminosity is only a small fraction (typically $\\ll 1\\%$)  of the bolometric luminosity. If this scaling persists to highly luminous quasars then low velocity flows are unlikely to lead to substantial feedback effects. Our results indicate that such flows show an interesting physical similarity among outflows in different objects suggesting that the gas microphysics in those systems is essentially the same. Different manifestation of such flows can be attributed, at least in part, to geometrical effects. Furthermore, we find that these flows are launched from parsec scales  and are unlikely to be associated with the inner parts of the accretion disk. Our study demonstrates the importance of high resolution and S/N X-ray grating observations  for understanding the physics of AGN flows."
    },
    "0812/0812.1909_arXiv.txt": {
        "abstract": "We obtained self-similar solutions of relativistically expanding magnetic loops taking into account the azimuthal magnetic fields. We neglect stellar rotation and assume axisymmetry and a purely radial flow. As the magnetic loops expand, the initial dipole magnetic field is stretched into the radial direction. When the expansion speed approaches the light speed, the displacement current reduces the toroidal current and modifies the distribution of the plasma lifted up from the central star. Since these self-similar solutions describe the free expansion of the magnetic loops, i.e., $Dv/Dt=0$, the equations of motion are similar to those of the static relativistic magnetohydrodynamics. This allows us to estimate the total energy stored in the magnetic loops by applying the virial theorem. This energy is comparable to that of the giant flares observed in magnetars. ",
        "introduction": "Soft gamma-ray repeaters (SGRs) are believed to be a young neutron star with strong magnetic fields ($\\sim 10^{15}~\\mathrm{G}$), namely the magnetar (see, e.g., \\citealt{2006csxs.book..547W}; \\citealt{2008A&ARv..15..225M}, for review). The magnetic fields inside the magnetar are amplified by the dynamo mechanism at the birth of the neutron star. The Lorentz force stressing the crust of the magnetar balances with the rigidity of the crust. When the critical twist is accumulated, the magnetic twist injected into the magnetar magnetosphere will trigger the expansion of magnetic loops \\citep{2001ApJ...561..980T}. The magnetic reconnection taking place inside the expanding magnetic loops can be responsible for SGR flares \\citep{2001ApJ...552..748W, 2006MNRAS.367.1594L}.  Recently, relativistic simulations have been performed to study the dynamics of the magnetospheres of neutron stars (\\citealt{2002MNRAS.336..759K}; \\citealt{2005PASJ...57..409A}; \\citealt{2006ApJ...648L..51S}; \\citealt{2006MNRAS.367...19K}). \\citet{2005KITP...CONF..HP} reported the results of 2-dimensional relativistic force-free simulations of the magnetar flares triggered by the injection of the magnetic twists at the footpoints of the loops. When the critical twist is accumulated, the magnetic loops expand relativistically. \\citet{2007Chiba...1..1} carried out 2-dimensional relativistic force-free simulations of expanding magnetic loops and showed that the Lorentz factor defined by the drift velocity $\\bmath v_d = c(\\bmath E \\times \\bmath B)/B^2$ exceeds 10 \\citep[see,][for the definition of the drift velocity]{1997PhRvE..56.2181U}. These simulations indicate that the magnetic loops expand self-similarly.  Assuming relativistic force-free dynamics, \\citet{2003astro.ph.12347L} obtained self-similar solutions of the spherically expanding magnetic shell. \\citet{2005MNRAS.359..725P} found self-similar solutions of the relativistic force-free field. In these studies of force-free dynamics, gas pressure and inertial terms are neglected. In the framework of the relativistic magnetohydrodynamics (MHD), \\citet{2002PhFl...14..963L} found self-similar solutions of the spherically expanding magnetic shells. \\cite{1982ApJ...261..351L} obtained non-relativistic self-similar MHD solutions of the expanding magnetic loops in solar flares or supernovae explosion by assuming axisymmetry. Subsequently, \\citet{1984ApJ...281..392L} extended his model to the case including toroidal magnetic fields and applied it to solar coronal mass ejections (CMEs). The latter model was employed by \\cite{1992ApJ...388..415S} as a test problem to check the validity and accuracy of axisymmetric MHD codes. In magnetar flares, the magnetic loops may be twisted by the shear motion at the footpoints of the loops. The shear motion generates Alfv\\'en waves propagating along the field lines. Such twisted magnetic loops expand by the enhanced magnetic pressure by the toroidal magnetic fields. Thus we should include the toroidal magnetic field to study the evolution of magnetic loops during magnetar flares. Also the relativistic effects should be included. The characteristic wave speed in the magnetar magnetosphere approaches the light speed because of the strong magnetic fields. Thus our aim is to obtain relativistic self-similar MHD solutions of expanding magnetic loops taking into account the toroidal magnetic fields by extending the non-relativistic solutions found by \\cite{1982ApJ...261..351L}.  This paper is organized as follows; in $\\S$ \\ref{ssmhd}, we present the relativistic ideal MHD equations and introduce a self-similar parameter which depends on both radial distance from the centre of the star and time. In $\\S$ \\ref{sss}, we obtain self-similar solutions. The physical properties of these solutions are discussed in $\\S$ \\ref{physprop}.  We summarize the results in $\\S$ \\ref{summary}.  \\section[]{Self-similar MHD Equations}\\label{ssmhd} In the following, we take the light speed as unity. The complete set of relativistic ideal MHD equations is \\begin{equation} \\frac{\\partial }{\\partial t}(\\gamma \\rho)+\\nabla\\cdot(\\gamma\\rho \\bmath v)=0, \\label{eq:MHDcont} \\end{equation} \\begin{equation} \\rho\\gamma\\left[\\frac{\\partial}{\\partial t}+(\\bmath v \\cdot \\nabla) \\right](\\xi\\gamma\\bmath v)= -\\nabla p +\\rho_e \\bmath E+\\bmath j \\times \\bmath B-\\frac{G M \\rho \\gamma }{r^2} \\bmath e_r,\\label{eq:MHDeom} \\end{equation} \\begin{equation} \\left[\\frac{\\partial}{\\partial t}+(\\bmath v \\cdot \\nabla)\\right]\\left(\\ln\\frac{p}{\\rho^\\Gamma}\\right)=0, \\label{eq:MHDentropy} \\end{equation} \\begin{equation} \\nabla \\cdot \\bmath E=4\\pi \\rho_e, \\label{eq:Gauss} \\end{equation} \\begin{equation} \\nabla \\cdot \\bmath B = 0, \\label{eq:nomonopole} \\end{equation} \\begin{equation} \\frac{\\partial \\bmath B}{\\partial t}+\\nabla \\times \\bmath E=0,\\label{eq:Faraday} \\end{equation} \\begin{equation} \\frac{\\partial \\bmath E}{\\partial t}=\\nabla \\times \\bmath B-4\\pi\\bmath j,\\label{eq:Ampere} \\end{equation} \\begin{equation} \\bmath E=-\\bmath v\\times \\bmath B,\\label{eq:mhdcond} \\end{equation} where $\\bmath E, \\bmath B,\\bmath j, \\bmath v,\\gamma,\\rho_e, \\rho, p,\\Gamma $ are the electric field, the magnetic field, the current density, the velocity, the Lorentz factor, the charge density, the mass density, the pressure and the specific heat ratio, respectively. The vector $\\bmath e_r$ is a unit vector in the radial direction. We include the gravity by a point mass $M$ as an external force. Here $G$ is the gravitational constant, and $r$ is the distance from the centre of the star. The relativistic specific enthalpy $\\xi$ is defined as \\begin{equation} \\xi=\\frac{\\epsilon+p}{\\rho}=1+\\frac{\\Gamma}{\\Gamma-1}\\frac{p}{\\rho}, \\end{equation} where $\\epsilon$ is the energy density of matter including the photon energy coupled with the plasma. In SGR outbursts, since the luminosity much exceeds the Eddington luminosity, radiation energy density can exceed the thermal energy of the plasma. In the following pressure $p$ includes the contribution from the radiation pressure.  In this paper, we consider relativistic self-similar expansions of magnetic loops which started expansion at $t=t_\\mathrm{s}$ by loss of dynamical equilibrium and entered into a self-similar stage at $t=t_0 > t_\\mathrm{s}$. We do not consider the evolution of the loops before $t=t_0$.  For simplicity, we ignore the stellar rotation and assume axisymmetry. We can express the axisymmetric magnetic field in terms of two scalar functions $\\tilde{A}$ and $B$ as \\begin{equation} \\bmath B=\\frac{1}{r \\sin\\theta}\\left(\\frac{1}{r}\\frac{\\partial \\tilde{A}}{\\partial \\theta}, -\\frac{\\partial \\tilde{A}}{\\partial r}, B\\right),\\label{def:A} \\end{equation} in the polar coordinates $(r,\\theta,\\phi)$. The scalar function $\\tilde{A}(t,r,\\theta)$ denotes the magnetic flux, whose contours coincide with magnetic field lines projected on to the $r-\\theta$ plane.  We further assume that the fluid flow is purely radial; \\begin{equation} \\bmath v=v(t,r,\\theta)\\,\\bmath e_r. \\label{def:v} \\end{equation} Equations (\\ref{eq:MHDcont}), (\\ref{eq:MHDeom}), (\\ref{eq:MHDentropy}), and (\\ref{eq:Faraday}) are then expressed as \\begin{equation} \\frac{\\partial (\\rho\\gamma)}{\\partial t}+\\frac{1}{r^2}\\frac{\\partial (r^2 \\rho\\gamma v)}{\\partial r}=0,\\label{eq:cont} \\end{equation} \\begin{equation} \\rho\\gamma\\left[\\frac{\\partial}{\\partial t}+v\\frac{\\partial }{\\partial r}\\right](\\xi \\gamma v)= -\\frac{\\partial p}{\\partial r}-\\frac{1}{4\\pi r^2\\sin^2\\theta}\\left\\{ \\frac{\\partial \\tilde{A}}{\\partial r}\\left[\\left(\\hat{\\mathcal{L}}_{(r,\\theta)}\\tilde{A}\\right) +\\frac{\\partial}{\\partial t}\\left(v\\frac{\\partial \\tilde{A}}{\\partial r}\\right)\\right]+B\\left[\\frac{\\partial B}{\\partial r}+\\frac{\\partial (vB)}{\\partial t}\\right]\\right\\}-\\frac{G M \\rho\\gamma}{r^2},\\label{eq:eomr} \\end{equation} \\begin{equation} 4\\pi r^2 \\sin^2\\theta\\frac{\\partial p }{\\partial \\theta}+(1-v^2)B\\frac{\\partial B}{\\partial\\theta}+\\frac{\\partial \\tilde{A}}{\\partial \\theta}\\left[\\left(\\hat{\\mathcal{L}}_{(r,\\theta)}\\tilde{A}\\right)+\\frac{\\partial}{\\partial t}\\left(v\\frac{\\partial \\tilde{A}}{\\partial r}\\right)\\right]-vB^2\\frac{\\partial v}{\\partial \\theta}=0,\\label{eq:eomtheta} \\end{equation} \\begin{equation} (1-v^2)\\frac{\\partial \\tilde{A}}{\\partial r}\\frac{\\partial B}{\\partial \\theta}-\\frac{\\partial \\tilde{A}}{\\partial \\theta}\\left[\\frac{\\partial B}{\\partial r}+\\frac{\\partial (vB)}{\\partial t}\\right]-vB\\frac{\\partial \\tilde{A}}{\\partial r}\\frac{\\partial v}{\\partial \\theta}=0,\\label{eq:eomphi} \\end{equation} \\begin{equation} \\left[\\frac{\\partial}{\\partial t}+v \\frac{\\partial}{\\partial r}\\right]\\left(\\ln\\frac{p}{\\rho^\\Gamma}\\right)=0, \\label{eq:MHDentropy2} \\end{equation} \\begin{equation} \\frac{\\partial \\tilde{A}}{\\partial t}+v\\frac{\\partial \\tilde{A}}{\\partial r}=0,\\label{eq:A} \\end{equation} \\begin{equation} \\frac{\\partial B}{\\partial t}+\\frac{\\partial (vB)}{\\partial r}=0,\\label{eq:B} \\end{equation} where we used the MHD condition given by (\\ref{eq:mhdcond}) and introduced the operator \\begin{equation} \\hat{\\mathcal{L}}_{(r,\\theta)}\\equiv \\frac{\\partial^2}{\\partial r^2}+\\frac{\\sin\\theta}{r^2}\\frac{\\partial}{\\partial\\theta}\\left(\\frac{1}{\\sin\\theta}\\frac{\\partial}{\\partial\\theta}\\right). \\end{equation} Since our aim is to obtain self-similar solutions of these relativistic MHD equations, we assume that the time evolution is governed by the self-similar variable: \\begin{equation} \\eta=\\frac{r}{Z(t)},\\label{def:eta} \\end{equation} where $Z(t)$ is an arbitrary function of time. We further assume that the flux function $\\tilde{A}$ depends on time $t$ and the radial distance $r$ through the self-similar variable $\\eta$, as \\begin{equation} \\tilde{A}(t,r,\\theta)=\\tilde{A}(\\eta,\\theta). \\label{def:Aself} \\end{equation} When equation (\\ref{def:Aself}) is satisfied, the radial velocity $v$ has a form \\begin{equation} v=\\eta \\dot Z, \\label{def:sv} \\end{equation} from equation (\\ref{eq:A}). Here dot denotes the time derivative. Equation (\\ref{def:sv}) implies that the radial velocity $v$ does not depend on the polar angle $\\theta$. It then follows from equations (\\ref{eq:cont}) and (\\ref{eq:B}) that \\begin{equation} \\rho(t,r,\\theta)\\gamma(t,r)=Z^{-3}(t) D(\\eta,\\theta),\\label{eq:DZ} \\end{equation} \\begin{equation} B(t,r,\\theta)=Z^{-1}(t) Q(\\eta,\\theta),\\label{eq:QZ} \\end{equation} where $Q$ and $D$ are arbitrary functions of $\\eta$ and $\\theta$. These relations indicate that the magnetic flux and the total mass are conserved. Next we take the pressure $p$ as $p(t,r,\\theta)=Z^{l}P(\\eta,\\theta)$. Substituting this equation into equations (\\ref{eq:eomtheta}) and (\\ref{eq:MHDentropy2}), we obtain \\begin{equation} 4\\pi\\eta^2Z^{l+4}\\sin^2\\theta\\frac{\\partial P}{\\partial \\theta}+\\frac{\\partial \\tilde{A}}{\\partial \\theta}\\left[\\hat{\\mathcal{L}}_{(\\eta,\\theta)}\\tilde{A}+\\left(\\eta Z \\ddot Z-2\\eta\\dot{Z}^2\\right) \\frac{\\partial \\tilde{A}}{\\partial \\eta}-\\eta^2\\dot{Z}^2\\frac{\\partial^2 \\tilde{A}}{\\partial \\eta^2}\\right]+(1-\\eta^2 \\dot{Z}^2)Q\\frac{\\partial Q}{\\partial\\theta}=0,\\label{eq:wk1}\\\\ \\end{equation} \\begin{equation} \\frac{\\Gamma \\eta^2 Z \\ddot Z }{1-\\eta^2\\dot Z^2}+(3\\Gamma+l)=0,\\label{eq:MHDentropy3} \\end{equation} where we introduced an operator $\\hat{\\mathcal{L}}_{(\\eta,\\theta)}$: \\begin{equation} \\hat{\\mathcal{L}}_{(\\eta,\\theta)}\\equiv \\frac{\\partial^2 }{\\partial\\eta^2}+\\frac{\\sin\\theta}{\\eta^2}\\frac{\\partial}{\\partial\\theta}\\left(\\frac{1}{\\sin\\theta}\\frac{\\partial}{\\partial\\theta}\\right) =\\frac{1}{Z^2}\\hat{\\mathcal{L}}_{(r,\\theta)}. \\end{equation} To satisfy these equations, $p$, $Z$ and $\\Gamma$ should have forms \\begin{equation} p(t,r,\\theta)=Z^{-4}P(\\eta,\\theta),\\label{eq:PZ} \\end{equation} \\begin{equation} Z(t)=t, \\label{sol:Z} \\end{equation} and \\begin{equation} \\Gamma = \\frac{4}{3}.\\label{sol:shr}\\\\ \\end{equation} This adiabatic index corresponds to the radiation pressure dominant plasma. Thus our model can describe the evolution of a fireball confined by magnetic fields.  Equations (\\ref{eq:DZ}) and (\\ref{eq:PZ}) indicate that the magnetic loops expand adiabatically. By using equations (\\ref{def:sv}), (\\ref{eq:DZ}), (\\ref{eq:QZ}), (\\ref{eq:PZ}), (\\ref{sol:Z}) and (\\ref{sol:shr}), equations (\\ref{eq:eomr}), (\\ref{eq:eomphi}) and (\\ref{eq:wk1}) are expressed as \\begin{equation} D(\\eta,\\theta)=\\frac{\\eta^2}{GM}\\left\\{\\frac{4\\eta P}{1-\\eta^2}-\\frac{\\partial P}{\\partial \\eta}\\right. \\left.-\\frac{1}{4\\pi \\eta^2 \\sin^2\\theta}\\left[\\frac{\\partial \\tilde{A}}{\\partial \\eta} \\left(\\hat{\\mathcal{L}}_{(\\eta,\\theta)}\\tilde{A} -\\frac{\\partial}{\\partial \\eta} \\left(\\eta^2 \\frac{\\partial\\tilde{A}}{\\partial \\eta}\\right)\\right) +Q \\frac{\\partial}{\\partial \\eta}\\left(Q(1-\\eta^2) \\right)\\right]\\right\\},\\label{eq:eomr2} \\end{equation} \\begin{equation} (1-\\eta^2)\\frac{\\partial \\tilde{A}}{\\partial \\eta}\\frac{\\partial Q}{\\partial \\theta}-\\frac{\\partial \\tilde{A}}{\\partial\\theta}\\frac{\\partial}{\\partial \\eta}\\left[(1-\\eta^2)Q\\right]=0,\\label{eq:eomphi2} \\end{equation} \\begin{equation} 4\\pi\\eta^2\\sin^2\\theta\\frac{\\partial P}{\\partial\\theta}+\\frac{\\partial \\tilde{A}}{\\partial \\theta}\\left[\\hat{\\mathcal{L}}_{(\\eta,\\theta)}\\tilde{A} -\\frac{\\partial}{\\partial \\eta}\\left(\\eta^2\\frac{\\partial \\tilde{A}}{\\partial \\eta}\\right) \\right]+(1-\\eta^2)Q\\frac{\\partial Q}{\\partial\\theta}=0.\\label{eq:eomtheta2} \\end{equation} From equation (\\ref{eq:eomphi2}), a formal solution of  $Q$ is obtained as \\begin{equation} Q(\\eta,\\theta)=\\frac{\\mathcal{G}(\\tilde{A})}{1-\\eta^2},\\label{sol:Qformal} \\end{equation} where $\\mathcal{G}$ is an arbitrary function.  Self-similar solutions can be constructed as follows. First we prescribe an arbitrary function $\\tilde{A}(\\eta,\\theta)$ (or $Q(\\eta,\\theta)$). Then, equation (\\ref{eq:eomphi2}) determines the function $Q(\\eta,\\theta)$ (or $\\tilde{A}(\\eta,\\theta)$). Functions $\\tilde{A}$ and $Q$ determine the pressure $P(\\eta,\\theta)$ according to equation (\\ref{eq:eomtheta2}). Finally, the density function $D(\\eta,\\theta)$ is obtained by equation (\\ref{eq:eomr2}).  Note that from the equation (\\ref{def:sv}) and (\\ref{sol:Z}), the radial velocity has a simple form as \\begin{equation} v=\\frac{r}{t}. \\label{sol:v} \\end{equation} Since the time derivative of the velocity becomes zero, i.e., $D\\bmath v/D t=0$, equations (\\ref{eq:eomr2})-(\\ref{eq:eomtheta2}) describe the freely expanding solution. This means that there is a reference frame that all forces balance. By substituting equations (\\ref{eq:DZ}), (\\ref{eq:QZ}), (\\ref{eq:PZ}), (\\ref{sol:Z}), and (\\ref{sol:v}) into the equations of motion (\\ref{eq:MHDeom}), we obtain \\begin{equation} \\frac{\\Gamma}{\\Gamma-1}\\frac{\\gamma^2 v^2 p}{r}\\bmath e_r-\\nabla p+\\rho_e \\bmath E+\\bmath j \\times \\bmath B-\\frac{GM\\rho\\gamma}{r^2}\\bmath e_r=0.\\label{eq:reducedeom} \\end{equation} The first term on the left hand side comes from the inertia. For convenience, we call this term as a thermal inertial term throughout this paper. When we neglect the terms of order $\\left(v/c\\right)^2$, equation (\\ref{eq:reducedeom}) reduces to the equations of the force balance in non-relativistic MHD. ",
        "conclusions": "\\label{summary} By extending the self-similar solutions derived by \\citet{1982ApJ...261..351L}, we derived self-similar solutions of relativistically expanding magnetic loops taking into account the toroidal magnetic fields. The dipolar solution derived in $\\S$ \\ref{sss1} gives us an insight into the relativistic expansion of the magnetic loops because of its simplicity. However, the shell and flux rope solutions derived in $\\S$ \\ref{sss2}, \\ref{sss3} have more physically interesting properties such as an enhanced magnetic pressure at the shells and flux rope structures. Such configurations might be more probable for SGR flares.  The equations of motion in the self-similar stage are similar to those of the static equilibrium state except the existence of the relativistic thermal inertial term and the electric field. This fact allows us to evaluate the non-kinetic part of the total energy in the magnetic loops by using the virial theorem. The magnetic loops with shell or flux rope structures carry more energy than the simple dipole solution. The energy is comparable to the observed energy of the SGR giant flares.  In relativistically expanding magnetic loops, the effect of the displacement current becomes important. In dipolar solution, the displacement current becomes larger than the real current $\\nabla \\times \\bmath B/(4\\pi)$ for faster expansion speed ($V_\\mathrm{max}>a_*$). This effect reduces the toroidal current and weakens the magnetic tension force. To balance the reduced magnetic tension force, the pressure decreases.   We found that the energy is transferred to $r>R(t)$ in dipolar solutions. In the shell and flux rope solutions, the energy is transferred to $r>R(t)$ unless condition (\\ref{sol2:knowave}) is satisfied. The condition can be interpreted as that for the total reflection of the MHD waves in the shell. Dipolar solutions always have leakage (transmission of Poynting flux to the region $r\\gid R(t)$) because $B_\\theta \\ne 0$ at $r=R(t)$. The shell and flux rope solutions have perfectly reflecting solutions in which the total energy in $r<R(t)$ is conserved. It means that the solutions are energy eigenstates of the system. The eigenstates can be obtained by adjusting the parameter $k$.  In this paper, we obtained solutions for freely expanding magnetic loops, i.e., $Dv / Dt = 0$. We assumed that the magnetic loops have sufficiently large energy to drive the expansion. When the flux function $\\tilde{A}$ increases with time, the toroidal magnetic fields will also increase with time. The toroidal magnetic fields will then affect the dynamics through the magnetic pressure. Such solutions can describe the accelerating magnetic loops.   Magnetic fields can be expressed as the sum of the Fourier modes in the polar angle. The modes and their amplitudes should be determined at the boundary where the magnetic twist is injected on the surface of the star. It is not shown but we can construct more complex solutions that the poloidal magnetic fields are expressed by the sum of the Fourier modes, i.e., $\\tilde A \\propto \\sin^n \\theta$. In actual explosion, the opening angle of the expanding magnetic loops depends on the location at which the magnetic twist is injected on the surface of the central star. Such a solution may be expressed as the sum of the Fourier modes for the poloidal and toroidal magnetic fields. We should note that SGR flares are not necessarily axisymmetric. Models including the non-axisymmetrically expanding magnetic loops will be a subject of future works."
    },
    "0812/0812.0308_arXiv.txt": {
        "abstract": "{Motivated by the recent PAMELA and ATIC data, one is led to a scenario with heavy vector-like dark matter in association with a hidden $U(1)_X$ sector below GeV scale. Realizing this idea in the context of gauge mediated supersymmetry breaking (GMSB), a heavy scalar component charged under $U(1)_X$ is found to be a good dark matter candidate which can be searched for direct scattering mediated by the Higgs boson and/or by the hidden gauge boson.  The latter turns out to put a stringent bound on the kinetic mixing parameter between $U(1)_X$ and $U(1)_Y$: $\\theta \\lesssim 10^{-6}$. For the typical range of model parameters, we find that the decay rates of the ordinary lightest neutralino into hidden gauge boson/gaugino and photon/gravitino are comparable, and the former decay mode leaves displaced vertices of lepton pairs and missing energy with distinctive length scale larger than 20 cm for invariant lepton pair mass below 0.5 GeV. An unsatisfactory aspect of our model is that the Sommerfeld effect cannot raise the galactic dark matter annihilation by more than 60 times for the dark matter mass below TeV.} ",
        "introduction": "It is now firmly established that 23\\% of the energy density of the universe consists of an unknown particle called dark matter. Discovering the nature of dark matter would be one of the most important tasks in current and future theoretical and experimental investigations.  Numerous searches for galactic dark matter have been made to observe direct signals of dark matter scattering with nuclei and indirect evidences of dark matter annihilation to various Standard Model particles. Recently, a number of stimulating results toward indirect signals have been announced.  The Payload for Antimatter Matter Exploration and Light-nuclei Astrophysics (PAMELA) collaboration has reported an excess in the positron fraction, $e^+/(e^++e^-)$, but no excess in anti-proton fraction $\\bar{p}/p$~\\cite{PAMELA}. The observed spectrum departed from the background calculation of the cosmic-ray secondary positron spectrum~\\cite{e+Background} for energies 10--100 GeV. This recent result is comparable with less certain excesses observed in HEAT~\\cite{HEAT} and AMS-01~\\cite{AMS-01}. The PPB-BETS balloon experiment and the Advanced Thin Ionization Calorimeter (ATIC) instrument have also reported an excess in the $e^++e^-$ energy spectrum of $300-800$ GeV~\\cite{PPB-BETS, ATIC}.  The origin of such excesses could be astronomical objects such as a pulsar which is expected to emit energetic electron-positron pairs~\\cite{Pulsars,GemingaPulsar,Hooper11}. On the other hand, excessive antiparticle fluxes in galactic cosmic rays have been considered as a primary way for indirect detection of dark matter. It is encouraging that galactic dark matter annihilation can easily reproduce the amplitudes and spectral shapes of the PAMELA and ATIC data. However, there appear two puzzling aspects. First, the dark matter annihilation cross-section  is required to be enhanced by order of $100-1000$ compared with what is allowed by the standard thermal relic abundance analysis. According to a recent analysis based on $\\Lambda$CDM $N$-body simulations, however, the boost factor from clumpy matter distribution can hardly larger than about 10~\\cite{BfactorLimit1_Nsimulation}. Second, no distinctive excess of antiprotons in the PAMELA data disfavors most conventional dark matter candidates. A new analysis of the $\\overline{p}/p$ ratio compared to the PAMELA data puts stringent limits on possible enhancements of the $\\overline{p}/p$ flux from dark matter annihilation~\\cite{BfactorLimit2_AntiprotonAnaysis}. Many ideas and models of dark matter annihilation or decay have already been suggested for the explanation of the recent observations~\\cite{Neutralino, SUSYDM1, MinimalDM, ModelIndep, DecayingU(1)x, ModelIndep2, ModelIndep3_Sommerfeld, TwoDM, ArkaniHamed:DM, SecludedDM, TwoDM2_1metastable, Non-MinimalDM, NWeinerLightBoson, AxionPortal, 4fermi-interaction, Stueckelberg:PAMELA, LeptonicallyDecay, SUSYDM2, singletUED, LeptophilicDM, U(1)x_U(1)B-L, DecayingDM, U(1)_Lmu-Ltau, ATIC_PAMELA, ATIC_PAMELA_DecayingDM, Multi-Zurek}. However, it appears to be a hard task to understand the nature of enhanced hadrophobic annihilation in a consistent framework with thermal dark matter.  \\medskip  In this paper, we elaborate the idea of dark matter charged under an extra $U(1)_X$ gauge group suggested in Refs.~\\cite{ArkaniHamed:DM, SecludedWIMP, ArkaniHamed:LHC}. The presence of $U(1)_X$ broken at the sub-GeV scale is motivated to explain the two intriguing features of the PAMELA data. First, the DM annihilation to extra light gauge bosons can be enhanced significantly by non-perturbative Sommerfeld effect \\cite{Sommerfeld-DM, Sommerfeld-Hisano, Sommerfeld-Profumo, Sommerfeld-Cirelli, Sommerfeld-HiggsPortal}. Second, the hadrophobic nature of the $U(1)_X$ gauge boson decay can be understood by kinematics when its mass is below the GeV scale~\\cite{ArkaniHamed:DM, NWeinerLightBoson}. We will consider the extra $U(1)_X$ gauge boson hidden from  the Standard Model sector except for a small kinetic mixing  between $U(1)_X$ and $U(1)_Y$, through which the $U(1)_X$ gauge bosons can decay mostly into lepton-antileptons. It was shown long ago that kinetic mixing can exist between  two $U(1)$ gauge bosons without violating the gauge-invariance and renormalizability \\cite{KineticMixing}. The mixing between an unbroken extra $U(1)$ and $U(1)_{em}$ has been used to explain other anomalous astronomical observation such as galactic 511 keV $\\gamma$-rays~\\cite{511millicharged}.     Breaking $U(1)_X$ below the GeV scale can arise naturally by radiative mechanism, in particular, in the context of gauge mediated supersymmetry breaking (GMSB) \\cite{ArkaniHamed:LHC}.  We will first present a concrete way of realizing such a scheme, which predicts a $U(1)_X$ charged scalar field as a dark matter candidate. The dark matter mass in the range of 600--1000 GeV is preferred for the simultaneous explanation of the PAMELA and ATIC data. Unfortunately, the Sommerfeld enhancement factor of our scenario is found to be about 40--60 for this mass range. The assumption of some clumpy distribution of dark matter would be a reasonable additional source for further increase of the boost factor. Our dark matter candidate can yield observable signals for direct detection through elastic scatterings mediated by the hidden gauge boson $X$ or Higgs bosons. In the first case, we draw a severe constraint on the kinetic mixing parameter: $\\theta\\lesssim 10^{-6}$ for $m_X\\sim 0.4$ GeV.  One of the interesting consequences of our scenario is that the ordinary lightest supersymmetric particle (OLSP), typically neutralino, decays instantaneously to the $X$ boson and its superpartner $\\tilde{X}$ through one-loop diagram with heavy dark matter superfields in the loop.  Then, the produced $X$ boson subsequently decays to two leptons with a decay length typically larger than 10 cm, which may be observed at the LHC. ",
        "conclusions": "The presence of TeV dark matter associated with a hidden $U(1)_X$ gauge symmetry broken below the GeV scale has been motivated by the recent results from PAMELA and ATIC. Elaborating this idea in the framework of gauge mediated supersymmetry breaking, we obtain some interesting constraints and prospects of the scenario.  Explaining the heaviness of dark matter in relation to the resolution of the $\\mu$ and $B\\mu$ problem of GMSB models, it follows that the vector-like scalar particle charged under $U(1)_X$ becomes a good dark matter candidate for the gauge coupling $\\alpha_X \\sim \\alpha$ and the dark matter mass $\\sim$ TeV. The other `hidden' particles, namely the $U(1)_X$ gauge boson and gaugino, higgs boson and higgsino, obtain masses below the GeV scale through a radiative mechanism with properly chosen particle contents and Yukawa couplings among them. The hidden sector communicates with the visible (Standard Model) sector either by the Yukwa terms generating the Higgsino and dark matter masses  or by the kinetic mixing between $U(1)_X$ and $U(1)_Y$.  It turns out that the kinetic mixing parameter $\\theta$ receives a strong upper limit $\\theta \\lesssim 10^{-6}(m_X/0.4 \\,\\mbox{GeV})^2$ coming from the experimental data on direct detection of dark matter.  In this scenario, positive signals for direct dark matter detection in future experiments may arise either from the vector interaction controlled by $\\theta$ or from the scalar interaction exchanging the Higgs boson.  Another interesting consequence for the LHC experiment is that the ordinary LSP decay leaves quite distinctive signals of missing energy and energetic lepton pairs with invariant mass below GeV and decay gaps larger than  about  10 cm. Furthermore, the rates for this decay and the usual decay to photon and gravitino, a characteristic of GMSB models, are found to be comparable in typical parameter range of the model. Thus, the scenario of the sub-GeV $U(1)_X$ sector in GMSB models may be tested at the LHC by observing both signatures of lepton pairs plus missing energy and photons plus missing energy.   However, we find that the boost factor for the galactic dark matter annihilation required by the PAMELA and ATIC data can not be obtained solely by the Sommerfeld effect in our scenario. Typical Sommerfeld enhancement factor is in the range of $40-60$ for the dark matter masses of $600-1000$ GeV.  Thus, we may have to invoke additional sources of the boost factor like the clumpy distribution of dark matter."
    },
    "0812/0812.1718_arXiv.txt": {
        "abstract": "We present large-area maps of the CO $J$=3-2 emission obtained at the James Clerk Maxwell Telescope for four spiral galaxies in the Virgo Cluster. We combine these data with published CO $J$=1-0, 24 $\\mu$m, and H$\\alpha$ images to measure the CO line ratios, molecular gas masses, and instantaneous gas depletion times. For three galaxies in our sample (NGC 4254, NGC4321, and NGC 4569), we obtain molecular gas masses of $7\\times 10^8 - 3\\times 10^9$ M$_\\odot$ and disk-averaged instantaneous gas depletion times of 1.1-1.7 Gyr. We argue that the CO $J$=3-2 line is a better tracer of the dense star forming molecular gas than the CO $J$=1-0 line, as it shows a better correlation with the star formation rate surface density both within and between galaxies. NGC 4254 appears to have a  larger star formation efficiency(smaller gas depletion time), perhaps because it is on its first passage through the Virgo Cluster. NGC 4569 shows a large-scale gradient in the gas properties traced by the CO $J$=3-2/$J$=1-0 line ratio, which suggests that its interaction with the intracluster medium is affecting the dense star-forming portion of the interstellar medium directly. The fourth galaxy in our sample, NGC 4579, has weak CO $J$=3-2 emission despite having bright 24 $\\mu$m emission; however, much of the central luminosity in this galaxy may be due to the presence of a central AGN. ",
        "introduction": "Molecular gas provides the fuel for star formation and therefore plays an important role in both the current appearance and future evolution of galaxies. Early surveys for molecular gas in galaxies revealed a wide range in total gas content, from the rarely detected elliptical galaxies to gas-rich massive spirals such as M51 \\citep{ys91,b93,y95}. For spiral galaxies, the variations in gas content roughly track the star formation rate for different systems, such that the global star formation efficiency (the inverse of which is the gas depletion time, $t_{gas} = M_{H_2} / SFR$) does not depend strongly on morphology when far-infrared luminosity is used to trace the star formation rate \\citep{ys91}. Detailed studies of individual galaxies provided evidence for enhanced star formation efficiencies (implying shorter instantaneous gas depletion times) at certain positions within spiral arms \\citep{v88} or along them \\citep{k92,k96} compared to the disk-averaged value. A focused survey of spiral galaxies in the Virgo cluster showed that the molecular gas content of HI-deficient cluster spirals is similar to that of their more gas-rich counterparts \\citep{ky89}. \\citet{kenn07} find that the star formation rate surface density correlates with the molecular but not the atomic gas surface density.  To date, most large surveys of molecular gas in galaxies have been carried out using the CO $J$=1-0 transition; this line is sensitive to the vast majority of molecular gas in the cold interstellar medium in galaxies, since the $J$=1 level is only 5.5 K above the ground state and the $J$=1-0 transition has a critical density of $1.1\\times 10^3$ cm$^{-3}$ in optically thin gas. The detailed connection between CO emission and molecular gas mass and ultimately the star formation process depends rather critically on the value of the CO-to-H$_2$ conversion factor \\citep{s88} and any dependencies of that conversion factor on temperature, density, or metallicity \\citep{dss85,mb88}. While individual galaxies have been studied in higher transitions, primarily CO $J$=2-1 and CO $J$=3-2, the only large surveys available to date in the CO $J$=3-2 transition are observations of the central regions of galaxies \\citep{m99,hs03,y03,komugi07}. Since the CO $J$=3 level is  33 K above the ground state and the $J$=3-2 transition has a critical density of $2.1\\times 10^4$ cm$^{-3}$ in optically thin gas, the $J$=3-2 transition naturally traces gas that is on average warmer and/or denser than the CO $J$=1-0 line.  Examining the CO $J$=3-2/1-0 line ratio both within and between galaxies can give information on any large-scale variations in the physical conditions in the molecular gas \\citep{w97,pw00}. In addition, the CO $J$=3-2 emission may trace the molecular gas that is more directly connected with star formation; indeed, \\citet{i08} have found that the global CO $J$=3-2 luminosity correlates nearly linearly with the global star formation rate over five orders of magnitude. \\citet{b08} also find a nearly linear relation between the H$_2$ surface density traced by the CO $J$=2-1 line and the star formation rate surface density. These relations are similar to the linear correlation of HCN with far-infrared luminosity seen by \\citet{gao04} and suggests that these CO transitions may also be effective tracers of the densest, star forming gas. In contrast, the CO $J$=1-0 luminosity  has a much shallower correlation with star formation rate  with a slope $0.62 \\pm 0.08$ \\citep{g05}. This shallower slope implies that there is not a one-to-one relation between the amount of molecular gas (as traced by the CO $J$=1-0 line) and the amount of star formation in a galaxy. The shallower slope compared to the CO J=3-2 data also indicates that the two lines are not tracing exactly the same gas components. For example, the CO $J$=1-0 line can include emission from lower opacity, more diffuse molecular gas \\citep{ww94,rb05}. Examining the correlation between molecular gas and star formation rate as a function of the gas density may yield new insights into the star formation process in galaxies.  An ideal laboratory for the acquisition of resolved molecular maps in galaxies is the nearby Virgo Cluster.  At a distance of 16.7 Mpc \\citep{mei07}, it is the largest concentration of galaxies \\citep[cluster mass $4.2 \\times 10^{14}$ M$_\\odot$;][]{mcl99} within 35 Mpc. It has been the target of numerous studies over the years and plays an important role in our understanding of how galaxies evolve in dense environments. The most extensive optical survey was carried out by \\citet{b85} and resulted in the Virgo Cluster Catalog (VCC) containing $\\sim 2100$ certain or possible cluster members over an area of 140 deg$^2$. The Virgo Cluster is particularly interesting for studying the impact of the dense environment on the ISM in galaxies, as properties such as ram-pressure stripping, tidal interactions, etc., become important in environments with a massive intracluster medium.  Recently, the Virgo Cluster has been included within the area being mapped in the HI 21 cm line by the ALFALFA survey \\citep{gio05}. Comparison with optical catalogs shows $\\sim 200$ detections associated with optical counterparts as well as 10 detections associated with tidal debris tails \\citep{gav08}. High-resolution HI observations with the VLA have identified seven spiral galaxies with long HI tails pointing away from M87 \\citep{chu07}, which suggests that interaction with the cluster as a whole has created these tails. Other evidence for environmental modification includes ram-pressure stripped extraplanar gas tails and disturbances in radio morphologies \\citep{k08,m08}.  Spiral galaxies in the Virgo Cluster were first surveyed for molecular emission by \\citet{s86}, who detected 23 out of 47 spirals observed with a beam size of 100$^{\\prime\\prime}$. The primary reference for the molecular gas content of spiral galaxies in the Virgo Cluster remains the work of \\citet{ky86,ky88a,ky88b,ky89}. Using a complete blue-magnitude selected sample, they found that the molecular gas content of HI-deficient cluster spirals in Virgo is similar that of to their more gas-rich counterparts. Interferometric CO $J$=1-0 observations of 15 Virgo galaxies show a wide range of molecular gas morphology as well as higher molecular gas concentrations in the galaxies with deeper gravitational potential wells \\citep{s03a,s03b}. High resolution CO $J$=1-0 images for a few galaxies in the Virgo cluster are also presented in \\citet{h03} and \\citet{k07}. \\citet{hs03} made low resolution observations of the CO $J$=3-2 and $J$=2-1 lines with the KOSMA 3~m telescope of the 20 strongest sources from the CO $J$=1-0 survey of \\citet{s86}. They detected 18 and 16 spiral galaxies, respectively, in the two lines and measured CO $J$=3-2/$J$=1-0 line ratios ranging from 0.14 to 0.35  and CO $J$=2-1/1-0 line ratios from 0.5 to 1.1.    The star formation properties of Virgo cluster galaxies have also been studied extensively. \\citet{ky04a} found that half of the bright Virgo spirals have truncated H$\\alpha$ disks even though they have relatively undisturbed stellar disks. This truncation suggests that ICM-ISM stripping is primarily responsible for the reduced star formation rates of bright Virgo spirals \\citep{g02,ky04b}. \\citet{fg08} found that cluster members which have lost even a moderate amount of their atomic gas also have lower molecular gas content and star formation rates.   In this paper, we present some first results from the JCMT Nearby Galaxies Legacy Survey (NGLS). We focus on new, wide-area observations of the CO $J$=3-2 emission for the four large spiral galaxies in the Virgo Cluster (\\object{NGC 4254}, \\object{NGC 4321}, \\object{NGC 4569}, and \\object{NGC 4579}) that are part of the Spitzer Infrared Nearby Galaxies Survey (SINGS) sample \\citep{k03}. We briefly describe the structure and goals of the NGLS in \\S\\ref{sec-ngls} and review previous observations of the interstellar medium in the particular Virgo Cluster galaxies in our sample in \\S\\ref{sec-bkgd}. We discuss our observations and data reduction in \\S\\ref{sec-obs}. In \\S\\ref{sec-lineratios}, we combine the data with CO $J$=1-0 maps from \\citet{k07} to estimate the CO $J$=3-2/$J$=1-0 line ratio and the molecular gas mass. In \\S\\ref{sec-tgas}, we combine these new data with H$\\alpha$ data and reprocessed 24 $\\mu$m data from the SINGS survey \\citep{k03} to obtain maps of the star formation rate \\citep{c07} and gas depletion time. We discuss the usefulness of the CO $J$=3-2 line as a tracer for dense star forming gas and the gas and star formation properties of the four galaxies in the context of their internal structure and broader environment in \\S\\ref{sec-discuss}. We give our conclusions in \\S\\ref{sec-concl}.  \\subsection{The JCMT Nearby Galaxies Legacy Survey}\\label{sec-ngls}  The JCMT Nearby Galaxies Legacy Survey (NGLS)\\footnote{http://www.jach.hawaii.edu/JCMT/surveys/} focuses on the molecular gas and dust content of galaxies within 25 Mpc using three primary samples: (1) galaxies from the SINGS \\citep{k03} sample \\citep[see][for first results from the NGLS for additional SINGS galaxies]{w08}; (2) an HI flux limited sample of galaxies from the Virgo cluster; and (3) an HI flux limited sample of galaxies in the field. The first phase of the NGLS consists of  CO $J$=3-2 observations out to $D_{25}/2$ for 21 galaxies from the SINGS sample as well as the complete Virgo sample; this phase of the survey is currently complete. The second phase consists of CO $J$=3-2 observations for the entire sample and will begin in 2009 if additional observing time is allocated to the survey. The third phase, for which observing time has been allocated already, will consist of SCUBA-2 observations of the entire sample and should begin in late 2009.  Specific science goals for the NGLS include: searching for evidence of cold dust and measuring its mass fraction in galaxies of different types; measuring the amount of warm, dense molecular gas associated with star formation using the CO $J$=3-2 line; comparing galaxies with similar morphologies and luminosities in the field and in the Virgo cluster to determine the effects of cluster membership; using CO rotation curves to trace the dark matter distribution and the frequency of occurrence of nuclear gas concentrations which may feed central starbursts or black holes; and measuring the local submillimeter luminosity function and dust mass function to luminosities up to 100 times fainter than previous studies \\citep[e.g.,][]{d00}.      \\subsection{Properties of four Virgo spiral galaxies}\\label{sec-bkgd}  Four large spiral galaxies (\\object{NGC 4254}, \\object{NGC 4321}, \\object{NGC 4569}, and \\object{NGC 4579}) are the only Virgo Cluster galaxies included in the SINGS survey \\citep{k03} and there is a large collection of information on their dust and interstellar medium properties. Some basic properties of the galaxies are given in Table~\\ref{tbl-props}. The global 1-850 $\\mu$m spectral energy distributions (SEDs) for these four galaxies are very similar \\citep{d05}, although the $\\alpha$ parameter in the SED fit ranges from a high 3.8 in NGC 4579 to a low of 2.4 in NGC 4254. \\citet{d07} calculate the dust mass and gas to dust ratio for 65 galaxies in the SINGS sample. Of the four Virgo spirals discussed in our paper, only NGC 4579 stands out as having a factor of 3 larger gas to dust ratio; this is the second highest gas to dust ratio in the entire sample analyzed by \\citet{d07}. This high gas-to-dust ratio may be a by-product of the difficulty in determining the 850 $\\mu$m flux from the relatively low signal-to-noise SCUBA map (Bendo, private communication). Morphologically, of the four Virgo spirals, NGC 4569 is the most centrally concentrated at 24 $\\mu$m, while NGC 4254 is both the least centrally concentrated and the most asymmetric \\citep{b07}. The star formation distribution based on H$\\alpha$ emission is classified as normal for NGC 4254 but truncated/normal for the other three galaxies \\citep{ky04a}. Below we give some important details on each of the individual galaxies.  Recent HI observations of NGC4254 (M99) show a long tidal tail extending to the north; its lopsided morphology, high HI content, and high relative velocity suggest that this galaxy is undergoing its first encounter with the center of the Virgo cluster \\citep{h07}. The galaxy has a normal molecular gas to total gas fraction \\citep{n06}.   NGC 4321 (M100) is classified as HII/LINER \\citep{h97}. High resolution CO $J$=1-0 observations show two spiral arms and gas in the nucleus as well as faint emission bridging the arms and the nucleus \\citep{gb05}. The HI emission is truncated near the edge of the optical disk \\citep{c90,braun07}. NGC 4321 is interacting with NGC 4322, as shown by the bridge between them that is visible in optical light but not in HI \\citep{k93}. It also has a rather prominent circumnuclear ring with much enhanced star formation \\citep{k95a,k95b}.  NGC 4569 (M90) has been called an anemic spiral because of the small extent of its HI emission relative to its optical size \\citep{s99}, although \\citet{ky04a} prefer a designation of truncated/normal because its star formation rate per unit area is relatively normal. \\citet{b07} note that there is almost no 24 $\\mu$m emission from the extended disk. The CO $J$=1-0 emission shows a compact bar oriented roughly north-south \\citep{s99} with complex kinematics indicating non-circular motions in the central arcminute \\citep{j05}. This galaxy has a higher molecular gas to total gas fraction than NGC 4254, perhaps due to ram pressure stripping of HI even within the stellar disk or a higher external pressure due to the intracluster medium \\citep{n06}. However, the fact that the only remaining gas is in the inner disk, which tends to have a high molecular gas fraction, could also explain this result. \\citet{ky04a} find that the galaxy has some HII regions that appear to be outside the plane. They also note that rotation plus forward motion relative to the ICM can create an asymmetric ram pressure, which in simulations can produce a dominant extraplanar gas arm emerging from a truncated gas disk. \\citet{v04} present also evidence for ram pressure stripping by comparing the kinematically distinct western arm with dynamical simulations. \\citet{b06} argue that active gas stripping, rather than simply stopping gas inflow, is required to explain the star formation history in NGC 4569. \\citet{j05} suggest that NGC 4569 is in the early stages of bar or tidally driven gas inflow.  NGC 4579 (M58) is classified as LINER 1.9/Seyfert 1.9 \\citep{gb05}. High resolution CO $J$=2-1 observations show two spiral arms approaching to within a few arcseconds of the AGN \\citep{gb05} and weak CO emission at the AGN itself \\citep{gb08}. At lower resolution, weak CO emission is detected throughout the disk \\citep{n06}. Despite the weak emission, this galaxy has a higher molecular gas to total gas fraction than NGC 4254, and so may also have been affected by ram pressure stripping or a higher external pressure \\citep{n06}.  In summary, this small sample of four spiral galaxies probes a large fraction of the characteristics seen for luminous spirals in the Virgo cluster such as HI deficiency, star formation rate, and presence of nuclear activity. Thus, a detailed study of the molecular gas and star formation properties of these galaxies should provide insight into the effects on the ISM in galaxies and galaxy structure and evolution that can be addressed using the complete Virgo sample from the NGLS. ",
        "conclusions": "We have used the JCMT to map the CO $J$=3-2 emission from four spiral galaxies in the Virgo cluster. These galaxies are included in the SINGS survey \\citep{k03} and so have a wealth of information on their dust and interstellar medium properties. These galaxies also allow us to probe a range of galaxy properties including HI deficiency, star formation rate, and nuclear activity.  Bright CO $J$=3-2 emission is detected over the extended disks of NGC 4254, NGC4321, and NGC 4569; NGC 4579 is weakly detected (3$\\sigma$) at one position in the disk and has only an upper limit in the central region. Combining our data with CO $J$=1-0 maps from \\citet{k07}, we derive disk-averaged CO $J$=3-2/$J$=1-0 line ratios of 0.33 in NGC 4254, 0.36 in NGC 4321, and 0.25 in NGC 4569. The line ratio in the center of NGC 4321 is significantly larger (0.79), while NGC 4569 shows a gradient in the line ratio from 0.06 in the south to 0.53 in the north. The very weak CO $J$=3-2 emission in NGC 4579 suggests a line ratio in the disk similar to the lowest value seen in NGC 4569. These line ratios are towards the low end of the range typically measured in galaxies \\citep{m99}. However, previous observations have focused on the extreme central regions, which our data suggest may have enhanced line ratios in many galaxies.  We combine H$\\alpha$ and 24 $\\mu$m images from the SINGS survey \\citep{k03} to obtain maps of the star formation rate following the prescription in \\citet{c07}. We combine these star formation rate images with maps of the molecular gas surface density derived from the CO $J$=3-2 data to obtain maps of the instantaneous gas depletion timescale (the inverse of the star formation efficiency). Of the three CO-bright galaxies, the disk-averaged gas depletion time is shortest for NGC 4254 (1.1 Gyr), roughly 50\\% less than for NGC 4321 (1.7 Gyr) and NGC 4569 (1.6 Gyr). The fourth galaxy, NGC 4579, has weak CO $J$=3-2 emission and small gas depletion times in the disk, 4-6 times smaller than in the other three galaxies. These times, which are substantially shorter than a Hubble time, are a measure of the instantaneous gas depletion time and do not take into account possible replenishing of the dense molecular gas from more diffuse molecular or atomic gas.  We have compared the gas depletion times obtained using the CO $J$=3-2 emission with those obtained using the CO $J$=1-0 emission, which is the more commonly used tracer of molecular gas. We find that using the CO $J$=3-2 line results in more uniform instantaneous gas depletion times both within galaxies and from one galaxy to another. We argue that the CO $J$=3-2 line is a better tracer of the dense molecular gas involved in star formation than the CO $J$=1-0 line. Alternatively, real variations in the star formation efficiency may result in variation in the amount of heating per unit gas, leading to more uniform apparent gas depletion times using the CO $J$=3-2 line. These two scenarios may be testable using radiative transfer codes and multiple CO emission lines to measure the average physical conditions in the molecular gas.  Looking at the internal structure and external environment, we suggest that the smaller gas depletion time (or larger star formation efficiency) in NGC 4254 may be due to the fact that it seems to be encountering the dense portion of the Virgo cluster for the first time \\citep{h07}. In NGC 4569, the gradient in the CO $J$=3-2/$J$=1-0 line ratio (and more weakly in the gas depletion time) from south to north suggests that active ram pressure in the northern half of the galaxy may be producing an increase in the average density or temperature of the molecular gas."
    },
    "0812/0812.3237_arXiv.txt": {
        "abstract": "Utilizing deep Hubble Space Telescope imaging from the two largest field galaxy surveys, the Extended Groth Strip (EGS) and the COSMOS survey, we examine the structural properties, and derive the merger history for $21,902$ galaxies with M$_{*} > 10^{10}$ \\solm at $z < 1.2$. We examine the structural CAS parameters of these galaxies, deriving merger fractions, at $0.2 < z < 1.2$, based on the asymmetry and clumpiness values of these systems.   We find that the merger fraction between $z = 0.2$ and $z = 1.2$ increases from roughly $f_{m} = 0.04\\pm0.01$ to $f_{m} = 0.13\\pm0.01$.  We furthermore detect, at a high significance, an abrupt drop in the merger fraction at $z < 0.7$, which appears relatively constant from $z = 0.7$ to $z = 1.2$. We explore several fitting formalisms for parameterising the merger fraction, and compare our results to other structural studies and pair methods within the DEEP2, VVDS, and COSMOS fields.  We also examine the basic features of these galaxies, including our selection for mergers, and the inherent error budget and systematics associated with finding mergers through structure. We find that for galaxies selected by M$_{*} > 10^{10}$ \\solm, the merger fraction can be parameterised by $f_{m} = f_{0} \\times (1+z)^{m}$ with the power-law slope $m = 2.3\\pm0.4$.  By using the best available $z = 0$ prior the slope increases to $m = 3.8\\pm0.2$, showing how critical the measurement of local merger properties are for deriving the evolution of the merger fraction.  We furthermore show that the merger fraction derived through structure is roughly a factor of 3-6 higher than pair fractions. Based on the latest cosmological simulations of mergers we show that this ratio is predicted, and that both methods are likely tracing the merger fraction and rate properly.  We calculate, utilising merger time scales from simulations, and previously published merger fractions within the Hubble Deep and Ultra Deep Fields, that the merger rate of galaxies with M$_{*} > 10^{10}$ \\solm increases linearly between $z = 0.7$ and $z = 3$. Finally, we show that a typical galaxy with a stellar mass of M$_{*} > 10^{10}$ \\solm  undergoes between 1-2 major mergers at  $z < 1.2$. ",
        "introduction": "Unlike the case for stars, the relative size and separation between galaxies is small, suggesting that galaxy interactions and mergers could be a major process in driving galaxy formation.  Observational attempts to determine the role of galaxy mergers in the local universe have been limited to studies of how star formation and black hole growth are induced by the merger process (e.g.,  Barton et al. 2000; Ellison et al. 2008).  It is however clear that ongoing major galaxy merging is rare in the nearby universe, and is likely not a dominant mode for the future evolution of galaxies (e.g., Patton et al. 2000).    If galaxy mergers play a major role in galaxy formation they must have occurred earlier in the history of the universe.  Attempts to trace the merger history in the early universe has been carried out using two primary methods.  The oldest method involves measuring the fraction of galaxies which are in physical pairs at various look-back times, or redshifts (e.g., Patton et al. 2002; Le Fevre et al. 2000; Lin et al. 2004; Lin et al. 2008; Bluck et al. 2008). A newer method, with a considerable background history (e.g., Holmberg 1941; Vorontsov-Velyaminov 1959; Arp 1966), uses the structures of galaxies to determine the merger history (e.g., Conselice 2003; Conselice et al. 2003; Lotz et al. 2008).    It is, however, not yet known if the two methods are revealing the same merger history for galaxies, although at $z \\sim 0$, the agreement between the two methods appears good (e.g., De Propris et al. 2007).  The merger history of galaxies is still largely only beginning to be measured with some certainty.  Two issues that still need to be addressed are the reliability of the current methods for determining merger histories, and the relatively small areas thus far used to find mergers.  Both of these issues limit our ability to accurately measure the merger history. The reliability argument has been addressed in great detail in papers such as Conselice (2003), Conselice et al. (2003), Conselice et al. (2005), and Lotz et al. (2004). Briefly, the merger history can only be measured using special techniques and data, such as complete redshift surveys and/or deep Hubble Space Telescope (HST) imaging.    It turns out that when examining galaxies that fit a structural merger criteria, nearly all are ongoing major mergers, as judged by eye and through kinematics (e.g., Conselice 2000a,b; Conselice 2003; Conselice et al. 2008). This hold for low and high redshift galaxies (e.g., Conselice et al. 2003), and is also seen when examining N-body models of the merging process (Conselice 2006; Lotz et al. 2008b).  Using structure to determine the merger history requires the use of HST imaging or adaptive optics (AO).  However, the small fields of view used in HST and AO surveys make it difficult to measure the merger history with a high certainty, simply due to the small number of galaxies that have been examined.  Another issue is that it is difficult to measure the merger history at higher redshifts due to the fact that there is very little high resolution deep near-infrared imaging, which is needed to directly study the rest-frame optical structures of $z > 1.2$ galaxies.  This is required, as the rest-frame optical appearance of a galaxy often differs significantly from the rest-frame ultraviolet morphology (e.g., Windhorst et al. 2002; Taylor-Mager et al. 2007).  Ideally in the future, using multiple band data, it is desirable to measure parameters and structures on stellar mass images (e.g., Lanyon-Foster et al. 2009).  For our merger measurements, we utilise the CAS method (e.g., Conselice 2003) for determining the presence of galaxy mergers in the galaxy population at $z < 1$.  We first re-evaluate the use of the asymmetry index (Conselice et al. 2000a), and the CAS method itself for locating major mergers.  We furthermore use our derived merger fractions to characterise the major galaxy merger evolution for M$_{*} > 10^{10}$ \\solm galaxies at $z < 3$.  We discuss various ways in which this merger evolution can be parameterised, investigating power-law, exponential, and the combination of the two forms.    We conclude that the use of a $z = 0$ prior in fitting the merger fraction is a significant factor in the determination of the power-law slope $m$.  We show that this value of $m$ can vary between $m = 1.5$ and $m = 4$ depending on how the merger fraction is fit (cf. Bluck et al. 2008 for ultra massive galaxies).  In this paper we use the two largest HST Advanced Camera for Surveys (ACS) imaging data sets - the Extended Groth Strip (EGS) and the Cosmic Evolution Survey (COSMOS) to determine, using $> 20,000$ massive galaxies, the detailed merger history at $z < 1.2$.  This is an important epoch for measuring the merger history, and there is still considerable uncertainty regarding the merging history during this epoch, which traces the last half of the universe. By using ACS F814W band images we are also tracing the structures of these galaxies in the rest-frame optical out to $z \\sim 1$ and do not need to consider the sometimes considerable morphological $k-$corrections when imaging galaxies in the ultraviolet (e.g., Taylor-Mager et al. 2007).  Our general conclusion is that the merger fraction increases with higher redshifts to $z \\sim 1.2$. We use the latest model based time-scales for galaxies in pairs to merge, and for galaxies to remain symmetric during mergers, to calculate the number of mergers a typical M$_{*} > 10^{10}$ \\solm galaxy will undergo at $z < 3$, as well as the galaxy merger rate. We conclude that a typical  M$_{*} > 10^{10}$ \\solm galaxy will undergo 1-2 major mergers at $z < 1$. We also investigate how structural merger fractions compare with pair fractions, finding that at a given redshift, within the same population of galaxies, the CAS merger fraction is 3-6 times higher than the pair fraction.  We show that the ratio of the time-scales for CAS mergers and 20 kpc pairs is nearly the same as the merger-pair fraction ratio. Both methods therefore appear to trace the same merging population at different phases. This is further evidence that we are indeed measuring correctly the merging properties of galaxies.  This paper is organised as follows: \\S 2 includes a discussion of the data sources we use in this paper, and how we select our sample, including a description of our morphological and structural analyses, and the stellar masses we utilise, \\S 3 is a discussion of our results, including a description of the merger history up to $z \\sim 1.2$, and \\S 4 is our summary and conclusions.   We use a standard cosmology of H$_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, and $\\Omega_{\\rm m} = 1 - \\Omega_{\\lambda}$ = 0.3 throughout. ",
        "conclusions": "We have carried out an analysis of the structural CAS parameters for a sample of $>$ 20,000 galaxies with stellar masses M$_{*} > 10^{10}$ \\solm within the EGS and COSMOS fields, between $z = 0.2$ and $z = 1.2$.  We explore in this paper the merger fraction history, including various parameterisations, comparison of structural and pair merger fractions, the merger rate and role of mergers in galaxy formation, as well as systematics which are possibly playing a role in the derived merger fraction.  The primary results from this paper are:  \\noindent I.  We find through the CAS structural method that the merger fraction slightly increases from $z = 0.2$ to $z = 1.2$ with the merger fraction increasing from $f_{m} = 0.04\\pm0.01$ to $f_{m} = 0.13\\pm0.01$.  \\noindent II. We compare our derived merger fractions to previously published merger fractions from Lotz et al. (2008a) and pair studies from Kartaltepe et al. (2007), Lin et al. (2008) and de Ravel et al. (2008). Our results are comparable with the Gini/M$_{20}$ derived merger fractions from Lotz et al. (2008a), but we find our merger fractions are between 3-6 times higher than those derived from pairs, even within the same stellar mass selection.  We argue this in some detail through the use of a M$_{*} > 10^{11}$ \\solm sample of galaxies in the EGS.  We also show that this ratio of structural mergers and pair fractions is however predicted in the latest N-body models of galaxy mergers from Lotz et al. (2008b) and is due to differing merger time-scales.  \\noindent III. We investigate various methods, including the use of priors, for parameterising the merger history through the CAS method.  We show that the power-law formalism, whereby the merger fraction is parameterised by $f (z) = f_{0} \\times (1+z)^{m}$, is a poor fit to the merger fractions at $z < 1$.  We further show that the value of the power-law slope, $m$, can vary depending on whether a prior is used to set the local $z = 0$ merger fraction.  If we fit only our $z = 0.2-1.2$ merger fractions from COSMOS+EGS we fit a merger fraction evolution of $f(z) = 0.025\\pm0.005(1+z)^{2.3\\pm0.4}$. However, by using a prior of $f_0 = 0.009$ from de Propris et al. (2007), the slope of the power-law fit is $m = 3.8\\pm0.2$  \\noindent IV. We find that a combined power-law exponential is a better fit to the merger fraction at $z < 3$ than just a simple power-law.  We fit this form to the merger fraction using a $z = 0$ prior, and without, finding a similar form that predicts a turnover in the merger fraction for M$_* > 10^{10}$ \\solm galaxies at z = 2.  \\noindent V. We calculate the merger rate by using the latest measures of merger time-scales from merger simulations, and number densities for galaxies with M$_* > 10^{10}$ \\solm from Conselice et al. (2007b).  Although the errors on measuring the merger rate are large, resulting from uncertainties in the merger fraction, the merger time-scale, and the galaxy density, it is the ultimate quantity, and can reveal the role of mergers in forming galaxies.  We find that the merger rate is roughly constant up to $z = 0.7$, and increases at $z > 1$ by a factor of 10. We also show that the merger rate for  M$_* > 10^{10}$ \\solm galaxies can be parameterised up to $z \\sim 2$ as a linear increasing function.    We thank the staff at the Palomar and Keck observatories for invaluable assistance in collecting the data used for the Extended Groth Strip survey. Funding to support this effort came from a National Science Foundation Astronomy \\& Astrophysics Fellowship, and grants from the Science Technology Facilities Council (STFC). Support for the ACS imaging of the EGS in GO program 10134 was provided by NASA through NASA grant HST-G0-10134.13-A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555."
    },
    "0812/0812.1835_arXiv.txt": {
        "abstract": "Detailed models of galactic disk formation and evolution require knowledge about the initial conditions under which disk galaxies form, the boundary conditions that affect their secular evolution and the micro-physical processes that drive the multi-phase interstellar medium and regulate the star formation history. Unfortunately, up to now, most of these ingredients are still poorly understood. The challenge therefore is to, despite this caveat, construct realistic models of galactic disks with predictive power. This short review will summarize some problems related to numerical simulations of galactic disk formation and evolution. ",
        "introduction": "The radial surface density profiles of galactic disks are determined by the gravitational potential of the galaxy that is dominated in the outer parts by dark matter and the specific angular momentum distribution of the infalling gas that dissipates its potential and kinetic energy while settling into centrifugal equilibrium in the inner regions of a dark matter halo. In addition one has to consider the secular evolution of galactic disks. Viscous angular momentum redistribution and selective gas loss in galactic winds strongly affects the evolution of disks making it difficult to infer the initial conditions from their presently observed  structure.  Angular momentum is an important ingredient in order for galactic disks to form. It is generally assumed that, before collapse, gas and dark matter are well mixed and therefore acquire a similar specific angular momentum distribution (Peebles 1969; Fall \\& Efstathiou 1980; White 1984). If angular momentum would be conserved during gas infall, the resulting disk size should be directly related to the specific angular momentum $\\lambda'$ of the surrounding dark halo where (Bullock et al. 2001)  \\begin{equation} \\lambda' = \\frac{J}{\\sqrt{2} M_{vir} V_{vir} R_{vir}} \\end{equation}  \\noindent with $R_{vir}$ and $V_{vir}^2=GM_{vir}/R_{vir}$ the virial radius and virial velocity of the halo, respectively, and $M_{vir}$ its virial mass. Adopting a flat rotation curve, the disk scale length is (Mo et al. 1998; Burkert \\& D'Onghia 04)  \\begin{equation} R_d \\approx 8 \\left( \\frac{\\lambda'}{0.035} \\right) \\left( \\frac{v_{max}}{200 km/s} \\right) kpc \\end{equation}  \\noindent where $v_{max}$ is the maximum rotational velocity in the disk.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=3.4in]{fig1.eps} \\caption{The observed scale lengths versus the maximum rotational velocities of galactic disks are shown for the Courteau (1997) sample. The sold line shows the theoretically predicted correlation for $\\lambda' = 0.035$. The dashed curve corresponds to $\\lambda'=0.025$.} \\label{fig1} \\end{center} \\end{figure}  Figure 1 shows the correlation between the disk scale length $R_{disk}$ and the maximum rotational velocity $v_{max}$ for massive spiral galaxies (Courteau 1997) which is consistent with a mean value of $\\lambda' \\approx 0.025$. The observationally derived average peak value of $\\lambda' = 0.025$ is somewhat smaller than the theoretically predicted value of $\\lambda' = 0.035$,  indicating that the gas could on average have lost some amount of angular momentum during the phase of decoupling from the dark component and settling into the equatorial plane.  This result is promising. The situation is however far less satisfactory when we consider more detailed numerical models of gas infall and accumulation in galactic disks. Many simulations of galactic disk formation suffer from catastrophic angular momentum loss which leads to disks with unreasonably small scale lengths and surface densities that are too large. The origin of this problem has been attributed to a strong clumping of the infalling gas which looses angular momentum by dynamical friction within the surrounding dark matter halo (Navarro \\& Benz 1991, Navarro \\& Steinmetz 2000). Other possibilites are low numerical resolution (Governato et al. 2004, 2007), the effect of substantial and major mergers (d'Onghia et al. 2006) or artificial secular angular momentum transfer from the cold disk to its hot surrounding (Okamoto et al. 2003). Over the years many groups have tried to solve this problem by including star formation and energetic feedback (e.g. Sommer-Larsen et al. 2003, Abadi et al. 2003, Springel \\& Hernquist 2003, Robertson et al. 2004, Oppenheimer \\& Dave 2006, Dubois \\& Teyssier 2008). The results are however confusing. First of all the origin of the angular momentum problem is not clearly understood. Secondly, no reasonable, universally applicable feedback prescription has been found that would lead to the formation of large-sized, late-type disks, not only for one special case, but in general.  Progress in our understanding of the cosmological angular momentum problem has recently been achieved by Zavala et al. (2008) who confirmed that the specific angular momentum distribution of the disk forming material follows closely the angular momentum evolution of the dark matter halo. The dark matter angular momentum grows at early times as a result of large-scale tidal torques, consistent with the prediction of linear theory and remains constant after the epoch of maximum expansion. During this late phase however angular momentum is redistributed within the dark halo with the inner dark halo regions loosing up to 90\\% of their specific angular momentum to the outer parts. The process leading to this angular momentum redistribution is not discussed in details. It is however likely that substantial mergers with mass ratios less than 10:1 that are expected to occur frequently even at late phases during galaxy formation perturb the halo and generate global asymmetries in the mass distribution that are known to be an efficient mechanism for angular momentum transfer. Small satellite infall is probably of minor importance. It would be interesting to study the role of major and minor mergers in this process in greater details.  It is then likely that any gas residing in the inner regions during such an angular momentum redistribution will also loose most of its angular momentum, independent of whether the gas resides already in a protodisk, is still confined to dark matter substructures or is in an extended, diffuse distribution. Zavala et al. (2008) (see also Okamoto et al., 2005 and Scannapieco et al., 2008) show that efficient heating of the gas component can prevent angular momentum loss, probably because most of the gaseous component resides in the outer parts of the dark halo during its angular momentum redistribution phase. The gas would then actually gain angular momentum  rather than loose it and could lateron settle smoothly into an extended galactic disk in an ELS-like (Eggen, Lynden-Bell \\& Sandage 1962) accretion phase.  Unfortunately little is known about the energetic processes that could lead to such an evolution. Obviously, star formation must be delayed during the protogalactic collapse phase in order for the gas to have enough time to settle into the plane before condensing into stars. However star formation is also required in order to heat the gas preventing it from collapsing prior to the angular momentum redistribution phase. Scannapieco et al. (2008) show that their supernova feedback prescription is able to regulate star formation while at the same time pressurizing the gas. Their models are however still not efficient enough in order to produce disk-dominated, late-type galaxies. Large galactic disks are formed. The systems are however dominated by a central, massive, low-angular momentum stellar bulge component. This is in contradiction with observations which indicate a large fraction of massive disk galaxies with bulge-to disk ratios smaller than 50\\% (Weinzirl et al. 2008) that cannot be produced currently by numerical simulations of cosmological disk formation. ",
        "conclusions": "We are currently living in a very exciting time where the various complex processes that can affect galactic disk formation and evolution are being uncovered and studied observationally and theoretically. Combined with the now well established cold dark matter structure formation scenario the time seems ripe for self-consistent models of galaxy formation with predictive power. Given the high capacities of present-day supercomputers it is understandable that one tries to including as many processes as possible, most of which being however not well understood. These models not only suffer from a large number of free parameters. They also do not necessarily lead to insight as they are so complex and depend on so many different implemented physical aspects that it is impossible to clearly understand what in the end the origin of a certain result will be.  I wonder whether one needs a high complexity in order to understand important questions of galactic disk evolution, like the two KS laws, the origin of turbulence in the diffuse interstellar medium or in molecular clouds or the origin of the strong correlation between the viscous timescale and the star formation timescale.  Let us try to solve simple questions first before we focus on the complex puzzles that involve many processes that are not well understood yet."
    },
    "0812/0812.0002_arXiv.txt": {
        "abstract": "Using the linear theory of perturbations in General Relativity, we express a set of consistency relations that can be observationally tested with current and future large scale structure surveys. We then outline a stringent model-independent program to test gravity on cosmological scales. We illustrate the feasibility of such a program by jointly using several observables like peculiar velocities, galaxy clustering and weak gravitational lensing. After addressing possible observational or astrophysical caveats like galaxy bias and redshift uncertainties, we forecast in particular how well one can predict the lensing signal from a cosmic shear survey using an over-lapping galaxy survey. We finally discuss the specific physics probed this way and illustrate how $f(R)$ gravity models would fail such a test. ",
        "introduction": "To understand the origin of the accelerating expansion of the universe is one of the salient question in contemporary cosmology \\cite{Riess:2004nr,Astier:2005qq,Eisenstein:2005su,Cole:2005sx,Tegmark:2006az,Komatsu:2008hk,Kilbinger:2008gk}. It is commonly attributed to the existence of some extra unknown physics loosely labelled Dark Energy (DE) \\cite{Peebles:2002gy}. Alternatively, it might cast some doubts on the theoretical foundation of  current cosmology, that is the theory of General Relativity (GR). From a theorist point of view, as most efforts to construct self-consistent DE models within GR seem unsuccessful, there is no a priori reason to rule out a modification to gravity on cosmological scales as an explanation \\cite{Durrer:2008in}. From an observer point of view, whereas GR passes direct tests probing the Solar System scales ($10^{11}$m) down to the laboratory scales ($10^{-3}$m) \\cite{adelberger03}, testing gravity on cosmological scales is challenging and overly model dependant so far. In this work, we propose to remedy these issues by defining a set of simple model independent tests of GR on cosmological scales, although more complicated models of GR (clumping dark energy or interacting dark energy models) are not considered in this paper (the complete set of test including both exotic dark energy models is being prepared as the second step of testing gravity). The tests we devise rely on simple and testable consistency relations based on linear GR and make use of the appropriate combinations of various observables of large scale structures. Were any of these relations proved to be violated, it would be a clear sign of a break-down of GR and it would highly motivate theoretical work on alternatives to GR. Reversely, before any such conclusion can be reached, it requires an absolute demonstration that observational systematic effects are well under controlled. A firm believer in GR might look at these states as simple test of systematic effects. Note that throughout this work, we will assume the validity of standard cosmology, that is to say the validity of the Copernican principle \\cite{Quercellini:2008ty}. If this was not the case, then the DE problem could be solve without new physics and more specific tests would have to be designed \\cite{Goodman:1995dt,Caldwell:2007yu,Uzan:2008qp}.  Several alternative approaches to test GR have been proposed so far, some parametric, some non-parametric. The parametric approaches can be classified into several categories according to their goal: i) to separate the geometrical and growth signatures in the $w$ $\\&$ $w_a$ plane~\\cite{Ishak:2005zs,Song:2005gm,Wang:2007fsa}, ii) to define the probed $w(z)$ range in a model independant manner \\cite{Knox:2006fh,Mortonson:2008qy,SongDore08}, iii) to look for an anomalous linear growth rate using the $\\gamma$ parameter~\\cite{Linder:2005in,Linder:2007hg,Huterer:2006mva,Polarski:2007rr,DiPorto:2007ym,Mortonson:2008qy,Thomas:2008tp,Dent:2008ia}, iv) to parametrize the metric perturbations~\\cite{uzan01,Song:2006sa,Bertschinger:2006aw,Koivisto:2005mm,Amendola:2007rr,Uzan:2006mf,Hu:2007pj,Hu:2008zd,Zhao:2008bn,Bertschinger:2008zb,Daniel:2008et,Caldwell:2007cw,Dore:2007jh}. Alternatively, non-parametric approaches were developed by~\\cite{Jain:2007yk,Song:2008vm,Zhang:2007nk,Zhang:2008ba}.  In this paper, we advocate the use of a non-parametric approach -- based on a set of GR based consistency relations -- to test gravity on cosmological scales. We will use explicit alternative to GR for illustrative purpose only. To do so, we jointly use various observables, \\ie velocities, galaxy clustering and weak-gravitational lensing~\\cite{Jain:2007yk}. More precisely, by using a combination of the two first, we will be able to compare it to the later. In this work, extending the work of \\cite{Song:2008vm} we focus closely on potential observational biases. We propose in particular various ways to deal with the galaxy bias, redshift space distortion and spectroscopic and photometric redshifts.  We first begin by introducing in Sec.~\\ref{sec:conv} a general framework to test GR on cosmological scales.  We then detail how to build faithful probes of matter perturbations in Sec.~\\ref{sec:tracing_matter} before detailing one test of GR in Sec.\\ref{sec:test}. We conclude and discuss our findings in Sec.~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  In this paper, we advocate the use of consistency relations to test gravity on cosmological scales. We did so by using a combination of observables. This approach remains model independent and does not rely on any specific parametrization. We focused in particular on the joint use of galaxy surveys and weak-lensing observables. We showed how using the former to predict the latter, we can build a strong self-consistency test for GR. We also showed how large-scale redshift surveys can be extremely valuable in such an endeavor. The test we proposed seems particularly appealing since any weak-lensing survey is also by nature a galaxy clustering survey and we illustrated how even a photo-z survey can be used to build a strong self-consistency test. We thus expect this test to be applied in the near future to the CFHTLS survey \\footnote{{\\texttt http://www.cfht.hawaii.edu/Science/CFHLS/}} and others like DES \\footnote{{\\texttt https://www.darkenergysurvey.org/}}, the proposed JDEM \\footnote{{\\texttt http://jdem.gsfc.nasa.gov/}} and Euclid \\footnote{{\\texttt http://sci.esa.int/science-e/www/object/index.cfm?fobjectid=42266}} space missions.  We address in this work the key observational and astrophysical systematics like the galaxy bias and redshift uncertainties but many more survey specific ones would have to be carefully studied before any conclusion can be drawn from any particular data-sets, e.g. $n(z)$ errors, PSF correction and variation over large scales. It is quite likely that this kind of test will be ultimately limited by the level of control of systematics. In particular, we purposefully worked on linear scales. Non-linear effects in real and redshift space will limit the useful scales for such a program. If we assume that we can handle non-linear effects up to k=0.3 $h$/Mpc, we can see in the left panel of Fig.~\\ref{fig:contour_bias} that our program is still tractable. Besides, weak-gravitational lensing has now been measured in the linear regime where this test could be applied \\cite{Fu:2007qq}. Another unexplored sides of our work lie in the possible degeneracies with underlying cosmological model parameters but we leave this study for future work.  More positively, although, our survey parametrization is quite generic, these surveys could certainly be customized to enhance the discriminatory power of this test and others. We encourage future missions to include the feasibility of such tests  as an extra criteria in the optimization of their design.  Assuming that all systematics are well under control, we can wonder what new physics can be probed with this test. As an illustration, we plotted in Fig.~\\ref{fig:cl_wl_wl_error} the predictions for a $f(R)$ theory where a non-linear function $f$ of the Ricci scalar $R$ is added to the Einstein-Hilbert action \\cite{Song:2006ej}. We considered typical value of the dimensionless Compton wavelength parameter $B_0=1$ (see Eq. 17 in \\cite{Song:2006ej}). In $f(R)$ gravity models, the Poisson equation is modified so that $G_N \\rightarrow G_N/(1+df/dR)$ and a non-zero anisotropy stress is introduced to break the simplest GR model anisotropy condition $\\Phi=-\\Psi$. As visible in Fig.~\\ref{fig:cl_wl_wl_error}, the predicted $\\bar{C}_l^{dd}$ for a $f(R)$ gravity model using GR assumptions would differ from the observed one at a level well exceeding the various potential biases. This discrepancy originates from the GR assumption in Eq.~\\ref{recon_m} since any non-zero anisotropic stress invalidates the relation $2\\Phi=\\Phi-\\Psi$. However, the reduced mass scale affecting $f(R)$ theories modify the Poisson equation. It can thus be probed directly using the cross-correlation technique detailled in Sec.~\\ref{sec:cc_with_vel}. Note that since the difference between this $f(R)$ model and  GR increases when the redshift decreases, so does the discrepancy between the observed and predicted $C_l^{dd}$ for a $f(R)$ theory. From the Fig.~\\ref{fig:alpha_error} it can be seen that constraints on $B_0$ of order $1$ or less are in principle feasible with such a test. Other aspects of $f(R)$ theories are discussed in \\cite{Cooney:2008wk,Gannouji:2008wt,Nojiri:2008nt}.   More generally, it can be argued that any theory with non-minimal interaction in the dark sector or dark energy clumping would fail this test since it either modifies the Poisson equation or introduce an anisotropic stress \\cite{Amendola:2007rr,Kunz:2006ca}. Reversely, it is important to note that some types of modified gravity models would pass the test we detailed here. For example, the potentials following from a DGP model \\cite{dvali00} are modified in such a way that the light path tracing potential differences is similar to the standard $\\Lambda$CDM predictions. To detect such models, we need to implement the program proposed in~\\cite{Song:2008vm}.  The tests presented above are only two of examples of the possible one highlighted in Sec.~\\ref{sec:conv} and we plan to extend our studies to include other consistency relations as well as other observables like the integrated Sachs-Wolfe effect \\cite{Sachs:1967er} and the lensing of the cosmic microwave background. We believe however that the observational program sketched in this paper already offers a stringent observational test of GR and we thus plan to put it into action using current data."
    },
    "0812/0812.2836_arXiv.txt": {
        "abstract": "We examine the angular momentum loss and associated rotational spin-down for magnetic hot stars with a line-driven stellar wind and a rotation-aligned dipole magnetic field. Our analysis here is based on our previous 2-D numerical MHD simulation study that examines the interplay among wind, field, and rotation as a function of two dimensionless parameters, W(=Vrot/Vorb) and 'wind magnetic confinement', $\\eta_\\ast$ defined below. We compare and contrast the 2-D, time variable angular momentum loss of this dipole model of a hot-star wind with the classical 1-D steady-state analysis by Weber and Davis (WD), who used an idealized monopole field to model the angular momentum loss in the solar wind. Despite the differences, we find that the total angular momentum loss averaged over both solid angle and time follows closely the general WD scaling $\\dot {J} \\sim \\dot {M} \\Omega R_A^2$. The key distinction is that for a dipole field Alfv\\`en radius $R_A$ is significantly smaller than for the monopole field WD used in their analyses. This leads to a slower stellar spin-down for the dipole field with typical spin-down times of order 1~Myr for several known magnetic massive stars. ",
        "introduction": "An outflowing wind carries away angular momentum and thus spins down the stellar rotation. Winds with magnetic fields exert a braking torque that is significantly larger than for non-magnetic cases, due to the larger lever arm of magnetic field lines that extend outward from the stellar surface. A seminal analysis of this process was carried out by Weber \\& Davis (WD, 1967), who modelled the angular momentum loss of the solar wind for the idealized case of a simple monopole magnetic field from the solar surface. In terms of the surface angular velocity $\\Omega$ and wind mass loss rate $\\dot M$, Weber \\& Davis concluded that the total angular momentum loss rate scales as: \\begin {equation} \\dot J= \\frac {2}{3} \\dot M \\Omega R_A^2 \\, , \\label{jdot-eqn} \\end {equation} with $R_A$ the Alfv\\'en radius,defined by where the radial components of the field and flow have equal energy density. However, WD did not provide a prescription for computing such an Alfv\\'en radius even for the idealized monopole field geometry.  For any radius $r$, the energy density ratio between radial field and flow is given by \\begin{equation} \\eta(r) \\equiv \\frac{B_{r}^{2}/8\\pi}{\\rho v_{r}^{2}/2} \\, . \\label{eta-eqn} \\end{equation} The Alfv\\'{e}n radius is then defined implicitly by $\\eta(R_{A})\\equiv 1$. We can derive approximate {\\em explicit} expressions in terms of fixed  values for the equatorial field strength $B_{eq}$ at the surface radius $R_{\\ast}$, and for the wind mass loss rate ${\\dot M}$ and terminal flow speed $v_\\infty$. Specifically, following ud-Doula \\& Owocki (paper 1, 2002) and ud-Doula et al. (paper 2, 2008a), if we define here a wind {\\em magnetic confinement parameter}, \\begin{equation} \\eta_{\\ast} \\equiv \\frac {\\Beq^2 \\, \\Rstar^2} {\\dot{M} \\, \\vinf} \\, , \\label{esdef} \\end{equation} then we can write the energy density ratio in the form \\begin{equation} \\eta(r) =  \\eta_{\\ast} \\, \\left [ \\frac{r}{\\Rstar} \\right ]^{2-2q} \\, \\frac{\\vinf}{v_{r}(r)} \\end{equation} where $q$ is the power-law exponent for radial decline of the assumed magnetic field and terminal speed $v_\\infty$. It turns out that in the strong magnetic confinement limit $\\eta_{\\ast} \\gg 1$, for a monopole field $R_{A} \\sim \\sqrt{\\eta_{\\ast}}$ whereas for a dipole field the scaling is significantly weaker $R_{A} \\sim {\\eta_{\\ast}}^{1/4}$. ",
        "conclusions": ""
    },
    "0812/0812.2007_arXiv.txt": {
        "abstract": "Photometry data were collected from the literature and analyzed for supernovae that are thought to have a gamma--ray burst association. There are several gamma--ray bursts afterglow light curves that appear to have a supernova component. For these light curves, the supernova component was extracted and analyzed. A supernova light curve model was used to help determine the peak absolute magnitudes as well as estimates for the kinetic energy, ejected mass and nickel mass in the explosion. The peak absolute magnitudes are, on average, brighter than those of similar supernovae (stripped--envelope supernovae) that do not have a gamma--ray burst association, but this can easily be due to a selection effect.  However, the kinetic energies and ejected masses were found to be considerably higher, on average, than those of similar supernovae without a gamma--ray burst association. ",
        "introduction": "A significant effort was made, over a number of years, to find the origin of gamma--ray bursts (GRBs). The discovery of optical afterglows in long--duration GRBs allowed for the determination of the GRB's redshift \\citep[for a review, see][and references therein]{vanparadijs00}. This led to the realization that long--duration GRBs are at cosmological distances \\citep[e.g.,][]{vanparadijs97,metzger97}. In this paper, long--duration, cosmological GRBs will be referred to simply as GRBs.  The light curve of the optical afterglow of GRB~980326 showed a rebrightening at around two weeks after the burst \\citep{bloom99}. This rebrightening, or bump, in the light curve was attributed to an underlying supernova (SN). This was the first observational indication of an association between GRBs and SNe.  During April of 1998, an unusual gamma--ray burst, GRB~980425, was discovered with its gamma--ray energy several orders of magnitude lower than typical GRBs \\citep{soffitta98,woosley99}. Optical observations of the region led to the discovery of an unusual supernova, SN~1998bw, within the error radius of the GRB \\citep{galama98}. Spectra of SN~1998bw showed it to be a Type~Ic SN with exceptionally broad lines (Ic--BL). This SN was also unusually bright for a Type~Ic SN \\citep{richardson06}. It was thought that GRB~980425 might simply be a normal GRB, but one where the jet was viewed off axis \\citep[e.g.][]{nakamura99,ioka01,yamazaki03}. However, according to late--time radio observations \\citep{soderberg04}, even if it was viewed off--axis it would still be an unusual GRB. There may be more dim, peculiar GRBs, like GRB~980425, than observations would indicate; due to the fact that those at large redshifts are difficult to detect.  The afterglow spectra of GRB~030329 were remarkably similar to that of SN~1998bw \\citep{stanek03}. The SN associated with this GRB was then named SN~2003dh. From this, it was clear that at least some long--duration gamma--ray bursts are connected to supernovae. Since then, two other SNe~Ic--BL have been spectroscopically confirmed to have a GRB connection: SN~2003lw/GRB~031203 \\citep{malesani04} and SN~2006aj/GRB~060218 \\citep{sollerman06,mazzali06b,modjaz06,mirabal06}. Not all SNe~Ic--BL have a convincing GRB connection; as is the case for SN~1997ef \\citep{iwamoto00} and SN~2002ap \\citep{mazzali02}. However, this does not rule out the possibility of an association with a weak (possibly off--axis) GRB. It also seems that there are some long--duration gamma--ray bursts with no apparent associated supernova; such as GRBs 060505 and 060614 \\citep{fynbo06}.  Similar to the  bump found in the optical afterglow light curve of GRB~980326 (mentioned above), there have been several other GRB afterglow light curves with just such a bump \\citep[e.g.,][]{reichart99,galama00,sahu00,castrotirado01}. This provides us with a number of possible GRB/SNe to consider, even without spectroscopic confirmation. These light curves are typically analyzed using a composite model \\citep[e.g.][]{zeh04, bloom04, stanek05}, discussed below in Section~5. Most of the light curves considered here are of this type. For alternative explanations for this rebrightening see \\cite{fynbo04} and references therein. For a discussion on the possible progenitors of long--duration GRBs, see \\cite{fryer07}. The observed data are discussed in Section~2. Section~3 covers the analysis of the light curves while the results are presented in Section~4. The conclusions are discussed in Section~5. ",
        "conclusions": "The average peak absolute magnitude was found to be higher for the GRB/SNe than for the SE SNe. However, in view of possible selection effects, the difference may not be significant. There were two GRB/SNe at the dim end of the distribution, but well within the distribution of SE SNe.  Most of the SN data analyzed here were taken from GRB afterglow light curves. The GRB light was treated as being independent of the SN light and was removed so that the resulting light curve could be analyzed as a SN light curve. These resulting light curves fit quite well with a SN light curve model (Figure~4).  GRB afterglow light curves that are suspected of having a SN component are usually analyzed by a different method. Usually the observed data are fit to a composite model, where all of the components are represented; GRB, SN and host galaxy \\citep[for example]{zeh04}. The assumptions made are similar to those made with the separation method of this paper. The main difference between the two methods is in the way the SN is treated. The composite method starts with a light curve of SN~1998bw (adjusted for the GRBs redshift). Two of the five free parameters in the composite model are then used in the overall fit to describe the SN; one for the brightness and one for the width of the SN contribution. While the separation method uses SN~1998bw for K--corrections, a general SE SN model is then fit to the resulting light curve. The two methods are similar; however, the separation method used here allows for closer analysis of the the SN. Reasonable estimates of the kinetic energy, ejected mass, nickel mass and rise times are found. However, the values of kinetic energy and ejected mass from other studies \\citep{nomoto06} tend to be larger by a factor of approximately two.  The coincidence of the GRB date and the derived date of the SN explosion is another point of interest. In all but two cases the difference between the two dates was less than a week. For GRB~050525A, the SN explosion date is estimated to have occurred about 10 days before the GRB date, but by looking at the model fit and the observed data in Figure~4, we see that the model was not able to simultaneously reproduce the narrow peak and the bright, late--time data point. It appears that a more accurate explosion date for the SN would bring it closer to the GRB date. For GRB~011121, the SN explosion date is estimated to have occurred about 9 days before the GRB date. Other studies have given similar results \\citep{bloom02}, but the lack of pre--peak data for this SN makes the SN explosion date difficult to accurately pin down. Thus there is no clear evidence for real differences between the times of the GRBs and the SNe.  The nickel masses determined here range from 0.05 to 0.99 M$_{\\odot}$. These are similar to the values found for SE SNe with no GRB association \\citep{richardson06}.  However, the kinetic energy values determined here range from about one to 31 foe and the ejected mass values range from about one to six M$_{\\odot}$.  These values are, on average, considerably higher than those found for SE SNe with no GRB association."
    },
    "0812/0812.2211.txt": {
        "abstract": "In this paper we show how the covariant gauge invariant equations for the evolution of scalar, vector and tensor perturbations for a generic $f(R)$-gravity theory can be recast in order to exploit the power of dynamical system methodology. In this way, recent results describing the dynamics of the background FRW model can be easily combined with these equations to reveal important details pertaining to the evolution of cosmological models in fourth order gravity. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The observational evidence for the accelerated expansion rate of the universe, and the introduction of the concept of Dark Energy has put theoretical cosmology into crisis. This is due to the fact that despite an increasing amount and quality of data, no model has been proposed thus far that is able to give a completely satisfactory theoretical explanation of these observations.  Among the many different ways to achieve cosmic acceleration, the modification of gravity and in particular higher order gravity has recently gained much attention \\cite{revnostra,Odintsov, Carroll,star2007,Sotiriou:2008rp}. The reason for this popularity is due to the fact that these models provide a somewhat more natural explanation of the cosmic acceleration: this effect is due to corrections to Einstein gravity which are directly related to the characteristic properties of the gravitational interaction. Most investigations of higher order gravity have focused on Fourth Order Gravity (FOG), i.e., on gravitational Lagrangian's in which the corrections are at most of order four in the metric and in what follows we will also focus on these models.  Because the field equations resulting from FOG are extremely complicated,  difficult conceptual and technical issues arise which need to be resolved in order to uncover the detailed physics of these models. Consequently it is important  to develop new methods which are able to assist in resolving these problems. Two such approaches, the {\\it dynamical systems approach to cosmology} and the {\\it covariant approach to cosmological perturbations}, have proved particularly useful in this respect.  The dynamical system approach, first developed by Collins \\cite{collins} and extensively reviewed in the book by Ellis and Wainwright  \\cite{ellisbook}  has proved to be an important tool in the understanding of cosmology of $f(R)$-gravity models. In fact, studying cosmological models using these techniques has the advantage of providing a relatively simple method for obtaining exact solutions, which appear as fixed points of the system and obtaining a global picture of the dynamics of these models. Consequently, such an analysis allows for an efficient preliminary investigation of these theories, allowing one to identify specific models which merit further investigation. In a series of recent  papers, a  wide range of features of  $f(R)$ cosmology have been presented, ranging from an analysis of the standard Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) models \\cite{Amendola,GenDynSys,RnGrav} to a discussion of the properties of the Einstein universe and the isotropization of Bianchi models \\cite{Leach:2006br, Barrow:2006xb,Goheer:2007wu}.  The 1+3 covariant approach was developed originally to study the evolution of linear perturbations of FRW models in General Relativity in\\cite{EllisCovariant,EB,EBH,BDE,DBE,BED,DBBE,conserved} with great success. For our purposes this approach has two major advantages. Firstly,  a specific recasting of the field equations allows one to easily make extensions to a wide range of modified gravity theories including Braneworlds \\cite{brane} and FOG \\cite{SantePertSca,StructForm}. Secondly the structure of the formalism is such that,  unlike other approaches , it is possible to keep track of the physical meaning of the equations at any stage of the calculations, which can be crucial when one modifies the theory of gravity.  In particular, in a number of recent papers \\cite{SantePertSca,StructForm,PRL} we derived the evolution equations for scalar and tensor perturbations of a subclass of fourth order theories of gravity characterized by an action which is a general analytic function of the Ricci scalar. The results obtained in \\cite{SantePertSca,StructForm,PRL} clearly demonstrated that the evolution of scalar perturbations is determined by a fourth order differential equation rather than a second order one. This implies that the evolution of the density fluctuations contains, in general, four modes rather than two and can give rise to a more complex evolution than what is obtained in General Relativity. It was also found that the perturbations depend on the scale for any value of the equation of state parameter of standard matter (while in GR the evolution of the dust perturbations are not scale dependent) and that there is a characteristic scale-dependent signature in the matter power spectrum \\cite{PRL}. This means that in FOG the evolution of super-horizon and sub-horizon  perturbations are different. It also turns out that the growth of large density fluctuations can occur also in backgrounds in which the expansion rate is increasing in time. This surprising result is strikingly different with what one finds in GR and could provide a strong constraint on some FOG theories using the ISW effect. In addition to that the structure of the general fourth order perturbation equations and the analysis of scalar perturbations lead to the discovery of a characteristic signature of fourth order gravity in the matter power spectrum \\cite{StructForm}, the details of which have not seen before in other works in this area. This could provide a crucial test for fourth order gravity on cosmological scales.  The aim of this paper is to combine dynamical systems methods with linear perturbation theory in such a manner that one is able to apply directly the results coming from the former to the latter and is able to gain a semi-qualitative idea of both the behavior of the background and that of the first order perturbations in a general $f(R)$-gravity theory. In order to achieve this we will express the coefficients of the perturbation equations in terms of the dynamical system variables in such a way that at any fixed point it will correspond a set of perturbation equations and as a consequence an evolution law for the linear fluctuations.  The paper is divided as follow. In Section 2 we give some basic equations. In Section 3 we present a formalism that allows one to investigate a large class of  $f(R)$-gravity models. In Section 4 we present the equations for the evolution of the scalar, vector and tensor perturbations in a generic $f(R)$-gravity theory written in terms of the dynamical system variables. In Section 5 we consider two simple examples. Finally, Section 6 is devoted to the conclusions.  Unless otherwise specified, natural units ($\\hbar=c=k_{B}=8\\pi G=1$) will be used throughout this paper, Latin indices run from 0 to 3. The symbol $\\nabla$ represents the usual covariant derivative and $\\partial$ corresponds to partial differentiation. We use the $-,+,+,+$ signature and the Riemann tensor is defined by \\begin{equation} R^{a}{}_{bcd}=W^a{}_{bd,c}-W^a{}_{bc,d}+ W^e{}_{bd}W^a{}_{ce}- W^f{}_{bc}W^a{}_{df}\\;, \\end{equation} where the $W^a{}_{bd}$ are the Christoffel symbols (i.e. symmetric in the lower indices), defined by \\begin{equation} W^a_{bd}=\\frac{1}{2}g^{ae} \\left(g_{be,d}+g_{ed,b}-g_{bd,e}\\right)\\;. \\end{equation} The Ricci tensor is obtained by contracting the {\\em first} and the {\\em third} indices \\begin{equation}\\label{Ricci} R_{ab}=g^{cd}R_{acbd}\\;. \\end{equation} Finally the Hilbert--Einstein action in the presence of matter is given by \\begin{equation} {\\cal A}=\\int d x^{4} \\sqrt{-g}\\left[\\frac{1}{2}R+ L_{m}\\right]\\;. \\end{equation}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  In this paper we have discussed the connection between the dynamical system approach and the covariant-gauge invariant theory of perturbations, presenting a method to calculate directly the evolution of the scalar, vector and tensor perturbations at a fixed point of the phase space of a generic $f(R)$-gravity theory. Within the limitations of the dynamical system approach one is then able to obtain an idea of the evolution of the perturbations in any $f(R)$ model.  Because of the non-linearity and the peculiar structure of the dynamical system formalism, the concept of fixed points in $f(R)$-gravity is more subtle than the one of GR. In particular, one can have fixed points which do not correspond to solutions which satisfy the cosmological equations. This is due to the fact that the fixed point conditions \\rf{FPcond} constitute additional constraints that can be incompatible with the cosmological equations and that the exact solution associated to the fixed points seems to depend more on the definition of the variables than the model itself.  This has profound implications on the interpretation of the results of the dynamical system approach. Specifically it suggests that this kind of approach offers only partial information on the actual evolution of the cosmology in this framework and that this information depends on the formalism used. In particular (i) one might not be able to see all the fixed points in the phase space because of the form of the variables  and (ii) some of these fixed points might not correspond to actual solutions of the system with obvious problems in the interpretations of the orbits which have these points as attractors. As a consequence one has to be very careful in using these tools to derive general conclusions about the dynamics for these cosmological models.  These peculiarities also have consequences on the results of the perturbations equations. In particular, we found that the evolution laws obtained using this form of the equations do not in general coincide with those that one would obtain by simply subsittuting in the background corresponding to the fixed point in the dynamical system equations. This can be explained when one considers that the additional conditions associated with the fixed points, further constrain the perturbation equations and therefore lead to different results. On the other hand such a fact also implies that one can associate an evolution law for first order perturbations to a specific fixed point,  which  in some sense may be  considered  as  a ``fixed point\" for the perturbation theory. Thus one can use the results of the equations given above in the same way in which one used the solution at the fixed points: gaining qualitative information about the behavior of the perturbations along an orbit. Such a feature, confirmed by the direct calculations presented in \\cite{StructForm}, can help with the understanding of the behavior of the perturbations in models for which a direct numerical integration is to complex or to resource intensive to be performed. It is also worth noticing that since the peculiarities in the predictions of the equations presented above derive ultimately from the additional condition on the fixed points, they only apply when one deals with fixed points and will not be present when one considers the dynamical system variables as functions of time.  In conclusion, in spite of these difficulties the dynamical system approach, when combined with the covariant gauge invariant formalism, represents an extremely powerful tool for the study of $f(R)$-gravity cosmological models. We believe that a careful use of these methods could be invaluable in determining the relevance of these models on cosmological scales and to constrain them using current and future observations."
    },
    "0812/0812.0366_arXiv.txt": {
        "abstract": "The silhouette cast by the horizon of the supermassive black hole in M87 can now be resolved with the emerging millimeter very-long baseline interferometry (VLBI) capability.  Despite being $\\sim 2\\times 10^3$ times farther away than Sgr A* (the supermassive black hole at the center of the Milky-Way and the primary target for horizon-scale imaging), M87's much larger black hole mass results in a horizon angular scale roughly half that of Sgr A*'s, providing another practical target for direct imaging.  However, unlike Sgr A*, M87 exhibits a powerful radio jet, providing an opportunity to study jet formation physics on horizon scales.  We employ a simple, qualitatively correct force-free jet model to explore the expected high-resolution images of M87 at wavelengths of $1.3\\,\\mm$ and $0.87\\,\\mm$ ($230\\,\\GHz$ and $345\\,\\GHz$), for a variety of jet parameters.  We show that future VLBI data will be able to constrain the size of the jet footprint, the jet collimation rate, and the black hole spin.  Polarization will further probe the structure of the jet's magnetic field and its effect on the emitting gas.  Horizon-scale imaging of M87 and Sgr A* will enable for the first time the empirical exploration of the relationship between the mass and spin of a black hole and the characteristics of the gas inflow/outflow around it. ",
        "introduction": "\\label{sec:intro} With the advent of millimeter VLBI, for the first time it has become possible to resolve the horizon of a black hole \\citep{Doel_etal:08,Doel:08}. Most of the theoretical and observational efforts on this front have been directed towards imaging Sagittarius A* (Sgr A*), the radio source associated with the supermassive black hole at the center of the Milky Way \\citep[see, e.g.,][]{Falc-Meli-Agol:00,Brod-Loeb:05,Brod-Loeb:06a,Brod-Loeb:06b,Miyo_etal:08,Doel_etal:08}. This is because Sgr A* subtends the largest angle of any known black hole candidate (roughly $55\\,\\muas$ in diameter).  Nevertheless, there is at least one other promising target for millimeter VLBI imaging, namely M87's supermassive black hole (hereafter labeled `M87' for brevity).  At a distance of $16\\,\\Mpc$ and with a black hole mass of $3.4\\times10^9\\,\\Ms$, the apparent diameter of M87's horizon is approximately $22\\,\\muas$, appearing about half as large as Sgr A*. Like Sgr A*, M87 is many orders of magnitude under-luminous in comparison to its Eddington luminosity.  However, M87 is different from Sgr A* in five important respects.  First, M87 is located in the Northern sky, and thus is more amenable to imaging by telescope arrays comprised of existing millimeter and submillimeter observatories, which exist primarily in the Northern hemisphere.  Thus, observations of M87 will generally have lower atmospheric columns and longer observing opportunities.  Arrays of existing telescopes (see \\S\\ref{sec:PT} for details), spanning projected baselines of $10^4\\,\\km$, are already sufficient to resolve the silhouette at $1.3\\,\\mm$ ($230\\,\\GHz$) and $0.87\\,\\mm$ ($345\\,\\GHz$) with beam sizes of $17\\,\\muas$ and $11\\,\\muas$, respectively, along the minor axis.  The inclusion of new facilities over the next few years will improve baseline coverage substantially, resulting in a roughly circular beam with these resolutions.  Second, M87's mass implies a dynamical timescale of roughly $5\\,\\hr$, and a typical orbital timescale at the innermost-stable orbit of 2-18 days, depending upon the black hole spin.  Thus, unlike Sgr A*, the structure of M87 may be treated as fixed during an entire day, one of the underlying assumptions of Earth aperture synthesis. This means that it will be possible to produce sequences of images of dynamical events in M87, occuring over many days.  Third, M87 exhibits a powerful jet, providing a unique opportunity to study the formation of relativistic jets on horizon scales. Realistic {\\em ab initio} computations of jet formation are not presently possible, requiring jet simulations to make a variety of assumptions about the underlying plasma processes.  Examples include assumptions about the dissipation scale, cooling rate, jet mass-loading mechanism, acceleration of the observed non-thermal electrons and even the applicability of the magnetohydrodynamic (MHD) prescription \\citep{DeVi-Hawl-Krol:03,McKi-Gamm:04,Komi:05,McKi:06,Hawl-Krol:06,Tche-McKi-Nara:08,Igum:08,McKi-Blan:08,Nish_etal:05}. Thus, horizon-resolving images of M87's jet will critically inform our understanding of the formation, structure and dynamics of ultra-relativistic outflows.  Fourth, unlike Sgr A*, one is primarily driven towards submillimeter wavelengths by resolution requirements.  There is no evidence for an analogous interstellar scattering screen that would blur high resolution images  of M87 \\citep[cf.][]{Bowe_etal:06}. Nor is there evidence that M87 is optically thick at or about millimeter wavelengths \\citep[cf.][]{Brod-Loeb:06a,Brod-Loeb:06b}. Thus, modeling and interpreting images of M87 is likely to be somewhat more straightforward.  Fifth, despite being under-luminous, M87 is similar to other radio-loud quasars.  The same cannot be said of Sgr A*, which is nearly dormant and is visible only due to its extraordinary proximity. As such, the faint low-mass black hole in Sgr A* is much more representative of the population of black holes that are accessible to the {\\em Laser Interferometer Space Antenna} (LISA \\footnote{http://lisa.nasa.gov/}).  Careful comparison between horizon-scale imaging of Sgr A* and M87 will be crucial as a touchstone for comparing the astrophysical properties of these two populations, and how the black hole mass and spin affect the electromagnetic and kinetic properties of the surrounding plasma.  It is not surprising, therefore, that already high-resolution images of M87 has been aggressively pursued.  The best known of these are the $7\\,\\mm$ VLBI images obtained by \\citet{Juno-Bire-Livi:99} \\citep[and more recently by][]{Ly-Walk-Wrob:04,Walk-Ly-Juno-Hard:08}.  $3\\,\\mm$ and $2\\,\\mm$ images have also been reported \\citep{Kric_etal:06,Kova-List-Homa-Kell:07}.  These represent the highest resolution observations of a relativistic jet ever produced, probing down to roughly 5 Schwarzschild radii.  Nevertheless, they are presently insufficient to directly address many outstanding questions in jet-formation theory.  Among these questions are the role of the black hole spin in jet production, the structure of the jet forming region, the importance of magnetic fields and the particle content of the jet itself.  However, all of these uncertainties may be settled by pushing to submillimeter wavelengths.  In this sense, millimeter imaging promises to provide a qualitative change, not simply an incremental improvement.  The reason for this is simple: the minimum relevant physical scale for a black hole of mass $M$ is $GM/c^2$, and thus resolving this scale provides full access to the physical processes responsible for jet formation and accretion. In addition, the observation of the characteristic silhouette associated with the black hole settle the present ambiguity in indentifying its position relative to the observed radio core \\citep{Marc_etal:08}.  Motivated by the above arguments, we present in this paper predicted theoretical images of M87 for a class of force-free jet models which fit all existing observations.  In particular, we assess the ability of possible millimeter \\& submillimeter VLBI experiments to distinguish different critical jet parameters.  In \\S\\ref{sec:PT} we review how a millimeter array can be created out of existing telescopes.  In \\S\\ref{sec:MM} we describe the jet-disk model we employ for M87.  In \\S\\ref{sec:JI} and \\S\\ref{sec:JP} we present images and polarization maps, respectively, for a variety of jet models.  Finally, \\S\\ref{sec:C} summarizes our main conclusions. ",
        "conclusions": "\\label{sec:C}  M87 provides a promising second target for the emerging millimeter and submillimeter VLBI capability.  Its presence in the Northern sky simplifies its observation and results in better baseline coverage than available for Sgr A*.  In addition, its large black-hole mass, and correspondingly long dynamical timescale, makes possible the use of Earth aperture synthesis, even during periods of substantial variability.  Despite being dominated by a jet at millimeter wavelengths, the black hole silhouette, cast against the base of the counter-jet, is still clearly present, and may just be resolved at $0.87\\,\\mm$ ($345\\,\\GHz$).  This requires that both the jet and the counter-jet are moving slowly near the black hole, and thus relativistic beaming does not preclude its observation.  If, however, the counter-jet is either relativistic even at its base or simply not present the silhouette will not be detectable.  Nonetheless, obtaining a sub-mm image of M87 appears to be more promising than efforts to image the silhouette of Sgr A* due to the lower opacity of the surrounding synchrotron emitting gas.  Unfortunately, it is unlikely that the silhouette alone will be sufficient to constrain the spin without additional information.  Better constraints may originate from exploring the effect of spin upon the jet image (via a change in the angular velocity profile of the jet footprint) or by carefully studying variability of hot spots in the accretion flow and jetted outflow. The latter possibility is feasible owing to the low optical depth to synchrotron self-absorption and the long dynamical timescales in M87.  Direct imaging of M87's jet on horizon scales will necessarily provide a wealth of information crucial to modern numerical efforts to model jet formation and propagation.  In particular, such observations will be sensitive to the jet footprint size, collimation rate, and (if the jet is truly a magnetically dominated outflow) black hole spin. Polarized imaging will further shed light upon the structure of the jet's magnetic field and its relation to the emitting gas.  Less well constrained will be the jet orientation.  Finally, comparisons of M87 and Sgr A* will provide the first opportunity to relate the gravitational properties of quasars, about which much is known, to the lower-mass, dormant black holes that LISA will access.  Ultimately, such measurements will associate the black hole demographics measured by LISA with the existing, extensively studied set of known quasars."
    },
    "0812/0812.2363_arXiv.txt": {
        "abstract": "BP~Vulpeculae is a bright eclipsing binary system showing apsidal motion. It was found in an earlier study that it shows retrograde apsidal motion which contradicts theory. In this paper we present the first $BV$ light curve of the system and its light curve solution as well as seven new times of the minima from the years 1959-1963. This way we could expanded the baseline of the investigation to five decades. Based on this longer baseline we concluded that the apsidal motion is prograde agreeing with the theoretical expectations and its period is about 365 years and the determined internal  structure constant is close to the theoretically expected one. ",
        "introduction": "The eclipsing nature of the $10^{th}$ magnitude star BP Vulpeculae was discovered by \\citet{illesalmar60} and she gave a period value of $1.938$ days which was slightly corrected by \\citet{huth65}. Then the star was neglected until the end of the 20th century when \\citet{lacy92} published $UBV$ colours of the system at certain phases. {Later \\citet{lacyetal03} presented more than 5000 data points in one colour ($V$) and they have solved their light curve and determined very accurately absolute dimensions of the system. This $V$ light curve was again solved by a simple, but automatized code \\citep{devor05} and the result was similar to that of \\citet{lacyetal03}.  The system has an age of about 1 Gyr consisting of A7m V + F2m V spectral type components and is slightly eccentric ($e=0.0345$, \\citealp{lacyetal03}). In eccentric binary star systems the apsidal line is revolving which is called apsidal motion and therefore the time difference between a primary minimum and the subsequent secondary minimum is variable. The apsidal motion is usually caused by two effects: the tidal forces of the components and the effects of general relativity. In the case of a more or less ordinary eccentric binary, i.e. no very close third companion in the system, and no extreme stellar rotation, both phenomena predict that the apside shows a prograde motion. The total apsidal motion is a sum of the classical and relativistic contributions. It is worthy to mention that there are several systems, like DI~Herculis or AS~Camelopardalis where the observed apsidal motions highly differ from the theoretically predicted ones \\citep[see e.g.][and references therein]{claret98}. It seems that the mentioned other effects, like third bodies etc. cause these pecularities \\citep[for an overview see][]{borkovitsetal07}.  Regarding the case of BP~Vul, \\citet{lacy03} established a very short apsidal motion period ($77\\pm22$yr) and they found that it is a retrograde one. If BP~Vul really would have had a retrograde apsidal motion it would be a new representative of the systems which confronts theory and requires further study because a retrograde apsidal motion is in contradiction with theory. But the observational window in the work of \\citet{lacy03} was about one decade only therefore this rough estimation for the period of apsidal motion should be refined. For this refinement we present $BV$ observations of BP~Vul which were obtained 45-49 years ago and were not published until now. ",
        "conclusions": "\\label{Summary}  BP~Vulpeculae is an eccentric eclipsing binary star showing the so-called periastron-precession effect. \\citet{lacyetal03} concluded that this effect causes a retrograde motion of the semi-major axis and it has a period of $77\\pm22$ years based on their about one decade long observational material. However, retrograde motion contradicts theory. Their explanation was that a possible third body in the system could perturb the orbit yielding the observed peculiar periastron precession. Nevertheless no spectroscopic evidence was found by them for such a third body.  BP~Vul was observed at the Konkoly Observatory more than forty years before the work of \\citet{lacyetal03} by one of the authors of this study. This observational material allowed us to determine the times of six primary and one secondary minima. With these early observations the baseline could be expanded to approximately five decades which was enough to refine the apsidal motion period determined by \\citet{lacyetal03}. Our $O-C$ analysis based on the extended time-line showed that an unseen, dark third body in the system cannot be extracted from the presently available minima observations. All of these makes very unlikely the presence of a third body with a mass and orbit which would cause a peculiar periastron precession.  First time two-colour light curves were presented by us for BP~Vul.  The light curve solution -- using the Wilson-Devinney Code -- yielded very similar results comparing to \\citet{devor05}'s one and \\citet{lacyetal03}'s one . In addition, the $O-C$ analysis in this paper showed that the apsidal motion period in BP~Vul is prograde and it has a period of about 365 years -- it is in agreement with the value determined with less accuracy from the light curve. The prograde motion means that BP Vul is not a representative of problematic cases and it is in agreement with theoretical expectations. Nevertheless we concluded that the negative apsidal motion rate determined by \\citet{lacyetal03} is only a consequence of their short observational window.  We also calculated the $k_2$ internal structure of BP~Vul and found it being close to the theoretical value. The slight difference should be refined in the future with a better observed $O-C$ diagram without gaps. Therefore the minima observations of BP~Vulpeculae in the future are needed.  The comments on the first version of the manuscript by Drs J. Jurcsik and K. Ol\\'ah is acknowledged.  \\begin{figure} \\begin{center} \\includegraphics[width=12cm]{figure1.eps} \\end{center} \\caption{The $B$ (top) and $V$ (bottom) light curves of BP~Vul obtained in the years 1959-1963. Bottom is the B curve while top curve is the V one which is shifted by 0.3 magnitudes for the sake of clarity.} \\label{fig:lightcurve} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=12cm]{figure2.eps} \\end{center} \\caption{The $O-C$ diagram of BP Vulpeculae. Filled and open squares represent primary and secondary minima, respectively. Lines show the fits to the $O-C$ residuals for the primary and secondary minima, respectively. The weights of the different kind of minima can be found in the text.} \\label{fig:O-C} \\end{figure}"
    },
    "0812/0812.3906.txt": {
        "abstract": "{} {We present a parameter study of the magnetohydrodynamical dynamo driven by cosmic rays in the interstellar medium (ISM) focusing on the efficiency of magnetic field amplification and the issue of energy equipartition between magnetic, kinetic and cosmic ray (CR) energies.} {We perform numerical CR-MHD simulations of the ISM using the extended version of ZEUS-3D code in the shearing box approximation and taking into account the presence of Ohmic resistivity, tidal forces and vertical disk gravity. CRs are supplied in randomly distributed supernova (SN) remnants and are described by the diffusion-advection equation, which incorporates an anisotropic diffusion tensor.} {The azimuthal magnetic flux and total magnetic energy are amplified in majority of models depending on a particular choice of model parameters. We find that the most favorable conditions for magnetic field amplification correspond to magnetic diffusivity of the order of $3\\times 10^{25} \\cm^2\\s^{-1}$, SN rates close to those observed in the Milky Way, periodic SN activity corresponding to spiral arms, and highly anisotropic and field-aligned CR diffusion. The rate of magnetic field amplification is relatively insensitive to the magnitude of SN rates in a rage of spanning 10\\% up to 100\\% of realistic values. The timescale of magnetic field amplification in the most favorable conditions is 150 Myr, at galactocentric radius equal to 5 kpc, which is close to the timescale of galactic rotation. The final magnetic field energies reached in the efficient amplification cases fluctuate near equipartition with the gas kinetic energy. In all models CR energy exceeds the equipartition values by a least an order of magnitude, in contrary to the commonly expected equipartition. We suggest that the excess of cosmic rays in numerical models can be attributed to the fact that the shearing-box does not permit cosmic rays to leave the system along the horizontal magnetic field, as it may be the case of real galaxies.} {} ",
        "introduction": "An attractive idea of {\\em fast galactic dynamo} has been proposed by Parker (1992). The idea relies on two ingredients: (1) cosmic rays (CR) continuously supplied to the disk by supernova (SN) remnants and (2) fast magnetic reconnection which operates in current sheets and allows to dissipate and relax the random magnetic field components in the limit of vanishing resistivity. Over the last decade we have investigated the different elements, physical properties and consequences of Parker's idea and scenario by means of analytical calculations and numerical simulations (Hanasz \\& Lesch 1993, 1997, 1998, 2000, 2001, 2003a, 2003b, Hanasz et al. 2002, 2004, 2006, Lesch \\& Hanasz 2003; Otmianowska-Mazur et al. 2003, 2007; Kowal et al. 2003, 2005)   The first complete 3D numerical model of the CR-driven dynamo has been demonstrated by Hanasz et al. (2004, 2006). In this paper we perform a parameter study of the CR-driven dynamo model by examining the dependence of magnetic field amplification on magnetic diffusivity, supernova rate determining the CR injection rate, temporal modulation of SN activity, grid resolution, and CR diffusion coefficients.   The principle of action of the CR-driven dynamo is based on the cosmic ray energy supplied continuously by SN remnants. Due to the anisotropic diffusion of cosmic rays and the horizontal magnetic field configuration, cosmic rays tend to accumulate within the disc volume. However, the configuration stratified by vertical gravity is unstable with respect to the Parker instability.  Buoyancy effects induce vertical and horizontal motions of the fluid and formation of undulated patterns -- magnetic loops in frozen-in, predominantly horizontal magnetic fields. The presence of rotation in galactic disks implies a coherent twisting of the loops by means of the Coriolis force, which leads to the generation of small-scale radial magnetic field components. The next phase is merging of small-scale loops by the magnetic reconnection process to form the large scale radial magnetic fields. Finally, the differential rotation stretches the radial magnetic field to amplify the large-scale azimuthal magnetic field component. The coupling of amplification processes of radial and azimuthal magnetic field components results in an exponential growth of the large scale magnetic field. The timescale of magnetic field amplification, resulting from the action of the CR-driven dynamo, has been found (Hanasz et al. 2004, 2006) to be equal to 140- 250 Myr, depending on the galactocentric radius, which is comparable to the galactic rotation period.   The CR-dynamo experiments reported in the afore mentioned papers rely on energy of CRs accelerated in SN remnants.  Recently Gressel et al. (2008a,b) reported a series of non-ideal MHD simulations demonstrating dynamo action resulting from the SN-driven turbulence, in absence of CRs. Using a similar set of galactic disk parameters, with angular velocity enlarged 4 times the  value of $0.025 \\Myr^{-1}$, typical for the galactocentric radius of Sun, these authors found amplification of the large scale magnetic fields on a time scale of $250 \\Myr$. This indicates some similarity between the CR-driven dynamo and the dynamo driven by thermal energy output from supernovae. The similarity is presumably related to the buoyancy effect which can be commonly attributed to the excess of both the thermal and cosmic ray energies in the disk volume.   The magnitudes of galactic magnetic fields are usually estimated from measurements of the radio synchrotron emission arising from the acceleration of cosmic ray electrons in the magnetic field. To interpret the radio emission spectrum, it is usually assumed that the energy density in the magnetic field is of the same order of magnitude as the energy density in cosmic ray protons (which are assumed to outnumber the electrons by 100:1, as they do in our Galaxy). There is however no compelling evidence of energy equipartition. Since the equipartition or minimum energy assumption is one of the few ways to calculate radio source parameters, it is important to determine how reasonable the approach generally is. Recently Strong et al. (2007) and Snodin et al. (2005) raised again the question about the applicability of the equipartition argument.  From the observational point of view (Fitt and Alexander 1993; Vallee 1995) the equipartition assumption seems to hold. In particular, Vallee's comparison of three different methods determining galactic magnetic field strengths (Faraday rotation method, equipartition method and Cosmic ray equipartition) shows that the equipartition fields are in a quite good agreement. On the other hand, Beck and Krause (2005) considered in detail a problem which was raised by Chi and Wolfendale (1993), that the commonly used classical equipartition or minimum-energy estimate of total magnetic field strengths from radio synchrotron intensities is of limited practical use because it is based on the hardly known ratio K of the total energies of cosmic ray protons and electrons and also has inherent problems. They present a revised formula using the {number density ratio} {K} for which they give estimates. For particle acceleration in strong shocks K is about 40 and increases with decreasing shock strength. Their revised estimate for the field strength gives larger values than the classical estimate for flat radio spectra with spectral indices of about 0.5-0.6, but smaller values for steep spectra and total fields stronger than about 10 $\\muG$. In very young supernova remnants, for example, the classical estimate may be too large by up to 10. On the other hand, if energy losses of cosmic ray electrons are important, K increases with particle energy and the equipartition field may be underestimated significantly.  From a more global point of view the assumption of equipartition seems to be a very natural one. A thermodynamical system will always distribute the free energy to all degrees of freedom available, if the system has time to do so. In cases of accelerated particles diffusing in the large-scale and turbulent magnetic fields one can expect that at least the turbulent magnetic field (since it represents three degrees of freedom) is somehow virialized with respect to any other pressure term, like the cosmic ray pressure. This may not be true for the ordered magnetic fields, which are supposed to be amplified by the combined action of differentially rotating shear flows in the disk and some helical upward and downward motion driven either by cosmic rays pressure or any activity in the disk. In any case the connection of star formation activity accompanied by enhanced flux of cosmic rays and the amplification of large-scale magnetic fields, inherently raises the expectation that magnetic fields should not exhibit higher pressures than the cosmic rays. Moreover, one would rather expect to find magnetic fields whose pressure is somewhat lower than the pressure of the cosmic rays, if the cosmic rays present a source for the galactic dynamo.  The paper is organized as follows: in Sect.~\\ref{sect:model} we describe the CR driven dynamo model and its numerical implementation, in Sect.~\\ref{sect:simul} we present  our simulation setup and describe parameters used in numerical simulations, in Sect.~\\ref{sect:results} we describe results, focusing on the effect of each parameter on magnetic field amplification rate. We discuss the final saturated states of models in terms of equipartition between kinetic, magnetic and CR energies. Finally, in Sect.~\\ref{sect:summary} we conclude our paper. ",
        "conclusions": "}  In the present paper we described an extensive series of simulations and presented parameter study of the CR-driven dynamo in a galaxy, characterized by the essential parameters typical for the Milky Way at galactocentric radius $R_G=5\\kpc$. In the presented study the magnetic diffusivity, as well as the parallel and perpendicular CR diffusion coefficients have been considered as free parameters, and dedicated simulation series have been performed to investigate their influence on the efficiency of the CR-driven dynamo process. The results of the parameter study can be summarized as follows:  (1) The magnitude of magnetic diffusivity influences essentially the efficiency of magnetic field amplification. The most favorable value of magnetic diffusivity is $100 \\pc^2\\Myr^{-1}  \\simeq 3 \\times 10^{25} \\cm^2\\s^{-1}$, the value comparable, although lower than the the value of turbulent diffusivity of ISM deduced from observational data.  (2) The efficiency of magnetic field amplification is enhanced by temporal modulation of the CR supply. An effect of this kind may be associated with the periodicity of star formation and supernova activity induced by spiral arms. The enhancement is apparent at lower values of magnetic diffusivity and is less significant at the optimal value of $\\eta = 100 \\pc^2\\Myr^{-1}$.  (3) Magnetic field amplification rate is relatively insensitive to the magnitude of SN rate within the range of SN rates spanning one decade below the value $f_{SN} = 130 \\kpc^{-2} \\,\\Myr^{-1}$ typical for galactocentric radius $R_G=5\\kpc$. We note also that magnetic field amplification saturates at the level of equipartition of magnetic and kinetic energies for all supernova rates for which amplification holds. Magnetic field is no longer amplified, if SN rate is further enhanced by factors 2 and 4 with respect to the realistic value, while other quantities (like e.g. gas column density) remain fixed.  (4) Magnetic field amplification in the CR-driven dynamo relies on anisotropic diffusion of cosmic rays. From the limited set of simulations of series E, one can deduce that the magnetic field amplification is possible only for $K_\\perp \\leq 3\\times 10^3\\pc^2 \\Myr^{-1} \\simeq  10^{27} \\cm^2 s^{-1} $, and for all considered values of the parallel diffusion coefficient $K_\\paral$ in the range $3\\div30\\times 10^{27} \\cm^2 s^{-1}$. Therefore, the 5\\% ratio of the perpendicular to parallel diffusion coefficients postulated by Giaccalone \\& Jokipii (1998) falls within the amplification range.  (5) By varying various parameters: magnetic diffusivity, supernova rate and the CR diffusion coefficients, we have found that the favorable conditions for magnetic field amplification correspond to approximately equal energies of the vertical and azimuthal magnetic field components in the case of buoyancy driven dynamo. An excess or deficit of vertical magnetic field with respect to the azimuthal one corresponds to significantly less efficient amplification or even decay of magnetic field.  (6) We note the problem indicated previously by Snodin et al (2005), that in all simulations the CR energy in the computational domain exceeds the turbulent kinetic energy and magnetic energy by more than one order of magnitudes. Moreover, the lowest ratios of CR to kinetic energies emerge for the largest values of the parallel diffusion coefficient.  It seems not plausible, however, that an enlargement of diffusion coefficients up to fully realistic values will reduce the excess of cosmic rays energy in the disk. It seems also, that the ratio of CR to other forms of energy in the ISM is not yet well restricted on observational grounds (Strong et al 2007). On the other hand, the currently used shearing box approximation does not permit CRs to leave the disk by means of diffusion along the predominantly horizontal magnetic field. Therefore, we suggest that subsequent work on the CR-driven dynamo, aimed to solve this problem, should be done in the framework of global galactic disk simulations.  %We end up with the problem, that the nonlinear dynamics of the interstellar %medium favors a strong coupling of all dynamically relevant participants like %turbulence, cosmic rays and magnetic fields. On the other hand, the details of %an eventually completely balanced network of all pressure components within the %interstellar medium is still a matter of debate."
    },
    "0812/0812.0199_arXiv.txt": {
        "abstract": "The operation of an interface dynamo (as has been suggested for the Sun and other stars with convective envelopes) relies crucially upon the effective transport of magnetic flux between two spatially disjoint generation regions. In the simplest % models communication between the two regions is achieved solely by diffusion. Here we incorporate a highly simplified % anisotropic transport mechanism in order to model the net effect of flux conveyance by magnetic pumping and by magnetic % buoyancy. We investigate the influence of this mechanism on the efficiency of kinematic dynamo action. It is found that the % effect of flux transport on the efficiency of the dynamo is dependent upon the spatial profile of the transport. Typically, % transport hinders the onset of dynamo action and increases the frequency of the dynamo waves. However, in certain cases % there exists a preferred magnitude of transport for which dynamo action is most efficient. Furthermore, we demonstrate the % importance of the imposition of boundary conditions in drawing conclusions on the role of transport. ",
        "introduction": "\\label{intro}  The solar magnetic field is maintained by the action of a hydromagnetic dynamo. The well-known eleven year solar activity % cycle corresponds to a twenty-two year magnetic cycle, the origin of which lies deep within the solar interior (see, for example, % Charbonneau 2005). The key to understanding the generation of the solar magnetic field lies in establishing the form of both the % large- and small-scale fluid flows in the interior. Helioseismology, the study of acoustic fluctuations within the Sun, has % provided a map of the solar differential rotation (Schou et al.~1998), which shows that the surface latitudinal dependence is % largely maintained throughout the convection zone while the solar interior rotates as a solid body. These two profiles are % matched in a thin, stably stratified layer of strong differential rotation (shear) at the base of the convection zone -- a region % known as the tachocline. A discussion of the dynamics of the tachocline can be found in Tobias (2005) and Hughes, Rosner % \\& Weiss (2007).  This region of pronounced shear is believed to be of paramount importance in the generation (via the so-called % $\\omega$-effect), storage and eventual instability of a strong toroidal magnetic field. The regeneration of poloidal field from % toroidal field is less well understood. It is typically ascribed to the $\\alpha$-effect of mean field electrodynamics, but the % precise origin and location of this mechanism remain unclear. A number of possibilities have been suggested (see the review of % Ossendrijver 2003). The most natural is to appeal to convective turbulence in the presence of rotation (Parker 1955; Steenbeck, % Krause \\& R\\\"adler 1966). There are however some difficulties for this approach both in the linear (Cattaneo \\& Hughes 2006, % Hughes \\& Cattaneo 2008) and nonlinear regimes (see the reviews by Diamond, Hughes \\& Kim 2005; Brandenburg \\& % Subramanian 2005). An alternative is to appeal to magnetic buoyancy instabilities of the tachocline field interacting with rotation % (Schmitt, Sch\\\"ussler \\& Ferriz-Mas 1996; Thelen 2000; see also Cline, Brummell \\& Cattaneo 2003 for a related model). A % third possibility places the generation of poloidal field well away from the tachocline, the assumed location for the generation of % toroidal field. In this paradigm (Dikpati \\& Charbonneau 1999; Dikpati \\& Gilman 2001) poloidal flux is regenerated via the % decay of active regions at the solar surface (Babcock 1961; Leighton 1969). In these models meridional circulation is invoked in order to % transport the new poloidal flux to the tachocline. This model too is not without difficulties (see the review by Tobias \\& Weiss % 2007).  In this paper we shall consider models of dynamos operating at an interface between regions of turbulent rotating convection % and velocity shear. The idea of an interface dynamo --- with neighbouring regions of $\\alpha$- and $\\omega$-effects --- was % introduced by Parker (1993) expressly to avoid the difficulties caused by the strong suppression of turbulent magnetic diffusion % by a weak mean field. In Parker's model, strong toroidal field is generated in the tachocline by radial shear. There, and only % there, does it suppress the diffusion, leading naturally to confinement of the toroidal field within that layer. The $\\alpha$-effect % and a healthy turbulent diffusion operate relatively unimpeded  in the bulk of the convection zone. Parker suggested that % provided the diffusivity is not suppressed too strongly, observed solar fluxes and the 22-year solar period could be % accommodated. Parker's original kinematic model has been extended into the nonlinear regime to investigate explicitly the % quenching not only of the turbulent diffusion but also of the $\\alpha$-effect (Tobias 1996; Charbonneau \\& MacGregor 1997; % Markiel \\& Thomas 1999).  A crucial ingredient for efficient interface dynamo action is the transport of magnetic flux between the two spatially disjoint % generation regions. For this reason, dynamo action is, typically, more difficult to excite in such models than in those with % spatially-coincident generation mechanisms (Parker 1993). Indeed, if such a transport mechanism is absent, this places a severe % constraint on the applicability of interface dynamos (Dikpati et al.\\ 2005). Within the Sun there are however a number of % possible mechanisms that could transport magnetic field. For example, convective pumping (e.g.\\ Drobyshevski \\& Yuferev % 1974) is believed to transport flux downwards from the convection zone into the tachocline (e.g.\\ Tobias et al.\\ 1998, 2001; % Dorch \\& Nordlund 2001). Turbulent diffusion, which is believed to act throughout the convection zone but could be % significantly reduced in the tachocline (Parker 1993), acts to reduce gradients in field. Mean flows, such as the meridional % circulation (Choudhuri et al. 1995), which is observed to be polewards at and just below the solar surface but is undetermined % at the base of the convection zone, will act to advect mean flux. The only transport mechanism that is unambiguously directed % outwards is magnetic buoyancy, which is believed to play a key role in the emergence of active regions at the surface (see, for % example, the reviews by Sch\\\"ussler 2005 and Hughes 2007). These competing effects may occur on different timescales and % depend on the local strength of magnetic field. Indeed, the general picture we have of the solar dynamo is of the field being % confined and maintained at the base of the convection zone, with occasional eruptions.  In this paper we consider the effect of incorporating non-diffusive flux transport in an interface dynamo model. In keeping with % the simple yet illustrative nature of Parker's original model, we include magnetic transport via an effective velocity. Indeed, such % a velocity arises naturally in mean field electrodynamics as the `$\\gamma$-effect', resulting from the anti-symmetric part of the % $\\alpha$-tensor (R\\\"adler 1968). It should be noted that mean field electrodynamics is not without problems, particularly in the % high magnetic Reynolds number regime applicable to the Sun. Nonetheless, our hope is that it is possible to capture some % features of the interaction of generation, transport and diffusion of magnetic fields via a simple mean-field model.  In this paper we shall investigate the role of the $\\gamma$-effect on the onset of the dynamo instability for two kinematic $\\alpha\\omega$-dynamo models. A natural starting point is to extend Parker's 1993 model, in which the interface % separates two spatially unbounded regions. It is though important to consider the more realistic case of vertically bounded % domains, which is the aim of our second model. In \\S2 we discuss those characteristics that are common to both models. The % effect of introducing flux transport is investigated, for the two models, in \\S3 and \\S4. The possible implications for an % interface dynamo in the Sun are discussed in \\S5. ",
        "conclusions": "\\label{sec:4}   We have considered the role of anisotropic transport on the efficiency of an interface dynamo in the linear regime, and its effect in modifying the temporal structure of the solution, through the frequency of the dynamo wave and the phase relation between the radial and azimuthal fields. One of the main findings of the investigation is that the conclusions to be drawn depend crucially upon the geometry of the model.  The first model considered was of infinite vertical extent. In the absence of transport, the preferred wavenumber for the onset % of the dynamo instability is $k=0$, where the critical parameter $(\\alpha G/\\eta_1^2)_c=0$. Although, clearly, this cannot be % decreased by incorporating the $\\gamma$-effect, for all forms of transport investigated it is found that $\\textrm{min}_k % [(\\alpha G/\\eta_1^2)_c]$ increases as the magnitude of the transport increases. Thus dynamo action is less readily achieved. % Also, the period of the dynamo wave decreases as the magnitude of the $\\gamma$-effect increases. Thus although the dynamo is more difficult to excite, it more quickly % reverses the polarity of the field. Perhaps most noteworthy is that for this model the two cases of the $\\gamma$-effect directed % towards the interface from both above and below and of the transport directed away from the interface in both regions have % identical critical dynamo numbers and corresponding frequencies  (\\S\\ref{sec:m2_bidirectional}).  The effect of the transport mechanism in the more realistic Model~II is more complex. The model geometry is the same as that of Model~I except that the domain of the problem is bounded in the vertical direction. Here it is found that a uniform transport velocity throughout the convection zone and tachocline delays the onset of the dynamo instability and increases the frequency of the dynamo waves. However, transport directed towards the interface in both $z>0$ and $z<0$ decreases both the critical dynamo number and the frequency. Thus transport aids the onset of the % instability in this case, although it then takes longer to reverse the field polarity. Restricting the transport mechanism to the % convection zone reveals a preferred strength for dynamo action.  It is of interest to relate our studies to previous investigations of the effects of transport on dynamo models. { Brandenburg et % al.~(1993) and Gabov et al.~(1996) considered models of galactic dynamos and investigated the role of turbulent diamagnetism % on the dynamo instability via the additional term $\\nabla \\times (-\\nabla \\beta/2 \\times \\mathbf{B})$ in the induction equation, % where $\\beta$ is the turbulent diffusivity. Brandenburg et al.~(1993) illustrated that a galactic wind can enhance the instability by % transporting field into regions of stronger generation, while turbulent diamagnetism can act to oppose this effect and hence % hinder dynamo action. Similarly, Gabov et al.~(1996) illustrated that the critical dynamo number is larger when the effects of % turbulent diamagnetism are present. This was attributed, in part, to an increase in the effective diffusion coefficient. It was then % shown that the effective dynamo number $D_c/\\tilde\\beta^2$ (where $\\tilde\\beta$ is the mean value of the turbulent diffusion % coefficient) decreases as $\\tilde\\beta$ is increased. It was further illustrated that for highly supercritical dynamo numbers the % linear growth rate is larger when turbulent diamagnetism is present than in its absence. In our study we have concentrated on the % effects of transport on marginally stable solutions. Nightingale (1983) incorporated the $\\gamma$-effect into a spherical % $\\alpha^2$-dynamo model and found that the dynamo growth rate is smaller when $\\gamma$ is non-zero than when there is no % transport. R\\\"adler (1986) considered spherical $\\alpha\\omega$- and $\\alpha^2$-dynamos and showed that although radial % transport mainly hinders dynamo action, certain strengths of $\\gamma$ can promote the instability.}  Finally, we remark that interface dynamo models relying purely upon diffusive communication are sometimes rejected owing to % their inability to reproduce either the observed direction of propagation of the dynamo belts (exemplified in the butterfly % diagram) or the observed phase relation between the radial and azimuthal fields. In this regard, it is of interest to note that the % study by Brandenburg et al. (1992) showed that a sufficiently strong downward pumping can lead to an equatorward migration % of the magnetic field even when the conventional rule for the sign of $\\alpha \\partial \\Omega/\\partial r$ yields a poleward % migration. Also, here we have shown that the transport mechanism changes the ratio of the phase difference to the period of the % dynamo wave. It is therefore entirely possible that tuning the $\\gamma$ parameter could enable interface dynamo models to % reproduce the observed behaviour."
    },
    "0812/0812.4933.txt": {
        "abstract": "Radio interferometry probes astrophysical signals through incomplete and noisy Fourier measurements. The theory of compressed sensing demonstrates that such measurements may actually suffice for accurate reconstruction of sparse or compressible signals. We propose new generic imaging techniques based on convex optimization for global minimization problems defined in this context. The versatility of the framework notably allows introduction of specific prior information on the signals, which offers the possibility of significant improvements of reconstruction relative to the standard local matching pursuit algorithm CLEAN used in radio astronomy. We illustrate the potential of the approach by studying reconstruction performances on simulations of two different kinds of signals observed with very generic interferometric configurations. The first kind is an intensity field of compact astrophysical objects. The second kind is the imprint of cosmic strings in the temperature field of the cosmic microwave background radiation, of particular interest for cosmology. ",
        "introduction": "\\label{sec:Introduction}Radio interferometry is a powerful technique for aperture synthesis in astronomy, dating back to more than sixty years ago \\citep{ryle46,blythe57,ryle59,ryle60,thompson04}. In a few words, thanks to interferometric techniques, radio telescope arrays synthesize the aperture of a unique telescope of the same size as the maximum projected distance between two telescopes on the plane perpendicular to the pointing direction of the instrument. This allows observations with otherwise inaccessible angular resolutions and sensitivities in radio astronomy. The small portion of the celestial sphere accessible to the instrument around the pointing direction tracked during observation defines the original real planar signal or image $I$ to be recovered. The fundamental Nyquist-Shannon theorem requires a signal to be sampled at a frequency of twice its bandwidth to be exactly known. The signal $I$ may therefore be expressed as a vector $x\\in\\mathbb{R}^{N}$ containing the required number $N$ of sampled values. Radio-interferometric data are acquired in the Fourier plane. The number $m$ of spatial frequencies probed may be much smaller than the number $N$ of discrete frequencies of the original band-limited signal, so that the Fourier coverage is incomplete. Moreover the spatial frequencies probed are not uniformly sampled. The measurements are also obviously affected by noise. An ill-posed inverse problem is thus defined for reconstruction of the original image.  Beyond the Nyquist-Shannon theorem, the emerging theory of compressed sensing aims at merging data acquisition and compression \\citep{candes06a,candes06b,candes06c,donoho06,baraniuk07}. It notably relies on the idea that a large variety of signals in Nature are sparse or compressible. By definition, a signal is sparse in some basis if its expansion contains only a small number of non-zero coefficients. More generally it is compressible if its expansion only contains a small number of significant coefficients, i.e. if a large number of its coefficients bear a negligible value. Compressed sensing theory demonstrates that a much smaller number of linear measurements is required for accurate knowledge of such signals than is required for Nyquist-Shannon sampling. The sensing matrix must simply satisfy a so-called restricted isometry property. In particular, a small number of random measurements in a sensing basis incoherent with the sparsity or compressibility basis will ensure this property with overwhelming probability, e.g. random Fourier measurements of a signal sparse in real or wavelet space. Consequently, if compressed sensing had been developed before the advent of radio interferometry, one could probably not have thought of a much better design of measurements for sparse and compressible signals in an imaging perspective.  In this work we present results showing that the theory of compressed sensing offers powerful image reconstruction techniques for radio-interferometric data. These techniques are based on global minimization problems, which are solved by convex optimization algorithms. We also emphasize on the versatility of the scheme relative to the inclusion of specific prior information on the signal in the minimization problems. This versatility allows the definition of image reconstruction techniques which are significantly more powerful than standard deconvolution algorithm called CLEAN used in the context of radio astronomy.  In Section \\ref{sec:Radio-interferometry}, we pose the inverse problem for image reconstruction from radio-interferometric data and discuss the standard image reconstruction techniques used in radio astronomy. In Section \\ref{sec:Compressed-sensing-perspective}, we concisely describe the central results of the theory of compressed sensing regarding the definition of a sensing basis and the accurate reconstruction of sparse or compressible signals. In Section \\ref{sec:Applications}, we firstly comment on the exact compliance of radio interferometric measurements with compressed sensing. We then study the reconstruction performances of various compressed sensing imaging techniques relative to CLEAN on simulations of two kinds of signals of interest for astrophysics and cosmology. We finally conclude in Section \\ref{sec:Conclusion}.  Notice that a first application of compressed sensing in astronomy \\citep{bobin08} was very recently proposed for non-destructive data compression on board the future Herschel space observatory% \\footnote{http://herschel.esac.esa.int/% }. The versatility of the compressed sensing framework to account for specific prior information on signals was already pointed out in that context. Moreover, the generic potential of compressed sensing for interferometry was pointed in the signal processing community since the time when the theory emerged \\citep{donoho06,candes06b,mary07,levanda08}. It was also very recently acknowledged in radio astronomy \\citep{cornwell08b}. The present work nonetheless represents the first application of compressed sensing for the definition of new imaging techniques in radio interferometry. A huge amount of work may be envisaged along these lines, notably for the transfer of the proposed techniques to optical and infrared interferometry. The extension of these techniques from the plane to the sphere will also be essential, notably with regard to forthcoming radio interferometers with wide fields of view on the celestial sphere \\citep{cornwell08a,mcewen08}, such as the future Square Kilometer Array (SKA)% \\footnote{http://www.skatelescope.org/% } \\citep{carilli04}. ",
        "conclusions": "\\label{sec:Conclusion}Compressed sensing offers a new framework for image reconstruction in radio interferometry. In this context, the inverse problem for image reconstruction from incomplete and noisy Fourier measurements is regularized by the definition of global minimization problems in which a generic sparsity or compressibility prior is explicitly imposed. These problems are solved through convex optimization. The versatility of this scheme also allows inclusion of specific prior information on the signal under scrutiny in the minimization problems. We studied reconstruction performances on simulations of an intensity field of compact astrophysical objects and of a signal induced by cosmic strings in the CMB temperature field, observed with very generic interferometric configurations. The $\\textnormal{BP}_{\\epsilon}$ technique provides similar reconstruction performances as the standard matching pursuit algorithm CLEAN. The inclusion of specific prior information significantly improves the quality of reconstruction.  Further work by the authors along these lines is in preparation. In particular, a more complete analysis is being performed to estimate the lowest string tension down to which compressed sensing imaging techniques can reconstruct a string signal in the CMB, in more realistic noise and Fourier coverage conditions. In this case, given the compressibility of the magnitude of the gradient of the string signal itself, TV norm minimization also represents an interesting alternative to the $\\textnormal{SBP}_{\\epsilon}$ problem proposed here."
    },
    "0812/0812.1010_arXiv.txt": {
        "abstract": "\\normalsize{ We present a de-trending algorithm for the removal of trends in time series. Trends in time series could be caused by various systematic and random noise sources such as cloud passages, changes of airmass, telescope vibration, CCD noise or defects of photometry. Those trends undermine the intrinsic signals of stars and should be removed. We determine the trends from subsets of stars that are highly correlated among  themselves. These subsets are selected based on a hierarchical tree  clustering algorithm. A bottom-up merging algorithm based on the departure from normal distribution in the correlation is developed to identify subsets, which we call clusters. After identification of clusters, we determine a trend per cluster by weighted sum of normalized light-curves. We then use quadratic programming to de-trend all individual light-curves based on these determined trends. Experimental results with synthetic light-curves containing artificial trends and events are presented. Results from other de-trending methods are also compared. The developed algorithm can be applied to time series for trend removal in both   narrow and wide field astronomy.} ",
        "introduction": "Small-aperture telescopes have detected a large number of   exo-planet transits \\citep{Alonso2004,Bakos2004,McCullough2005,Bakos2007,Burke2008,Pal2008,Pollacco2008}. A large number of variable stars have also been detected by surveys that use  such telescopes \\citep{Schmidt1991,Akerlof2000,Pojmanski2005,Schmidt2007,Pigulski2008,Szczygiel2008}. A weakness in these surveys is that the signal to noise ratio (S/N) is lower than   the S/N obtained by larger-aperture telescopes. The low S/N can be  attributed not only to the small-aperture size but also to noise in CCD images such as non-uniform illumination, or to local weather changes throughout the field (especially in the case of wide field surveys). To improve the S/N and thus improve the detectability of variability,   these noise sources should be minimized.  Some of these  noise sources are strongly correlated between light-curves of different stars. For example, if a star appears fainter, other stars near it may appear fainter at the same time. We call such coherent changes through parts of the field {\\em trends}. These trends could be caused by local weather patterns such as thin cloud passages or airmass changes \\citep{Howell1986, Kjeldsen1992, Gilliland1988} throughout the night. The conventional approach for trend removal is differential photometry with a {\\em reasonable} selection of template stars near the star of interest   \\citep{Young1991, Everett2001}. With the help of modern CCDs, it is not   hard to select a sufficient number of bright stars as a template set. However, the de-trended results are then sensitive to the selection of template stars. If the template stars contain intrinsic variables, the determined trends will be different from  the true trends. Therefore, excluding such intrinsically variable stars from template stars is essential. Furthermore, because there is no guarantee that trends are the same for all stars throughout the entire field,  the template selection method should be able to handle localization of trends in large fields of wide field surveys.   In this paper, we propose a new de-trending method   (hereafter PDT, for  Photometric De-Trending algorithm), which incorporates a systematic template   selection algorithm that can solve the problems mentioned above and consequently shows superior de-trended results. Experiments with simulated light-curves show that PDT correctly reproduces localization.  We present details of PDT in Sec. \\ref{sec:Algorithm}. In Sec. \\ref{sec:Test}, we show de-trended results for synthetic   light-curves containing artificially added trends and events. In addition, comparison results with the Trend Filtering Algorithm (hereafter, TFA) \\citep{Kovacs2005} are  also presented. In Sec. \\ref{sec:TAOSandMMT}, we show two examples of astronomical datasets and their de-trended results. We outline future work in Sec. \\ref{sec:NoteFutureWork}.  We summarize our  conclusions  in Sec. \\ref{sec:Conclusion}. ",
        "conclusions": "\\label{sec:Conclusion}   In this paper, we presented the Photometric De-Trending algorithm (PDT), a new de-trending algorithm. We first determined the trends by constructing a hierarchical tree based on the similarity matrix. Elements of the similarity matrix are the Pearson correlation values of all pairs of light-curves. After that, a bottom-up merging algorithm was applied to the constructed tree in order to identify subsets of light-curves that we call clusters. At each step of the merging process, we tested the normality of the subsets and determined where to stop. By means of the normality test, we could select reliable clusters of trends. For each cluster, we determined one representative master-trend by weighted sum  of the normalized light-curves. This procedure greatly constrained the number of free parameters to be calculated, and thus  showed less significant signal depression than other de-trending algorithms such as Trend Filtering Algorithm (TFA). Finally, in order to remove the trends from individual light-curves, we used quadratic programming to minimize the residual  between each target light-curve and the determined master-trends. Note that PDT is designed to remove only the fluctuations that are common among stars. If the fluctuations are unique to an individual star, the fluctuations will be preserved.  We performed several simulations of synthetic light-curves with different initial parameters such as  total duration of observation, transit duration, field of view, exposure time etc,  to test PDT and showed some of the simulation results in this paper.   First, we tested PDT with $\\sim$500 synthetic light-curves  that contain the first order atmospheric extinction (airmass), artificial trends, Poisson noise and events (three transits and two eclipsing binaries). We applied PDT to these synthetic light-curves in order to determine trends  and to regenerate the inserted events. PDT successfully identified multiple clusters of different  trends which were the mixture of different trends and noise. These identified clusters well represented the overall characteristic of the trends through the field. We compared de-trended results of PDT with one another de-trending algorithm (TFA). PDT results were an improvement over TFA results, especially when 1) the dataset contains intrinsic variables  that would be included in template set of TFA or 2) the rms contribution from the intrinsic signals is significant.   We also tested PDT with $\\sim$500 synthetic light-curves that contain color dependent second order extinction and Poisson noise. Trends appearing in light-curves can be slightly different due to differences in color. We found that PDT can identify clusters according to color. However, in realistic scenarios, it is not easy to isolate only the second order extinction  because, even if one correctly removes the first order extinction for all stars, there exist other various noise sources which dilute trends caused by the second order extinction.   In the case of dataset of random fluctuations (e.g. pure Poisson noise), which does not have any trends, we do not need to de-trend the dataset. PDT can distinguish the light-curves of random fluctuations using the characteristics of the distribution of correlation coefficients. Therefore, PDT does not de-trend these light-curves and thus preserves any intrinsic signals.  Examples of two astronomical datasets are also presented. They show multiple trends in the field caused by various noise sources such as airmass, cloud passages, telescope vibration, defects of photometry and so on. PDT performed well and removed trends that appeared in the datasets.  In this paper, we show the simulation results of wide field data only. However, PDT can be applied to narrow field data as well if there are enough stars in the field ($\\sim$ a few hundreds). In addition, PDT is useful to extract global trends that can represent the overall  characteristics of a dataset. The extracted trends can give a general idea  of how much the data are contaminated by the trends.  The software package of PDT is be provided at \\href{http://timemachine.iic.harvard.edu/detrending01/}{http://timemachine.iic.harvard.edu}."
    },
    "0812/0812.2365.txt": {
        "abstract": "%context {The source \\src\\ is one of the most enigmatic high-mass X-ray binaries. In spite of intensive searches, X-ray pulsations have not been detected in the time range $10^{-3}$-$10^{3}$ s. A cyclotron line at $\\sim$ 30 keV has been suggested by various authors but never detected with significance. The stellar wind of the optical companion is abnormally slow. The orbital period, initially reported to be 9.6 days, disappeared and a new periodicity of 19.25 days emerged.} %aims {The main objective of our RXTE monitoring of \\src\\ is to study the X-ray orbital variability of the spectral and timing parameters. The new long and uninterrupted RXTE observations allow us to search for long ($\\sim$1 hr) pulsations for the first time.  } %methods {We divided the $\\sim$7-day observation into five intervals and obtained time-averaged energy spectra and power spectral density for each observation interval. We also searched for pulsations using various algorithms.} %results {We have discovered 5560-s pulsations in the light curve of \\src. Initially detected in RXTE data, these pulsations are also present in INTEGRAL and EXOSAT observations. The average X-ray luminosity in the energy range 2--10 keV is $1.5 \\times 10^{35}$ erg s$^{-1}$ with a ratio $F_{\\rm max}/F_{\\rm min}\\approx 5$. This ratio implies an eccentricity of $\\sim$0.4, somewhat higher than previously suggested. The power spectrum is dominated by red noise that can be fitted with a single power law whose index and strength decrease with X-ray flux. The source also shows a soft excess at low energies. If the soft excess is modelled with a blackbody component, then the size and temperature of the emitting region agrees with its interpretation in terms of a hot spot on the neutron star surface. } %conclusions {The discovery of X-ray pulsations in \\src\\ confirms the neutron star nature of the compact companion and definitively rules out the presence of a black hole. The source displays variability on time scales of days, presumably due to changes in the mass accretion rate as the neutron star moves around the optical companion in a moderately eccentric orbit. If current models for the spin evolution in X-ray pulsars are correct, then the magnetic field of \\src\\ at birth must have been $B\\simmore 10^{14}$ G.} ",
        "introduction": "The X-ray source \\src\\ is one of the most enigmatic high-mass X-ray binaries (HMXB). Although it has been studied by numerous ground- and space-based observatories, many questions about the nature of its variability patterns remain unanswered. In the optical and UV bands, the emitting spectrum is complex \\citep{negu01} and far from that of a Be star. Initially classified as a B0-2 star \\citep{stei84}, the optical counterpart defies any standard spectral analysis. The absence of any obvious long-term trends in the evolution of the \\ha\\ line spectral parameters and the lack of any correlation between the \\ha\\ line equivalent width and infrared magnitudes and colours  rules out a Be classification. However, the strength and shape of the \\ha\\ line do not resemble those seen in supergiant stars either. Although the presence of an O9.5V star is virtually secured, the He abundance is abnormally strong \\citep{blay06}.  Without a circumstellar disc around the donor, the material needed for accretion and production of X-rays must come from the stellar wind. However, the expected X-ray luminosity using the typical wind velocities of an O9.5V star is about two orders of magnitude lower than the observed X-ray luminosity $L_{\\rm X}~\\sim 10^{35}$ erg s$^{-1}$. The distance to \\src\\ is estimated to be 2.6 kpc \\citep{blay06}. \\citet{ribo06} found that the wind terminal velocity of \\src\\ is abnormally slow ($\\sim$350 km s$^{-1}$). With such a low velocity, an eccentric orbit and using the Bondi-Hoyle formalism it is possible to reproduce the average X-ray luminosity of the system as well as the orbital variability seen in the {\\it RXTE}/ASM light curve between 1996 and 2005. \\src\\ is the only permanent wind-fed HMXB with a main-sequence donor \\citep{ribo06}.  Equally enigmatic is the nature of the compact companion, which has not been settled beyond doubt. Although a neutron star is the most likely scenario, the low X-ray luminosity, compared to other persistent HMXB, has been used as an argument to suggest the presence of a white dwarf \\citep{sara92,corb01}, while the lack of pulsations and the similarities with LS 5039 do not allow us to rule out a  black hole companion \\citep{negu01} -- the most recent results show, however, that LS 5039 is probably a non-accreting neutron star \\citep{ribo08}.  Broad-band X-ray observations favour the presence of a neutron star \\citep{torr04}. When modelled with Comptonization models, the resulting seed photons have $kT>1$ keV, while the bolometric luminosity is low ($L_{\\rm X}\\approx 10^{35}$ erg s$^{-1}$). The only way to reconcile these two results is by invoking a small emitting area ($< 2$ km), like that of a polar cap in a neutron star \\citep{mase04}, also ruling out the presence of an accretion disc \\citep{torr04}. The presence of a magnetic field with $B\\sim 2 \\times 10^{12}$ G is implied by the possible detection of a cyclotron resonance scattering feature at $\\sim$30 keV. Although the statistical significance of this feature is only marginal, it has been observed by three different observatories, namely {\\it RXTE}, {\\it BeppoSAX} and {\\it INTEGRAL} \\citep{torr04,mase04,blay05}.  %If the cyclotron line is finally confirmed then the black-hole scenario %would be ruled out.  Another debated question is the orbital period of the system. \\citet{corb01} reported an orbital period of 9.6 days. This value was based on observations made with the {\\it RXTE} ASM instrument over 5 years. This orbital period was refined to 9.5591 days by \\citet{ribo06} using 9 years worth of data of the same instrument. However, \\citet{corb07} could not find the 9.6-day modulation in the light curves of {\\it SWIFT} BAT and {\\it RXTE} ASM for observations carried out from 2004 to 2005. Instead, they found a strong modulation at 19.25 d, that is, consistent with twice the 9.6-day period (or 19.12 d).  In this paper we investigate the X-ray variability of \\src\\ on time scales of hours to days. The main objective is to search for $\\sim$1-hr pulsations and study possible spectral and timing orbital variability.   %-----------------------------fig 2--------------------------------------------- \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{./10950f2.eps}} \\\\ \\caption[]{Portion of $\\sim 6$ hours of the light curves of \\src\\ corresponding to each one of the observation intervals. The axis scales were left the same in all panels to allow easy comparison of the variability amplitude. The bin size is 10 seconds. Intervals C and D are more affected by observational gaps. } \\label{lc_int} \\end{figure} %------------------------------------------------------------------------------ %----------------------------------fig 3-------------------------------------- \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{./10950f3.eps}} \\\\ \\caption[]{Hardness-intensity diagram of \\src\\ for various X-ray bands. Count rates correspond to PCU2. Bin size is 512 s.} \\label{hr} \\end{figure} %------------------------------------------------------------------------------ ",
        "conclusions": "We have finally unveiled the nature of the compact companion in the high-mass X-ray binary \\src. The {\\it RXTE} light curve of \\src\\ appears to be modulated with a period of $\\sim$1.54 hours that we interpret as the spin period of the neutron star. Evidence for this modulation is also found in {\\it INTEGRAL} and {\\it EXOSAT} data. Thus we rule out that the periodicity is due to the {\\it RXTE} orbit. \\src\\ becomes the  slowest rotating neutron star in a high-mass X-ray pulsar with a main-sequence companion. The X-ray flux and the spectral index of the power-law components that fit  energy and power spectra exhibit variability on time scales of days. The amplitude of variability of the X-ray flux changed by a factor of 5 on time scales of days. The power-law index is larger (i.e. steeper spectrum) during low-flux states in the energy spectra and during high-flux states in the power spectra. We attribute this variability to changes in the mass accretion rate as the neutron star orbits the O-type companion in a moderately eccentric orbit.  The time span covered by the observations is shorter than any of the two suggested orbital periods, hence further monitoring campaigns covering a larger fraction of the orbit are needed to pin down the correct orbital period. However, the fact that the maximum of the PCA light curve occurs at phase zero, i.e., in coincidence with the ASM maximum when folded onto the 9.6-day period and the detection of a strong signal at $\\sim$815000 s in the {\\it RXTE} periodogram would favour the shorter period. The detection of a soft component in the energy spectrum has been interpreted as emission from the polar caps. The plot of the blackbody emitting radius and temperature as a function of X-ray luminosity constitutes a useful tool to distinguish the type of accretion mechanism at work in the X-ray binaries that show a soft excess.   %No agreement is found between our PCA data and the ASM light curve if the %orbital period is $2P_{\\rm orb}$.  %That is, the PCA light curves is consistent with the modulation of the %folded ASM light curve \\citep[see][]{corb07}."
    },
    "0812/0812.1226_arXiv.txt": {
        "abstract": "We present the most likely optical counterparts of 113 X-ray sources detected in our \\chandra survey of the central region of the Small Magellanic Cloud (SMC) based on the OGLE-II and MCPS catalogs. We estimate that the foreground contamination and chance coincidence probability are minimal for the bright optical counterparts (corresponding to OB type stars; 35 in total). We propose here for the first time 13 High-Mass X-ray Binaries (HMXBs), of which 4 are Be X-ray binaries (Be-XRBs), and we confirm the  previous classification of 18 Be-XRBs. We estimate that the new candidate Be-XRBs have an age of $\\sim15-85$ Myr, consistent with the age of Be stars. We also examine the ``overabundance'' of Be-XRBs in the  SMC fields covered by \\chandrasimple, in comparison with the Galaxy. In luminosities down to $\\sim10^{34}$\\ergs\\,, we find that SMC Be-XRBs are $\\sim1.5$ times more common when compared to the Milky Way even after taking into account the difference in the formation rates of OB stars. This residual excess can be attributed to the lower metallicity of the SMC. Finally, we find that the mixing of Be-XRBs with other than their natal stellar population is not an issue in our comparisons of Be-XRBs and stellar populations in the SMC. Instead, we find indication for variation of the SMC XRB populations on kiloparsec scales, related to local variations of the formation rate of OB stars and slight variation of their age, which results in different relative numbers of Be stars and therefore XRBs. ",
        "introduction": "X-ray binaries (XRBs) are stellar systems consisting of a compact object (neutron star, black hole, white dwarf) accreting material from a close companion star. They are the end points of stellar evolution, and thus by studying them we can set constraints on stellar evolution and compact object formation models. They constitute a numerous class of X-ray bright objects, with typical luminosities of $\\sim10^{36}-10^{38}$\\lunit\\, when in outburst. They are divided into two classes depending on the mass of the donor star, with M $\\leq1$\\msun for the Low-Mass XRBs (LMXBs), and M $\\geq10$\\msun for the High-Mass XRBs (HMXBs).  The companion star in HMXBs is of O or B spectral type, with optical bolometric luminosity usually exceeding that of the accretion disk (e.g. van Paradijs \\& McClintock 1995). HMXBs can further be divided into two groups, the supergiant X-ray binaries (hereafter SG-XRBs), and the Be/X-ray binaries (hereafter Be-XRBs). In the SG-XRBs the primary is a supergiant of spectral type earlier than B2, or an Of star, and it has evolved away from the main sequence (MS), while in the Be-XRBs, the primary is an Oe or Be star lying close to the MS. The optical spectra of Be-XRBs are characterized by emission lines (mostly of the Balmer series of hydrogen; for a review of HMXBs and their properties see for example van Paradijs \\& McClintock 1995).   Different mechanisms are believed to be responsible for the mass transfer in these two groups of HMXBs. The most luminous SG-XRBs usually have an accretion disk fueled via Roche-lobe overflow, while the less luminous systems may be fed by a supersonic stellar wind. In the Be-XRBs, Be stars are characterized by a low velocity, high density equatorial wind, resulting in periodic accretion episodes during the passage of the compact object through the decretion disk of the donor.   Be-XRBs show pulsations and have hard 1-10 keV spectra (i.e. with a power-law energy index of $\\Gamma\\sim0-1$; e.g. White, Nagase \\& Parmar 1995, Yokogawa \\etal 2003), which are signatures of accretion onto strongly magnetized neutron stars. In contrast, SG-XRBs have generally softer spectra, and do not always have a pulsar as the compact object.   Be-XRBs are the most numerous sub-class of HMXBs. In the Milky Way as well as in the Large Magellanic Cloud (LMC), they constitute 60-70\\% of all HMXBs (Sasaki, Pietsch \\& Haberl 2003), while in the Small Magellanic Cloud (SMC) only one out of 92 known or probable HMXBs (Liu, van Paradijs \\& van den Heuvel 2005) is a supergiant system.   The SMC is an excellent laboratory to study the Be-XRB populations. It is the second nearest star-forming galaxy, and it has a large number of Be-XRBs seen through moderate Galactic foreground absorption ($N_{H}\\simeq6\\times10^{20} cm^{-2}$; Dickey \\& Lockman 1990), and well mapped extinction (Zaritsky \\etal 2002). In addition, its well measured distance (60 kpc; e.g. van den Bergh 2000, Hilditch, Howarth \\& Harries 2005), small line-of-sight depth of the young populations at its main body ($<$10 kpc; e.g. Crowl \\etal 2001; Harries, Hilditch \\& Howarth 2003), well studied recent star-formation (SF) history (Harris \\& Zaritsky 2004), and generally uniform metallicity of the young populations facilitate the interpretation of the results. Therefore, one can study in great detail and in an homogeneous way the faint end of the XRB populations.   Several studies of the optical characteristics of the XRB systems in the SMC, have been published in the last few years. Haberl \\& Sasaki (2000), based on the ROSAT surveys of the SMC, identified 25 X-ray sources as new Be-XRBs. Haberl \\& Pietsch (2004) extended this work presenting 65 SMC HMXBs, of which 45 are associated with an emission-line object indicating that they are Be-XRBs. In the latest census of HMXBs by Liu \\etal (2005) 62 Be-XRBs are listed (for 92 confirmed or proposed HMXBs). In addition out of the 38 currently known X-ray pulsars in the SMC, 34 have identified optical counterparts (Coe \\etal 2005a).   In this paper we study in a systematic way the young SMC XRB populations. The SMC ROSAT survey reached a non-uniform detection limit of $\\sim5\\times10^{34}-10^{35}$\\ergs\\, (e.g. Kahabka \\& Pietsch 1996, Haberl \\etal 2000, Sasaki \\etal 2000). Here we are using data down to $\\sim4\\times10^{33}$\\ergs\\, from the \\chandra survey of the SMC (Zezas \\etal 2003) that allow to investigate the faintest of the HMXB populations.    \\subsection{The \\chandra survey of the SMC}\\label{Chandrasurvey}   The SMC as a prime target to study the faint end of the XRB populations (with typical ${\\rm L_{x}}<10^{34}$\\lunit), overcoming in this way the inherent problems of observations in the Milky Way (e.g. distance determination, and obscuration). For this reason we initiated a \\chandra survey, consisting of five observations of the central part of the SMC performed between May and October 2002 with the \\chandra ACIS-I detector (Advanced CCD Imaging Spectrometer; Garmire \\etal 2003). The survey covers an area of $1280\\,arcmin^{2}$ along the central, most actively star forming, region of the SMC (also referred to as the SMC ``Bar'', although it bears no relation to a dynamical bar; van den Bergh, 2000).  In Figure \\ref{DSS} we plot the footprints of the \\chandra and OGLE-II fields (for the latter see \\S \\ref{opticalcat}) overlaid on a Digitized Sky Survey (DSS) optical image of the main body of the SMC.   The survey yielded a total of 158 sources, detected on the ACIS-I CCDs, down to a limiting luminosity of $\\sim4\\times10^{33}$\\lunit, which is within the luminosity range of quiescent HMXBs (typical ${\\rm L_{x}}\\sim10^{32}-10^{34}$\\lunit\\,; van Paradijs \\& McClintock 1995). Each of the 5 fields contained between 8 and 16 sources with fluxes at least of $3\\sigma$ significance (assuming Gehrels statistics; Gehrels 1986) above their local background. A brief description of the survey and the first results are presented in Zezas \\etal (2003). The final source-list (including the X-ray luminosity functions) and the spectral properties of the sources are presented in Zezas \\etal (2008, in prep.). Because of the large off-axis angle (and therefore large positional uncertainty) of sources on the ACIS-S3 and ACIS-S4 CCDs we do not include them in the present work.   The absolute astrometric accuracy of on-axis sources is dominated by the boresight error of \\chandra ($<0.6\\arcsec$ at the 90\\% confidence level\\footnotemark)\\footnotetext{http://cxc.harvard.edu/cal/docs/cal\\_present\\_status.html\\#abs\\_spat\\_pos}. However, at larger off-axis angles the Point Spread Function (PSF) becomes broader and asymmetric, resulting in larger positional uncertainties (see \\chandra POG, 2005). To estimate the positional errors of the off-axis sources, we used the empirical formula of Kim \\etal (2004), which gives positional uncertainties for sources with 20 and 100 counts as a function of their off-axis angle (up to a maximum of $10\\arcmin$). For sources with net number of counts within the range we linearly interpolated between these two estimates, while for sources with less than 20 or more than 100 counts we used the appropriate branch of the equation.   Five sources out of the 158, which were at large off-axis angles ($\\rm{D_{off-axis}>10}\\arcmin$), yielded high positional errors and were discarded from further consideration in this paper. The \\chandra field IDs, and the coordinates of their centers are listed in Table \\ref{tablesurvey} (Columns (1)-(3)). In Column (4) we present the number of all detected X-ray sources in each \\chandra field, while in parenthesis we give the number of sources with fluxes at $3\\sigma$ or greater level above their local background. In Column (5) we give the number of \\chandra sources for which we searched for counterparts (i.e. detected in $\\rm{D_{off-axis}<10}\\arcmin$).   The identification and classification of the optical counterpart of an X-ray source allow us to identify the interlopers, while it is the only way to confirm the different types of XRBs (e.g. Be-XRBs, SG-XRBs, LMXBs). Although the classification of the counterparts is only secure via optical spectroscopic observations, color-magnitude and/or color-color diagrams (hereafter CMDs and 2-CDs, respectively) can be used to obtain rough classifications for the optical counterparts, and at least identify HMXBs. From the appropriate isochrones and the stellar tracks we can also obtain a feel for the mass and the age of the donor star, if it dominates the optical emission of the X-ray sources.   In this paper we have used the \\chandra survey of the SMC (Zezas \\etal 2003), and searched for optical counterparts of the X-ray sources included in this survey. In particular, we confirm or suggest new counterparts of the above X-ray sources, based on their optical properties, and when possible we give a tentative classification. In Section \\ref{opticalcat} we present the optical catalogs used in this study, and in Section \\ref{cc} the counterparts of the X-ray sources. We also present their finding charts (Section \\ref{charts}), and we examine the photometry of the optical data (Section \\ref{photometry}). In Section \\ref{CMDs} we describe the data used in the construction of the $V$, $B-V$ CMD for the classification of the sources, and we estimate the expected contamination by foreground stars, while in Section \\ref{chance} we estimate the chance coincidence probability. In the discussion (Section \\ref{discussion}) we present the criteria used to classify the X-ray sources, and distinguish among multiple matches the most likely counterpart. We also discuss the properties of the optical counterparts, and the relative numbers of Be-XRBs in the SMC, LMC and Milky Way. In the Appendix we present notes on individual sources. ",
        "conclusions": "In this paper we report the results for the optical counterparts of 153 X-ray sources detected in the \\chandra survey of the central region of the SMC. For 120 sources we find optical matches in the OGLE-II and/or MCPS catalogs. Using these photometric data, we propose the most likely optical counterpart for 113 \\chandra sources, while for 7 sources the candidate counterparts are equally likely. By combining their optical and X-ray properties (i.e. spectral and timing properties), we classify 34 \\chandra sources of which 13 were previously unclassified. In particular:  \\begin{enumerate}  \\item{We find that 52 X-ray sources have a single counterpart, 68 have two or more matches, while 33 sources do not have counterparts in either OGLE-II or MCPS catalog within the $1\\sigma$ search radius (7 out of these 33 sources have bright optical matches ($M_{{V}_{o}}\\leq-0.25$) between the $1\\sigma$ and $2\\sigma$ search radii, but none of them is associated with an OB star).}  \\item{Early-type counterparts (of OB spectral type) have been identified for 35 X-ray sources (with a chance coincidence probability $\\leq19\\%$). Based on spectroscopic observations of stars in the SMC we define the photometric locus of early-type (OB) stars in the $M_{{V}_{o}}$ vs. $(B-V)_{o}$ CMD. Based on the position on the CMD and their X-ray spectral properties, we propose: (a) 4 new Be-XRBs, and (b) 7 new HMXBs, and 2 new candidate HMXBs. Two of the new HMXBs have hard X-ray spectra and early-type (OB) counterparts within their $2\\sigma$ search radius.}  \\item{Based on the isochrones of Lejeune \\& Schaerer (2001), we estimate the age of the 13 new (confirmed and candidate Be-XRBs and HMXBs) to be $\\sim15-85$ Myr, in agreement with the age of Be stars.}   \\item{Twenty X-ray sources with bright counterparts ($M_{{V}_{o}}\\leq-0.25$) could not be classified, because we do not have X-ray spectral information for them and/or they have other than OB spectral-type counterparts.}   \\item{We find that the mixing of Be-XRBs with other than their natal stellar population is not an issue in our comparisons of Be-XRBs and stellar populations in the SMC, because the SF activity across the SMC ``Bar'' is generally uniform in scales larger than the size of the \\chandra fields. Instead we find indication for variation in the XRB populations across the ``Bar'' and in scales of $\\sim1$ kpc: fields with higher SF rate at ages which are more prone to produce Be stars show increased number of Be-XRBs.}   \\item{Using the catalogs of Liu \\etal (2005, 2006) we find that for luminosities down to ${\\rm L_{x}}\\sim10^{34}$\\ergs\\, (chosen as a moderate minimum cut-off between different SMC and Galactic X-ray surveys) the \\chandra SMC fields contain $\\sim1.5$ times more Be-XRBs in comparison to the Milky Way (including the new confirmed and candidate Be-XRBs from this study), even after taking into account the different OB star formation rates in the two galaxies. This residual excess can be explained when we account for the different metallicity of the SMC and the Milky Way. This is in good agreement with theoretical predictions derived from population synthesis studies (Dray 2006), and previous observational (both photometric and spectroscopic) works (e.g. Wisniewski \\& Bjorkman 2006; Martayan \\etal 2006, 2007).}  \\end{enumerate}"
    },
    "0812/0812.1156_arXiv.txt": {
        "abstract": "{We observed the blazar \\object{0716+714} with the VLBA during its active state in 2003-2004. In this paper we discuss multi-frequency analysis of the inner jet (first 1~mas) kinematics. The unprecedentedly dense time sampling allows us to trace jet components without misidentification and to calculate the component speeds with good accuracy. In the smooth superluminal jet we were able to identify and track three components over time moving outwards with relatively high apparent superluminal speeds (8.5-19.4\\,$c$), which contradicts the hypothesis of a stationary oscillating jet in this source. Component ejections occur at a relatively high rate (once in two months), and they are accompanied by mm-continuum outbursts. Superluminal jet components move along wiggling trajectories, which is an indication of actual helical motion. Fast proper motion and rapid decay of the components suggest that this source should be observed with the VLBI at a rate of at least once in one or two months in order to trace superluminal jet components without confusion.} ",
        "introduction": "The blazar S5 \\object{0716+714} is one of the most active sources of its class. It is highly variable on time scales from hours to months at all observed wavelengths from radio to $\\gamma$-rays. Intra-day variability (IDV) was detected in the source in the optical, millimetre, and radio bands \\citep[e.g.][and references therein]{montagni2006, agudo2006}. The IDV in the different bands of electromagnetic radiation is strongly correlated \\citep{quirrenbach1991, wagner1996, stalin2006, fuhrmann2008}, which suggests its intrinsic origin. Redshift of the source is 0.31, as recently estimated by \\citet{Nilsson_z2008}, who observed the host galaxy in the I-band during the quiescent state of the active nucleus.  Very long baseline interferometry studies (VLBI) show a core-dominated jet pointing to the north \\citep{jorstad2001, bach2005}, and VLA data show a halo-like jet misaligned with it by $\\sim$~90$^{\\circ}$ \\citep{wagner1996}. There are several scenarios for the milliarcsecond scale jet kinematics. The most recent studies propose two very different interpretations for the long time series ($\\sim$10~yr) of centimetre VLBI observations. One scenario consists of fast superluminal components moving along the jet with speeds ranging from 5\\,$c$ to 16\\,$c$ \\citep{jorstad2001, bach2005}. An alternative interpretation suggests that instead of travelling along the jet, components are ``oscillating'' around their mean position due to the jet precession. Component speeds, calculated taking position angle changes into account, range from 5\\,$c$ to 10\\,$c$ \\citep{britzen2006}.  This paper is the first one in a series of papers that describe changes in the parsec-scale structure of \\object{0716+714} during the active period in 2003-2004. Here we present the kinematics of the inner jet of \\object{0716+714}. In Paper II we will discuss the suitability of different data imaging methods for restoration of the source structure in the case of a very smooth brightness distribution along the jet. Source polarisation, quantitative analysis of the flux density evolution and spectral properties will be published in the subsequent papers. Here we use a standard cosmological model with a flat universe and values of Hubble constant $H_{o}$=71\\,km~s$^{-1}$Mpc$^{-1}$, energy density of matter $\\Omega_{m}$=0.3, and cosmological constant energy density $\\Omega_{\\Lambda}$=0.7. ",
        "conclusions": "We studied the kinematics of the inner jet of \\object{0716+714} at 43~GHz. Even though the jet has a very smooth brightness distribution, we identified four components: three moving and one stationary. Components are well-defined at several frequencies and traceable over time. Components move along the jet with superluminal speeds ranging from 8.5\\,$c$ to 19.4\\,$c$, which contradicts the model of an oscillating jet.  Ejections of new superluminal components occur at a relatively high rate, once in two months, and a component life time is approximately ten months. Therefore, future observers have to note that for accurate structure evolution studies, \\object{0716+714} has to be observed at a rate of once in one or two months. Superluminal component ejections coincide with outbursts in a millimetre continuum.  We also noticed that components move along wiggling trajectories in the jet, which might be a signature of their helical motion. Careful quantitative tests of models that predict physically helical component motion (e.g. two-fluid model and helical jet model) should be performed in the future in order to understand nature of the jet in this object."
    },
    "0812/0812.3153_arXiv.txt": {
        "abstract": "We study the axisymmetric and non-axisymmetric, time-dependent hydrodynamics of gas that is under the influence of the gravity of a super massive black hole (SMBH) and the radiation force produced by a radiatively efficient flow accreting onto the SMBH. We have considered two cases: (1)~the formation of an outflow from the accretion of the ambient gas without rotation and (2)~that with \\textcolor{black}{weak} rotation. The main goals of this study are: (1)~to examine if there is a significant difference between the models with identical initial and boundary conditions but in different dimensionality (2-D and 3-D), and (2)~to understand the gas dynamics in AGN.  Our 3-D simulations of a non-rotating gas show small yet noticeable non-axisymmetric small-scale features inside the outflow.  The outflow as a whole and the inflow do not seem to suffer from any large-scale instability. In the rotating case, the non-axisymmetric features are very prominent, especially in the outflow which consists of many cold dense clouds entrained in a smoother hot flow. The 3-D outflow is non-axisymmetric due to the shear and thermal instabilities. In both 2-D and 3-D simulations, gas rotation increases the outflow thermal energy flux, but reduces the outflow mass and kinetic energy fluxes. Rotation also leads to time variability and fragmentation of the outflow in the radial and latitudinal directions. The collimation of the outflow is reduced in the models with gas rotation. The time variability in the mass and energy fluxes is reduced in the 3-D case because of the outflow fragmentation in the azimuthal direction.  The virial mass estimated from the kinematics of the dense cold clouds found in our 3-D simulations of rotating gas underestimates the actual mass used in the simulations by about 40~\\%.  The opening angles ($\\sim30^{\\circ}$) of the bi-conic outflows found in the models with rotating gas are very similar to that of the nearby Seyfert galaxy NGC~4151 ($\\sim33^{\\circ}$).  The radial velocities of the dense cold clouds from the simulations are compared with the observed gas kinematics of the narrow line region of NGC~4151. ",
        "introduction": "\\label{sec:Introduction}  Active Galactic Nuclei (AGNs) are powered by accretion of matter onto a super massive ($10^{6}$--$10^{10}\\,\\Msun$) black hole (SMBH), and produce a large amount of energy (e.g., \\citealt{Lynden-Bell:1969}) as electromagnetic radiation ($10^{10}$--$10^{14}\\Lsun$), over a wide range of wavelengths (from radio to hard X-rays, and even to TeV photons).  The strong radiation from AGNs influences the physical properties (e.g., the ionization structure, gas dynamics and density distribution) of their vicinity, their host galaxies, and even of the inter-galactic material of galaxy clusters to which they belong (e.g., \\citealt*{Quilis:2001}; \\citealt{DallaVecchia:2004}; \\citealt{McNamara:2005}; \\citealt{Zanni:2005}; \\citealt{Fabian:2006}; \\citealt{Vernaleo:2006}). The importance of the radiation-driven outflows from AGNs as a feedback process is recognized in many of the galaxy formation/evolutionary models (e.g., \\citealt{ciotti:1997}, \\citeyear{ciotti:2001}, \\citeyear{ciotti:2007}; \\citealt{king:2003}; \\citealt{Hopkins:2005}; \\citealt*{Murray:2005}; \\citealt{Sazonov:2005}; \\citealt*{Springel:2005}; \\citealt{Brighenti:2006}; \\citealt{Fontanot:2006}; \\citealt*{Wang:2006}, \\citealt*{Tremonti:2007}; Ciotti et al.~2008, in preparation).  The formation of AGN outflows, of course, can be caused by some mechanisms other than radiation pressure, e.g., magnetocentrifugal force (e.g., \\citealt{Blandford:1982}; \\citealt*{Emmering:1992}; \\citealt{Konigl:1994}; \\citealt{Bottorff:1997}), \\textcolor{black}{ Poynting flux/magnetic towers (e.g., \\citealt{Lovelace:1987}; \\citealt{Lynden-Bell:1996,Lynden-Bell:2003}; \\citealt{Li:2001}; \\citealt{Kato:2004}; \\citealt{Nakamura:2006}; \\citealt{Kato:2007}), } and thermal pressure (e.g., \\citealt{Weymann:1982}; \\citealt*{Begelman:1991}; \\citealt{Everett:2007}).  However, the highly blueshifted line absorption features often seen in the observed UV and optical spectra of AGNs can be best described by the radiation-driven wind models (e.g.,~\\citealt{Murray:1995b}; \\citealt{Proga:2000}; \\citealt{Proga:2004}), provided that the ionization state of the gas is appropriate.  In reality, these forces may interplay and contribute to the dynamics of the outflows in AGNs in somewhat different degrees (e.g., \\citealt{Konigl:2006}; \\citealt{Proga:2007a}, and references therein).  The AGN environment on relatively large scales ($10^{2}-10^{3}$~pc) is a mixture of gas and dust (e.g.~\\citealt{Antonucci:1984}; \\citealt{Miller:1990}; \\citealt{Awaki:1991}; \\citealt{Blanco:1990}; \\citealt{Krolik:1999}).  The radiation pressure on dust can drive the dust outflows, and their dynamics is likely to be coupled with the gas dynamics (e.g., \\citealt{Phinney:1989}; \\citealt{Pier:1992}; \\citealt{Emmering:1992}; \\citealt{Laor:1993}; \\citealt{Konigl:1994}; \\citealt{Murray:2005}). On much smaller scales ($<\\sim10$~pc), the dust is less likely to survive because the temperature of the environment is too high ($>10^{4}\\,\\Kelvin$); hence, the studies of the radiation-driven outflow dynamics using only gas component (e.g.,~\\citealt*{Arav:1994}; \\citealt{Proga:2000}) would be justified in those cases.  In the first paper of this series (\\citealt{Proga:2007b}, hereafter Paper~I), the initial phase of our gas dynamics studies of AGNs on sub-parsec and parsec scales was set.  Since the problem is complex, as it involves many aspects of physics such as multi-dimensional fluid dynamics, radiative processes, and magnetic processes, our approach was to set up simulations as simple as possible. The study focused on exploring the effects of X-ray heating (which is important in the so-called preheated accretion; e.g., \\citealt{Ostriker:1976}; \\citealt{Park:2001, Park:2007}) and radiation pressure on gas that is gravitationally captured by a black hole (BH). We adopted the numerical methods developed by \\citet{Proga:2000} for studying radiation-driven disk winds in AGNs. Our simulations covered a relatively unexplored range of the distance from the central BH, i.e., the outer boundary of our simulations coincides with the inner boundary of many galaxy models (e.g., \\citealt{Springel:2005}; \\citealt{ciotti:2007}), and our inner boundary starts just outside of the outer boundary of many BH accretion models (e.g., \\citealt{Hawley:2002}; \\citealt{Ohsuga:2007}). The effect of gas rotation was not included in Paper~I.   In the second paper in this series (\\citealt{Proga:2008}, hereafter Paper~II), the effect of gas rotation, position-dependent radiation temperature, density at large radii, and uniform X-ray background radiation were explored. As in the non-rotating case considered in Paper~I, the rotating flow settles into a configuration with two components: (1)~an equatorial inflow and (2)~a bipolar inflow/outflow with the outflow leaving the system along the polar axis. However, with rotation the flow does not always reach a steady state.  In addition, rotation reduces the outflow collimation and the outward fluxes of mass and kinetic energy. Moreover rotation increases the outward flux of the thermal energy, and it can lead to fragmentation and time-variability of the outflow. It is also shown that the position-dependent radiation temperature can significantly change the flow solution, i.e.,  the inflow in the equatorial region can be replaced by a thermally driven outflow. As it \\textcolor{black}{has} been discussed and shown in the past (e.g., \\citealt{ciotti:2007}; Ciotti et al.~2008, in preparation), the self-consistently determined preheating/cooling from the quasar radiation can significantly reduce the mass accretion rate of the central BH. Our results clearly demonstrated that quasar radiation can drive non-spherical, multi-temperature and very dynamic flows. This effect becomes dominant for the systems with luminosity in excess of 0.01 times the Eddington luminosity.   The work presented here is a direct extension of the previous axi-symmetric models of Paper~I and Paper~II to a full 3-D model, and is an extended version of the 3-D models presented in \\citet{Kurosawa:2008} to which we have added the radiation force due to spectral lines and the radiative cooling and heating effect.  Here, we consider two cases from Paper~I and Paper~II: (1)~the formation of \\textcolor{black}{relatively large scale ($\\sim 10$pc)} outflows from the accretion of the ambient gas with no rotation and (2)~that with rotation, in 3-D. We note that our work is complimentary to the work by \\citet{Dorodnitsyn:2008a, Dorodnitsyn:2008b} who studied the hydrodynamics of axisymmetric torus winds in AGNs.  The main goals of this study are (1)~to examine if there is a significant difference between two models with physically identical conditions but in different dimensionality (2-D and 3-D), (2)~to study if the radiation driven outflows that were found to be stable in the previous studies in 2-D (Paper~I; Paper~II) remain stable in 3-D simulations, and (3)~to understand gas dynamics in AGNs, in particular the dynamics of the narrow line regions (NLR) by comparing our simulation results with observations.   In the following section, we describe our method and model assumptions. We give the results of our hydrodynamical simulations in \\S~\\ref{sec:Results}. Discussions on virial mass estimates and comparisons with the observations of Seyfert galaxies will be given in \\S~\\ref{sec:Discussions}. The summary and conclusions are in \\S~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  We have presented the dynamics of gas under the influences of the gravity of a SMBH and the radiation force from the luminous accretion disk around the SMBH. This is a direct extension of the previous axi-symmetric models of Paper~I and Paper~II to a full 3-D model, and is an extended version of the models presented in \\citet{Kurosawa:2008} to which we have added the radiation force due to line processes and the radiative cooling and heating effect. We have considered two cases from Paper~I and Paper~II: (1)~the formation of outflow  from the accretion of the ambient gas with no rotation and (2)~that with \\textcolor{black}{weak} rotation. The models have been considered in both 2-D and 3-D hence, in total, four models have been presented. Our first main goal is to examine if there is a significant difference between two models with identical initial and outer boundary conditions but in different dimensionality (2-D and 3-D). In particular, we examine whether the radiation driven outflows that were found to be stable in the previous studies in 2-D (Paper~I; Paper~II) still remain stable in 3-D. Our second main goal is to gain some insights into the gas dynamics in AGNs and Seyfert galaxies by comparing the simulation results with observations. In the following, we summarize our main findings through this investigation.  1.~For non-rotating gas cases, the outflow occurs in very narrow cones (with the opening angles $\\sim5^{\\circ}$) in polar directions. Overall density and temperature of the both 2-D and 3-D models (Models~I and III) are very similar to each other (Figs.~\\ref{fig:Density-3d} and \\ref{fig:Density-Temp-Map-I-III}). Small but noticeable differences are seen in the narrow outflow regions.  2.~Rotation of gas significantly changes the morphology of the outflows (Models~II and IV in Figs.~\\ref{fig:Density-3d} and \\ref{fig:Density-Temp-Map-II-IV}). The centrifugal force pushes the outflow away from the polar axis and forms much wider outflows (with the opening angles $\\sim30^{\\circ}$). The outflow occurs mainly on and near bi-conic surfaces, and relatively low values of density are found in the polar directions, unlike the outflows in the non-rotating cases. The models with gas rotation show cold clouds (clumps) in their outflows in their 2-D density and temperature maps (Fig.~\\ref{fig:Density-Temp-Map-II-IV}). Although the overall density and temperature structures of the flows of the 2-D and 3-D models are similar to each other, the outflows in 3-D occur in much less organized manner. We find that the cloud-like structures seen in the 2-D model (Model~II), which are rings if the density is expanded in 3-D using the axisymmetry (Fig.~\\ref{fig:Density-3d}), are not stable in full 3-D simulations due to the shear and thermal instabilities. The rings break up into smaller pieces, and fully 3-D clouds are formed in Model~IV.   3.~The mass and energy fluxes plotted as a function of radius for the 3-D non-rotating case are almost identical to those of the non-rotating 2-D case (Fig.~\\ref{fig:mass-energy-flux}). For the rotating cases, the bumps seen in the mass-inflow rate and the net mass flux at the outer radii ($r' \\gtrsim 10^{4}$) for the 2-D model (Model II) are smoothed out in the 3-D model (Model~IV) due to the fragmentation of the ring structures in the 3-D model. While the kinetic power dominates at all radii for the non-rotating cases, the thermal power contributes significantly to the outflow driving force for the rotating cases. In spite of the differences in the flow geometries, the rotating models in both 2-D and 3-D show very similar values of the mass accretion and outflow rates at the outer and inner boundaries (Table~\\ref{tab:Model-Summary}). In other words, AGN feedback due to radiation is similar in the 2-D and 3-D cases as far as the time-averaged mass and energy fluxes are concerned.  4.~For the non-rotating cases, the amount of variability in the mass flux at the outer boundary is higher in the 2-D model than that in the 3-D model, but the opposite is true for the rotating cases (Fig.~\\ref{fig:mdot-evolution}).  5.~\\textcolor{black}{In the 3-D models}, the deviations from the axisymmetry are observed in both rotating and non-rotating cases (Figs.~\\ref{fig:rho-temp-vr-Model-I-III} and \\ref{fig:rho-temp-vr-Model-II-IV}). The amounts of the azimuthal angle variations of the density, temperature, and radial velocity (Fig.~\\ref{fig:phi-variation}) are relatively small for the non-rotating case (Model~III), but they are relatively large at all radii for the rotating case (Model~IV).  6.~The gas properties of the 2-D and 3-D models are very similar to each other for both non-rotating and rotating cases (Figs.~\\ref{fig:phase-T-xi}, \\ref{fig:phase-vr-xi} and \\ref{fig:phase-vr-temp}). The majority of the outflowing gas for the rotating cases (Models~II and IV) has relatively large values of the photoionization parameter ($\\xi>10^{6}$) while for the non-rotating cases, it has relatively small values of the photoionization parameters ($\\xi<10^{2}$) (Fig.~\\ref{fig:phase-vr-xi}). This is due to the difference in the dominant outflow mechanisms between the non-rotating and the rotating cases, i.e., the outflow is mainly radiatively driven for the non-rotating cases while the thermal pressure significantly contributes to the outflows of the rotating cases (cf., Fig.~\\ref{fig:mass-energy-flux}). For the rotating models, the majority of the outflowing gas has relatively high ($T>10^{6}$~K) temperature while for the non-rotating cases, it has relatively low ($T<10^{5}$~K) temperature. The higher $\\xi$ values seen in the rotating cases are mainly from the \\textcolor{black}{low-density} hot outflowing gas in between the outflowing cold clouds.  7.~For Model~IV, we find the average speed and the radial position of the cold cloud elements (\\S~\\ref{sub:Virial-Mass}) as $V=285\\,\\mathrm{km\\, s^{-1}}$ and $R=1.00\\times10^{19}\\,\\mathrm{cm}$. The corresponding viral mass is $M_{\\mathrm{vir}}=1.22 \\times10^{41}\\,\\mathrm{g}$ which is about 40~\\% smaller than the actual mass of the BH used in the simulation, i.e., $M_{\\mathrm{BH}}=1.989\\times10^{41}\\,\\mathrm{g}$. This is in general agreement with the previous studies (e.g.~\\citealt{Peterson:2000}; \\citealt{Krolik:2001}; \\citealt{Marconi:2008}) which predict that the virial mass estimated without considering the effect of the radiation force underestimates the actual mass of the SMBH.  8.~The opening angles ($\\sim30^{\\circ}$) of the bi-conic outflows found in the the rotating models (Models~II and IV) are very similar to that of the nearby Seyfert galaxy NGC~4151 ($33^{\\circ}$) determined by \\citet{das:2005}. Although the physical size of the long slit observations of NGC~4151 by \\citet{das:2005} is in much lager scale ($\\sim50$ times larger) than that of our model, their radial velocities as a function of the position along the slit (see their Figs.~5 and 6) show a similar pattern as in our model (Fig\\@.~\\ref{fig:v-projected-das}). An important difference between the observation of \\citet{das:2005} and our models is the lack of clearly decelerating clouds at larger radii in our models. However, we note that the clouds found in our simulations reach a constant velocity near the outer boundary of our simulations, and show a hint of deceleration. This puzzling outflow deceleration seen in the observations might be due to the inflow that interacts with the polar outflows. The reason for no clear cloud deceleration seen in our model may be simply due to the relatively small simulation box size we used, and the issue could be resolved in a lager scale simulation. Spectroscopic studies of the NLR of Seyfert galaxies by \\citet{Komossa:2008} also favor a scenario in which the NLR clouds are traveling in decelerating wind. The hot outflowing gas, on the other hand, does show deceleration at the larger radii in our models (cf.~Fig.~\\ref{fig:phi-variation}).  To perform a better comparison of our models with observations hence to constrain the model parameters, in future studies, we need to increase the size of the simulation box to match the physical sizes of the NLR of Seyfert galaxies. It would take the outer radius of the computational domain to be expanded by one or even two orders of magnitude compared to the one used here. The dust is very likely important in the dynamics of the outflow in the larger scale simulations since the temperature becomes low enough for the dust survival and formation in the larger radius. We showed in Paper~II, relatively high density set at the outer boundary promotes formation of cold clouds. Therefore, we plan to explore the effects of dust and \\textcolor{black}{outer boundary} density.  To compare the model results directly with observations, we would need to compute the radiative transfer models of the important emission lines (e.g., {[}\\ion{O}{3}]~$\\lambda5007$, H$\\beta$ and \\ion{C}{4}~$\\lambda1549$), which will be the topic of our future paper. Taking these steps will allow a quite strict test of our results against observations of Seyfert galaxies and AGN. It would also be interesting to check if our models could reproduce large-scale outflows in quasars, for example the high-velocity outflow components seen in \\ion{C}{4} and \\ion{Mg}{2} quasar absorption-line systems (e.g., see a recent work by \\citealt{Wild:2008}) which \\textcolor{black}{would} provide an additional \\textcolor{black}{constraint} on our wind models."
    },
    "0812/0812.0293_arXiv.txt": {
        "abstract": "{The Gaia mission is expected to provide highly accurate astrometric, photometric, and spectroscopic measurements for about $10^9$ objects. Automated classification of detected sources is a key part of the data processing. Here a few aspects of the Gaia classification process are presented. Information from other surveys at longer wavelengths, and from follow-up ground based observations will be complementary to Gaia data especially at faint magnitudes, and will offer a great opportunity to understand our Galaxy. ",
        "introduction": "ESA's Gaia mission, to be launched in 2011, is meant to obtain accurate position, parallax and  proper motion for 10$^9$ object all over the sky, up to magnitude $G$=20 with an astrometric accuracy at the $\\mu$arcsec level. The low-dispersion spectroscopy obtained in the BP and RP passbands (330--1000 nm, resolution of 3--30 nm) will be used not only to correct the astrometry for color effects, but also to obtain a characterization of the sources themselves. The Radial Velocity Spectrograph (RVS, 840--890 nm, R$_{\\rm P}$=11\\,500) will measure radial velocities, with a precision of few km s$^{-1}$, up to $G$=16. Gaia will observe the whole sky for five years (plus a possible year extension) achieving a mean of $\\sim$ 80 observations for each source. The final catalog will be available in 2020, preceded by an early data release.  The data reduction is a great challenge: the size of Gaia related data will be about 10$^{15}$ bytes, while the final data delivered to the community would be of about 20 TB. The estimated computation size will be of the order of $10^{21}$ Flops \\citep[see][]{mignard}. ",
        "conclusions": ""
    },
    "0812/0812.0546_arXiv.txt": {
        "abstract": "The Short GAmma Ray Front Air Cherenkov Experiment (SGARFACE) uses the Whipple 10~m telescope to search for bursts of $\\gamma$ rays. SGARFACE is sensitive to bursts with duration from a few ns to $\\sim$20~$\\mu$s and with $\\gamma$-ray energy above 100~MeV. SGARFACE began operating in March 2003 and has collected 2.2 million events during an exposure time of 2267~hours. A search for bursts of $\\gamma$ rays from explosions of primordial black holes (PBH) was carried out. A Hagedorn-type PBH explosion is predicted to be visible within 60~pc of Earth. Background events were caused by cosmic rays and by atmospheric phenomena and their rejection was accomplished to a large extent using the time-resolved images. No unambiguous detection of bursts of $\\gamma$ rays could be made as the remaining background events mimic the expected shape and time development of bursts. Upper limits on the PBH explosion rate were derived from the SGARFACE data and are compared to previous and future experiments. We note that a future array of large wide-field air-Cherenkov telescopes equipped with a SGARFACE-like trigger would be able to operate background-free with a 20 to 30 times higher sensitivity for PBH explosions. ",
        "introduction": "} SGARFACE is designed to detect short bursts of $\\gamma$ rays with energies above 100 MeV using the air-Cherenkov technique \\cite{Sgarface_Krennrich2000}. The design allows detection of bursts lasting up to $\\sim$20 $\\mu$s and has a directional accuracy of up to a few arc-minutes \\cite{Sgarface_Lebohec2005}. Bursts of this type could be produced by the evaporation of primordial black holes (PBH), from emission associated with giant radio pulses from pulsars \\cite{Lyutikov2008}, or by a high-energy component accompanying very short gamma-ray bursts.  Primordial black holes would have been the first objects formed by the gravitational collapse of inhomogeneities in the early universe \\cite{Zeldovich1967,Hawking1971,Carr_Hawking1974}. Inhomogeneities may have resulted from density fluctuations \\cite{Carr1975} present after inflation \\cite{Gilbert1995}, the presence of cosmic strings \\cite{Polnarev1991,MacGibbon_Brandenberger1998}, collapse of domain walls \\cite{Rubin2001}, or from bubble collisions during a phase transition \\cite{Hawking1982}. Consequently, even without a detection, limits on the PBH density provide important insight into many aspects of the physics of the early universe \\cite{Carr2005}, including: cosmological phase transitions \\cite{Kapusta2007} and references therein, structure formation \\cite{Afshordi2003,Khlopov2005}, mass clustering \\cite{Chisholm2006}, and magnetic monopoles \\cite{Stojkovic2005}.  Three effects have been used to search for PBHs:  (1) Gravitational effects: Black holes of mass $\\gtrsim 10^{-8}$ $M_\\odot$ ($2\\times10^{25}$ g) may be detectable by their gravitational effects in microlensing observations \\cite{Paczynski1986,Alcock1996} or from mass clustering measured with the Ly$\\alpha$ forest \\cite{Afshordi2003}. A significant number of heavier PBHs may have formed binary systems \\cite{Nakamura1997} and may be detected from the gravitational radiation emitted during coalescence. An upper limit on the BH density in the mass range from 0.2 to 1.0 $M_\\odot$  has been reported by the LIGO collaboration \\cite{Abbott2005}.  (2) Hawking radiation: Black holes of mass less than $10^{14}$~g may be detected by their emitted Hawking radiation \\cite{Hawking1974}. Black holes of mass, $M$, emit real particles with a black-body energy spectrum at temperature $T = 1.06 (10^{13} \\mathrm{g}/M)$ GeV \\cite{MacGibbon_Webber1990}. As mass accretion by PBHs is thought to be negligible \\cite{Custodio2002}, PBHs will evaporate completely in $10^{64}(M/M_\\odot)^3$ years, where $M_\\odot\\approx 2\\times10^{33}$g is the Solar mass. At the present time, black holes with an initial mass of $4.5 \\times 10^{14}$~g \\cite{MacGibbon1991} would be reaching the final phase of their evaporation which might be explosive. PBHs that have already evaporated, would have left an imprint on the cosmic particle backgrounds and set stringent limits on the PBH density; a discussion is given in Sect.~\\ref{sec:PBH_searches}. Lastly, PBHs may also be detectable as transient gamma-ray sources in the solar system \\cite{Page_Hawking1976,MacGibbon_Carr1991}.  (3) Accretion effects: The accretion onto light black holes may produce distinct observable radiation in the relatively dense environments of interstellar space, in binary systems, or near planets~\\cite{Trofimenko1990}.  The types of particles emitted by Hawking radiation and their energy spectra are well understood below the QCD confinement temperature of $T_c \\approx 175^{+5}_{-18}$ MeV \\cite{Braun_Munzinger2004}. Above this temperature, it is unknown whether individual elementary particles will be emitted in the strong gravitational field, or if a phase transition of the surrounding vacuum occurs and a thermodynamic model of the emission is valid. If the standard particle physics model (SM) is valid at all temperatures, the emission spectra can be determined from the quark-gluon jet fragmentation functions convolved with the Hawking emission function \\cite{MacGibbon_Webber1990,MacGibbon1991}. This would result in a slow evaporation at high temperatures: the final explosion would last about 1 s with $\\gamma$-rays of energies above 400 GeV being emitted \\cite{Halzen1991,Connaughton1998}. 22\\% of the total energy would be emitted in $\\gamma$ rays with a flux spectrum peaking at $\\sim$ 100 MeV and falling off as $E^{-3}$ at higher energies \\cite{MacGibbon1991}.  On the other hand, in thermodynamic models of the QCD phase transition point, the large number of hadronic resonances implies a large phase space and therefore the final evaporation becomes extremely rapid, resulting in an explosion. The most extreme thermodynamic model is that of \\cite{Hagedorn1965,Hagedorn1970}, referred to as the Hagedorn model (HM), where the density of states, $\\rho_s$, increases exponentially with mass: $\\rho_s \\propto m^{-5/2} e^{m/T_c}$. In this picture, the ground state of the vacuum contains elementary and composite particles, which are evaporated by the gravitational field of the PBH. During the final explosion, 10-30\\% of the mass at $T_c$ is converted into gamma-rays between 100 MeV and 1 GeV; the explosion lasts $\\sim$100~ns~\\cite{Page_Hawking1976}.  Instead of completely evaporating, PBHs may leave behind Planck mass relics, see \\cite{Carr2005} and references therein. Because of their small mass, Planck relics do not affect $\\gamma$-ray emission, but may possibly contribute to dark matter \\cite{MacGibbon1987,Barrau2004}. If extra dimensions exist below the currently measured scale of gravity, $\\sim$~10~$\\mu$m \\cite{Ferrari2006}, black hole evaporation would release less detectable energy~\\cite{Argyres1998,Nouicer2007}.  For SGARFACE, we consider the detection of PBHs via their final-stage $\\gamma$-ray radiation in the Hagedorn model. The nominal burst is assumed to come from a black hole of mass $6.5\\times 10^{13}$ g, corresponding to $T$ = 160 MeV. It is assumed that 20\\% of the mass-energy, or $7.5\\times 10^{45}$ eV, is emitted in the form of 1~GeV $\\gamma$ rays. Realistically, the $\\gamma$-ray spectrum would not be a $\\delta$-function, but a product of the direct $\\gamma$ emission with the decays from hadrons. It is straightforward to scale the upper limits on the PBH explosion rate presented in Section~\\ref{sec:bursts} with the total burst energy and the average energy of $\\gamma$ rays, see Section~\\ref{sec:volume}.  \\subsection{Summary of PBH Searches \\label{sec:PBH_searches}} Limits on the PBH explosion rate have been placed by direct searches for the final stage emission and require a particular particle physics model. Limits on the PBH density that have been derived from cosmic-particle background spectra require an assumed initial PBH mass spectrum. To compare the results, assumptions need to be made on PBH clustering in the galactic halo \\cite{Halzen1991,Wright1996,Chisholm2006} and, in the case of charged particle backgrounds, the enhancement due to confinement by the Galactic magnetic field \\cite{MacGibbon_Carr1991}.  Air-Cherenkov telescopes (ACT), due to their large collection area, but small field of view of $\\sim$4\\degr, are well suited to search for individual explosions. The sensitivity to detect standard model PBH explosions is limited to distances of 0.4 pc \\cite{Connaughton1998}, and to about 200 pc for the HM bursts. Upper limits on the explosion rate are shown in Fig.~\\ref{fig:intro_limits}, where Hagedorn burst models have been assumed at time scales shorter than 1 $\\mu$s and SM bursts at longer time scales. The first search for HM bursts was carried out by \\cite{Porter_Weekes1977,Porter_Weekes1978,Porter_Weekes1979}, who arrived at an upper limit for bursts of 150 ns duration. The results of their two experiments were carefully reanalyzed using air-shower simulations to characterize the trigger threshold across the field of view. The revised value of their more sensitive, short-baseline, experiment is shown in Fig.~\\ref{fig:intro_limits}.  The first search for SM bursts of duration between 0.01~s and 1~s and energy above 7~TeV was carried out by \\cite{Fegan1978} with a pair of scintillation particle detectors separated by 250~km. Several groups, using single or closely spaced detectors, have set limits at other energies and duration \\cite{Bhat1980_gamma,Nolan1990,Alexandreas1993,Amenomori1995,Connaughton1998,Linton2006}. Some of the limits shown in Fig.~\\ref{fig:intro_limits} were standardized to the same SM burst by \\cite{Halzen1991}. Satellites, though they have a small collection area, are useful in searching for PBH explosions because they are essentially background-free and have a large field of view. EGRET data was searched for multiple $\\gamma$-ray tracks within its trigger window of 600$\\pm$100~ns \\cite{Fichtel1994}. Using BATSE data, a possible association of PBH explosions was made with some short (6-200~ms) X-ray bursts on the basis of the spectral evolution and spatial isotropy \\cite{Cline1997}.  The evaporation of PBHs over the lifetime of the universe would have left an imprint on the observed cosmic particle spectra: $\\gamma$, $e^{\\pm}$, $p$, $\\bar{p}$, and $\\nu$.% \\footnote{We use the following constants: $H_0$=71~km/s/kpc, $13.6\\times10^9$~yr as the age of the universe, and a critical density of $1\\times10^{-29}$~g/cm$3$.} The underlying assumption in placing limits on the PBH density using particle backgrounds is that the initial PBH mass distribution arose from scale-invariant density fluctuations \\cite{Page_Hawking1976}. Using the measured anisotropy in the EGRET $\\gamma$-ray background above 100~MeV, \\cite{Wright1996} determined the rate of PBH explosions to be less than 0.07-0.42~pc$^{-3}$yr$^{-1}$, the confidence level of this upper limit is not given. Of the charged particles, $\\bar{p}$ are the most effective in setting limits on the PBH density \\cite{MacGibbon_Carr1991,Carr_MacGibbon1998}, however additional uncertainties on the leakage time out of the Galaxy are introduced. The limit on the PBH explosion rate from the $\\bar{p}$ data is (0.011-0.067)~pc$^{-3}$~yr$^{-1}$, where $\\bar{p}$ are produced by jet-fragmentation. Here, the systematic error due to uncertainties in the halo enhancement may be as large as a factor of two \\cite{Carr_MacGibbon1998}. The most recent and stringent upper limit comes from \\cite{Maki1996}, who use a 3D Monte Carlo code to model the $\\bar{p}$ propagation in the local Galactic environment and compare it to the BESS 1995 $\\bar{p}$ data \\cite{Yoshimura1995}. We scale their result to the 99\\% CL as $3.4\\times 10^{-2}$~pc$^{-3}$~yr$^{-1}$, assuming Poisson statistics. Another method to detect the evaporation of PBH relies on the neutrino background. The flux emitted in neutrinos is at least 2 to 3 times larger than the $\\gamma$-ray flux~\\cite{Carr1976}, but so far only the low-energy $\\nu$-background has been used to constrain the PBH mass-distribution~\\cite{Bugaev2002}.  Another idea to detect PBH explosions was put forward by \\cite{Rees1977}, who suggested that a HM burst produces a sphere of simultaneously emitted charged particles, but see the arguments by~\\cite{Heckler1997} and~\\cite{MacGibbon2008}. These generate radio and optical emission by interaction with the ambient magnetic field. Searches for optical pulses have been carried out by \\cite{Porter_Weekes1977_optical,Porter_Weekes1978} resulting in an upper limit of 0.3~pc$^{-3}$~yr$^{-1}$ at the 99\\% CL. A dedicated search for coincident optical flashes between two optical telescopes with duration between 1 and 30~ms by \\cite{Bhat1980_optical} set an upper limit at $0.028$~pc$^{-3}$~yr$^{-1}$.  Searches for radio pulses from PBH evaporation are summarized in \\cite{Halzen1991} ; the best limit being $<7\\times10^{-10}$~pc$^{-3}$~yr$^{-1}$. However, as the physical interaction model of \\cite{Rees1977} is highly speculative, these upper limits should not be regarded as definite. \\begin{figure}[!ht] \\includegraphics[width=0.48 \\textwidth,clip]{ul_intro.eps} \\caption{Limits on the density of PBH explosions. Direct-detection limits are at 99\\% CL, except for \\cite{Cline1997}, and use $\\Omega_m = 0.06$. The limits from particle background measurements have possibly large systematic uncertainties associated with them. } \\label{fig:intro_limits} \\end{figure}  \\subsection{Description of SGARFACE} SGARFACE~\\cite{Sgarface_Lebohec2005} is now operating as part of the Whipple 10~m telescope, an imaging air-Cherenkov telescope (ACT) located in southern Arizona \\cite{Crab_Nebula_Weekes89,Whipple_Kildea07}. The Whipple 10~m telescope is located at 31.6804\\degr\\ latitude, 110.8790\\degr\\ W longitude, and 2312~m a.s.l.  The camera of the Whipple 10~m telescope consists of 379 close-packed PMTs covering a 2.4\\degr\\ field of view. The PMT signals for SGARFACE are split off the 10~m signal cables via passive couplers. To reduce cost and complexity of the SGARFACE system, clusters of seven nearest-neighbor pixels are summed into 55 channels, each viewing $\\sim$0.36\\degr\\ of the sky. This pixelation is sufficient to image the extended Cherenkov images produced by bursts of $\\gamma$ rays. The channels are sampled by flash analog-to-digital (FADC) converters at an interval of 20 ns and with memory depth of 35.04 $\\mu$s (1752 samples). A trigger, described below, causes the FADC buffers to be stored to disk. The FADC input was designed with a time-constant of about 25 ns to preserve the accumulated charge between samples. SGARFACE achieves background rejection of cosmic rays and atmospheric phenomena by the time-resolved imaging of the Cherenkov wavefronts.  The air-Cherenkov technique has a detection area of $\\sim 1\\times 10^5$ m$^2$ for single TeV $\\gamma$-ray primaries~\\cite{Mohanty98}. At sub-GeV energies, the Cherenkov emission from one $\\gamma$ ray is not detectable, but a wavefront of near-simultaneous sub-GeV $\\gamma$-rays produces a detectable signal~\\cite{Sgarface_Krennrich2000}. The time dispersion of Cherenkov photons from an instantaneous burst is 10~ns (35~ns) on-axis (3\\degr\\ off-axis) due to the shower geometry, while a 6.5~ns time-spread is introduced by the geometry of the 10~m telescope. The angular distribution of Cherenkov photons of these low-energy $\\gamma$-ray wavefronts is smooth and broad, as shown in Fig.~\\ref{fig:Cherenkov_distribution} for a burst of 1 GeV $\\gamma$ rays at zenith. The widths of the angular distributions, shown in the lower part of Fig.~\\ref{fig:Cherenkov_distribution}, are a combination of Compton and Coulomb scattering, and the geomagnetic bending of $e^{\\pm}$. Taking into account the number of radiation lengths available for electrons until they reach the threshold for Cherenkov emission in the atmosphere, the total multiple Coulomb scattering angle is roughly proportional to $p^{-1/2}$, where $p$ is the initial electron momentum~\\cite{PDG}. Similarly, the width of the photon angular distribution arising from Compton scattering is proportional to $p^{-1/2}$, where $p$ is the  photon momentum ~\\cite[pg. 354-358]{LandauQED}. As the primary $\\gamma$-ray energy decreases towards the Cherenkov threshold, relatively less Cherenkov radiation is produced from the total available energy; the threshold for $e^{\\pm}$ is 80~MeV at 20~km altitude, the height of shower maximum for 100~MeV $\\gamma$ rays. Therefore, the Cherenkov-light density on the ground depends strongly on the primary energy. For a fixed total energy, Figure~\\ref{fig:sensitivity_energy} shows the Cherenkov photon yield with primary energy relative to 1 GeV $\\gamma$ rays. \\begin{figure}[!ht] \\includegraphics[width=0.4 \\textwidth,clip ]{Cherenkov_distribution.eps} \\caption{Angular distribution of Cherenkov light received at 2300~m altitude from a plane wave front of 1~GeV $\\gamma$ rays at zenith with fluence of 4~$\\gamma$/m$^2$. The x-axis is perpendicular to the Earth's magnetic field. The \\emph{solid} line shows the photon distributions along the x-axis; the \\emph{dashed} line along the y-axis.  } \\label{fig:Cherenkov_distribution} \\end{figure} \\begin{figure}[!ht] \\includegraphics[width=0.45 \\textwidth,clip]{sensitivity_energy.eps} \\caption{Variation of the Cherenkov photon yield divided by the primary energy, $E_0$, versus $E_0$, corresponding to the yield of Cherenkov photons for a burst with fixed total energy. The yield has been normalized relative to $E_0$ = 1~GeV. The yield decreases as the primary $\\gamma$-ray energy approaches the threshold for Cherenkov emission. } \\label{fig:sensitivity_energy} \\end{figure}  The SGARFACE trigger consists of two systems~\\cite{Sgarface_Lebohec2005}. A multi-time scale discriminator (MTSD) and a pattern-sensitive unit. The MTSD integrates the signal and triggers on six time scales: 60~ns, 180~ns, 540~ns, 1620~ns, 4,860~ns, and 14,580~ns. The trigger threshold requirement is tested at three equally spaced time intervals during the trigger window, requiring pulses to be at least as long as the trigger window. The thresholds were set conservatively above the night-sky background and occasionally adjusted to reflect changes in the telescope throughput.  The pattern-sensitive coincidence unit determines if a sufficient number of nearest neighbor channels have triggered the MTSD. The majority of observation time was taken with a 7-fold nearest-neighbor coincidence requirement. Once an event triggers, the 55 FADC stacks are read out after a 16.28~$\\mu$s delay. This data is stored together with information on telescope pointing, PMT voltage settings, and trigger information. ",
        "conclusions": "The direct search for PBH explosions with SGARFACE covered a broad range of $\\gamma$-ray energies and burst durations compatible with a Hagedorn-type black hole evaporation. The experiment is collecting data since March 2003 and makes use of time-resolved images of Cherenkov radiation to distinguish the candidate bursts of $\\gamma$ rays from the much more numerous cosmic-ray initiated showers. In the analysis, the detector was modeled in detail and Monte-Carlo simulations of bursts of $\\gamma$ rays were used to determine the astrophysical volume over which PBH explosions could be seen. Using only a single imaging telescope, this is the first time such a detailed search has been carried out and a high background suppression of 98.4\\% was achieved. The majority of the remaining cosmic-ray background events are indistinguishable from burst events with a single telescope. Figure~\\ref{fig:limits} shows the upper limits that could have been achieved in the absence of the cosmic-ray background, i.e. if a second, identical detector had been placed at a large distance and operated in coincidence. Originally, SGARFACE was designed when the camera on the 10~m camera had a 4.8\\degr\\ field of view. With this larger field of view an almost three-fold increase in PBH detection volume would have been achieved.  The GLAST satellite is sensitive to bursts of $\\gamma$ rays above 0.03~GeV energy. While the detector has a dead-time of 26~$\\mu$s, multiple photons can be detected within the $\\sim$ 1~$\\mu$s coincidence gate width, enabling a direct search for bursts of $\\gamma$ rays. In determining the fluence sensitivity of GLAST to sub-$\\mu s$ burst of $\\gamma$-rays, we require a minimum of 10 photons to be collected. Assuming that no events are detected, we show the 99\\% CL upper limits on the PBH explosion rate for a 2 year exposure in Fig.~\\ref{fig:limits}. Shown in Fig.~\\ref{fig:glast_limits}, is the 99\\% upper limit over the entire energy range of GLAST, where the total burst energy was taken as before as $7.5\\times 10^{45}$ eV. GLAST is most sensitive to burst of $\\gamma$ rays at $\\sim$100 MeV due to the relation between sampled volume and $\\gamma$-ray energy at a fixed total burst energy.  In comparing the upper limits shown in Fig.~\\ref{fig:limits}, it should be kept in mind that different regions of space are sampled by the different detection techniques. For example, GLAST is limited in distance to 14 pc to detect our nominal burst of $\\gamma$ rays, while SGARFACE probes PBH explosions out to 60 pc. Limits from cosmic particle backgrounds measure the PBH distribution on the scale of the Galaxy, while relying on the assumed mass distribution function, galactic halo enhancement, and in some cases the cosmic-ray propagation model, to set limits on the current rate of PBH explosions. These different techniques complement each other in searching for all possible signatures of PBH evaporation.  In this paper we considered bursts of $\\gamma$ rays to be produced during the final, explosive evaporation of PBHs. In itself, a burst detected with SGARFACE would not be sufficient to claim the detection of a PBH explosion, as bursts might be produced by the emission of $\\gamma$ rays contemporaneous with giant radio pulses~\\cite{Lyutikov2008} or perhaps via some other, as yet undiscovered, mechanism. The smoking gun to claim the detection of a Hagedorn-type PBH explosion would include measurements of the neutrino and graviton emission, as well as radio and optical observations~\\cite{Rees1977,Heckler1997}, but see \\cite{MacGibbon2008} for a refutal on such emission. Future searches for bursts should strive towards such a multi-messenger approach. \\begin{figure}[!ht] \\includegraphics[width=0.45 \\textwidth,clip]{glast_limits.eps} \\caption{ Upper limits on the PBH explosion rate that GLAST can set in 2 years of observation versus energy of the emitted $\\gamma$ rays, $E_0$. The total burst energy released in $\\gamma$ rays is $7.5\\times 10^{45}$ eV. } \\label{fig:glast_limits} \\end{figure}"
    },
    "0812/0812.3918_arXiv.txt": {
        "abstract": "We explore the observational characteristics of jet-driven supernovae by simulating bipolar-jet-driven explosions in a red supergiant progenitor.  We present results of four models in which we hold the injected kinetic energy at a constant $10^{51}$ ergs across all jet models but vary the specific characteristics of the jets to explore the influence of the nature of jets on the structure of the supernova ejecta.  We evolve the explosions past shock-breakout and into quasi-homologous expansion of the supernova envelope into a red supergiant wind.  The simulations have sufficient numerical resolution to study the stability of the flow.  Our simulations show the development of fluid instabilities that produce pristine helium clumps in the hydrogen envelope.  The oppositely-directed, nickel-rich jets give a large-scale asymmetry that may account for the non-spherical excitation and substructure of spectral lines such as H$\\alpha$ and He I 10830\\AA.  Jets with a large fraction of kinetic to thermal energy punch through the progenitor envelope and give rise to explosions that would be observed to be asymmetric from the earliest epochs, inconsistent with spectropolarimetric measurements of Type II supernovae.  Jets with higher thermal energy fractions result in explosions that are roughly spherical at large radii but are significantly elongated at smaller radii, deep inside the ejecta, in agreement with the polarimetric observations.  We present shock breakout light curves that indicate that strongly aspherical shock breakouts are incompatible with recent {\\it GALEX} observations of shock breakout from red supergiant stars.  Comparison with observations indicates that jets must deposit their kinetic energy efficiently throughout the ejecta while in the hydrogen envelope.  Thermal energy-dominated jets satisfy this criterion and yield many of the observational characteristics of Type II supernovae. ",
        "introduction": "The observations of core-collapse supernovae are replete with evidence of asphericity.  Hints of aspherical structure have been found in many supernovae (SNe) of various types, including both Type Ib/c  \\citep{maund07c, maund07b, Modjaz:08, WW08} and Type II events \\citep{Fassia:98, Leonard:01, Elmhamdi:03, Chugai:05, Leonard:06}.  Challenges to the classic spherical model of supernovae (SNe) have been building since the observation of Type II SN 1987A, wherein asphericity was inferred from net polarization measurements \\citep{Jeffery:91} and from substructure in spectral lines \\citep[i.e., the `Bochum' event,][]{Hanuschik:88, Lucy:88, Phillips:89}.  Mixing one or two nickel clumps outward in the ejecta has been shown to reproduce the substructure of the H$\\alpha$ line in SN 1987A \\citep{Spyromilio:90, Utrobin:95}.  With the Hubble Space Telescope (HST), we have gained late-time, resolved images of the structure of the nascent SN remnant that dramatically reveal aspherical, elongated ejecta \\citep{wang02}.  Supernova 1987A, while especially observable, was a single event, and with an atypical blue supergiant progenitor; however, typical Type II-P supernovae (SNeII-P), thought to arise from a red supergiant progenitor, show similar signs of asphericity and asymmetry.  \\citet{Fassia:98} present infrared spectra of Type II-P SN 1995V and discuss implications of nonspherical elemental mixing.  They find that in order to fit the He I 10830-\\AA\\ emission, significant amounts of $^{56}$Ni must be ``dredged-up'' to high velocities in order to excite the He I line through radioactive decay.  Fassia et al. also conclude that about 10\\% of the He mass must be contained in pure He clumps that are not microscopically mixed with hydrogen.  When mixed with hydrogen, collisional ionization of hydrogen \\citep[Penning ionization,][]{Bell:70, Chugai:91} is the dominant de-excitation mechanism for He in the hydrogen envelope.  This analysis suggests the need for a mechanism that produces significant outward propagation of nickel and the presence of fluid instabilities that lead to helium clumping.  The detailed shape of spectral lines provides indirect evidence of asymmetry and asphericity of supernova ejecta.  We can gain a direct measure of the shape of the supernova photosphere and line-forming regions from continuum and line polarization \\citep[see][]{WW08}.   While the sample size of Type II-P SNe for which polarization measurements have been made is modest, the data to-date suggest asphericity which depends on the depth of the photosphere.  Leonard et al. (2001) made spectropolarimetric observations of SN 1999em and find a net continuum polarization of about 0.2\\% on day 7 after maximum light, increasing to 0.5\\% at day 165.  Citing the calculations of polarization from electron scattering spheroids by \\citet{hoflich91}, \\citet{Leonard:01} estimate the asphericity of SN 1999em to be $\\sim$7\\% on day 7 and $\\sim$10\\% on day 165.  As they discuss, the late time polarization may be reduced due to the decreased optical depth to electron scattering in the SN atmosphere, implying the intrinsic asphericity of SN 1999em may be greater than their estimates based on the electron scattering calculations of \\citet{hoflich91}.  Further hints at the asphericity of the ejecta in SN 1999em are discussed by \\citet{Elmhamdi:03}.  They report substructure in the H$\\alpha$ and He I lines that is reminiscent of the `Bochum' event in SN 1987A.  The spectral modeling of \\citet{Elmhamdi:03} shows that an asymmetric $^{56}$Ni distribution may account for the structure of these lines.  A striking, and provocative, example of time-dependent SNeII-P polarization is found in SN 2004dj.  In this event, the continuum polarization was practically zero, within the errors, during the plateau phase \\citep{Leonard:06}, but during the transition from photospheric to nebular phases, as the photosphere receded to reveal the core of the SN, the polarization increased dramatically to 0.56\\% on day 91, implying an asphericity of $\\sim$30\\%, and then slowly decreased until no polarization was detected on day 271.  \\citet{Leonard:06} show that the decay of the polarization signal is roughly consistent with the expected decrease in electron scattering opacity of the receding photosphere.  This seems to imply that the core of the explosion was significantly aspherical but was masked by a roughly spherical envelope.  As discussed by \\citet{Leonard:06}, and first argued by \\citet{Wang:01}, when combined with data from core-collapse SNe of other types (i.e., Ib/c), this may indicate that the core-collapse SN mechanism is similar across all spectral types and is intrinsically aspherical.  Similar to SN 1999em, the spectral lines of SN 2004dj imply an asymmetric structure in the ejecta, inferred from the secondary peaks to the red and blue of the central peak.  Spectral modeling including the energy deposition from radioactive decay of $^{56}$Ni show that a {\\it bipolar} nickel distribution in a spherical hydrogen envelope can reproduce the substructure of the H$\\alpha$, [O I], and [Ca II] lines due to asymmetric line excitation from the decay of $^{56}$Ni \\citep{Chugai:05}.    Furthermore, \\citet{Chugai:06} shows that a bipolar nickel distribution in a spherical envelope can also approximately account for the time evolution of the decreasing continuum polarization during the nebular phase.  Based on the preceding discussion, explosion models for SNeII-P must account for ``dredge-up'' of nickel to higher expansion velocities than expected from spherically symmetric explosions, clumps of pristine helium in the hydrogen envelope, an asymmetric and possibly bipolar distribution of nickel, and roughly spherical electron-scattering photospheres during the plateau phase that become remarkably aspherical as the photosphere recedes below the hydrogen envelope.  The jet-driven model of core-collapse SNe \\citep[e.g.,][]{kho99, hof01, maeda03} has features that may produce some of the necessary observed characteristics.  In this model, the explosion is driven by nonrelativistic bipolar jets that send strong bow shocks into the progenitor envelope.  These bows shocks intersect near the equatorial plane of the star and establish an equatorial outflow that complements the polar outflows and completes the explosion.  The material that comprises the jets is elongated along the jet axis and is expected to be nickel-rich owing to the jets' origin in the inner core of the explosion.  Fluid instabilities at the interface between the helium core and hydrogen envelope may produce clumps of helium in the hydrogen envelope, as seen in non-jet-driven supernova calculations  \\citep[e.g.,][]{kifon03, kifon06}.  The character of such an unstable flow, however, has not been investigated for the case of jet-driven SNe.  It is also unclear that a jet-driven explosion in a red supergiant progenitor can produce a photosphere that is nearly spherical early-on and becomes aspherical later \\citep[see][]{hof01}.  In this work, we present a numerical study of the structure of jet-driven SNe with adequate resolution to explore the nature of any unstable flows and time evolution sufficient to make a direct comparison with observations.   We assume that bipolar jets are formed as a result of the core-collapse process and do not self-consistantly simulate the jet production process.   Collapsing stellar cores could produce bipolar jets  in the presence of rotation and magnetic field amplification \\citep{Wheeler:00, Wheeler:02, Akiyama:03}.  The process of magnetic jet formation in collapsing stars has been studied in previous simulations \\citep{LeBlanc:70, Shibata:06, Obergaulinger:06b, Obergaulinger:06a,  Burrows:07, kom07}, however, the rotation rate and initial magnetic field structure of the progenitor stars, two very critical properties, are not well understood from first principles.  Therefore, we vary the parameters of the jets in four, physically-plausible jet models in order to explore the importance of the nature of the jets on the final structure of the SN.  The progenitor model in all cases is a red supergiant star of 15 $M_\\odot$, roughly analogous to Betelgeuse, and the models are evolved to nearly six days after the initiation of the jets, several days after shock breakout.  Many previous studies have focused on highly energetic jets in connection with Gamma-Ray Bursts (GRBs) \\citep[e.g.,][]{MacFadyen:01, Zhang:03} and explosions of Pop III stars \\citep[e.g.,][]{maeda03, Tominaga:07}.  We focus explicitly on explosions of typical supernova energies ($10^{51}$ erg) driven by {\\it nonrelativistic} jets.  We use a progenitor star of solar metallicity that is not extremely massive.  The goal of this work is to explore the observational features of jet-driven supernovae in common Type II events.  This work is organized as follows.  In \\S\\ref{sec:sims} we describe our numerical approach and discuss the details of our jet models.   In \\S\\ref{sec:dynamics} we provide a narrative of the general evolution of the explosions, from shock propagation in the core to eruption from the progenitor's surface.  The transport of nickel to radii and velocities greater than those of lighter elements is a feature of a jet-driven explosion model in which the jets contain a significant amount of nickel.  In \\S\\ref{sec:clumps} we discuss the significance of this source of overturn, as well as the development of fluid instabilities that lead to the mixing of elements, such as nickel and helium.  In \\S\\ref{sec:geometry} we discuss the resultant shapes of the explosions and discuss the evolution of the surfaces of constant electron scattering optical depth.  In \\S\\ref{sec:breakout} we calculate the bolometric shock breakout light curves of our models and compare them to {\\it GALEX} observations of supernova shock breakout from red supergiant stars.  We discuss the implications of our study and list our conclusions in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have performed 2D axisymmetric, high-resolution simulations of stellar explosions driven by bipolar jets.  We have followed the evolution of four different jet models to near free expansion.  The energy budgets of  three of our jet models (v3m12, v5m06, and v6m04) are dominated by kinetic energy whereas that of one jet model (v1m12) contains a large fraction of thermal energy, and all jet models inject a total of $10^{51}$ erg.  This distinction in the four models drastically affects the resulting structures of the models.  The kinetic energy models remain highly elongated as they expand into the wind surrounding our progenitor model, whereas model v1m12 produces an effectively spherical outer envelope with large departures from symmetry in the core.   These shapes correspond to how efficiently the jet kinetic energies are isotropized.  Only model v1m12 results in an effectively spherically-symmetric kinetic energy distribution (see Fig. \\ref{fig:ekv1}).  An anisotropic distribution of kinetic energy may have implications for the measurement of supernova energies from observations.   While radiative transport calculations of our results would be required to reproduce typical spectral lines used in estimating supernova energies, it is likely that the kinetic energy models, with their asymmetric energy distributions, would yield viewing angle-dependent kinetic energy estimates.  If a spherical explosion model were assumed when analyzing the observations of an aspherical supernova, errors in the estimate of the kinetic energy could arise.  We resolve the growth of fluid instabilities in the explosions.  We show that the jet fluid is unstable and develops several RT fingers that, when reverse-shocked in the H envelope, serve as initial perturbations for rapid RM growth.  The RM fingers travel at nearly the free-fall velocity and are able to overtake the shock in the kinetic energy models.  The edges of the jets are also KH unstable, but the growth rate depends sensitively on the speed and density of the jet fluid.  Model v3m12 does not show significant KH growth, but all three other models do.  The equatorial outflow in all four models is also subject to KH growth along the top and bottom surfaces, as predicted by \\citet{Wheeler:08}.  In comparison with \\citet{kho99}, model v3m12 shows a similar large-scale structure up to the time the shocks depart the helium core (see their Figure 2).  The extension in the jet direction is similar and the development of the equatorial outflow is comparable.  In contrast, the higher resolution employed in our model ($\\Delta r_{\\rm min} \\sim 3\\times10^6$ cm versus $\\Delta r_{\\rm min} \\sim 4\\times10^7$ cm in Khokhlov et al.) allows us to see the development and growth of fluid instabilities, particularly in the jet fluid (see the right panel of Figure \\ref{fig:v3m12core}).  We have compared the late-time geometries of our models with observations by calculating the electron scattering optical depths under the assumption of adiabatic expansion (see Figs. \\ref{fig:tau_v1} and \\ref{fig:tau_v3}).  We show that the kinetic energy-dominated models (v3m12, v5m06 and v6m04) are too aspherical at early times to match polarimetry of SNeII-P.  Model v1m12 produces a photosphere that is nearly spherical for the first 50 days of expansion and increases in prolate elongation to axis ratios that are in close agreement with those derived for SN 2004dj \\citep{Leonard:06}.  Our simplified calculation of optical depths neglect several important radiative effects, however a full radiative transfer calculation is beyond the scope of this work and is unlikely to qualitatively change our results.  The development of instabilities may be important to the overall understanding of core-collapse supernovae in a number of ways.  Substantial observational evidence shows that the transport of heavier material to radii ahead of lighter material is a ubiquitous phenomenon in core-collapse supernovae, especially SNeII-P.  This ``overturn'' of the original ``onion-skin'' configuration of the progenitor has been explained in neutrino-driven explosions by the growth of fluid instabilities \\citep[see][]{kifon03,kifon06}.  The substructure of certain strong lines (such as H$\\alpha$) indicates large, fast-moving clumps of nickel in the ejecta.  Similar to the calculations of \\citet{kifon03}, many of the instabilities that develop in our models are of very small scale and it is unlikely that such small instabilities can explain the substructure of spectral lines, however, the jets are effectively a large, single-mode asymmetry capable of producing high-velocity bipolar clumps.  The jet fluid, which we assume to be nickel, is transported ahead of lighter materials, providing significant ``overturn.''  In our models, a large amount of nickel mass is transported to radii where it can excite the H$\\alpha$ and He I lines, as is required by spectral modeling of SNeII-P \\citep{Fassia:98, Elmhamdi:03, Chugai:05, Chugai:06}.   \\citet{kifon06} show that the mode of the instabilities that grow in neutrino-driven explosions is sensitive to the nature of the neutrino emission.  In jet-driven models, the jets create a robust, large-scale asymmetry that is not very sensitive to the detailed nature of the jets.  We show that He clumps form due to RT instabilities in the transition layer between the H and He in the explosion ejecta in all four jet models.  These clumps may contribute the pristine helium required by \\citet{Fassia:98} to fit the strength of the He I 10830 \\AA\\ line, though the amount of clumping we find is below that required by Fassia et al.  As we have noted, the extra nickel at high velocities in the jet models (about a factor of two greater than that assumed by Fassia et al.) may compensate for the low helium clump fraction by providing excess gamma-ray excitation.  Evolving our calculations beyond shock breakout allows us to study the breakout characteristics of the explosion models.  In Figure \\ref{fig:breakout} we present the bolometric shock breakout light curves for models v1m12 and v3m12.  The nearly-spherical breakout of model v1m12 produces a light curve that is very similar to spherically-symmetric model light curves.  The dramatically aspherical breakouts of the kinetic models show a first peak in the light curve associated with the breakout of the jets followed by a roughly constant luminosity period created by the expanding and cooling jet ``cocoons'' and then a large peak as the equatorial shocks break out.  The early peak and plateau in the light curves of the kinetic models would have been evident in {\\it GALEX} observations of SNLS 04D2dc, if they were present.  Their absence argues against the likelihood of a significantly aspherical break out for this SNeII-P, as does the comparison to polarimetry.  The light curve of model v1m12 agrees with the observations for shock breakout for the two SNeII-P caught by {\\it GALEX} \\citep{Schawinski:08, Gezari:08}.   In Table \\ref{tab:obs}, we summarize the observational characteristics of SNeII-P and the four jet models.  We also include the characteristics a neutrino-driven SN model based loosely on that of \\citet{kifon03, kifon06} for comparison.  A multi-dimensional neutrino-driven model that includes fluid instabilities can be expected to account for transport of heavy elements to high ejecta velocities as well as the formation of pristine helium clumps in the hydrogen envelope; however, it is not clear that neutrino-driven models can consistently produce fast-moving, bipolar nickel clumps in the ejecta.  Likewise, it is unlikely that these models can account for the time dependence of SNeII-P polarization.  Significantly aspherical neutrino-driven explosions may result from the standing accretion shock instability \\citep[SASI;][]{Blondin:03, Marek:07} or the acoustic mechanism \\citep{Burrows:06, Burrows:07a}, however the large-scale structure of supernovae resulting from these mechanisms has not been explored.  As we have discussed, the kinetic energy-dominated jet models produce explosions that are too aspherical at early times and are, thus, incompatible with polarization and shock breakout observations of SNeII-P.   In contrast, our thermal energy-dominated model (v1m12) shows good agreement with all of the observations we have discussed here.  The comparison of observational evidence with our simulations consistently indicates that jets capable of retaining a large fraction of their initial kinetic energy while crossing the hydrogen envelope of a red supergiant star are unlikely to be present in SNeII-P explosions.  This is evidence that if jets are present, or are the driving mechanism in core-collapse supernovae, they must be jets that are capable of transferring their kinetic energy to the hydrogen envelope of a RSG.  This may not require thermal energy-dominated jets, as we have presented here, but it does seem to rule out fast, dense jets that remain ballistically detached from the surrounding ejecta and punch out of the progenitor with a large fraction of their initial kinetic energy intact.  \\citet{hof01} and \\citet{kho01} have noted that if jets efficiently penetrate the core of a stripped envelope supernova, but are mostly trapped in the envelope of extended hydrogen-rich supernovae, then average stripped-envelope supernovae should have lower kinetic energies than average SNeII-P. Since this is not the case, H\\\"{o}flich deduced that any jet energy must already be fairly efficiently deposited in the cores of supernova progenitors by dint of any jets being rather wide and slow. The conclusions of our current study, that thermal jets are favored over kinetic energy jets, is roughly in accord with Hoeflich's supposition. We note, however, that even our thermal jets that ultimately deposit their energy rather uniformly in the hydrogen envelope are less efficient in depositing energy in the helium core. The energy deposition cannot be completely isotropized in the helium core without violating the observations from spectropolarimetry that the inner core is substantially aspherical.  How these two constraints of energetics versus asymmetry play off against one another needs to be studied in more detail."
    },
    "0812/0812.1589_arXiv.txt": {
        "abstract": "The bright Type II-plateau supernova (SN) 2004dj occurred within the young, massive stellar cluster Sandage-96 in a spiral arm of NGC~2403. New multi-wavelength observations obtained with several ground-based and space-based telescopes are combined to study the radiation from Sandage-96 after SN~2004dj faded away.  Sandage-96 started to dominate the flux in the optical bands starting September 2006 ($\\sim$800~d after explosion). The optical fluxes are equal to the pre-explosion ones within the observational uncertainties. An optical Keck spectrum obtained $\\sim$900~d after explosion shows the dominant blue continuum from the cluster stars shortward of 6000~\\AA\\ as well as strong SN nebular emission lines redward.  The integrated spectral energy distribution (SED) of the cluster has been extended into the ultraviolet region by archival {\\it XMM-Newton} and new {\\it Swift} observations, and compared with theoretical models. The outer parts of the cluster have been resolved by the {\\it Hubble Space Telescope}, allowing the construction of a color-magnitude diagram. The fitting of the cluster SED with theoretical isochrones results in cluster ages distributed between 10 and 40 Myr, depending on the assumed metallicity and the theoretical model family.  The isochrone fitting of the color-magnitude diagrams indicates that the resolved part of the cluster consists of stars having a bimodal age distribution: a younger population at $\\sim$10--16 Myr, and an older one at $\\sim$32--100 Myr.  The older population has an age distribution similar to that of the other nearby field stars.  This may be explained with the hypothesis that the outskirts of Sandage-96 are contaminated by stars captured from the field during cluster formation.  The young age of Sandage-96 and the comparison of its pre- and post-explosion SEDs suggest $12 \\lesssim M_{\\rm prog} \\lesssim 20$~M$_\\odot$ as the most probable mass range for the progenitor of SN~2004dj.  This is consistent with, but perhaps slightly higher than, most of the other Type II-plateau SN progenitor masses determined so far. ",
        "introduction": "The theory of stellar evolution predicts that massive stars ($M \\gtrsim 8~M_\\odot$) end their lives as core-collapse supernovae (CC~SNe, e.g., \\citealt{woo02}). In particular, after the main-sequence phase the most massive stars undergo heavy mass loss, become stripped stellar cores and explode as Type Ib/c supernovae (SNe~Ib/c; see \\citealt{fil97} for a discussion of SN spectral classification).  Stars close to the lower mass limit of CC are thought to produce Type II-plateau SNe (SNe~II-P). Recent observations to detect the progenitors of CC~SNe support this scenario.  Currently there are 10 SNe II (1987A, 1993J, 1999ev, 2003gd, 2004A, 2004et, 2005cs, 2006ov, 2008bk; the last seven are Type II-P) whose progenitors have been directly identified in pre-explosion images (see \\citealt{hen06, li07, mat08, leonard08, smartt08}, and references therein), and the mass estimates are $M \\lesssim 15$--20 ~M$_\\odot$ for all of them. Moreover, upper mass limits were derived for a number of other SNe~II from nondetections of their progenitors \\citep{vand03, ms05,leonard08}, and the highest upper limit was found to be $M \\approx 20$~M$_\\odot$. These observations have led to the conclusion that SNe~II-P likely originate from ``low-mass'' massive progenitors with $M \\lesssim 20$~M$_\\odot$ \\citep{li06, li07, smartt08}, and the fate of stars with $M \\gtrsim 20$~M$_\\odot$ may be a SN~Ib/c explosion.  On the other hand, the progenitors of SNe~Ib/c (even the brightest and closest ones) have escaped direct detection so far \\citep{cr07a}.  The most promising candidate is SN~2007gr, which occurred in a compact, massive stellar cluster in NGC~1058 that has been detected with the {\\it Hubble Space Telescope (HST)} prior to explosion \\citep{cr07b}.  SN~2004dj, the closest ($\\sim 3.5$ Mpc, \\citealt{v06}) and one of the brightest SNe since SN~1987A, was a Type~II-P event. It occurred within a young, massive cluster, Sandage-96 (S96) in NGC~2403. SN~2004dj has been extensively studied through multi-wavelength observations; see \\citet{v06} (hereafter Paper~I) for references. In particular, several attempts were made to infer the mass of the progenitor by comparing the pre-explosion magnitudes and colors of S96 with theoretical spectral energy distributions (SEDs) to determine the age, hence the turnoff mass, of the cluster. These resulted in a range of possible progenitor masses from $M \\approx 12$~M$_\\odot$ to $M \\gtrsim 20$~M$_\\odot$ depending on the assumed metallicity and/or reddening \\citep{maiz04, wang05, v06}. However, all of these studies suffered from the age-reddening and age-metallicity degeneracy \\citep{rebuz86}, because the available pre-explosion observations covered only the optical and near-infrared (NIR) bands.  The main aim of this paper is to derive further constraints on the progenitor mass of SN~2004dj from post-explosion observations of S96, made after its reappearance from the dimming light of the SN. The cluster has been successfully redetected both by our ground-based and space-based observations, showing no significant change in its optical light level with respect to the pre-explosion level.  We have extended the wavelength coverage of the observed SED into the ultraviolet (UV) with new {\\it Swift}/UVOT observations. Moreover, the cluster has been partially resolved by our new $HST$/ACS observations, which provides an additional opportunity to infer age constraints on its stellar population.  The paper is organized as follows. In \\S 2 we present our new multi-wavelength observations made from the ground and space.  The SED of the cluster and its color-magnitude diagram are constructed and fitted with theoretical model predictions in \\S 3. We discuss the results in \\S 4 and present our conclusions in \\S 5.  \\begin{figure*} \\begin{center} \\epsscale{1.0} \\plottwo{f1a.eps}{f1b.eps} \\caption{Left panel: color-combined $B$, $V$, and H$\\alpha$ frames of NGC~2403 obtained with the 2.3~m Bok telescope at Steward Observatory on 2007 Jan. 28 (UT dates are used throughout this paper).  Right panel: {\\it Swift}/UVOT image of NGC 2403 obtained on 2007 Dec. 2 ($u$, uvw1, and uvw2 filters were selected as red, green, and blue channels).  The field of view of both frames is $10.0' \\times 7.2'$; north is up and east is to the left.  The position of S96/SN~2004dj is marked. } \\label{fig-bok} \\end{center} \\end{figure*} ",
        "conclusions": "We have presented late-time photometry of SN~2004dj and the surrounding cluster S96, extending the time coverage of the observational sample up to $\\sim 1000$ d after explosion. In the optical, the continuum flux from SN~2004dj faded below the level of the integrated flux of S96 in 2006 Sep., $\\sim 800$ d after explosion. The pre- and post-explosion SEDs of S96 show no significant differences in the range 2000--9000~\\AA. The nebular spectrum of SN~2004dj at $\\sim900$ d after explosion was dominated by the blue continuum from S96 shortward of 6000~\\AA, and by strong H$\\alpha$, [\\ion{O}{1}] $\\lambda\\lambda$6300, 6363, and [\\ion{Fe}{2}] $\\lambda7155$ emission line, characteristic of a typical nebular spectrum of a SN~II-P.  We have examined the multi-wavelength observations of S96 by different methods, in order to derive constraints on the cluster age and evolutionary status.  The fitting of the cluster SED (using the average of pre- and post-explosion fluxes) results in cluster ages distributed between $\\sim 8$ and $\\sim 40$ Myr, with the best-fitting solutions being within 10--20 Myr.  The observed reddening is $E(B-V) \\approx 0.10 \\pm 0.05$ mag; its uncertainty is greatly reduced compared with previous studies, due to the inclusion of the UV fluxes from {\\it Swift} and {\\it XMM-Newton}.  S96 appears to be partly resolved in images obtained with $HST$/ACS on 2005 August 28 ($\\sim 425$~d after explosion), although the light from SN~2004dj was still very strong at that time.  We have computed photometry of the ACS images obtained through the $F435W$, $F606W$, and $F814W$ filters, and combined the magnitudes of the detected stellar sources in color-magnitude diagrams. Theoretical isochrones fitted to the observed CMDs reveal that the resolved stars in the outskirts of the cluster have a bimodal age distribution. The younger population consists of stars with ages of $10 < t < 16$ Myr, while the members of the older one have $30 < t < 100$ Myr. The ages of the older population has a distribution that is similar to that of the field stars, not associated with S96. This similarity may suggest that about half of the cluster stars resolved by ACS were captured from the field population during the formation of S96.  The absence of a visible H$\\alpha$-emitting cloud around S96 implies a lower limit for the cluster age of $\\sim$8--10 Myr, in agreement with the other age estimates.  The 10 Myr age of S96 would imply a SN 2004dj progenitor mass of $M_{\\rm prog} \\approx 20$~M$_\\odot$, while the mass limit for core collapse (7--8 M$_\\odot$) would mean $t \\approx 60$ Myr for the age of the progenitor. This latter limit is consistent with the age of the older population within S96, leaving the possibility of a low-mass progenitor open. The age of the younger population (10--16 Myr) corresponds to $M_{\\rm prog} \\approx$ 12--15 M$_\\odot$, which seems to be the most probable mass estimate at present. We verified that even a 20~M$_\\odot$ progenitor would be consistent with the unobservable flux difference between the pre- and post-explosion SEDs. However, more observations, especially in the $JHK$ bands, would be essential to narrow the mass range of the progenitor."
    },
    "0812/0812.3290_arXiv.txt": {
        "abstract": "{ We explore the relationship between the hard X-ray photon index $\\Gamma$ and the Eddington ratio ($\\xi=L_{X}(0.5-25\\ \\rm keV)/L_{Edd}$) in six X-ray binaries (XRBs) with well-constrained black hole masses and distances. We find that different XRBs follow different anti-correlations between $\\Gamma$ and $\\xi$ when $\\xi$ is less than a critical value, while they follow the same positive correlation when $\\xi$ is larger than the critical value. This anti-correlation and positive correlation are also found in low luminosity active galactic nuclei (LLAGNs) and luminous QSOs respectively. The anti-correlation and positive correlation of different XRBs roughly converge to the same point ($\\log \\xi=-2.1\\pm0.2, \\Gamma=1.5\\pm 0.1$), which may correspond to the accretion mode transition, since that the anti-correlation and positive correlation between $\\Gamma$ and $\\xi$ are consistent with the prediction of advection dominated accretion flows (ADAFs) and standard disk/corona system respectively. We note that traditional low-hard state are divided into two parts by the cross point $\\log \\xi\\sim-2.1$, i.e., \\emph{faint-hard state} in the anti-correlation part and \\emph{bright-hard state} in the positive correlation part. The accretion process in the \\emph{bright-hard state} may be still the \\emph{standard accretion disk} as that in the high/soft state, which is consistent with that both the cold disk component and broad Fe K emission line are observed in some bright-hard state of XRBs (e.g., GX 339-4, Miller et al. 2006). The \\emph{ADAF} is only important in the \\emph{faint-hard state} XRBs.  Motivated by the similarities of the state transition and timing properties of the ultraluminous X-ray sources (ULXs) to that of XRBs, we then constrain the black hole masses for seven luminous ULXs assuming that their X-ray spectral evolution is similar to that of XRBs (i.e., photon index is only depend on the Eddington ratio). We find that the BH masses of these seven ULXs are around $10^{4}M_{\\odot}$, which are typical intermediate-mass BHs (IMBHs). Our results are roughly consistent with the BH masses constrained from the model fitting with a multi-color disk and/or the timing properties (e.g., QPO and break frequency).  }  \\FullConference{VII Microquasar Workshop: Microquasars and Beyond\\\\ September 1-5 2008\\\\ Foca, Izmir, Turkey}  \\begin{document} ",
        "introduction": "It is generally assumed that black hole (BH) X-ray binaries (XRBs), active galactic nuclei (AGNs) and other BH systems have a similar central engine, namely the central BH, the accretion flow and a relativistic jet (for reviews see Fender et al. 2007, and $\\rm K\\ddot{o}rding$ et al. in this proceeding). The similarities are supported by recent two ``fundamental planes\" of the BH activity on all mass scales, which include $L_{\\rm radio}-L_{\\rm Xray}-M_{\\rm BH}$ (e.g., Merloni et al. 2003; Falcke et al. 2004) and $L_{\\rm bol}-M_{\\rm BH}-T_{\\rm break}$ (McHardy et al. 2006). The hardness-intensity diagram and pattern of radio loudness in XRBs and AGNs also support this unification scenario (e.g., $\\rm K\\ddot{o}rding$ et al. 2006).  The low/hard state and high/soft state are two main states in XRBs. Typically, the X-ray spectrum in hard state can be described as a power law with a photon index $\\Gamma=1.5-2.1$, while the energy spectrum can be described with a thermal disk component and a power law with a photon index $\\Gamma=2.1-4.8$ in soft state (e.g., Remillard \\& McClintock 2006). The hard X-ray photon index is anti-correlated to the Eddington ratio in some low/hard state XRBs (e.g., Yamaoka et al. 2005) and a sample of low luminosity AGNs (LLAGNs, Gu \\& Cao 2008) when the Eddington ratio is less than a critical value. However, a positive correlation is found in XRBs with higher Eddington ratio (e.g., Kubota \\& Makishima 2004) and luminous QSOs (e.g., Shemmer et al. 2006).  Ultraluminous X-ray sources (ULXs) are pointlike, nonnuclear X-ray sources with X-ray luminosities between $10^{39}$ and $10^{41}\\rm ergs/s$, well in excess of the Eddington limit for a stellar mass black hole. The true nature of ULXs is still unclear, especially their BH mass, and current models include two main alternatives: (1) ``intermediate-mass BHs\" (IMBHs) with mass $M_{\\rm BH}\\simeq 10^{2}-10^{5} M_{\\odot}$ (e.g., Colbert et al. 1999); (2) stellar mass BHs (e.g., Mirabel \\& Rodriguez 1999). ",
        "conclusions": "\\subsection{The hard X-ray spectral evolution in XRBs and AGNs} We present the relation between the photon index, $\\Gamma$, and the Eddington ratio, $\\xi$, of the XRB sample in Figure 1a, where the photon index $\\Gamma$ is fitted in the 3-25 keV band. We define the Eddington ratio $\\xi=L_{X}(0.5-25\\ \\rm keV)/\\it L_{\\rm Edd}$, in which $L_{X}(0.5-25\\ \\rm keV)$ is extrapolated from the unabsorbed 3-25 keV band luminosity, since which was regarded as more represent the bolometric luminosity. It can be clearly seen that the X-ray photon index $\\Gamma$ is strongly correlated with the Eddington ratio $\\xi$ when $-2\\lesssim\\log \\xi\\lesssim0$. In the low Eddington ratio (e.g., $\\log \\xi\\lesssim -2$) part, the anti-correlation is generally present, which actually has already been found in several occasions (e.g., Yamaoka et al. 2005; Yuan et al. 2007). However, the most remarkable result in our work is that different sources follow the different anti-correlations (Fig. 1a). We fit the anti-correlation points with linear least-square method for each source. We also fit all data points at the positive correlation ($-2\\lesssim \\log \\xi \\lesssim -1$) region as a whole with the same method (short dashed line in Fig. 1a). Therefore, for each source, we obtained the cross point between the positive and anti-correlation part. Although the anti-correlation varies from source to source, we find the cross points of all XRBs roughly converge to the same point with small scatter ($\\log \\xi=-2.1\\pm0.2, \\Gamma=1.5\\pm 0.1$).  \\begin{figure} \\includegraphics[width=.99\\textwidth]{fig.eps} \\caption{(a) The relation between the X-ray photon index and the Eddington ratio in XRBs and AGNs. The short dashed lines are the linear least-square fits of XRBs,  and long-dashed lines are for LLAGNs (left) and QSOs (right) respectively. (b) The relation between the X-ray photon index and the Eddington ratio in ULXs. For comparison, two XRBs (XTE J1748-288 and 4U 1543-47) are also plotted.} \\label{fig1} \\end{figure}  For comparison, we also plot the best fits to the relation between the photon index and the Eddington ratio in LLAGNs (long dashed-line in Fig. 1a, Gu \\& Cao 2008) and luminous QSOs (long dashed-line in Fig. 1a, Shemmer et al. 2006), We also include the Seyfert galaxy NGC 4051, which show strong X-ray spectrum and flux variation (green solid circles in Fig. 1a). We find that the photon index is anti-correlated to the Eddington ratio in LLAGNs, while a positive correlation is present in QSOs and NGC 4051. The anti-correlation in LLAGNs and the positive correlation in luminous AGNs cross at point ($\\log \\xi=-2.78, \\Gamma=1.61$), where the Eddington ratio $\\xi$ is several times less than that of XRBs, which is mainly caused by the different bolometric correction factor (see Wu \\& Gu 2008 for details). Therefore, it seems that the hard X-ray spectral evolution are also similar in XRBs and AGNs, which can be formulized as: \\begin{equation} \\Gamma = \\kappa (\\xi+2.1^{+0.7}_{-0.2}) +1.5^{+0.1}_{-0.1} \\end{equation} where $\\kappa$ is the slope and $\\xi$ is the Eddington ratio.   \\subsection{The possible accretion process and the so-called ``hard state problem\" } Both the UV/optical bumps observed in QSOs and the soft X-rays observed in high state XRBs can be naturally interpreted as blackbody emission from the standard accretion disk (Shakura \\& Sunyaev 1976). ADAF is a hot, optically thin, geometrically thick accretion flow, which can successfully explain most features of the LLAGNs and low/hard-state XRBs (see Ho 2008; Narayan \\& McClintock 2008 for recent reviews). The qualitative behavior of the \\emph{anti-correlation} between the photon index and Eddington ratio is consistent with the predictions of the \\emph{ADAF} spectrum (e.g., Esin et al. 1997). The Comptonization of thermal synchrotron photons in the ADAF is the dominated cooling mechanism at low Eddington ratios. As the accretion rate increases, the optical depth of the ADAF also increases, increasing the Compton $y$-parameter, thereby leading to a harder X-ray spectrum. However, the \\emph{positive correlation} is consistent with the \\emph{standard disk-corona} model (e.g., Janiuk \\& Czerny 2000). As the accretion rate increases, both the fraction of accreting energy released to the corona and electron temperature decreases, the corona becomes weak and the \\emph{optical depth} decreases, reducing the y-parameter, and leading to a softer X-ray spectrum.  Recently, the observations of the cool disk component and the relativistic broad Fe K emission line in hard state of several XRBs (e.g., GX 339-4, Miller et al. 2006; Tomisic et al. 2008) have suggested that the standard disk may extend to the innermost stable circular orbit, which challenged the popular ADAF model in the hard state (so-called ``\\emph{hard state problems}\", see a review of Tomsick et al. in this proceeding). It is interesting to note that the traditional hard state are divided into two parts, i.e., \\emph{bright hard state} and \\emph{faint hard state}, by the cross point($\\log \\xi \\sim -2.1\\pm0.2$) of the anti-correlation and the positive correlation. The X-ray spectral evolution in the bright hard state is the same as that of high/soft state of XRBs and luminous QSOs. It implies that accretion process in this \\emph{bright-hard state} may be the \\emph{standard accretion disk}. It is consistent with the cool disk component and the relativistic broad Fe K emission line in the several bright hard state of XRBs (e.g., Miller et al. 2006). In contrast, the \\emph{ADAF model} becomes important in the \\emph{faint-hard state} of the XRBs with $\\log \\xi \\lesssim -2.1$. If this is the case, the cross point may correspond to the accretion mode transition, which, however, is different from the popular idea that the disk transition occurs at $\\Gamma=2.1$.   \\subsection{The X-ray spectral evolution in ULXs and constraints on their BH masses} It is found that the X-ray photon index of NGC 1313 X-1 is positively correlated to the luminosity, while these two parameters are anti-correlated in NGC 1313 X-2  (Feng \\& Kaaret 2006), which are similar to that of XRBs/AGNs. This phenomenon has also been found in many other ULXs (e.g., Roberts et al. 2004). The similarity between XRBs, ULXs and AGNs implies that there may be similar physics behind the phenomena, which motivates us to explore the properties of ULXs utilizing the unification of ULXs and XRBs (and AGNs). Specifically, the BH masses of ULXs can be constrained if we assume that the BH accretion is scale-free and that their X-ray spectral evolution is only determined by the Eddington ratio. Assuming the spectral evolution can be described by Eq. 3.1, we thus calculate the BH masses of ULXs through the least-squares linear fit on the available data for each ULX. Our results show that all the BHs in these luminous ULXs are IMBHs of around $(4-30)\\times 10^{3}M_{\\odot}$ (Table 1). We find that our estimated BH masses are consistent within a factor of two with that derived either using $\\nu_{\\rm break}-M_{\\rm BH}-L_{ \\rm Bol}$ relation for XRBs and AGNs (McHardy et al. 2006) or from the multi-color disk fittings (Miller et al. 2004) (see Table 1).  These consistency may support the validity of our method, which further implies that the assumption of similarity between ULXs and XRBs is likely reasonable.   \\begin{table} \\caption{BH mass of ULXs.} \\label{tab1} \\begin{tabular}{lcc|lcc} \\hline \\hline Source Name & Our BH mass & other BH mass$^1$ & Source Name & Our BH mass & other BH mass$^2$\\\\ \\hline M81 X-9      &  $6.6\\times10^{3}$   &  $5^{+7}_{-2}\\times10^{3}$   &  NGC 5408 X-1  &  $<4.0\\times10^{3}$  &  $2.5\\pm1\\times10^{3}$   \\\\ M82 X-1      &  $6.6\\times10^{3}$   &  $5^{+7}_{-2}\\times10^{3}$   &  NGC 4559 X-7  &  $1.4\\times10^{4}$   &  $7.6\\times10^{3}$     \\\\ NGC 1313 X-1 &  $7.9\\times10^{3}$   &  $4^{+2}_{-1}\\times10^{3}$   &  NGC 4559 X-10 &  $2.7\\times10^{4}$   & ...    \\\\ NGC 1313 X-2 &  $2.4\\times10^{4}$   &  ...                  \\\\ \\hline \\end{tabular} \\footnotesize (1)BH mass derived from the multicolor disk fittings; (2) BH mass derived from the timing properties in this work. \\end{table}"
    },
    "0812/0812.0806_arXiv.txt": {
        "abstract": "This paper summarizes Kormendy et al.~(2009, ApJS, in press, arXiv:0810.1681). We confirm that spheroidal galaxies have fundamental plane correlations that are almost perpendicular to those for bulges and ellipticals.  Spheroidals are not dwarf ellipticals.  They are structurally similar to late-type galaxies.  We suggest that they are defunct (``red~and~dead'') late-type galaxies transformed by a variety of gas removal processes. Minus spheroidals, ellipticals come in two varieties:~giant, non-rotating, boxy galaxies with cuspy cores and smaller, rotating, disky galaxies that lack cores.  We find a new feature of this ``E--E dichotomy'':~Coreless ellipticals have extra light at the center with respect to an inward extrapolation of the outer S\\'ersic profile.  We suggest that extra light is made in starbursts that swamp core scouring in wet mergers.  In general, only giant, core ellipticals contain X-ray gas halos.  We suggest that they formed in mergers that were kept dry by X-ray gas heated by active galactic nuclei. \\pretolerance=15000  \\tolerance=15000 ",
        "introduction": "Figure 1 shows examples of photometry of all known ellipticals in the Virgo cluster from Kormendy et al.~(2009; hereafter KFCB). The key to this paper is the accuracy and large dynamic range attained when data from many telescopes with different resolutions and field sizes are combined into composite profiles.  \\begin{figure}[h!] \\pretolerance=15000  \\tolerance=15000  \\vskip 2.2truein  \\special{psfile=./kormendy-figure1a.ps hoffset=-84 voffset=-75 hscale=49.5 vscale=49.5}  \\special{psfile=./kormendy-figure1b.ps hoffset=107 voffset=-75 hscale=49.5 vscale=49.5}  \\caption{\\lineskip=-6pt \\lineskiplimit=-6pt Major-axis brightness profiles of Virgo ellipticals from KFCB.  Surface brightness $\\mu$ is in $V$ mag arcsec$^{-2}$.  The curve is a S\\'ersic (1968) fit between the vertical dashes; the index $n$ and total galaxy absolute magnitude $M_{VT}$ are in the key. This figure illustrates the E--E dichotomy: NGC 4486 is a core galaxy with $n > 4$; NGC 4458 is an extra light galaxy with $n \\la 4$ (\\S\\thinspace3). } \\end{figure} \\eject ",
        "conclusions": ""
    },
    "0812/0812.3202_arXiv.txt": {
        "abstract": "In this letter, we suggest that a nearby clump of 600-1000 GeV neutralinos may be responsible for the excesses recently observed in the cosmic ray positron and electron spectra by the PAMELA and ATIC experiments. Although neutralino dark matter annihilating throughout the halo of the Milky Way is predicted to produce a softer spectrum than is observed, and violate constraints from cosmic ray antiproton measurements, a large nearby (within 1-2 kiloparsecs of the Solar System) clump of annihilating neutralinos can lead to a spectrum which is consistent with PAMELA and ATIC, while also producing an acceptable antiproton flux. Furthermore, the presence of a large dark matter clump can potentially accommodate the very large annihilation rate required to produce the PAMELA and ATIC signals. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.1800_arXiv.txt": {
        "abstract": "It is now well accepted that the galaxies are distributed in filaments, sheets and  clusters all of which form  an interconnected network known as the Cosmic Web. It is a big  challenge to quantify the shapes of the  interconnected structural elements  that  form this   network. Tools  like the Minkowski functionals  which use  global properties, though well suited for an isolated object like  a single sheet or filament, are not suited for an interconnected network of such objects. We consider the Local Dimension $D$, defined through $N(R)=A R^D$, where $N(R)$ is the galaxy number count within a sphere of comoving radius  $R$ centered on  a particular galaxy, as a tool to locally quantify  the shape  in the neigbourhood of different galaxies along the Cosmic Web. We expect $D \\sim 1,2$ and $3$ for a galaxy located in a filament, sheet and cluster respectively. Using LCDM  N-body simulations  we find  that it is possible to determine $D$ through a power law fit to $N(R)$ across the length-scales $2$ to $10 \\, {\\rm Mpc}$ for $\\sim 33 \\%$ of the galaxies. We have visually identified the filaments and sheets corresponding to many of the galaxies  with $D \\sim 1$ and $2$ respectively. In several other situations the structure responsible for the $D$  value could not be visually identified, either due to its being  tenuous or due to  other dominating structures in the vicinity. We also show that  the global distribution of the $D$ values can be used to visualize and interpret how the different structural elements are woven into the Cosmic Web. ",
        "introduction": "Filaments are  the most prominent  features visible in the galaxy distribution.  This finding dates back to a few papers in the seventies and eighties \\citep{joe, einas4,zel}. Subsequent work substantiates this (e.g. \\citealt{gel}, \\citealt{shect}, \\citealt{shand2}, \\citealt{bharadwaj2000}, \\citealt{mul}, \\citealt{basil}, \\citealt{doro2}, \\citealt{pimb}) and shows the filaments to be statistically significant \\citep{bharadwaj2004,pandey}. It is now well accepted that galaxies are distributed in an interconnected network of clusters, sheets and filaments encircling voids. This complicated  pattern is often referred to as the Cosmic Web.  Despite this progress, it still remains a challenge to quantify the Cosmic Web that is so distinctly visible in galaxy redshift surveys (eg. SDSS DR5, \\citet{adel}).  Statistical measures like  the  void probability function \\citep{White1979}, percolation analysis \\citep{shandarin1983} and  the genus curve \\citep{Gott1} each quantifies a different aspect of the Cosmic Web. The Minkowski functionals \\citep{mecke1994} are very effective to quantify the shapes of  individual structural elements like sheets or filaments. In $3$ dimensions there are $4$ Minkowski functionals, namely the volume, surface area, integrated mean curvature and integrated Gaussian curvature.  \\citet{Sahni1998} introduce   the 'Shapefinders', essentially ratios of the  Minkowski functionals,  as a very effective shape diagnostic. A $2$ dimensional version of Shapefinders \\citep{bharadwaj2000} has been  extensively used  to quantify  the filamentarity in the galaxy distribution (\\citealt{pandey3} and references therein).   \\begin{figure} \\includegraphics[width=0.4\\textwidth]{fig2.eps} \\caption{This shows an interconnected  sheet and filament. Also shown, a sphere of  radius $R$ centered on a galaxy located in the filament.} \\label{fig:exp1} \\end{figure}  Though the Minkowski functionals and the Shapefinders are very effective techniques to quantify the shapes of individual structural elements like sheets or filaments,  it is very  different when dealing with the  Cosmic Web which is an interconnected network of filaments, sheets and clusters. For example consider a sheet connected to a filament as shown in Figure \\ref{fig:exp1}.  The Minkowski functionals are global properties of the entire object {\\it ie.} the area is the sum of the areas of the sheet and the filament etc., and the fact that  object is actually a combination of two different elements would be lost. It is necessary to quantify the local shape at different points in the object in order to determine that it actually is a combination of a sheet and a filament. In this paper  we consider the ``Local Dimension'' as a means to quantify the local shape of the galaxy distribution at different positions along the Cosmic Web. ",
        "conclusions": "\\begin{figure} \\includegraphics[width=.45\\textwidth]{f6.eps} \\caption{The  galaxies within $10 \\, {\\rm Mpc}$ of a  center with $D \\sim 1$ are shown as points. The galaxy distribution was converted to a set of 1s and 0s on a grid of spacing $1 \\, {\\rm Mpc}$. Connected structures were identified using the Friend-of-Friend algorithm where adjacent 1s are identified as belonging to the same connected structure. Only the two largest structures have been shown, both pass very close to the center and they appear connected in the figure.  The $D$ values have been painted on these structures with red, blue and green representing $D \\sim 1,2$ and $3$ respectively. The $D$ value is undetermined for  the  parts of the structures  shown in black.} \\label{fig:fil1n} \\end{figure}  The number of centers for which it is possible to determine a Local Dimension varies between $69,017$ to $73,481$ {\\it ie.} less than $7 \\%$ variation across the $10$ different realizations. The $D$ values in the interval $1 \\pm 0.5$, $2 \\pm 0.5$ and $3 \\pm 0.5$ have respectively been binned as $D \\sim 1,2$ and $3$. We  have  inspected  the galaxy distribution in the vicinity of a few  centers in order to visually identify the structures  corresponding to the respective $D$ values. Figure \\ref{fig:fil1n} shows the galaxy distribution within $10 \\, {\\rm Mpc}$ of a center which has Local Dimension $ D \\sim 1$.  We expect this center to be located in a filament. We have used  the Friend-of-Friend algorithm with linking length $\\sqrt{3}  \\, {\\rm Mpc}$  to identify   connected patterns in the galaxy distribution. We note that all the galaxies in this region connect up into a single structure if the linking length is doubled to $2 \\times \\sqrt{3} \\, {\\rm Mpc}$.  Using $\\sqrt{3}  \\, {\\rm Mpc}$ yields several  disconnected  structures of which we show only the  largest two.  Both the structures pass  close to the center, and appear connected in the figure. Of the two structures, one  with two  well separated  tentacles appears below the center, whereas the other with two nearly connected tentacles appears  above it. The filamentary  nature of these structures is quite evident. Out of the total $79$ galaxies in these two structures, $D$ is undetermined for $58$. There are $8$ galaxies with $D\\sim 1$, all located in a contiguous region near the center,  and   $6$ and $7$ galaxies with $D\\sim 2$ and $3$ respectively spread out along the structure  at locations quite different  from the center. While it is relatively easy  to visually identify the  filament corresponding to the centers with  $D\\sim1$, we have not been able to visually identify the sheets and clusters corresponding to  the centers with $D\\sim 2$ and $3$ shown in   Figure \\ref{fig:fil1n}. We note that a  visual identification of the sheets and clusters  corresponding to $D \\sim 2$ and $3$ respectively has been  possible in other fields where the largest structure is predominantly sheetlike or clusterlike.  It is quite apparent that the largest structure in Figure \\ref{fig:fil1n} is predominantly filamentary.   \\begin{figure} \\includegraphics[width=.45\\textwidth]{all_shown.eps} \\caption{This  shows the distribution of $D$ values in a box of size $[40 \\, {\\rm Mpc}]^3$,  red, blue and green denote $D \\sim 1,2$ and $3$ respectively.} \\label{fig:allshown} \\end{figure}  Figure \\ref{fig:allshown} shows the  distribution of $D$ values over a large   region from one of the realizations. The distribution shows distinct patterns, there being regions with size ranging from a few ${\\rm   Mpc}$    to tens of ${\\rm Mpc}$ where the $D$ value is constant. The centers  with $D \\sim 1$ appear to have a more  dense and  compact  distribution compared to the centers with $D \\sim 2$, whereas the centers with $D \\sim 3$ appear to  have a rather diffuse distribution.  \\begin{figure} \\rotatebox{270}{\\includegraphics[height=0.45\\textwidth]{comp.ps}} \\caption{(A) The fraction of centers with a particular $D$ value, and (B) the volume fraction  containing a particular $D$ value. The bins in $D$ have size $\\pm 0.25$. The solid and dashed curves use $D$ values   determined over length-scales $2-10 \\, {\\rm Mpc}$ and $2-5 \\, {\\rm Mpc}$ respectively. } \\label{fig:hist1} \\end{figure}  \\begin{figure} \\rotatebox{270}{\\includegraphics[height=.45\\textwidth]{nuhist.ps}} \\caption{For each $D$ value we separately show  the fraction of centers  as a function of the  density environment. The latter is quantified through $\\nu=\\delta/\\sigma$ where $\\delta$ is the density contrast and $\\sigma^2$ its  variance. This is evaluated on a grid of spacing $1 \\,  {\\rm Mpc}$  and the density field is smoothened with a Gaussian of width $5 \\, {\\rm Mpc}$. } \\label{fig:hist2} \\end{figure}  Figure \\ref{fig:hist1} separately quantifies two different  aspects of the distribution of  $D$ values, (A.) the fraction  of centers with a particular  $D$ value, and (B.)  the fraction  of volume with a particular  $D$ value. For this analysis the $D$ values  were divided into bins of width $\\pm 0.25$ and the centers for which a $D$ value could not be determined were discarded. We find that the fraction of centers and the volume fraction both have  very robust statistics with a variation of $\\sim 10 \\%$ or less across the $10$ different realizations. We first consider the distribution of the fraction  of centers with different  $D$ values. The bin  at $D=1.5$ contains the maximum number of centers, and the two  bins  at $D=1.5$ and $2$ together contain more than $50 \\%$ of the centers. This indicates that the maximum number   of centers for which $D$ can be determined lie in  sheets and filaments. Assuming that this is representative of the entire matter distribution, we conclude that the bulk of the matter is contained in sheets and filaments, with the sheets dominating. We next consider the the fraction of  volume occupied by  a particular $D$ value. We have estimated this using a mesh with grid spacing $1 \\, {\\rm Mpc}$.  Each  grid  position was assigned the $D$ value of the centers within the corresponding  grid cell. Cells that contain centers with different   values of $D$ were discarded. It was possible to assign a $D$ value to $\\sim 35,000$ grid cells. We find that the distribution peaks at $D=2.5$, the peak is rather broad and more than $50 \\%$ of the volume  for which it was possible to assign a $D$ value  is occupied by  values in the range $2$ to $3$. Assuming that this  result is  valid for the entire volume, we conclude that the volume is predominantly filled with sheets and clusters.  While the matter is mainly contained in sheets and filaments, the filaments do not occupy a very large fraction of the volume. This leads to a picture where the filaments have relatively higher  matter  densities  as compared to clusters. To test this we  separately consider the centers with $D$ values $1,2$ and $3$, and for each value we determine the fraction of centers as a function of the  density environment. The density was calculated on a $1 \\, {\\rm Mpc}$ grid after smoothing with  a Gaussian filter of width $5 \\, {\\rm Mpc}$.  We find  (Figure \\ref{fig:hist2}) that in comparison  to the centers with $D \\sim 1$ and $2$, the centers  with  $D \\sim 3$  are more abundant  in under-dense regions  in preference to the over-dense regions.  This is consistent with the picture where  filaments and sheets located along the Cosmic Web  contain  the bulk of the matter and also  most of the high density regions, whereas the centers  with $D \\sim 3$  are predominantly located in under-dense regions away from the Cosmic Web, namely the voids.  The results presented above are specific to the length-scale range $2-10 \\, {\\rm Mpc}$.  It is quite  possible that the relative abundance  of clusters, sheets and filaments  would be quite different if the same analysis were carried out over a considerably different range of length-scales, say for example $0.1 - 1 \\, {\\rm Mpc}$.  Here we have repeated the entire analysis using a smaller range of length-scales $2-5 \\, {\\rm Mpc}$ for which the results are also shown in Figure \\ref{fig:hist1}. We find that the results are qualitatively similar to those obtained using  $2-10 \\, {\\rm Mpc}$, indicating that our  conclusions regarding the relative abundance of clusters, sheets and filaments  are quite robust and are not very sensitive to small changes in the range over which $D$  is determined.  This also indicates  that the $D$ values determined using length-scales $2-10 \\,{\\rm Mpc}$ are not severely biased by the presence of other structural elements with $5-10 \\, {\\rm   Mpc}$ of the one on which the center is located.   In conclusion we note that the Local Dimension provides a robust method to quantify the shapes and probe the distribution of the  different, interconnected   structural elements that  make up the Cosmic Web. We plan to apply this to the SDSS and other galaxy surveys in future."
    },
    "0812/0812.3491_arXiv.txt": {
        "abstract": "{ A complete map of the 3D distribution of molecular (CO) gas was constructed using a realistic dynamical model of the gas flow in the barred potential of the Milky Way. The map shows two prominent spiral arms starting at the bar ends connecting smoothly to the 4-armed spiral pattern observed in the atomic hydrogen gas in the outer Galaxy. Unlike previous attempts, our new map uncovers the gas distribution in the bar region of the Galaxy and the far side of the disk. For the first time, we can follow spiral arms in gas as they pass behind the galactic centre. ",
        "introduction": "Half a century ago, \\cite{Oort58} used the Leiden/Sydney 21 cm line survey to construct a map of the neutral atomic hydrogen gas distribution for the Milky Way (Fig.~\\ref{oort}). This map, the first large scale map of the Milky Way's gas distribution - mostly located in spiral arms, was disturbed by distance errors and excluded the inner part of the Galaxy as well as the region beyond. The apparent expansion of the Galactic centre region was speculated to be due to a bar, and this view has been generally accepted in the last decade. The method used by Oort, however, assumed circular rotation of the gas, since distance information was not known, and therefore was not applicable in the innermost $\\sim5\\,\\kpc$'s of the Galaxy. The map showed spiral arms, but also many fingers pointing towards earth (the ``Finger-of-God'' effect), and the pattern was incomplete in direction of the centre and anti-centre. There was also an apparent difference between spiral arms on the left and right side of the Galaxy; not a single spiral pattern could fit the observations.  \\begin{figure}[] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true,angle=90]{englmaier1.eps}} \\caption{ Map of neutral atomic hydrogen (21-cm line) published by \\protect{\\cite{Oort58}}; figure taken from the text book  \\cite{ElsaesserBook}. The Sun is in the upper part of the plot at $8\\,\\kpc$. } \\label{oort} \\end{figure}  Since then, many studies attempted to chart the spiral arm pattern of the Milky Way in several tracers, but often assuming circular rotation laws for translating radial velocity into distance. Some tracers follow a 4-armed spiral pattern, while others only follow 2 arms (see \\cite{Vallee95} for a review). The perhaps most successful attempt to chart spiral arms, was achieved by \\cite{GG76}, who used HII regions and also a circular rotation law for distance estimation, or more direct methods for nearby objects.  More recently, \\cite{NakanishiSofue} used the $^{12}$CO ($J=1-0$) survey data of \\cite{Dame01} to recover the 3D distribution of the molecular gas in the Milky Way. Again, a circular rotation law was assumed, and the area beyond the Galactic centre was excluded. The face-on view is compatible with a 4-armed spiral pattern.  In the outer disk, spiral arms have been traced by analyzing the  HI layer thickness \\citep{Levine06}, again finding at least four spiral arms. ",
        "conclusions": "The Milky Way has four symmetric spiral arms in the inner part, two of which, the near and far 3-kpc-arm, end inside corotation, two other arms start at $\\sim4\\,\\kpc$ on the major axis of the bar, continue through corotation, and branch at $\\sim7\\,\\kpc$ into four spiral arms which continue to  $\\sim20\\,\\kpc$.  The outer pattern very likely is a superposition of $m=2$ and $m=3$ spiral density waves, as has been observed in external galaxies similar to the Milky Way.  The observed symmetries might be useful in future attempts to invert the gas distribution. One of the spiral arm branches occurs near to us and might leave an imprint in the local stellar velocity distribution. Models of the gas dynamics should take the symmetry of the two 3-kpc-arms into account."
    },
    "0812/0812.1494_arXiv.txt": {
        "abstract": "{The detection of $\\gamma$-rays from dark matter (DM) annihilation is among the scientific goals of the Fermi Large Area Telescope (formerly known as GLAST) and Cherenkov telescopes.} {In this paper we investigate the existence of realistic chances of such a discovery selecting some nearby dwarf spheroidal galaxies (dSph) as a target, and adopting the DM density profiles derived from both astronomical observations and $N$-body simulations. We also make use of recent highlights about the presence of black holes and of a population of sub-subhalos inside the Local Group (LG) dwarfs for boost factor studies.} {We study the detectability with the Fermi-LAT of the $\\gamma$-ray flux from DM annihilation in four of the nearest and highly DM-dominated dSph galaxies of the LG, namely Draco, Ursa Minor, Carina, and Sextans, for which the state-of-art DM density profiles were available. We assume the DM is made of Weakly Interacting Massive Particles such as the Lightest Supersymmetric Particle (LSP) and compute the expected $\\gamma$-ray flux for optimistic choices of the unknown underlying particle physics parameters. We then compute the boost factors due to the presence of DM clumps and of a central supermassive black hole. Finally, we compare our predictions with the Fermi-LAT sensitivity maps.} {We find that the dSph galaxies shine above the Galactic smooth halo: e.g., the Galactic halo is brighter than the Draco dSph only for angles smaller than 2.3 degrees above the Galactic Center. We also find that the presence of a cusp or a constant density core in the DM mass density profile does not produce any relevant effect in the $\\gamma$-ray flux due to the fortunate combination of the geometrical acceptance of the Fermi-LAT detector and the distance of the galaxies and that no significant enhancement is given by the presence of a central black hole or a population of sub-subhalos.} { We conclude that, even for the most optimistic scenario of particle physics, the $\\gamma$-ray flux from DM annihilation in the dSph galaxies of the LG would be too low to be detected with the Fermi-LAT.} ",
        "introduction": "\\label{sec:introduction}  Since the first evidence of the presence of dark matter (DM) in the universe, scientists have worked to understand its nature and distribution. This investigation involves different fields of research such as particle physics, cosmology and, observational astronomy (e.g. \\cite{Bahcall1999, Spergel2003}).  The Fermi Large Area Telecope (Fermi-LAT) will test theories in which DM candidates are the Lightest Supersymmetric Particles (LSPs) such as the neutralinos, arising in Supersymmetric extensions of the Standard Model of particle physics (SUSY), or the Lightest Kaluza-Klein Particles (LKKPs) such as the $B^{(1)}$s, first excitation of the hypercharge gauge boson in theories with Universal Extra Dimensions (see \\cite{Bergstrom2000, Bertone2005}, and references therein). Typical values for the mass of these candidates range from about 50 GeV up to several TeV.  Cosmological models, mainly based on $N-$body simulations in a $\\Lambda$-Cold Dark Matter (CDM) framework, successfully reproduce relevant characteristics of the universe such as the cosmic microwave background anisotropy and the large scale structure of the universe. They also predict well-defined properties of DM haloes, whose radial mass density distribution follows a universal law, and it is described by a steep power law for a wide range of masses ranging from dwarf galaxies to galaxy clusters (see, e.g., \\cite{Navarro1996, Navarro1997, Navarro2004, Moore1998, Moore1999, Diemand2005}). However, the astronomical community is still debating whether DM haloes are characterized by a central density cusp. In fact, haloes with a constant density core are in most cases preferred to account for the observed kinematics of galaxies (see \\cite{Binney2004} for a review).  Generally speaking, the uncertainty in the choice of the density profile can result in several order of magnitude of indetermination on the $\\gamma$-ray flux prediction, which already suffers from the high uncertainties arising by the unknown underlying particle physics (\\cite{Fornengo2004}). For this reason it would be important to derive the DM density profile of galaxies directly from the available kinematic data. Although data-sets for the very inner part of the galaxies are scarce and affected by large errors, the situation is not better in $N-$body simulations, whose resolution goes down to 0.05 times the virial radius at most. Using real data we have the advantage of deriving a flux prediction which takes into account the peculiarity of each galaxy, without any model-dependent generalization which would arise the astrophysical uncertainties.  The expected $\\gamma$-ray flux at the telescope from a given source is directly proportional to the DM density squared along the line-of-sight (LOS), and inversely proportional to the square of its distance. The best targets are therefore nearby dense objects such as the local dwarf spheroidal galaxies (see \\cite{Mateo1998} for a review). Indeed, in the last decade, the large collecting area of the 8-m class telescope and the use of multi-fiber spectrographs allowed astronomers to obtain high-resolution spectra of a large number of stars. This made possible to isolate the galaxy member stars, to measure their radial velocity with an accuracy of few \\kms\\ and to build accurate dynamical models of a number of such systems (see \\cite{Gilmore2007} and references therein).  Annihilation of $\\gamma$-rays in dSph galaxies would give a clean signal because of the absence of high astrophysical uncertainties in modeling the expected background and could hopefully be detected with upcoming experiments like the Fermi-LAT. Many authors studied the feasibility of such a detection, using a large variety of cuspy and cored universal density profiles, reflecting the theoretical as well as the experimental uncertainties. Different works (\\cite{Baltz2000,Tyler2002,Peirani2004,Pieri2004,Bergstrom2006}) found that only the presence of a spike and/or an enhancement due to clumpiness and/or a more favourable combination of the unknown particle physics parameters could make the Draco dwarf galaxy observable with the Fermi-LAT. \\cite{Strigari2007} are optimistic about the detection of Draco with the Fermi-LAT in 5 years. They adopted a King profile (\\cite{King1966}) to derive the surface density of the stellar luminosity. This was deprojected and converted into the stellar mass density by adopting the typical range of the mass-to-light ratio of dSphs. The luminous mass they derived is at the very least one order of magnitude below the mass of the DM halo. For the halo they assumed a NFW mass density profile (\\cite{Navarro1996, Navarro1997}) whose free degenerate parameters were constrained by marginalizing over the stellar velocity dispersion anisotropy parameter. \\cite{Colafrancesco2007} showed how diffuse radio emission would actually be a more promising process to look at in order to detect a DM signal. They also claimed that he presence of a supermassive black hole (SBH) at the centre of Draco, which could enhance the $\\gamma$-ray signal up to detection, is not actually excluded by experiments. Detection of annihilation $\\gamma$-rays from Draco has been recently excluded by \\cite{Sanchez2007} through the use of density profiles which are compatible with the latest observations.  In this paper we use the latest available astrophysical measurements for four of the nearest and highly DM-dominated dSph galaxies of the Local Group, namely  Draco, Ursa Minor, Carina, and Sextans. to compute the expected $\\gamma$-ray flux from DM annihilation. \\\\  In Sect.~\\ref{sec:flux} the most optimistic particle physics scenarios and the DM density profiles derived both from the available kinematic measurements and from $N$-body simulations are used to predict the expected $\\gamma$-ray flux from DM annihilation in Draco, Ursa Minor, Carina, and Sextans. In Sect.~\\ref{sec:glast} the predicted flux is compared with the experimental sensitivity of the Fermi-LAT. The presence of DM clumps and a central SBH could enhance the $\\gamma$-ray flux. But their effects have to be rescaled for the limits imposed on the extragalactic $\\gamma$-ray background (EGB) by the Energetic Gamma-Ray Experiment (EGRET) and on the $\\gamma$-ray flux in Draco by the measurements of the Major Atmospheric Gamma-Ray Imaging Cherenkov (MAGIC) telescope. Our conclusions are given in Sect.~\\ref{sec:conclusions}. \\\\  The main differences with the other papers that already discussed this argument are the following: we show that, even adopting a very lucky case for the unknown particle physics sector, the expected flux from the DM halo is about two orders of magnitude below the detectability limit of the Fermi-LAT experiment; we also show how the use of a cored or a cuspy profiles does not produce any relevant effect in the expected $\\gamma$-ray flux, because of a combination of the galaxy distance and the angular acceptance of the Fermi-LAT; we then show that the current limits on the mass of the supermassive black hole (SMBH) inside Draco lead to an insignificant boost factor due to the presence of such a SMBH; and we numerically compute the boost factor due to the presence of a population of subhaloes inside the dwarf galaxies, limiting the possible range of models for the sub-subhalo structure making use of the constraints imposed by the EGRET extragalactic measurements; the boost factor due to the presence of sub-subhaloes is computed in two ways: first, we obtain it integrating over the whole volume of the galaxy, as done, e.g., in \\cite{Strigari2007}. This gives the correct boost factor when considering cosmological haloes. But, if we consider the closer dwarf galaxies, we have to take into account that only the very inner part of the galaxy is observed within the angular resolution of the instrument. We therefore also compute the angular dependence of the boost factor due to sub-subhaloes. Although we find a huge enhancement of the flux far from the galaxy center, there is actually no enhancement along the LOS pointing toward the galaxy center. As a last improvement with respect to the other papers, we compare our predictions with an recently released detectability map for the Fermi-LAT which takes into account the response of the detector to different energies and incidence angles, as far as effective energy and angular resolution are concerned. ",
        "conclusions": "\\label{sec:conclusions}  The Fermi-LAT telescope was launched in June 2008 and is taking data on $\\gamma$-rays in the energy range between about 20 MeV and 300 GeV. Its all-sky survey operation mode will allow an unprecedented precise study of the $\\gamma$-ray sky, so that many DM models will be tested. The dSph galaxies of the Local Group will be primary targets for DM analysis with the Fermi-LAT, because of the low astrophysical background expected in their direction. Therefore, we studied the detection limit of the $\\gamma$-ray flux from DM annihilation in four of the nearest dSph galaxies, namely Draco, Ursa Minor, Carina, and Sextans.  State-of-art DM density profiles were available for these galaxies and we computed the expected $\\gamma$-ray flux from DM annihilation for different particle physics parameters. We varied the DM particle mass as well as the annihilation cross-section and branching ratios. We found that the presence of NFW-like cusp or constant density core in the DM mass density profile does not produce any relevant effect in the $\\gamma$-ray flux due to a combination of the geometrical acceptance of the Fermi-LAT detector, which is not able to resolve the very inner shells of the studied galaxies, and the distance of the sample galaxies.  In the case of Draco and Ursa Minor, we found that they would shine above the Galactic smooth halo for all but the smallest angles ($\\sim$ 2 degrees) above the Galactic Center. Yet, the upper values of the predicted flux were found to be about two orders of magnitude below the Fermi-LAT detection threshold as derived in \\cite{Baltz2008}. Such values were computed for the most optimistic particle physics scenario of a 40 GeV particle with  $\\sigma_{\\rm ann} v= 3 \\times 10^{-26}$ \\cms annihilating into $b \\bar b$. We have shown how the effect of the boost factor due the presence of a population of DM clumps inside the dSph galaxies, though possibly dramatic (of the order of $10^3$ when integrated over the whole galaxy), had to be rescaled for the limit on the EGB measured by EGRET. The overall maximal effect was reduced to a factor 70. The reader should keep in mind though that the calculation was made for a toy-model where the subhalo population of the dwarf galaxies is described by the same model used for the MW. In any case, since the closest dwarfs are not pointlike for the Fermi-LAT angular acceptance $\\Delta \\Omega$, the factor to be taken into consideration is the effect of the sub-subhaloes inside $\\Delta \\Omega$, which resulted in a decrease of the expected flux due to a redistribution of the DM inside the halo. The presence of a central SBH in agreement with the $M_{\\rm SBH}-\\sigma$ relation extrapolated to the observed low $\\sigma$ values resulted in a negligible boost factor. More extreme models would result in a much higher boost factor, though they are not theoretically supported.  Contrarily to previous papers which addressed the presence of subhaloes or of SMBHs to boost the signal, we have demonstrated that the boost factor must be searched for in some exotic extension or modification of the particle physics sector. Unless such a scenario happens, we conclude that the annihilation of DM inside the local dwarfs is unlikely to be detected with the Fermi-LAT.  As a further development, it will be interesting to repeat the present analysis of boost factors for the recently catalogued ultra faint dwarfs of the Local Group. \\cite{Strigari2008} have computed the expected $\\gamma$-ray flux from those sources, deriving the halo parameters from kinematical data. The inclusion of such galaxies in our study will improve the sensitivity of a joint multi-centred likelihood analysis.  It is worth noticing that,since the DM spectra would be the same for all the galaxies, such an analysis could be performed in order to maximise the detection efficiency and to allow portions of the particle physics phase-space to be explored."
    },
    "0812/0812.2424_arXiv.txt": {
        "abstract": "We report an analysis of twins of spectral types F or later in the {\\it 9th Catalog of Spectroscopic Binaries (SB9)}.  Twins, the components of binaries with mass ratio within 2\\% of 1.0, are found among the binaries with primaries of F and G spectral type. They are most prominent among the binaries with periods less than 43 days, a cutoff first identified by Lucy.  Within the subsample of binaries with P$<43$ days, the twins do not differ from the other binaries in their distributions of periods (median P$\\sim 7$ d), masses, or orbital eccentricities.   Combining the  mass ratio distribution in the SB9 in the mass range 0.6 to 0.85 \\msun ~with that measured by Mazeh et al. for binaries in the Carney-Latham high proper motion survey, we estimate that the frequency of twins in a large sample of spectroscopic binaries is about 3 \\%.  Current theoretical understanding indicates that accretion of high specific angular momentum material by a protobinary tends to equalize its masses.   We speculate that the excess of twins is produced in those star forming regions where the accretion processes were able to proceed to completion for a minority of protobinaries.  This predicts that the components of a young twin may appear to differ in age and that, in a sample of spectroscopic binaries  in a star formation region, the twins are, on average, older than the binaries with mass ratios much smaller than 1. ",
        "introduction": "Stars form by fragmentation and grow to their final mass by accretion from the molecular cloud and their immediate surroundings.  Binaries, and higher order multiples, are common products of star formation (see for example Patience et al. (2002), Delgado-Donate et al. (2004) and references therein). Broadly, the same processes of fragmentation and accretion also operate on a protobinary, with the additional possibilities of dynamical interactions with circumbinary disks and other protostars in the star forming region.  Bate et al. (2002) in their Fig. 2 list the sequence of events that produced some of their model binaries.  The variety of processes and their diverse sequences are bewilderingly complex. Nonetheless, there is a theoretical consensus that, for binaries with small separations, $< \\sim 10$ AU, the accretion processes favor the formation of binaries with mass ratios approaching 1  (see the pioneering work of Artymowicz, 1983, and the reviews by Bate, 2002, and Clarke, 2007).  We use mass ratio $q$ in the usual sense, $q$ = secondary mass/primary mass.   The distribution of mass ratios, $f(q)$, thus provides a fossil remnant of the processes that produced the final masses of their components.   There are two general approaches to the measurement of $f(q)$.  Mass ratios in angularly resolved binaries are determined by photometric estimates of the component masses.   This approach depends on, and is limited by, using a reliable mass-luminosity relation and knowing the distance to the binaries.  To get close to the epoch of star formation in well-defined young clusters, Patience et al. (1998, 2002) studied the visual binary population  of the Hyades, $\\alpha$ Persei, and Praesepe, determined masses photometrically, and measured $f(q)$ down to $q\\sim 0.4$. In the second approach spectroscopic velocity measurements of a double-lined spectrocopic binary (SB2) yield $q$ directly, thereby avoiding the uncertainties of the mass-luminosity relation and distances. The spectroscopic measurements are usually directed at binaries with much shorter periods than the visible binaries and that are angularly unresolved. The only statistically complete dynamical measurement of $f(q)$ over the full mass ratio range of a well-defined sample of binaries of which we are aware was obtained by Mazeh et al. (2003, M03). Bender (2006) and Bender and Simon (2008) describe the start of such a determination for the Hyades open cluster.  In this paper we are concerned with understanding two aspects of the dynamically measured  $f(q)$ of SB2s. First, M03 measured $f(q)$  of main sequence double-lined spectroscopic binaries (SB2s) drawn from the Carney-Latham survey of high proper motion stars  in the galactic disk (Carney  et al. 1994, CL94). They showed that $f(q)$  is approximately constant for $0.3 <q< 1.0$ in the mass range 0.6 to 0.85 \\msun. The second aspect is that of the stellar ``twins'', the primary and secondary components of SB2s with very  nearly the same mass, first identified in Lucy and Ricco's (1979) seminal work.  Tokovinin (2000, T00) showed that twins with $q>0.95$ were present at a small but statistically significant level in a sample of main sequence binaries with primaries of spectral type K0V to F5V and periods in the range 2 to 30 days. Tokovinin suggested that accretion from a circumbinary disk could be responsible for their formation. Lucy (2006, L06) confirmed and sharpened the result using the large sample in the 9th Catalog of the Elements of Spectroscopic Binaries (Pourbaix et al. 2004, SB9). Lucy found that twins in the SB9 are limited to narrow interval $0.98<q<1.00$~ and that the twin phenomenon is most prominent among the spectroscopic binaries with periods less than 43 days.   We wanted to understand why twins did not appear in M03's work.  We therefore sought answers to the following questions in this work:  \\parindent=0.0in  1) What are the properties of the low-mass twins?  How do they differ from other spectroscopic binaries?  2) What is their frequency?  3) What is their origin?  We divided the SB2s in the the SB9 into subsamples by mass that were large enough to provide statisically significant results (\\S 2) and describe the period spectral type, and mass ranges of the twins in \\S 3.   We estimate the frequency at which the twins are found in \\S4.  We discuss our findings and a possible formation mechanism in \\S 5. ",
        "conclusions": "\\subsection{Nature of the Twins}  Twins are found in the range of spectral type K2 to F8 in the HPS (Fig. 9).  That twins are also of masses greater than corresponding to F8 spectral type is indicated by the presence of twins of spectral type F0 and F2 in the HPS. These are included in the ``twin bin'' in Fig. 4.  Twins  with components outside the K2 to F0 range  are not found in the HPS.   We do not know whether this is the result of the difficulty in identifying SB2s and measuring their parameters precisely or of the formation mechanism of the twins. We believe the former is the more likely because Fig. 1 indicates that most of the binaries in the SN3 are of F and G spectral type, Fig. 6 indicates that the sample of K-star binaries available in the HPS is too small to measure their number reliably, and because we cannot identify a reason for a cutoff in twin formation outside this range.   The twins are distinguished from the other binaries with very precisely determined parameters only by the fact that they are found at periods less than 43 days, the demarcation identified by Lucy (2006) or, in our HPS, at periods less than 25 days.  In fact, half of the twins have periods less than 7 days. For a typical component mass of 0.75 \\msun, hence 1.5 \\msun ~total mass of the binary twin, the 7 and 43 day periods correspond to semi-major axes 0.08 and 0.28 AU, respectively. Theoretical modeling shows that the high mass ratio binaries form by the accretion of high specific angular momentum material from the molecular cloud and that the separation of the protobinary components increases as the accretion continues (Bate 2000 and earlier references therein).   Our empirical result indicates that this mechanism operates most effectively on the binaries that now have separations less than $\\sim 0.3$ AU and had smaller separations in their protobinary phase.  \\subsection{A Possible Formation Mechanism}  Within the sample of binaries with P$<43$ days, the twins are not different from the binaries with mass ratios in the range $0.84 \\le q < 0.98$ in their periods, eccentricities, and masses.   The frequency of twins in a large sample of binaries of diverse origins is small,  only about 3 \\%  if our linking the HPS and M03 samples is valid.   This suggests that the accretion processes acting to equalize the masses of the protobinary components proceed to completion at mass ratio 1 in only a small fraction of protobinaries. Hertzsprung-Russell diagrams of diverse star forming regions show that typically there is a spread of a few million years in the  ages of their stars.  It is plausible then that protobinaries do not have the same amount of time available to form their final masses; the most recently formed may experience accretion from the molecular cloud for the shortest time.  The form of the mass ratio distribution at the end of binary formation in a region thus contains information not only about the accretion processes but about their history.  These ideas are testable by observation.  First, consider a region where star formation is still continuing, in which astronomers have identified spectroscopic binaries, and have measured their mass ratios.  When stars are young, theoretical isochrones of different ages are well separated in the HR diagram.  It is thus possible to determine reliable estimates of the relative, if not absolute, ages of the stars in the region.  According to the accretion scenario applied to a protobinary, the initially  lower mass secondary preferentially gains in mass.    Our suggestion for twin formation thus predicts that if mass equalization proceeds to completion, the star that was the original secondary will reach its final mass later than the original primary and appear younger.  This may explain the few $\\times 10^5$ year difference in age  of the components of Parenago 1802, a pre-main sequence twins binary in the Orion Cluster (Stassun et al. (2008). On average, twins need more time to reach their final mass ratio than do the lower mass ratio binaries. Thus, for star forming regions in which a large sample of SB2s with a distribution of mass ratios is available,  our suggestion also predicts that the twins are older, on average, than the binaries with smaller mass ratios.   According to our scenario, the form of the mass ratio distribution, $f(q)$, in a cluster is determined by the time dependence of the rate of star formation and how the mass accreted is distributed between the secondary and primary during the evolution of the protobinary. Given that an ensemble of SB2s have such a distribution, it is not the existence of twins that is significant but rather their excess above the  average of $f(q)$ at $q < 1$. The excess in an open cluster is determined by the average mass accretion rate during the formation of its stars, and by the duration of star formation.  Calculation of model $f(q)$'s on this scenario is beyond the scope of this paper, but we note that Artymowicz(1983) described early analytical models for the formation of an excess of twins. The mass ratio distribution of the SB2s in an open cluster is a fossil record of its epoch  of star formation.  Measurements of $f(q)$  will require large samples of SB2s with well determined mass ratios.  For specificity, consider the necessary size of a well defined sample in order that it yield a statistically significant number of twins. If twins occur at the estimated frequency 3.4\\%,  a sample size of 250 would yield 9 twins which would be significant at the $3 \\sigma $ level.  The more stringent requirement that the number of twins must be significant as an excess over the average distribution at the $3 \\sigma $ level would require a sample of 560.  Samples such as these are far bigger than available now but the availability of fiber-fed multi-object spectrographs suggests that they are not out of reach."
    },
    "0812/0812.4344_arXiv.txt": {
        "abstract": "We develop analytic approximations of thermodynamic  functions of fully ionized nonideal electron-ion plasma mixtures. In the regime of strong Coulomb coupling, we use our previously developed analytic approximations for the free energy of one-component plasmas with rigid and polarizable electron background and apply the linear mixing rule (LMR). Other thermodynamic functions are obtained through analytic derivation of this free energy. In order to obtain an analytic approximation for the intermediate coupling and transition to the Debye-H\\\"uckel limit, we perform hypernetted-chain calculations of the free energy, internal energy, and pressure for  mixtures of different ion species and introduce a correction to the LMR, which allows a smooth transition from strong to weak Coulomb coupling in agreement with the numerical results. ",
        "introduction": "\\label{sect:intro}  We study the equation of state (EOS) of fully ionized nonideal electron-ion plasmas (EIP). In a previous work \\cite{CP98,PC00}, hypernetted chain (HNC) calculations were performed and analytic formulae were proposed for EOS calculations of EIP containing a single ion species. For \\emph{mixtures} of different ion species, the EOS was calculated using the linear mixing rule (LMR), whose high accuracy at $\\Gamma>1$ was previously confirmed in a number of studies \\cite{HTV77,ChabAsh,Rosenfeld,DWSC96,DWS03}. However, the LMR is inaccurate for weakly coupled plasmas. Some consequences of its violation were studied by Nadyozhin and Yudin \\cite{NadyozhinYudin05}, who showed that the  differences between the linear and nonlinear mixing at moderate Coulomb coupling ($0.1\\lesssim\\Gamma\\lesssim1$) can shift the nuclear statistical equilibrium at the final stage of a stellar gravitational collapse.  In this paper, we perform HNC calculations of the free energy, internal energy, and pressure for mixtures of various kinds of ions in the weak, intermediate, and strong coupling regimes, and suggest an analytic correction to the LMR.  In Sec.~\\ref{sect:plasmpar} we define the basic plasma parameters. In Sec.~\\ref{sect:mix} we calculate the EOS of ion mixtures and propose an analytic formula for the EOS of multicomponent EIP, applicable at any $\\Gamma$ values. The summary is given in Sec.~\\ref{sect:concl}. ",
        "conclusions": "\\label{sect:concl}  We have performed HNC calculations of the free energy, internal energy, and pressure for various ionic mixtures with different fractional abundances of the ion species in a broad range of $Z$ and $\\Gamma$ values. We have constructed an analytic approximation to the deviation from the LMR, which recovers the Debye-H\\\"uckel formula for multicomponent plasmas at $\\Gamma\\ll1$ and the LMR at $\\Gamma\\gg1$, and which describes our calculations at any $\\Gamma$ values, as well as the HNC and MC results for the internal energy of plasma mixtures, available in the literature, with an accuracy better than 2\\%.\\\\"
    },
    "0812/0812.1561_arXiv.txt": {
        "abstract": "We examine influence of the circum-nuclear disc (CND) upon the orbital evolution of young stars in the Galactic Centre. We show that gravity of the CND causes precession of the orbits which is highly sensitive upon the semi-major axis and inclination. We consider such a differential precession within the context of an ongoing discussion about the origin of the young stars and suggest a possibility that all of them have originated in a thin disc which was partially destroyed due to the influence of the CND during the period of $\\sim6\\mathrm{Myr}$. ",
        "introduction": "Recent observations of the central stellar cusp of the Milky Way have revealed a numerous population of massive young stars within the (projected) distance of $\\lesssim0.5\\mathrm{pc}$ from the supermassive black hole (e.g. Genzel et al.~1996, Paumard et al.~2006). The apparent youth of these stars is considered to be in contradiction with their vicinity to the black hole which is assumed to prevent stellar formation due to strong tidal forces acting on the parent gaseous clouds. Further analyses have shown that subset of stars above the projected radius $r\\approx0.04\\pc$ forms a coherently rotating `clockwise' disc (CWS). It has been proposed that these stars may have been born in a massive gaseous disc fragmenting due to it own gravity (Levin \\& Beloborodov~2003, Nayakshin~2006). In spite of its attractivity, this model, cannot explain origin of all young stars in the central parsec as more than one half of them have large inclinations with respect to the CWS. Even if we admit existence of another `counter-clockwise' disc of $\\lesssim 20$ young stars (Genzel et al.~2003, Paumard et al.~2006) which could also have been born via fragmentation of a gaseous disc, we will still be facing a problem of the origin of several tens of massive young stars that apparently do not belong to any disc-like structure in the Galactic Centre.  The aim of this contribution is to discuss dynamical processes that should be taken into considerations about the relation of the initial and current kinematical states of the young stars. In the following section, we present notes on the dynamics of stars in the idealised model of the Galactic Centre. In Section~\\ref{sec:results} we discuss consequences of the orbital evolution of the young stars with regard to the stability of a disc-like stellar system. We summarise our results in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} Within our idealised model, the orbits of stars in the central parsec of the Galaxy precess around the symmetry axis of the massive molecular torus. We have verified by means of numerical integration of the equations of motion that this process is robust and it is not substantially modified when the axial symmetry is broken and CND is represented by a finite number of discrete particles. We estimate that relaxation processes act on the young stars on time-scales considerably longer than their presumed lifetime. Therefore, we suggest that the precession due to the gravity of the CND played a dominant role in the orbital evolution of these stars during the past $\\sim6\\mathrm{Myr}$.  The rate of precession is sensitive upon the semi-major axis of the orbit and its inclination with respect to the CND. This process is of particular importance for the dynamics of a stellar disc -- the differential precession tends to destroy a disc-like system which consists of orbits within some range of semi-major axis an inclination. As an example, we have shown that for a system with $a \\in (0.04\\mathrm{pc},\\, 0.4\\mathrm{pc})$ and an opening angle of $\\approx 5\\degr$ which is nearly perpendicular to the CND, the precession will drag normal vectors of a considerable fraction of stars by more than $20\\degr$ away from their original positions.  We point out that, due to the differential precession, there is a nontrivial mapping between initial and current orientations of orbits of the young stars in the Galactic Centre. We suggest a partial destruction (warping) due to the gravity of the massive gaseous torus of an initially coherently rotating disc as an explanation of the origin of those young stars in the Galactic Centre that do not belong to the recently identified clockwise stellar disc. Our hypothesis of a warped disc predicts that normal vectors of all of its member stars should lie close to the circumference perpendicular to the axis of the CND. We have verified that this hypothesis is compatible with the observational data published by Paumard et al.~2006. Nevertheless, the lack of robust determination of the current values of the orbital elements form the observational data prevents a convincing identification of the warped stellar disc. Finally, let us remark that the hypothesis of the single warped stellar disc also cannot explain apparently random orientations of the orbits of the S-stars, as these are the least affected by the gravity of the CND.   \\ack This work was supported by the Research Program MSM0021620860 of the Czech Ministry of Education, the Centre for Theoretical Astrophysics in Prague and the Czech Science Foundation (ref.\\ 205/07/0052)."
    },
    "0812/0812.3564_arXiv.txt": {
        "abstract": "The free photon dispersion relation is a reference quantity for high precision tests of Lorentz Invariance. We first outline theoretical approaches to a conceivable Lorentz Invariance Violation (LIV). Next we address phenomenological tests based on the propagation of cosmic rays, in particular in Gamma Ray Bursts (GRBs). As a specific concept, which could imply LIV, we then focus on field theory in a non-commutative (NC) space, and we present non-perturbative results for the dispersion relation of the NC photon. ",
        "introduction": "Lorentz Invariance (LI) plays a central r\\^{o}le in relativity: it holds as a global symmetry in Special Relativity Theory, and as a local symmetry in General Relativity Theory. We refer to the former in order to be able to describe particle physics in terms of quantum field theory, where a field $\\Phi$ (scalar, vector, tensor or spinor) is transformed globally in some representation $D$ of the Lorentz group $SO(1,3)$, \\be \\Phi (x) \\to U(\\Lambda )^{\\dagger} \\Phi (x) U(\\Lambda ) = D(\\Lambda ) \\Phi (\\Lambda^{-1} x) \\ , \\quad U ~ {\\rm unitary} \\ , \\quad \\Lambda \\in SO(1,3) \\ . \\ee If a local field theory obeys LI, then it is also CPT invariant \\cite{CPT}. If LI is broken, CPT may be intact or not. If, however, CPT is broken, then a LI Violation (LIV) is inevitable \\cite{Greenberg}. Therefore LI can be probed indirectly by testing CPT predictions, such as the equivalence of the mass and magnetic moment of particles and their antiparticles. In particular the masses of $K^{0}$ and $\\bar K^{0}$ coincide to a relative precision $< 8 \\cdot 10^{-19}$ \\cite{databook}.  Specific direct LI tests are even more stringent. In the framework of the Standard Model Extension (SME) \\cite{SME}, local terms are added to the Standard Model Lagrangian, which can cause a spontaneous LIV, {\\it e.g.} \\be {\\cal L} = {\\rm i} \\bar \\psi \\gamma_{\\mu} \\partial^{\\mu} \\psi - g \\bar \\psi \\phi \\psi - {\\rm i} \\bar \\psi (g' G_{\\mu \\nu} + g'' H_{\\mu \\nu} \\gamma_{5}) \\gamma^{\\mu} \\partial^{\\nu} \\psi + \\dots \\ee Spontaneous Symmetry Breaking (SSB) of the Higgs field $ \\phi \\in \\C^{2}$ leads to the fermion mass. In analogy the new tensor fields $G$ and $H$ may also undergo SSB, which affects LI (unlike the SSB of the Higgs field). For a simple change of the observers frame, $G$ and $H$ are transformed as well, so that ${\\cal L}$ remains invariant; but the fermion $\\psi$ perceives them as an invariant background. LIV is then manifest for instance through a modified fermion dispersion relation, depending on the Yukawa-type couplings $g'$, $g''$. Kosteleck\\'{y} {\\it et al.} identified more than 100 parameters of this kind (for a review, see Ref.\\ \\cite{Bluhm}), which are claimed to be compatible with the usual properties of the Standard Model except for LI (and in part CPT) \\cite{ColKosCPT}. This corresponds to a low energy expansion, after lifting the usual LI condition. This collaboration interprets the photon as a Nambu-Goldstone boson \\cite{SSBinSME}. All LIV parameters emerging in this framework are experimentally bounded by $O(10^{-27})$ (in a dimensionless form, which factors out the Planck scale) \\cite{LIVbounds}. If some of these parameters are non-zero, they still allow for sizable LIV effects at tremendous energies, though this scenario poses a severe hierarchy problem \\cite{LIVfinetune}.  In an effective approach, Coleman and Glashow considered {\\em Maximal Attainable Velocities} (MAVs, attained in the high energy limit), which could slightly differ from the speed of light, depending on the particle type \\cite{ColGla}. They were most interested in the impact on cosmic rays (which attain the highest energies in the Universe). In particular the prediction that proton rays cannot exceed the Greisen-Zatsepin-Kuz'min (GZK) energy of $E_{\\rm GZK} \\approx 6 \\cdot 10^{19} ~ {\\rm eV}$ over a distance of more than $\\sim 100 ~ {\\rm Mpc}$ in the Cosmic Microwave Background (CMB) (mainly due to photopion production via a $\\Delta$-resonance) could be evaded by tiny MAV differences. However, recent observations favour the validity of the GZK energy cutoff \\cite{GZKholds}: up to $E_{\\rm GZK}$ the cosmic ray flux follows closely a behaviour $\\propto E^{-2.7}$, but beyond $E_{\\rm GZK}$ there is an extra suppression compared to this power law. This suggests that LI holds even under boosts of Lorentz factors like $\\gamma = E_{\\rm GZK}/ m_{\\rm proton} \\sim 10^{11}$, {\\it i.e.}\\ 6 orders of magnitude beyond the $\\gamma$-factors accessible in laboratories, which further constrains the MAV differences \\cite{ScuSte08}.  This issue is reviewed in Ref.\\ \\cite{cosmo}, covering theoretical and phenomenological aspects and the current status, along with MAV effects in cosmic rays with respect to decays, \\v Cerenkov radiation and neutrino oscillation \\cite{nuMAV} (these effects were addressed already in Ref.\\ \\cite{ColGla}, and reviewed earlier in Ref.\\ \\cite{JLM}). ",
        "conclusions": "Cosmic $\\gamma$-rays enable LI tests to extreme precision. So far LIV has not been detected anywhere --- we reviewed failed attempts based on GRB data.  As theoretical frameworks leading to LIV we discussed on one side low energy effective theories \\cite{SME,Bluhm,ColKosCPT,SSBinSME} and an effective ansatz which introduces particle specific Maximal Attainable Velocities \\cite{ColGla}. Another approach is based on NC geometry --- the presence of IR divergences in quantum field theory on such spaces clarifies its distinction from low energy effective theories.  In Section 4 we were concerned with the kinematics of a NC photon. Perturbation theory to one loop suggests its IR instability \\cite{LLT}, which looks like a disaster for NC field theory. The picture changes, however, in light of a non-perturbative consideration.  Our numerical study revealed a phase of spontaneously broken Poincar\\'{e} symmetry, which is not visible to perturbative calculations. In fact the limit to the continuum and to infinite volume at fixed non-commutativity (the Double Scaling Limit) leads to this phase of broken symmetry,and {\\em not} to the symmetric weak coupling phase. In the broken symmetry phase we observed stability for a number of observables \\cite{NCQED}. In particular the photon dispersion relation in a commutative sub-space appears linear. However, its behaviour at non-zero momentum components in the NC directions is not explored yet. In that respect, non-perturbative results could be confronted with phenomenological data --- in particular from GRBs --- to establish a robust bound on $\\Vert \\Theta \\Vert$ in Nature. On the phenomenological side, more stringent data are expected from new large-scale GRB observations, for instance in the projects discussed in Refs.\\ \\cite{Lamon,SVOM}.  Our non-perturbative results so far imply that the photon may survive in a NC world --- despite the IR disaster on the 1-loop level --- without the requirement to proceed to SUSY.  \\begin{acknowledgement}  It was my pleasure to attend the 4$^{\\rm th}$ EU RTN Workshop on ``Constituents, Fundamental Forces and Symmetries of the Universe'' in the beautiful city of Varna (Bulgaria), where this talk was presented. I would like to thank my collaborators for their contributions to Refs.\\ \\cite{BBT,NCphi4,NCQED,2dNCU1}, the authors of Ref.\\ \\cite{Ellis} for their permission to reproduce Fig.\\ \\ref{Ellisfig}, and F.W.\\ Stecker for a helpful comment. This work was supported by the Deutsche Forschungsgemeinschaft (DFG) through Sonderforschungsbereich SFB/TR55 ``Hadron Physics from Lattice QCD''.  \\end{acknowledgement}"
    },
    "0812/0812.4491_arXiv.txt": {
        "abstract": "Two kinds of difficulties have challenged the physics community for many years: (1) knowing nature's building blocks (particle physics) and (2) understanding interacting many-body systems (many-body physics). Both of them exist in the research of quark matter and compact stars. This paper addresses the possibility that quark clustering, rather than a color super-conducting state, could occur in cold quark matter at realistic baryon densities of compact stars, since a weakly coupling treatment of the interaction between quarks might not be reliable. Cold quark matter is conjectured to be in a solid state if thermal kinematic energy is much lower than the interaction energy of quark clusters. Different manifestations of pulsar-like compact stars are discussed, as well as modeled in a regime of solid quark stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.3093_arXiv.txt": {
        "abstract": "The disruption of the M33 galaxy is evident from its extended gaseous structure. We present new data from the Galactic Arecibo L-Band Feed Array HI (GALFA-HI) Survey that show the full extent and detailed spatial and kinematic structure of M33's neutral hydrogen. Over 18\\% of the HI mass of M33 (M$_{\\rm HI_{tot}}=1.4 \\times 10^9$ \\Msun) is found beyond the star forming disk as traced in the far-ultraviolet (FUV).  The most distinct features are extended warps, an arc from the northern warp to the disk, diffuse gas surrounding the galaxy, and a southern cloud with a filament back to the galaxy.   The features extend out to 22 kpc from the galaxy center (18 kpc from the edge of the FUV disk) and the gas is directly connected to M33's disk.  Only five discrete clouds (i.e., gas not directly connected to M33 in position-velocity space) are catalogued in the vicinity of M33, and these clouds show similar properties to Galactic and M31 halo clouds. M33's gaseous features most likely originate from the tidal disruption of M33 by M31 1-3 Gyr ago as shown from an orbit analysis which results in a tidal radius $<15$ kpc in the majority of M33's possible orbits.  M33 is now beyond the disruptive gravitational influence of M31 and the gas appears to be returning to M33's disk and redistributing its star formation fuel. M33's high mean velocity dispersion in the disk ($\\sim$18.5 \\kms) may also be consistent with the previous interaction and high rate of star formation. M33 will either exhaust its star formation fuel in the next few Gyrs or eventually become star formation fuel for M31.  The latter represents the accretion of a large gaseous satellite by a spiral galaxy, similar to the Magellanic Clouds' relationship to the Galaxy. ",
        "introduction": "Large spiral galaxies at low redshift are thought to gradually evolve through the accretion of material from their halos.  This halo material originates from the accretion of satellite galaxies and the cooling of the hot gaseous halo, with the hot halo most likely originating from a combination of the initial baryon collapse and galaxy feedback.   Continual gas accretion provides star formation fuel and allows for the range of stellar populations and metallicities we find in galaxies. Simulations that examine the methods galaxies obtain gas indicate the dominant mechanism depends on a galaxy's total mass and redshift \\citep{keres09, keres05, brooks09}.  At $z=0$ massive galaxies are  expected to accrete gas through satellites and the gradual accretion of hot halo gas, while less massive galaxies ($<$ few $\\times~10^{11}$ \\Msun) obtain star formation fuel primarily via direct accretion of cold gas along cosmic filaments.   Observational evidence for the various accretion modes remains limited.  The Galaxy and M31 have been shown to have numerous halo clouds (e.g., Muller et al. 1963; Putman et al. 2002; Thilker et al. 2004\\nocite{muller63, putman02, thilker04}), with their origin definitely partly from satellite accretion \\citep{mathewson74, putman03}, and possibly also from condensing halo clouds in the hot halo medium \\citep{peek08,maller04}.  There has been additional evidence for halo clouds that will fuel spiral galaxies through deep HI observations of nearby spirals \\citep{oosterloo07, fraternali01}.  In all of these cases, the amount of fuel currently found in the halo is not large considering the expected star formation rates and chemical evolution models. For instance, the Milky Way is expected to have an average infall of 1 \\Msun/yr \\citep{chiappini01, fenner03}, but this reservoir is not readily apparent from either the accretion of gas from small satellites \\citep{grcevich09} or the condensing halo clouds \\citep{peek08}. In the case of galaxies smaller than the Milky Way, observational evidence for halo clouds remains limited, and it is possible they will largely not be found, as the bulk of the accretion may have been completed for these systems at early times.   M33 is often considered a dwarf spiral galaxy of the Local Group, with a total mass of $\\sim5\\times10^{10}$ \\Msun (Corbelli 2003\\nocite{corbelli03}), which is $\\sim$20 times lower than the Milky Way and M31, and only $\\sim$2 times higher than the combined mass of the Magellanic Clouds. M33 is $\\sim730$ kpc from the Milky Way and $\\sim$ 200 kpc from M31 \\citep{brunthaler05}. Its stellar component indicates it is a fairly quiescent galaxy with no evidence for a recent interaction. Deep optical surveys have found a stellar halo component, but no clear evidence for the stream-like sub-structures found in the halo of the Milky Way and M31 (Ferguson et al. 2007; McConnachie et al. 2006\\nocite{ferguson07, mcconnachie06}). In neutral hydrogen there is evidence for a strong warp that has been examined by \\cite{corbelli97} (hereafter CS97) and Rogstad, Wright \\& Lockhart (1976).  Both authors concluded that the tidal force from M31 is unlikely to be sufficient to generate the warp.   On the other hand, recent work estimating the motion of M31 from indirect methods indicates that M33 and M31 are almost certainly gravitationally bound and M33 should show signatures of tidal perterbation \\citep{vandermarel08,loeb05}.   Sensitive large area neutral hydrogen surveys of M33's extended halo can investigate the methods that M33 is losing and/or obtaining star formation fuel, and the fate of this galaxy in the context of galaxy mergers at $z=0$. Recent observations taken as part of the GALFA-HI Survey provide these data at unprecedented resolution (3.4\\arcmin, 0.18 \\kms) and sensitivity (0.1 K ($3\\sigma$) over 5 \\kms).\\footnote{The data are publicly available through a link at http://sites.google.com/site/galfahi/.}  See also Grossi et al. (2008\\nocite{grossi08}) for data taken with a different spectrometer and reduced with different techniques as part of the ALFALFA survey. In this paper we present the GALFA-HI Survey observations of the HI environment of M33, complete simulations of the orbit of M33 relative to M31, and discuss the results in the context of fueling both M33 and M31.  The gaseous environment of M33 is an important tracer of the past and future evolution of the galaxies of the Local Group. ",
        "conclusions": "We have presented new HI data of M33 and its surroundings from the GALFA-HI Survey and have highlighted several distinct gaseous features that extend out to 18 kpc from the edge of the stellar component traced in FUV emission (Figure~\\ref{hi} \\& \\ref{hifuv}; Section 3.1).  The maximum extent of M33's HI is 28 $\\times$ 38 kpc, and includes features such as a northern arc of gas that extends from M33's warp back towards the disk, a southern cloud with a filament to the galaxy, and diffuse emission and several fingers to the east and west of the galaxy (Figures~1-6).   The HI gas beyond the star-forming disk contains 18\\% of M33's total HI mass of $1.4 \\times 10^9$ \\Msun\\ and the vast majority of the gas surrounding M33 is directly connected to the disk in position-velocity space.  The GALFA-HI data is compared to the warped disk model of CS97 to highlight the relationship of the extended gas to this model (Figures~\\ref{modchans} \\& \\ref{modvel}; Section 3.2). The following additional results are described in the paper. \\begin{itemize}  \\item The orbit of M33 relative to M31 is examined to investigate the feasibility of a tidal interaction between the two galaxies resulting in M33's gaseous features (Section~\\ref{sec-orbit}).  Consistent with previous analyses of the likely orbits, we find a previous close encounter ($<$100 kpc) between the two galaxies is highly likely ($> 75\\%$ of the orbits for the most likely V(M31)$_{tan} < 150$ \\kms), resulting in a tidal radius for M33 of $< 15$ kpc.   This suggests a satellite of M33 would not remain bound to it and the gaseous features were most likely tidally pulled directly from M33 in the past 3 Gyr.   \\item The combination of tidal and ram pressure forces acting on M33 as it passes through an extended halo of M31, is likely to cause the tidally extracted gas to lose angular momentum and fall back towards M33's disk as it moves away from M31.   This is consistent with the velocity sense of the observed gas and the elevated star formation in M33's disk.  \\item An automated object finder is run on the M33 cube to define truly discrete objects from those directly connected to the galaxy (Sections 3.3 \\& 5.1).  Besides the clouds associated with the Wright's cloud complex (which is unlikely to be associated with M33 given its offset and large size at M33's distance), there are only 5 small clouds catalogued, and all other gas is connected to M33.  Of the 5 clouds, 2 are very close to being connected to M33, and the other 3 are either in a filament going out 38 kpc from the center of M33 or part of a Galactic filament given their lower velocities.  These clouds show similar properties to Galactic HVCs and M31 halo clouds, including their velocity widths ($\\sim25$ \\kms) which are indicative of a warm cloud component.  \\item M33 will eventually become star formation fuel for M31, and if it does not substantially deplete its gas content before then, it will provide $\\sim$25\\% of M31's current HI mass.   This together with the Magellanic System, which hosts close to 50\\% of the Milky Way's current HI mass, indicates the contribution of gas from large mergers can play an important role in the ongoing fueling of galaxies.  The fact that satellite mergers that are gas-rich are not expected to be overly disruptive to the disk is also support for these types of mergers providing substantial fuel for a spiral galaxy.  \\item M33 will either run out of star formation fuel in the next 2 Gyr if it continues to form stars at its current rate, or it must accrete gas from an unknown Local Group medium.  Other notable characteristics of M33's gas and star formation are the high mean velocity dispersion across the disk ($\\sim18.5$ \\kms; Figure~\\ref{hiveldisp}) that may be consistent with the relatively high SFR and recent interaction, and the fact that M33 does not have star formation as traced in the FUV in its extended gaseous component, or beyond the contours at $3-6 \\times 10^{20}$ cm$^{-2}$ (Figure~\\ref{hifuv}).  \\end{itemize}"
    },
    "0812/0812.4399_arXiv.txt": {
        "abstract": "The huge size and uniformity of the Sloan Digital Sky Survey makes possible an exacting test of current models of galaxy formation. We compare the predictions of the {\\tt GALFORM} semi-analytical galaxy formation model for the luminosities, morphologies, colours and scale-lengths of local galaxies. {\\tt GALFORM} models the luminosity and size of the disk and bulge components of a galaxy, and so we can compute quantities which can be compared directly with SDSS observations, such as the Petrosian magnitude and the S\\'{e}rsic index. We test the predictions of two published models set in the cold dark matter cosmology: the \\citet{Bau05} model, which assumes a top-heavy initial mass function (IMF) in starbursts and superwind feedback, and the \\citet{Bow06} model, which uses AGN feedback and a standard IMF. The Bower et al model better reproduces the overall shape of the luminosity function, the morphology-luminosity relation and the colour bimodality observed in the SDSS data, but gives a poor match to the size-luminosity relation. The \\citeauthor{Bau05} model successfully predicts the size-luminosity relation for late-type galaxies. Both models fail to reproduce the sizes of bright early-type galaxies. These problems highlight the need to understand better both the role of feedback processes in determining galaxy sizes, in particular the treatment of the angular momentum of gas reheated by supernovae, and the sizes of the stellar spheroids formed by galaxy mergers and disk instabilities. ",
        "introduction": "There is a growing weight of evidence in support of the hierarchical paradigm for structure formation \\citep{Spr06}.  The principal process responsible for the growth of structure, gravitational instability, has been modelled extensively using large numerical simulations \\citep[e.g.][]{Spr05,Spr08}. The cold dark matter model gives an impressive fit to measurements of temperature fluctuations in the cosmic microwave background radiation \\citep{Hinshaw08}. When combined with other data, such as measurements of local large scale structure in the galaxy distribution or the Hubble diagram of type Ia supernovae, there is a dramatic shrinkage in the available range of cosmological parameter space \\citep{Per02,Ariel06,Kom08}. The spectacular progress made in constraining the background cosmology has allowed the focus to shift to trying to understand the evolution of the baryonic component of the Universe \\citep{Bau06}.  Over the same period of time there has been a tremendous increase in the quantity and range of observational data on galaxies at different redshifts. Observations at high redshift have uncovered populations of massive, actively star forming galaxies which were already in place when the Universe was a small fraction of its current age \\citep[e.g.][]{Sm97,St99,Blain02,Gia02}. Huge surveys of the local universe made possible by advances in multifibre spectrographs allow the galaxy distribution to be dissected in numerous ways \\citep[e.g.][]{Colless03,Adel08}. It is now possible to make robust measurements of the distribution of various intrinsic galaxy properties, such as luminosity, colour, morphology and size. The trends uncovered in these local surveys are influenced by a wide range of physical effects, such as star formation, supernova and AGN feedback, the cooling of gas and galaxy mergers, and hence provide as strict a test of galaxy formation models as that posed by the high redshift data.  In order to test current ideas about galaxy formation against observations, well developed theoretical tools are needed which can follow many complex processes concurrently. Most importantly, it is vital that the model predictions are produced in a form which can be compared as directly as possible with observations. Gas dynamic simulations typically follow galaxy formation in great detail for an individual dark matter halo \\citep[e.g.][]{Gov04,Takashi05,Gov07} or within some small volume \\citep[e.g.][]{Nag04,Sca06,Croft08}. In general, only limited output is available which can be compared directly with observations, such as, for example, estimates of galaxy luminosity.  Currently the closest contact with observations is made by semi-analytical models of galaxy formation (see \\citet{Bau06} for a review). In their most sophisticated form, these models can make predictions for the luminosity, colour, scale-length and morphology of galaxies in a wide range of environments \\citep[e.g.][]{Bow06,Cro06,Catt06,Monaco07,Claudia08}. These models necessarily have to treat the baryonic physics in a somewhat more idealized way than is done in the gas dynamic simulations. Physical processes are described using rules, some of which contain parameters whose values are set by comparing the model predictions with selected observations. As the range of data the model is compared against increases, the parameter space open to the models reduces. For example, the strength of supernova feedback (see Section 2) has an impact not only on the number of galaxies in the faint end of the luminosity function, but also on the size of galactic disks and even the morphological mix of galaxies.  Two key properties of semi-analytical models are their computational speed and modular nature. The impact of different processes on the nature of the galaxy population can be rapidly assessed by running models with different parameter choices. This makes the models ideal tools with which to interpret observational data. Any discrepancy uncovered between the model predictions and observations can help to identify physical ingredients which may either require better modelling or which may be missing altogether from the calculation. One clear example of how observations drive the development of the models is given by the recent efforts to reproduce the location and sharpness of the break in the galaxy luminosity function.  With the current best fitting cosmological parameters, galaxy formation models struggled to avoid producing too many bright galaxies \\citep{Ben03}. One solution to this problem was suggested by observations showing the apparent absence of cooling flows at the centres of rich clusters (e.g. \\citealt{Pet03}) which motivated the idea of taking into account the energy released by active galactic nuclei (AGN). This acts as a feedback process that heats the gas in massive haloes. The incorporation of this feedback mechanism suppresses the formation of galaxies in massive haloes, such that the right number of bright galaxies can be produced, and, furthermore, these galaxies have red colours to match those observed \\citep{Bow06,Cro06,Claudia08}.  In this paper we test two published galaxy formation models run using the {\\tt GALFORM} semi-analytical model against statistics measured from the Sloan Digital Sky Survey. The \\citet{Bau05} model (hereafter Baugh2005) invokes a ``superwind'' channel for supernova feedback, which ejects gas from low and intermediate mass haloes. This model assumes that stars are produced with a normal initial mass function (IMF) in quiescent disks but with a top-heavy IMF during merger-driven starbursts. The Baugh2005 model is able to reproduce the counts and redshift distribution of sub-millimetre selected galaxies, along with the abundance of Lyman-break galaxies and Lyman-$\\alpha$ emitters \\citep{LeD05,LeD06,Orsi08}. The \\citet{Bow06} model (hereafter referred to as Bower2006) incorporates AGN feedback, with the energy released by the accretion of mass onto the central supermassive black hole in halos with quasistatic hot gas atmospheres being responsible for stifling the cooling rate. The Bower2006 model gives a good match to the evolution of the K-band luminosity function and the inferred stellar mass function.  The paper is organized as follows. In Section 2, we outline the galaxy formation model {\\tt GALFORM} tested in the paper.  In Section 3 we describe how some additional galaxy properties are computed from the model output; these properties are needed to compare the model predictions directly with observations of SDSS galaxies.  Section 4 contains the comparisons between model predictions and SDSS data for the luminosity function, the distribution of morphological types, the colour distribution and the size distribution. In this section we also show the impact on the predictions of changing the strengths of various processes in the model. In Section 5 we present our conclusions.  The appendices discuss how certain photometric properties of galaxies have changed between SDSS data releases and compare different indicators of galaxy type. Finally, we note that magnitudes are quoted on the AB system assuming a Hubble parameter of $h=H_{0}/100$ km s$^{-1}$ Mpc$^{-1}$; the cosmological parameters adopted depend on the choice of semi-analytic model as explained in Section 2.   \\section[]{GALAXY FORMATION MODEL} We use the Durham semi-analytical galaxy formation model, {\\tt GALFORM}, introduced by \\citet{Col00} and developed in a series of subsequent papers \\citep{Ben03,Bau05,Bow06}. The model tracks the evolution of baryons in the cosmological setting of a cold dark matter universe. The physical processes modelled include: i) the hierarchical assembly of dark matter haloes; ii) the shock heating and virialization of gas inside the gravitational potential wells of dark matter haloes; iii) the radiative cooling of the gas to form a galactic disk; iv) star formation in the cool gas; v) the heating and expulsion of cold gas through feedback processes such as stellar winds and supernovae; vi) chemical evolution of gas and stars; vii) mergers between galaxies within a common dark halo as a result of dynamical friction; viii) the evolution of the stellar populations using population synthesis models; ix) the extinction and reprocessing of starlight by dust.  In this section we first give a comparison of the main features of the Baugh2005 and Bower2006 models, introducing the recipes used to model phenomena which are varied in Section 4. Similar discussions of these models can be found in \\citet{Al07}, \\citet{Al08} and \\citet{Violeta08}. In the second part of this section we review the model used to compute galaxy sizes, which was originally devised by \\citet{Col00} and tested by these authors for galactic disks and by \\citet{Al07} for spheroids.   \\subsection[]{A comparison of the Baugh2005 and Bower2006 models} The Baugh2005 and Bower2006 models represent alternative models of galaxy formation. The parameters which specify the models were set by the requirement that their predictions should reproduce a subset of the available observations of local galaxies together with certain observations of high redshift galaxies. Different solutions were found due to the use of different physical ingredients, as set out below, and because different emphasis was placed on reproducing particular observations. We refer the reader to the original papers for a full description of each model; the Baugh2005 model is also described in detail in \\citet{Lacey08}.  We now summarize the main differences between the two models.  \\begin{itemize} \\item {\\it Cosmology.} The Baugh2005 model adopts a $\\Lambda$CDM cosmology with a present-day matter density parameter, $\\Omega_{0}$=0.3, a cosmological constant, $\\Omega_{\\lambda}$=0.7, a baryon density, $\\Omega_{b}=0.04$ and a power spectrum normalization given by $\\sigma_{8}=0.93$. The Bower2006 model uses the cosmological model assumed in the Millennium simulation (\\citealt{Spr05}), where $\\Omega_{0}$=0.25, $\\Omega_{\\lambda}$=0.75, $\\Omega_{b}=0.045$ and $\\sigma_{8}=0.9$, which are in somewhat better agreement with the constraints from the anisotropies in the cosmic microwave background and galaxy clustering on large scales (e.g. \\citealt{Ariel06}).  \\item {\\it Halo merger trees.}  The Baugh2005 model uses a Monte-Carlo technique to generate merger histories for dark matter haloes, which is based on the extended Press-Schechter theory \\citep{Lacey93,Col00}. The formation and evolution of a representative sample of dark matter haloes is followed. In the Bower2006 model, the merger histories are extracted from the Millennium simulation (see \\citealt{Har06} for a description of the trees). The mass resolution of the simulation trees is $1.72 \\times 10^{10} h^{-1} M_{\\odot}$, which is a factor of three worse than that used in the Monte-Carlo trees. By comparing the output of models using Monte-Carlo and N-body merger trees, \\citet{Helly03} found very similar predictions for bright galaxies, with differences only becoming apparent below some faint magnitude, the value of which depends on the mass resolution of the N-body trees.  The resolution of the Millennium simulation yields a robust prediction for the luminosity function to around three magnitudes fainter than $L_{\\star}$, which is more than adequate for the comparisons presented in this paper.   \\item{\\it Quiescent star formation timescale.}  The quiescent star formation rate in disks, $\\psi$, is given by $\\psi = M_{\\rm gas}/\\tau_{\\ast}$, where $M_{\\rm gas}$ is the mass of cold gas and the time-scale, $\\tau_{\\ast}$, is parameterized differently in the two models. In both cases, the star formation timescale is allowed to depend upon some power of the circular velocity of the disk and is multiplied by an efficiency factor. In the Baugh2005 model, the efficiency factor is assumed to be independent of redshift, whereas in the case of the Bower2006 model, this factor scales with the dynamical time of the galaxy \\textbf{[$\\tau_{\\rm dyn}$]}, measured at the half-mass radius of the disk. Since the typical dynamical time gets shorter with increasing redshift, the star formation timescale in the Bower2006 model is shorter at high redshift than it would be in the equivalent disk in the Baugh2005 model. This has implications for the amount of star formation in merger-triggered starbursts (or following a disk becoming dynamically unstable in the Bower2006 model - see later).  In the Bower2006 model, disks at high redshift tend to be gas poor, with the gas being turned into stars on a short timescale after cooling, whereas in the Baugh2005 model, high redshift disks are gas rich.  \\item {\\it Initial mass function (IMF) for star formation.}  The Bower2006 model uses the \\citet{Ken83} IMF, consistent with deductions from the solar neighbourhood, in all modes of star formation. The Baugh2005 model also adopts this IMF in quiescent star formation in galactic disks.  However, in starbursts triggered by galaxy mergers, a top-heavy IMF is assumed. The yield of metals and the fraction of gas recycled per unit mass of stars formed are chosen to be consistent with the form of the IMF.  \\item {\\it Supernova (SN) feedback.} With each episode of star formation, a mass of cold gas is reheated and ejected from the disk by supernova explosions at a rate given by: \\begin{equation} \\dot{M}_{\\rm eject}=(V_{\\rm hot}/V_{\\rm disk})^{\\alpha_{\\rm hot}} \\psi, \\label{eqsnfeedback} \\end{equation} where $V_{\\rm disk}$ is the velocity at the disk half-mass radius, and $V_{\\rm hot}$ and $\\alpha_{\\rm hot}$ are parameters.  The SN feedback is stronger in the Bower2006 model ($V_{\\rm hot}=485 {\\rm km s}^{-1}$ and $\\alpha_{\\rm hot} = 3.2$ compared with $V_{\\rm hot}=300 {\\rm km s}^{-1}$ and $\\alpha_{\\rm hot} =2$ in the Baugh2005 model).   \\item {\\it AGN vs superwind feedback.}  Perhaps the most significant difference between models is the manner in which the formation of very massive galaxies is suppressed. In the Baugh2005 model, an additional channel or fate for gas heated by supernovae is invoked, called superwind feedback. In addition to the standard SNe feedback model described in the previous bullet point, some gas is assumed to be expelled completely from the halo due to heating by supernovae.  The amount of mass ejected is taken to be a multiple of the star formation rate, multiplied by a function of the circular velocity of the halo. The superwind is most effective in removing gas from low circular velocity haloes, with the mass of gas ejected falling with increasing circular velocity. The gas expelled in the superwind is not allowed to recool, even in more massive haloes at later times in the merger history.  This has the effect of reducing the cooling rate in massive haloes since these haloes have less than the universal fraction of baryons. Such winds have been observed in massive galaxies, with the inferred mass ejection rates found to be comparable to the star formation rate (e.g. \\citealt{Heckman90,Pett01,Wil05}).  In the Bower2006 model an AGN feedback model is implemented which regulates the cooling rate, effectively switching off the supply of cold gas for star formation in quasi-static hot gas haloes.  These are haloes in which the cooling time of the gas exceeds the free-fall time within the halo. The cooling flow is quenched by the energy injected into the hot halo by the central AGN. The growth of the black hole is followed using the model described by \\citet{Rowena07}  \\item {\\it Disk instabilities.}  In the Baugh2005 model the only process that leads to the formation of bulge stars is a galaxy merger. In the Bower2006 model, strongly self-gravitating disks are considered to be unstable to small perturbations, such as encounters with minor satellites or dark matter substructures. Such events can lead to the formation of a bar and eventually the disk is transformed into a bulge. The onset of instability is governed by the ratio \\begin{equation} \\epsilon =  \\frac{V_{\\rm disk}}{(GM_{\\rm disk}/r_{\\rm disk})^{1/2}}. \\label{eqinstab} \\end{equation} Disks for which $\\epsilon < \\epsilon_{\\rm disk}$, are considered to be unstable; in Bower2006, the threshold for unstable disks is set at $\\epsilon_{\\rm disk} = 0.8$. Any cold gas present when the disk becomes unstable is assumed to participate in a starburst. As with starbursts triggered by galaxy mergers, a small fraction of the gas involved in the burst is accreted onto the central black hole. This is an important channel for the growth of low and intermediate mass black holes in the Bower2006 model.  \\item {\\it Treatment of reheated cold gas.}  The fate of gas reheated by supernovae is different in the two models.  In the Baugh2005 model, as discussed above, there are two possible fates for the gas heated by supernovae: ejection from the disk to be reincorporated into the hot halo and ejection from the halo altogether, with no possibility of recooling at a later time. In the case of the first of these channels, the gas is added back into the hot halo when a new halo forms. This happens when the original halo has doubled in mass since its formation time. In the Bower2006 model, this timescale is instead taken to be some multiple of the dynamical time of the halo. Thus, gas can be reheated by supernovae, be added back into hot halo and cool again on a shorter timescale in the Bower2006 model than in the Baugh2005 model.  \\end{itemize}   \\subsection{Galaxy scale-lengths} For completeness, We now review the calculation of the sizes of the disk and bulge components of galaxies used in {\\tt GALFORM}, to complement the discussion of the size predictions presented in Section 4.  For a more detailed description we direct the reader to Cole et~al (2000).  In the following subsections, we outline how the scale-lengths of the disk and bulge are calculated. The scale-lengths are determined by solving for the dynamical equilibrium of the disk, bulge and dark matter together.  Dark matter haloes are assumed to have a \\citet{NFW97} density profile. The hot halo has a modified isothermal density profile with a core.  \\subsubsection{Disks}  The size of a galactic disk is determined by the conservation of the angular momentum of the gas cooling from the halo and the application of centrifugal equilibrium.  The disk is assumed to have an exponential surface density profile with half-mass radius $r_{\\rm disk}$. The half-mass radius of the disk is related to the exponential scale-length, $h_{\\rm D}$ by $r_{\\rm disk} = 1.68 h_{\\rm D}$.  The angular momentum of the dark matter halo is quantified by a dimensionless spin parameter which is drawn from a log-normal distribution, as suggested by measurements from high resolution N-body simulations \\citep{Col96}. This scheme is used in both models; the spin parameter of low-mass haloes can not be measured reliably from the Millennium simulation \\citep{Bett07}. Each newly formed halo in a halo merger tree is assigned a new spin parameter drawn from the distribution at random, independently of the previous value of the spin parameter.  As the halo collapses, the associated gas will shock and so be heated to approximately the virial temperature of the halo. Thereafter it will begin to cool via atomic processes. As gas cools it loses pressure support and flows to the center of the halo, where it is assumed to settle into a rotationally supported disk. The model assumes that the specific angular momentum of the cooling gas depends on the radius from which it originated, as described in \\citet{Col00}. The scale-length of the disk is dependent on the angular momentum of the halo gas which cools.  \\subsubsection{Bulges}  Spheroids result from galaxy mergers, and in the case of the Bower2006 model also from dynamical instabilities of disks. The size of the spheroid produced by these events is calculated by considering virial equilibrium and energy conservation. We assume that the projected density profile is well described by a de Vaucouleurs $r^{1/4}$ law (e.g. Binney \\& Tremaine 1987). The effective radius, $r_{\\rm e}$, of the $r^{1/4}$ law, i.e. the radius that contains half the mass in projection, is related to the half-mass radius in three dimensions, $r_{\\rm bulge}$, by $r_{\\rm bulge}=1.35 r_{\\rm e}$.  When dark matter haloes merge, the galaxy in the most massive halo is assumed to become the central galaxy in the new halo while the other galaxies become satellites. The orbits of the satellites gradually decay as energy and angular momentum are lost via dynamical friction. Eventually a satellite will merge with the central galaxy if the timescale for the orbit to decay is shorter than the halo lifetime.  Two types of merger are distinguished, major mergers and minor mergers, according to the ratio of the mass of the smaller galaxy to the larger galaxy $M_2/M_1$. If the ratio is $M_2/M_1\\geq f_{\\rm ellip}$, then a major merger is assumed to have taken place. Both stellar disks are transformed into spheroid stars and added to any pre-existing bulge. Any cold gas present takes part in a starburst. If the ratio is $M_{\\rm 2}/M_{\\rm 1}\\leq f_{\\rm ellip}$ then a minor merger is assumed and the stars of the smaller galaxy are added to the bulge of the central galaxy, which keeps its stellar disk.  If $f_{\\rm ellip} > M_{\\rm 2}/M_{\\rm 1}\\geq f_{\\rm burst}$ and the central galaxy is gas rich, then the minor merger may also be accompanied by a burst. The parameter $f_{\\rm ellip}=0.3$ in both models; $f_{\\rm burst}=0.05$ in Baugh2005 and $f_{\\rm burst} = 0.1$ in Bower2006. Disks are considered gas rich if 10\\% of the total disk mass in stars and cold gas is in the form of cold gas in Bower2006; this threshold is set higher, 75\\%, in Baugh2005.   In a merger, the two galaxies are assumed to spiral together until their separation equals the sum of their half-mass radii, which is the moment when they are considered to have merged. We estimate the radius of the merger remnant using energy conservation. Assuming virial equilibrium, the total internal energy $E_{\\rm int}$ of each galaxy (both for the merging components and the remnant of the merger) is related to its gravitational self-binding energy $U_{\\rm int}$ by $E_{\\rm int} = -\\frac{1}{2} U_{\\rm int}$, and so can be written as \\begin{equation} E_{\\rm int}=-\\frac{\\bar{c}}{2} \\frac{GM^2}{r}, \\label{Ebind} \\end{equation} where $M$ and $r$ are the mass and half-mass radius respectively and $\\bar{c}$ is a form factor which depends on the distribution of mass in the galaxy. For a de Vaucouleurs profile, $\\bar{c}=0.45$, while for an exponential disk, $\\bar{c}=0.49$. For simplicity, in the model we assume $\\bar{c}=0.5$ for all galaxies.  The orbital energy of a pair of galaxies at the moment of merging is given by \\begin{equation} E_{\\rm orbit}=-\\frac{f_{\\rm orbit}}{2}\\frac{GM_{1}M_{2}}{r_{1}+r_{2}}, \\label{Eorbit} \\end{equation} where $M_{1}$ and $M_{2}$ are the masses of the merging galaxies, $r_{1}$ and $r_{2}$ are their half-mass radii, and $f_{\\rm orbit}$ is a parameter which depends on the orbital parameters of the galaxy pair. A fiducial value of $f_{\\rm orbit}=1$ is adopted which corresponds to two point mass galaxies in a circular orbit with separation $r_{1}+r_{2}$. The galaxy masses $M_{1}$ and $M_{2}$ include the total stellar and cold gas masses and also some part of the dark matter halo. We assume that the mass of dark matter which participates in the galaxy merger in this way is a multiple $f_{\\rm dark}$ of dark halo mass $M_{i,\\rm dark}$ within the galaxy half-mass radius. We adopt a fiducial value $f_{\\rm dark}=2$. We thus have: \\begin{equation} M_{i}=M_{i,\\rm stellar+gas}+f_{\\rm dark} M_{i,\\rm dark}. \\label{DMcontrib} \\end{equation} Later on we will investigate the effect of varying $f_{\\rm orbit}$ and the dark matter mass contribution $f_{\\rm dark}$.  Assuming that each merging galaxy is in virial equilibrium, then their total energy equals one half of their internal energy. The conservation of energy means that: \\begin{equation} E_{\\rm int,new}=E_{\\rm int,1}+E_{\\rm int,2}+ E_{\\rm orbit}, \\label{Econservation} \\end{equation} and replacing $E_{\\rm int}$ with Eq.~\\ref{Ebind} and $E_{\\rm orbit}$ with Eq.~\\ref{Eorbit} leads to \\begin{equation} \\frac{(M_{1}+M_{2})^2}{r_{\\rm new}}=\\frac{M_{1}^2}{r_{1}}+\\frac{M_{2}^2}{r_{2}}+\\frac{f_{\\rm orbit}}{\\bar {c}}\\frac{M_{1}M_{2}}{r_{1}+r_{2}}, \\label{rnew} \\end{equation} where $r_{\\rm new}$ is the half-mass radius of the remnant immediately after the merger. These equations lead to the result that, in a merger of two identical galaxies, the half-mass radius of the new galaxy increases by a factor of $4/3$ which agrees reasonably well with the factor of 1.42 found in simulations of equal mass galaxy mergers by Barnes (1992).  In the case of a bulge produced by an unstable disk (which is only considered in the Bower2006 model), the considerations that lead to the remnant size are similar to those applied to a bulge produced by mergers, leading to: \\begin{eqnarray} \\frac{\\bar{c}_{\\rm B}(M_{\\rm disk}+M_{\\rm bulge})^2}{r_{\\rm new}} & = & \\frac{\\bar {c}_{\\rm B}M_{\\rm bulge}^2}{r_{\\rm bulge}}+\\frac{\\bar {c}_{\\rm D}M_{\\rm disk}^2}{r_{\\rm disk}}\\\\ & + & {f_{\\rm int}}\\frac{M_{1}M_{2}}{r_{1}+r_{2}}, \\nonumber \\label{rnewinstab} \\end{eqnarray} where $M_{\\rm bulge}$, ${r_{\\rm bulge}}$ and $M_{\\rm disk}$, ${r_{\\rm disk}}$ refer to the masses and half-mass radii of the bulge (if any) and disk respectively. As mentioned above, the form factors $\\bar {c}_{\\rm B}=0.49$ and $\\bar{c}_{\\rm D}=0.45$ for a bulge and disk respectively. The last term in Eq.~\\ref{rnewinstab} represents the gravitational interaction energy of the disk and bulge, which can be approximated for a range of $r_{\\rm bulge}/r_{\\rm disk}$ with $f_{\\rm int}=2$.   \\section[]{DERIVED GALAXY PROPERTIES}   \\begin{figure} \\includegraphics[width=8.5cm]{figs/BtvsPetBAU.ps} \\caption{ The ratio of Petrosian flux to total flux for a sample of {\\tt GALFORM} galaxies with the same selection as the SDSS, plotted as a function of the bulge to total luminosity ratio (B/T), measured in the $r-$band taking into account dust extinction. Points are colour-coded according to the ratio of disk to bulge scale-lengths, as indicated by the legend. } \\label{BTvsPF} \\end{figure}   In this section we describe how standard {\\tt GALFORM} outputs, such as disk and bulge total luminosities and half-mass radii are transformed into quantities which are measured for SDSS galaxies. This allows a direct comparison between the model predictions and the observations. We outline the calculation of Petrosian magnitudes (\\S3.1), the concentration index (Morgan 1958; \\S3.2) and the S\\'{e}rsic index \\citep{Ser68}. The latter two quantities are used as proxies for morphological type in analyses of SDSS data. We also illustrate how the Petrosian magnitude, concentration index and S\\'{e}rsic index depend on the bulge-to-total luminosity ratio and on the ratio of the disk and bulge radii.  \\subsection[]{Petrosian magnitude}  As a measure of galaxy flux, the SDSS team uses a modified definition of the Petrosian (1976) magnitude. The Petrosian radius $r_{\\rm Pet}$ is defined as the radius for which the following condition holds:  \\begin{equation} \\frac{\\int_{0.8r_{\\rm Pet}}^{1.25r_{\\rm Pet}}{\\rm d}r2\\pi rI(r)/[\\pi(1.25^2-0.8^2)r_{\\rm Pet}^2]}{\\int_{0}^{r_{\\rm Pet}} {\\rm d}r2\\pi rI(r)/[\\pi r_{\\rm Pet}^2]}=0.2, \\label{Petr} \\end{equation} where I(r) is the surface brightness profile. Defined in this way, $r_{\\rm Pet}$ is the radius where the local surface brightness averaged within a circular annulus centred on the Petrosian radius is 0.2 times the mean surface brightness interior to that radius. The Petrosian flux defined by the SDSS is then obtained within a circular aperture of radius $2r_{\\rm Pet}$. In the SDSS, the aperture used in all five bands is set by the profile of the galaxy in the r-band. I(r) is the azimuthally-averaged surface brightness measured in a series of annuli. In the case of {\\tt GALFORM} model galaxies, we calculate the disk and bulge sizes, and adopt an exponential profile for the disk with $I(r) \\propto \\exp(-1.68(r/r_{\\rm D}))$, where $r_{\\rm D}$ is the disk half-light radius, and a de Vaucouleurs profile for the bulge (assumed to be spherical), with $I(r) \\propto \\exp(-7.67(r/r_{\\rm B})^{1/4})$, where $r_{\\rm B}$ is the bulge half-light radius in projection (see Cole et~al. 2000). The total surface brightness profile for a galaxy is given by the sum of the disk and bulge profiles. The disk and bulge magnitudes include dust extinction. A random inclination angle is assigned to the galactic disk for calculating the dust extinction. The Petrosian flux within a circular aperture of $2r_{\\rm Pet}$ recovers a fraction of the total light of the galaxy which depends on its luminosity profile and hence its morphology. For a pure disk with an exponential profile, the Petrosian flux recovers in excess of 99\\% of the total flux. On the other hand, for a pure bulge with a de Vaucouleurs profile, the percentage of the total light recovered by the Petrosian magnitude is closer to 80\\%. Fig.~\\ref{BTvsPF} shows the ratio of Petrosian flux to total flux for model galaxies as a function the bulge-to-total luminosity ratio in the $r$-band. The limiting cases described above are apparent in the plot, which also shows the fraction of light recovered by the Petrosian definition for composite disk plus bulge systems, and for different ratios of the disk and bulge scale-lengths.    \\subsection{Concentration index} \\begin{figure} \\includegraphics[width=8.5cm]{figs/BtvscBAU.ps} \\caption{ The concentration index, $c$, plotted as a function of the bulge-to-total luminosity (B/T) in the $r-$band for the Baugh2005 model galaxies. The ratio of disk to bulge scale-lengths is indicated by the colour of the symbol as shown by the key. } \\label{BTvsc} \\end{figure}  The concentration index can be straightforwardly derived once the Petrosian flux and radius have been calculated. The concentration index is defined as $c = R_{90}/R_{50}$, where $R_{90}$ and $R_{50}$ correspond to the radii enclosing $90\\%$ and $50\\%$ of the Petrosian flux respectively in the $r$-band. Hence, the luminosity is dominated by the bulge for high concentration galaxies and is dominated by the disk for low concentration galaxies. In Fig.~\\ref{BTvsc} we plot the bulge-to-total luminosity versus the concentration index for model galaxies in the r-band.  We can see that pure disks have $c = 2.3$, pure bulges have $c = 3.3$, and intermediate values of concentration index correspond to galaxies with different combinations of $r_{\\rm disk}/r_{\\rm bulge}$ and B/T. Most galaxies lie in a narrow locus of B/T versus c, but different combinations of disk and bulge scale-lengths and B/T produce the scatter shown.  Observationally, the Petrosian concentration index is affected by seeing \\citep{Bla03}. The same galaxy can show different concentrations under different seeing conditions.   \\subsection{S\\'{e}rsic index}  \\begin{figure} \\includegraphics[width=8.7cm]{figs/BtvsnBAU.ps} \\caption{ The S\\'{e}rsic index $n$ plotted against the bulge-to-total luminosity ratio (B/T) for the Baugh2005 model galaxies. The ratio of disk to bulge scale-lengths is indicated by the colour of the symbol, as shown by the key. } \\label{B05BTvsN} \\end{figure}  The S\\'{e}rsic index describes the shape of a fit made to the surface brightness profile of a galaxy without prior knowledge of the scale-lengths of the disk and bulge components. The radial dependence of the profile is given by (S\\'{e}rsic 1968):  \\begin{equation} I(r)=I_{0}\\exp[-(r/r_{0})^{1/n}]. \\label{Sers} \\end{equation} Here $I_{0}$ is the central surface brightness, $r_{0}$ is the S\\'{e}rsic scale radius and $n$ is the S\\'{e}rsic index. The S\\'{e}rsic index has been shown to correlate with morphological type (e.g. Trujillo, Graham \\& Caon 2001). We can see that if $n=1$ we recover an exponential profile (used for pure disk galaxies) and if $n=4$ we recover a de Vaucouleurs profile, used to describe pure bulge galaxies. The S\\'{e}rsic index has been calculated in the New York University Value Added Catalogue (NYU-VAGC; \\citealt{Bl05}).  Here, we do not attempt to reproduce exactly the procedure that Blanton et~al. used to obtain the parameters in Eq.~\\ref{Sers} (which takes into account seeing and pixelization).  Since we know the full surface brightness profile of the model galaxies out to any radius, we want the S\\'{e}rsic profile that best reproduces the composite disk plus bulge profile. In order to determine the parameters of the S\\'{e}rsic profile, $I_{0}$, $r_{0}$ and $n$, we minimize a $\\chi^2$ function which depends on the difference between the S\\'{e}rsic profile and the sum of the disk and bulge surface brightness profiles, given the scale-lengths and luminosity ratio of these components:  \\begin{equation} \\chi^{2}=\\sum_{i}[\\log I_{{\\rm disk}+{\\rm bulge}}(r_{i})-\\log I_{{\\rm S\\acute{e}rsic}}(r_{i},r_{0},n)]^{2}W_{i}. \\label{Chi2} \\end{equation} The total luminosity of the fitted S\\'{e}rsic profile is constrained to be equal to that in the true disk + bulge profile. Here $r_{i}$ is a series of rings between $r=0$ to $r=R_{90,D+B}$, the radius enclosing the 90\\% of the disk plus bulge profile flux, and the weight $W_{i}$ is given by the luminosity in the ring containing $r_{i}$. Since the steepness of the S\\'{e}rsic profile is more evident at the centre of the galaxy, we assign half of the bins to the region within the bulge size $r_{\\rm bulge}$ (so as long as $r_{\\rm bulge} < R_{90,D+B}$). As a test to check the consistency of changing from a disk+bulge profile to the S\\'{e}rsic profile, we have compared $R_{50}$ (the radius enclosing 50\\% of the total luminosity) obtained from the two descriptions of the surface density profile and find very similar results.  If it is assumed that the distribution of light in real galaxies is accurately described by the S\\'{e}rsic profile, then quantities derived from fitting a S\\'{e}rsic profile have some advantages over the corresponding Petrosian quantities. (i) The total flux integrated over the fitted S\\'{e}rsic profile would equal the true total flux, unlike the Petrosian flux, which underestimates the true total value, especially for bulge-dominated galaxies, as shown in Fig.~\\ref{BTvsPF}. (ii) The effects of seeing can be included in the S\\'{e}rsic profile fitting, so that quantities obtained from the fit (total flux, scale size $r_0$ and S\\'{e}rsic index $n$) are in principle corrected for seeing effects, unlike the corresponding Petrosian quantities. However, since the Petrosian quantities are the standard ones used by the SDSS community, in the rest of this paper we work with Petrosian magnitude and radius, unless otherwise stated.  Fig.~\\ref{B05BTvsN} shows the correlation between S\\'{e}rsic index and the bulge-to-total luminosity ratio in the $r$-band. There is a considerable scatter between these two proxies or indicators of morphology, driven by the ratio of the scale-lengths of the disk and bulge components. For example, galaxies with a S\\'{e}rsic index of $n=4$, usually interpreted as a pure bulge light profile, can have essentially any value of bulge-to-luminosity ratio from $B/T=0.1$--$1$. A key feature of this plot is the distribution of disk to bulge size ratios generated by {\\tt GALFORM}. It is possible to populate other parts of the $n$ - $B/T$ plane in the cases of extremely large or small values of the ratio of disk to bulge radii. Without the guidance of a physical model, if a grid of $r_{\\rm d}/r_{\\rm b}$ was used instead, the distribution of points would be even broader than shown in Fig.~\\ref{B05BTvsN}. Note that only model galaxies brighter than $M_{r}-5 \\log h = -16$ are included in this plot.  A similar scatter is seen between these two quantities for real galaxies, as shown by Fig.~\\ref{ALLBTvsN} in the Appendix.      \\section[]{RESULTS}  The primary observational data set we compare the model predictions against is the New York University Value Added Catalogue (NYU-VAGC), which gives additional properties to those found in the SDSS database for a subset of DR4 galaxies \\citep{Bl05}. The NYU catalogue covers an area of 4783 deg$^{2}$ and contains 49968 galaxies with redshifts. The sample is complete to $r_{\\rm Pet} = 17.77$ over the redshift interval $0.0033 < z < 0.05$, and has a median redshift of $z=0.036$. The relatively low median redshift compared with the full spectroscopic sample is designed to provide a sample of large galaxy images, suitable for measurements of galaxy morphology.  Examples of the extra properties listed for galaxies, beyond the information available in the SDSS database, include the rest-frame (AB) absolute magnitude, the S\\'{e}rsic index and the value of $V_{\\rm max}$ (i.e. the maximum volume within which a galaxy could have been observed whilst satisfying the sample selection; this quantity is used to weight each galaxy in statistical analyses).  We run {\\tt GALFORM} with an output redshift of $z=0.036$ to match the median of the NYU-VAGC and derive properties from the output which can be compared directly against the observations, as described in \\S 3.  In \\S 4.1, we compare the model and observed luminosity functions, in \\S 4.2 we show the distribution of morphological types versus luminosity, in \\S 4.3 we examine the colour distributions and explore this further as a function of morphology in \\S 4.4. Finally in \\S 4.5 we test the size predictions against observations and assess the sensitivity of the model output to the strength of various processes.   \\subsection{Luminosity function: all galaxies and by colour}  \\begin{figure} \\includegraphics[width=8.5cm]{figs/Lumfunall.ps} \\caption{ The $r-$band luminosity function predicted by the Bower2006 (red) and Baugh2005 (blue) models. For comparison, we also plot the SDSS luminosity function estimated using the SWML estimator by Blanton et al (2005) from DR2 (solid histogram) and our result using the $1/V_{\\rm max}$ estimator from DR4 (dotted histogram). } \\label{LumFun} \\end{figure}  \\begin{figure*} \\includegraphics[width=16cm]{figs/Lumfun3colours.ps} \\caption{ The colour-dependent luminosity function, for three populations of galaxies defined by the $g-r$ colour as shown by the label on top of each panel. The black lines show the luminosity function estimated from DR4 and the red (Bower2006) and blue (Baugh2005) lines show the model predictions (see the legend in the left panel). } \\label{LumFunbyCol} \\end{figure*}   The local luminosity function plays a key role in constraining the parameters which specify a galaxy formation model. The comparison between the predicted and observed luminosity functions is hence a fundamental test of any model. The original papers describing the Baugh2005 and Bower2006 models showed the comparison of the model predictions with the observed local luminosity function in the optical and near-infrared. However a comparison with SDSS data was not made in those papers. Fig.~\\ref{LumFun} shows the luminosity function in the Petrosian $r$-band predicted by the Baugh2005 and Bower2006 models, compared with our estimate of the luminosity function from SDSS DR4 made using the values of $V_{\\rm max}$ from the NYU-VAGC catalogue. We also overplot the luminosity function estimated from the SDSS DR2 by \\citet{Bl05b} using the stepwise maximum likelihood (SWML) method. The SWML and $1/V_{\\rm max}$ estimates are in very good agreement, particularly for magnitudes brighter than $M_{r}-5\\log h=-17$. At fainter magnitudes, the $1/V_{\\rm max}$ estimator could be affected by very local large scale structure \\citep{Bl05b}.  Both models overpredict the abundance of bright galaxies.  The Bower2006 model produces a somewhat better match to the shape of the SDSS luminosity function.  This offset in the $r-$band luminosity function has also been noted by Cai et~al. (2008), who made the model galaxies in the Bower2006 model fainter by 0.15 magnitudes before using this model to make mock galaxy surveys.  It is worth noting that the Bower2006 model gives an excellent match to both the $b_{\\rm J}$-band luminosity function estimated from the 2dF galaxy redshift survey (\\citealt{Nor02}) and to the K-band luminosity function (e.g. Cole et~al. 2001; Kochanek et~al. 2001) without the need to shift the model magnitudes by hand.  We can study the impact on the luminosity function of different physical processes in more detail by using colour to separate galaxies into different samples. For the SDSS data we can compute the luminosity function of colour sub-samples using the $1/V_{\\rm max}$ estimator, bearing in mind that fainter than $M_{r}-5\\log h=-17$ this method gives an unreliable estimate of the luminosity function due to local large-scale structure. We use the Petrosian $g-r$ colour to split galaxies into blue ($g-r < 0.45$), red ($g-r > 0.65$) and intermediate ($0.45 < g-r < 0.65$) colour samples. In Fig.~\\ref{LumFunbyCol}, we can see that both models reproduce the intermediate colour population fairly well (middle panel). The Bower2006 model in particular matches the shape of the observed luminosity function closely, albeit with a shift to brighter magnitudes, similar to that seen in the case of the overall luminosity function in Fig.~\\ref{LumFun}. The models fare worst for blue galaxies, with both models overpredicting the number of bright blue galaxies.  This suggests that star formation is not quenched effectively enough in massive haloes or that the timescale for gas consumption in star formation is too long. For the case of red galaxies, the models do best brightwards of $L_{*}$, but get the number of faint red galaxies wrong, with the Baugh2005 model giving too many faint red galaxies and Bower2006 too few.   \\subsection{The distribution of morphological types}  \\begin{figure*} \\includegraphics[width=16cm]{figs/Morf2modelsNlow5.5.ps} \\caption{ The fraction of different morphological types as a function of magnitude $M_{r}$ for SDSS data (squares) and {\\tt GALFORM} (solid lines). The fraction of disk-dominated galaxies (as defined by a value of the S\\'{e}rsic index $n<2.5$) is shown in blue and bulge-dominated galaxies (i.e. those with $n>2.5$) are plotted in red. The left panel shows the Baugh2005 model and the right panel the Bower2006 model. } \\label{Morfn}  \\includegraphics[width=16cm]{figs/Morf2modelsCE.ps} \\caption{ The fraction of different morphological types as a function of magnitude $M_{r}$ for SDSS data (squares) and {\\tt GALFORM} (solid lines), using the Petrosian concentration index $c$ to define type. The fraction of disk-dominated galaxies (as defined by $c<2.86$) is shown in blue and bulge-dominated galaxies (i.e. those with $c>2.86$) are plotted in red. The left panel shows the Baugh2005 model and the right panel the Bower2006 model.  } \\label{Morfc}  \\includegraphics[width=16cm]{figs/Morf2modelsBT.ps} \\caption{ The fraction of different morphological types as a function of magnitude $M_{r}$ using the bulge-to-total luminosity ratio in the $r-$band to define type. Disk-dominated galaxies ($B/T<0.5$) are shown in blue and bulge-dominated ($B/T>0.5$) are plotted in red. The solid curves show the model predictions, according to the label above each panel. The shaded region shows an observational estimate made from SDSS data by Benson et~al. (2007). The extent of the shading shows by how much the fraction changes when a correction is applied to the observational estimates (see Benson et~al. for details). } \\label{MorfBT} \\end{figure*}  In this section we examine the mix of morphological types as a function of luminosity, using different proxies for galaxy morphology.  We first look at the mix of galaxies using the S\\'{e}rsic index.  Using a S\\'{e}rsic index value of $n=2.5$ (which is half-way between $n=1$ and $n=4$), we separate galaxies into two broad morphological classes, disk-dominated galaxies (late type galaxies) with $n<2.5$ and bulge-dominated galaxies (early type galaxies) with $n>2.5$. Fig.~\\ref{Morfn} shows the fraction of galaxies in each morphological type, as a function of Petrosian magnitude, $M_{r}$, for SDSS galaxies and the {\\tt GALFORM} model (Baugh2005 in the left panel and Bower2006 in the right panel). The trend found for SDSS galaxies is that the disk-dominated population is the more common at faint magnitudes, whereas bulge-dominated objects are in the majority brighter than $L_{*}$. Fig.~\\ref{Morfn} shows that both models follow the same general trend, but with the changeover from one population to the other occuring brighter than $L_{*}$ in the Baugh2005 model, whereas the Bower2006 model looks more similar to the observations.  Another way to morphologically classify galaxies using the profile shape is to use the Petrosian concentration index $c$. In Fig.~\\ref{Morfc}, we show the fraction of early and late type galaxies as a function of luminosity based on this, where we classify galaxies with $c<2.86$ as late type and $c>2.86$ as early type. The resulting plots look very similar to those based on S\\'{e}rsic index in Fig.~\\ref{Morfn}, though there are differences in detail, particularly at the bright end. The agreement between the models and the SDSS data is generally better using the concentration as the classifier, particularly for the Bower2006 model.  As a third approach to determining galaxy type, we consider the bulge to total luminosity ratio, $B/T$ measured in the $r-$band. \\citet{Ben07} fitted disk and bulge components to images of 8839 bright galaxies selected from the SDSS EDR. In fitting the disk and bulge components of each galaxy, they used the bulge ellipticity and disk inclination angle, $i$, as free parameters. The resulting distribution of $\\cos(i)$ showed an excess of face-on galaxies. This is due in part to the algorithm mistaking part of the bulge as a disk. Benson et~al. attempted to correct for the uneven distribution of inclination angles in the following way. Galaxies with $\\cos(i)<0.5$ are assumed to have been correctly fitted. Since a uniform distribution in $\\cos(i)$ is expected, for each galaxy with $\\cos(i)<0.5$, a galaxy with a similar bulge and face-on, projected disk magnitude but with $\\cos(i)>0.5$ is also selected.  The galaxies with $\\cos(i)>0.5$ which are left without a match are assumed to correspond to cases where the disk component has been used to fit some feature in the bulge. Benson et~al. assigned to these galaxies a value of B/T=1. The correction has a considerable impact on the fraction of bulge- and disk-dominated galaxies, as shown by the extent of the shaded region in Fig.~\\ref{MorfBT}.  The observational estimates of the mix of morphological types presented in Figs.~\\ref{Morfn}, \\ref{Morfc} and \\ref{MorfBT} are qualitatively the same, but show that the transition from disk-dominated to bulge-dominated depends on the choice of property used to define morphology. We note that the model predictions are very similar when we set the division at $c=2.86$ or at $B/T=0.5$. The model predictions made using the S\\'{e}rsic index, concentration and bulge-to-total luminosity ratio appear to be closer to each other than the corresponding observational measurements. This comparison gives some indication of the observational uncertainty in measuring fractions of different morphological types using the S\\'{e}rsic index, concentration and bulge-to-total luminosity ratio.   \\subsection{Colour distribution}  \\begin{figure*} \\includegraphics[width=14.7cm]{figs/2COLOURSPD.ps} \\caption{ The $(u-r)$ colour distribution as function of luminosity for the Baugh2005 model (blue histograms), the Bower2006 model (red histograms) and the SDSS data (yellow histograms). The centre of the magnitude bin used in each panel is given by the legend. All the histograms are normalized to have unit area. } \\label{Colorshort} \\end{figure*}  \\begin{figure*} \\includegraphics[width=12.5cm]{figs/ColorsvsM.W.r.WHITEb.ps} \\caption{ The galaxy distribution in the colour magnitude plane. Each column shows the distribution for a different colour. The top row shows the SDSS distributions, the middle row the predictions of the Baugh2005 model and the bottom row the Bower2006 model. Galaxies are weighted by $1/V_{\\rm max}$. The inner contour encloses 68\\% of the total number density of galaxies and the outer contour encloses 95\\% of the density. The colour shading reflects the square root of the number density. } \\label{ColorvsmW} \\includegraphics[width=12.5cm]{figs/ColorsvsM.LW.r.WHITEb.ps} \\caption{ The same as Fig.~\\ref{ColorvsmW}, but with each galaxy weighted by the product of its luminosity and $1/V_{\\rm max}$. } \\label{ColorvsmLW} \\end{figure*}  An important feature uncovered in the SDSS data is a bimodality in the galaxy colour-magnitude relation (e.g. \\citealt{Str01}; \\citealt{Bal04}).  In Fig.~\\ref{Colorshort} we plot the distribution of Petrosian $(u-r)$ colour in selected bins of magnitude $M_{r}$ for models and SDSS data. The SDSS data shows a dominant red population at bright magnitudes, with a blue population that becomes more important at fainter magnitudes.  Although we see blue and red populations in the Baugh2005 model predictions for intermediate magnitudes, the red population always dominates, even at the faintest magnitudes. The Bower2006 model, on the other hand displays a clear bimodality, with the red population dominating at bright magnitudes, comparable red and blue populations at intermediate magnitudes and a slightly more dominant blue population at faint magnitudes. The Bower2006 model shows the same behaviour as the SDSS data at bright magnitudes. At faint magnitudes, the Bower2006 model still shows a red population which is not apparent in the data. \\citet{Font08} argued that these faint red galaxies are predominantly satellite galaxies, which in the Bower2006 model have exhausted their cold gas reservoirs. In the \\citeauthor{Font08} model, which is a modified version of the Bower2006 model, the stripping of hot gas from satellites is incomplete, and so gas may still cool onto the satellite, fuelling further star formation and causing these galaxies to have bluer colours on average.  Since the SDSS photometry covers five bands ($u,g,r,i$ and $z$), we can investigate the bimodality further using different colours.  In Fig.~\\ref{ColorvsmW} we plot the abundance of galaxies in the colour magnitude plane, for $(u-g),(g-r),(r-i)$ and $(i-z)$ Petrosian colours against magnitude $M_{r}$. The top row of panels shows the distributions for the NYU-VAGC SDSS data, the middle row shows the predictions of the Baugh2005 model, and the bottom row gives those of the Bower2006 model. In this plot, each galaxy contributes $1/V_{\\rm max}$ to the density.  The contours in the plot indicate the regions containing 68\\% and 95\\% of the number density of galaxies in the samples. Note that the colour shading scales as the square root of the density.  For the case of the SDSS data in all colours except for $(i-z)$ we can see a bright red population and a fainter, bluer population as indicated by the splitting of the 68\\% density contour. The Baugh2005 model predicts a dominant red population at the brightest magnitudes. On moving to fainter magnitudes, bluer galaxies appear but red galaxies still dominate and there is no clear bimodality. The Bower2006 model displays a strong bimodality in colour, though with a bluewards shift in the locus of the colour-magnitude relation compared with the observations. \\citet{Font08} obtained better agreement of their model with the observed locus of red galaxies in the SDSS by increasing the assumed yield of metals by a factor of two relative to the Bower2006 model.  In Fig.~\\ref{ColorvsmLW} we plot a similar colour-magnitude distribution, but this time each galaxy contributes $L/V_{\\rm max}$ to the density. By doing this more emphasis is given to brighter galaxies. As a consequence, the bimodality in the SDSS data is less readily apparent. Intriguingly, the Baugh2005 model appears visually to be in better agreement with the observations when presented in this way.   \\subsection{Colour distribution by morphology}   \\begin{figure*} \\includegraphics[width=12.7cm]{figs/4MEDcolors.ps} \\caption{ The median colour $(u-g), (g-r), (r-i), (i-z)$ for the models (continuous lines) and SDSS galaxies (triangles) as a function of magnitude $M_{r}$. Different coloured lines and symbols correspond to different morphological types of galaxies, as given by the S\\'{e}rsic index ($n<2.5$ for the disk-dominated population and $n > 2.5$ for the bulge-dominated population). Each panel corresponds to a different colour. The panels each have the same range of colour on the y-axis. The results are plotted only when there are ten or more galaxies present in a bin. } \\label{Colorvsn} \\end{figure*}   The bimodality of the colour-magnitude relation seen for SDSS galaxies suggests that different populations or types of galaxy dominate at different magnitudes. We also saw in Section 4.2 that disk-dominated galaxies are more abundant at faint magnitudes and the bulge-dominated population is more prevalent at bright magnitudes. A correlation is therefore expected between morphology, colour and luminosity. To see this effect more clearly, we use the S\\'{e}rsic index, $n$, to separate galaxies into an ``early-type'' bulge-dominated population (with $n>2.5$) and a ``late-type'' disk-dominated population ($n<2.5$) and replot the colour-magnitude relation.  We calculate the median Petrosian colour for $(u-g),(g-r),(r-i),(i-z)$ in bins of magnitude $M_{r}$ for the two populations. The results are plotted in Fig.~\\ref{Colorvsn} for the models and the SDSS data.  We can see from the SDSS data that the different populations display different colour-magnitude correlations, confirming that the S\\'{e}rsic index is an effective morphological classifier. The bulge-dominated galaxies are redder than the disk-dominated galaxies, with the size of the difference decreasing as the effective wavelength of the passbands increases. Also, at fainter magnitudes, the colours of the two population tend to become more similar.  In the Baugh2005 model, Fig.~\\ref{Colorvsn} shows that the populations split by S\\'{e}rsic index have similar colours except for the brightest galaxies. Both populations are predicted to be too red at faint magnitudes. At brighter magnitudes ($M_{r}-5\\log h<-19$), bulge-dominated galaxies show similar behaviour to the SDSS data. Disk-dominated galaxies become bluer at the brightest magnitudes, which is opposite to the trend seen in the data.  The Bower2006 model predicts a clear separation in colour for populations classified by S\\'{e}rsic index, with blue disk-dominated galaxies even at faint magnitudes, which is in better agreement with SDSS data.  In both models, the faint bulge-dominated population is predicted to be too red.   \\subsubsection{What drives the colours? A look at the specific star formation rate and metallicity}  \\begin{figure} \\includegraphics[width=8.5cm]{figs/MSTARMET.ps} \\caption{ The top panel shows the specific star formation rate, i.e. the star formation rate per unit stellar mass, as a function of magnitude $M_{r}$ for the {\\tt GALFORM} models, as indicated by the key in the lower panel. The lower panel shows the stellar metallicity, weighted by the V-band luminosity as a function of magnitude $M_{r}$. Different colours correspond to different morphological types as set by the S\\'{e}rsic index, as shown by the label in the top panel. } \\label{sfrMetal} \\end{figure}   To identify which feature of the models is producing the differences in colour seen in Fig.~\\ref{Colorvsn}, we now examine the specific star formation rate SSFR (the star formation rate [SFR] per unit stellar mass) and stellar metallicity of galaxies in both models. The specific star formation rate quantifies how vigorously a galaxy is forming stars in terms of how big a contribution recent star formation makes to the total stellar mass. Galaxies with a high specific star formation rate will tend to have bluer colours and stronger emission lines than more ``passive'' galaxies. We use the S\\'{e}rsic index to separate the galaxies as before, into an ``early-type'' bulge-dominated population (with $n>2.5$) and a ``late-type'' disk-dominated population ($n<2.5$). In the top panel of Fig.~\\ref{sfrMetal} we plot the median of the SSFR as a function of magnitude $M_{r}$. In the bottom panel of this figure we plot the median of the V-band luminosity-weighted stellar metallicity.  The top panel of Fig.~\\ref{sfrMetal} shows that bulge-dominated galaxies have very low specific star formation rates in both models.  The disk-dominated galaxies have very different specific star formation rates in the two models. In the Bower2006 model, the disk-dominated galaxies are undergoing significant amounts of star formation, except at the brightest magnitudes. Although the strength of supernova feedback is stronger in the Bower2006 model than it is in the Baugh2005 model, reheated gas tends to recool on a shorter timescale because it is reincorporated into the hot halo faster. The drop in the specific star formation rate for the brightest galaxies in the Bower2006 model can be traced to the AGN feedback which shuts down gas cooling for these galaxies. Note that disk-dominated galaxies make up only a small fraction of the galaxies at these magnitudes.  Within a given model, the metallicities of the disk and bulge-dominated populations are similar. However, the metallicities in the Baugh2005 model are higher than in Bower2006, presumably because some fraction of the star formation in the former model occurs in starbursts with a top-heavy IMF, which correspondingly produces a higher yield of metals.  Hence, given this differences, one expects bluer galaxies at faint magnitudes in the Bower2006 model than in the Baugh2005 model.   \\subsubsection{Correlation between S\\'{e}rsic index, colour and magnitude}  \\begin{figure} \\includegraphics[width=8.5cm]{figs/NvscolmagLW.COLOR.ps} \\caption{ The luminosity-weighted density of galaxies in different projections of the S\\'{e}rsic index ($n$), (g-r) colour and magnitude plane. Contours indicate the regions containing $68\\%$ and $95\\%$ of the total density of galaxies. Top panels: SDSS data, intermediate panels: Baugh2005 model and bottom panels: Bower2006 model. } \\label{nvsM} \\end{figure}  To investigate further the correlation between the S\\'{e}rsic index, colour and absolute magnitude, we plot in Fig.~\\ref{nvsM} the luminosity weighted density in the various projections of the S\\'{e}rsic index ($n$), $(g-r)$ colour and $M_{i}$ absolute magnitude plane, both for SDSS data and {\\tt GALFORM} models. In the data we can see that disk-dominated galaxies (i.e. those with small $n$ values) tend to be bluer and also fainter, whereas the bulge-dominated galaxies (those with large $n$ values) tend to be redder and brighter. The predictions of both {\\tt GALFORM} models are peaked around S\\'{e}rsic indices of $n=1$ (nominally pure disc galaxies) and $n=4$ (pure bulge galaxies). Despite these density peaks, the numbers of galaxies in the different morphological classes are similar to the SDSS data (as shown in Fig.~\\ref{Morfn}) showing that at a broad-brush level, the distribution of $n$ predicted by the {\\tt GALFORM} models is in reasonable agreement with the observations.  Bearing in mind the level at which the models are able to match the distribution of S\\'{e}rsic indices, both models reproduce fairly well the behaviour seen in the SDSS observation, with a spike corresponding to a faint blue disk-dominated population, which changes to a red, bulge-dominated population at bright magnitudes. Compared with the SDSS, the Baugh2005 model overpredicts the number of red disk-dominated galaxies (around values $(g-r) \\sim 1$ and $n \\sim 1$) and the number of moderate luminosity bulge-dominated galaxies (around values $M_{i}-5\\log h \\sim -19$ and $n \\sim 4$). The Bower06 model predicts a distribution which agrees better with the observational data.    \\subsection{The distribution of disk and bulge sizes}  We now examine the model predictions for the linear size of the disk and bulge components of galaxies. We compare the model predictions with SDSS observations using the radius enclosing $50\\%$ of the Petrosian flux, $R_{50}$. The calculation of disk and bulge sizes was reviewed in Section 2.3 (see also Cole et al. 2000 and Almeida et al. 2007). We use the concentration index, $c$, to divide galaxies into two broad classes of disk-dominated and bulge-dominated samples. First we discuss the accuracy of the predictions for $R_{50}$ for disk-dominated galaxies (\\S~4.1) and then for bulge-dominated galaxies (\\S~4.2), before illustrating the sensitivity of the results to various physical ingredients of the models. The observations we compare against are our own analysis of the size distribution in the NYU-VAC constructed from DR4, as discussed below, and the results from Shen et~al.  (2003; hereafter Sh03), which were derived from a sample of $~140 000$ galaxies from DR1.  \\subsubsection{Disk-dominated galaxies}  \\begin{figure*} \\includegraphics[width=18cm]{figs/4SIZESPD.ps} \\caption{ The distribution of the Petrosian half-light radius, $R_{50}$, for early type galaxies (with concentration parameter $c>2.86$) in top six panels and for late type galaxies ($c<2.86$) in the lower six panels. Each panel corresponds to a different one magnitude wide bin, as indicated by the legend.  The {\\tt GALFORM} predictions are shown by unshaded histograms (Baugh2005 - blue; Bower2006 - red) and the SDSS data by the yellow shaded histogram. All of the histograms are normalized to have unit area.} \\label{AllSizesShort} \\end{figure*}   \\begin{figure*} \\includegraphics[width=17cm]{figs/TalkMEDvhotpaper.ps} \\caption{ A compilation of predictions for the Baugh2005 (left column) and Bower2006 (right column) models. The top row shows the median $R_{50}$ as a function of magnitude for bulge-dominated galaxies ($c>2.86$), the second row shows the median $R_{50}$ for disk-dominated galaxies ($c < 2.86$), the third row shows the $r-$band luminosity function of all galaxies and the bottom row shows the fraction of early-type galaxies as a function of magnitude. The predictions of the fiducial model in both cases are shown by the green lines. In this plot, we also show the impact of changing the strength of supernova feedback, rerunning the model with either half the fiducial value of $V_{\\rm hot}$ (blue curves; see Section 2.1) or twice the value (red curves).  In first and second rows, the open symbols show our determination of the median size from the NYU-VAGC; the filled symbols show the results obtained by Sh03. The black histogram in the third row shows our determination of the luminosity function in DR4 using the $1/V_{\\rm max}$ estimator. The squares in the bottom row show the fraction of early-types in the NYU-VAGC, defined according to concentration parameter $c > 2.86$, as a function of magnitude.  } \\label{Shenmedvhot} \\end{figure*}   \\begin{figure*} \\includegraphics[width=17cm]{figs/TalkMEDediscpaper.ps} \\caption{ Similar to Fig.~\\ref{Shenmedvhot}, but varying the disk instability threshold $\\epsilon_{\\rm disk}$ in Eq.~\\ref{eqinstab}. } \\label{Shenmedinstab} \\end{figure*}   \\begin{figure*} \\includegraphics[width=17cm]{figs/TalkMEDfopaper.ps} \\caption{ Similar to Fig.~\\ref{Shenmedvhot}, but varying the $f_{\\rm orbit}$ parameter in Eq.~\\ref{rnew} } \\label{Shenmedforbit} \\end{figure*}    \\begin{figure*} \\includegraphics[width=17cm]{figs/TalkMEDfdpaper.ps} \\caption{ Similar to Fig.~\\ref{Shenmedvhot}, but varying the $f_{\\rm dark}$ parameter in Eq.~\\ref{DMcontrib} } \\label{ShenmedDMc} \\end{figure*}  Following Sh03, we take $c<2.86$ to define the disk-dominated sample (recall that pure disk galaxies have $c \\sim 2.3$ and pure bulges have $c \\sim 3.3$, as shown by Fig.~\\ref{BTvsc}). Besides the selection introduced by use of the SDSS spectroscopic sample ($r_{\\rm Pet}<17.77$), further selection criteria are required in the size distribution analysis.  The size measurement for compact galaxies could be affected by the point-spread function of the image or because these objects could be misclassified as stars by the SDSS imaging processing software.  To minimize such contamination, Sh03 selected galaxies with angular sizes $R_{50} > 1.6''$ (this excludes only a few percent of the galaxies). Sh03 further restricted the sample to galaxies with surface brightness $\\mu_{50}\\leq 23.0$ mag arcsec$^{-2}$, and apparent magnitude in the range $r_{\\rm min} (\\theta,\\phi) \\leq r \\leq r_{\\rm max} (\\theta,\\phi)$ with, typically, $r_{\\rm min} \\sim 15.0$ and $r_{\\rm max} \\sim 17.77$, and redshift $z \\geq 0.005$.  We apply the same criteria as used by Sh03 to the low redshift NYU-VAGC catalogue. This requires us to recalculate the value of $V_{\\rm max}$ needed to construct statistical distributions, to take into account the bright magnitude limit, size cut and surface brightness cut.  Note that although the Sh03 sample is from the smaller DR1, it contains more galaxies than the NYU-VAGC sample used here because it extends to higher redshift.  The correlation of size with luminosity for disk-dominated galaxies is shown in the upper six panels of Fig.~\\ref{AllSizesShort}, in which we plot the distribution of $R_{50}$ in selected magnitude bins in the $r-$band.  The {\\tt GALFORM} predictions are plotted as unshaded histograms, with the Baugh2005 results in blue and the Bower2006 results in red. The observed distributions are shown by the yellow filled histograms. Except for the brightest two magnitude bins shown, the models tend to overpredict the size of disk-dominated galaxies, particularly in the case of the Bower2006 model. In the Baugh2005 model, the peak of the distribution shifts to larger sizes with brightening magnitude, reproducing the trend seen in the observations. On the other hand in the Bower2006 model, there is little dependence of disk size on luminosity. Both models display a larger scatter in sizes than is seen in the data. The panels for early-types are discussed in the next section.  To further quantify the trend of size with luminosity, we calculate the median value of $R_{50}$ and plot the results in Fig.~\\ref{Shenmedvhot}, where the continuous green line represents the fiducial {\\tt GALFORM} model (the left panel shows the results for the Baugh2005 model and the right panel for the Bower2006 model) and open symbols represent the NYU-VAGC data, where we overplot for comparison (and to check for consistency) the results from Sh03 (filled triangles).  Fig.~\\ref{Shenmedvhot} shows that our analysis of the size distribution in the NYU-VAGC is consistent with the results of Sh03. The apparent magnitude cut $r_{\\rm min}$ together with using a low-redshift sample in comparison with Sh03, removes the galaxies brighter than $M_{r}-5\\log h = -21.5$.  The SDSS data show an increase of over one decade in $R_{50}$ across the luminosity range plotted for disk-dominated galaxies. This increase is reproduced by the predictions of the Baugh2005 model. The behaviour of the Bower2006 model is quite different, with an essentially flat size-luminosity relation to $~L_{*}$, followed by a decrease in size for brighter galaxies. We investigate the impact of various processes on the form of the size predictions in Section 4.5.3.  \\subsubsection{Bulge-dominated galaxies}  We select a bulge-dominated sample by taking those galaxies with concentration index $c>2.86$. In the lower six panels of Fig.~\\ref{AllSizesShort}, we plot the distribution of sizes $R_{50}$ for bulge-dominated galaxies for a selection of magnitude bins. In general, the model predicts values of $R_{50}$ larger than observed, except for the $M_{r}-5 \\log h = -20.5$ bin for Bower2006 and  $M_{r}-5 \\log h = -21.25$ for Baugh2005. As for the case of disk-dominated galaxies, the predicted scatter in sizes is larger than observed.  We plot the median size of the bulge-dominated samples in the top row row of Fig.~\\ref{Shenmedvhot}. The predicted size-luminosity relation is flatter than observed, turning over at the brightest magnitudes plotted.  The brightest galaxies are three to five times than smaller than observed, confirming the conclusion reached by \\citep{Al07}.  As the DR4 data set we are working with covers a larger solid angle than the sample used by Sh03, the combined set of data measurements covers a wdier range of magnitudes than can be reached by either sample alone. Again where there is overlap, we find that our analysis of DR4 is consistent with the results of Sh03.  \\subsubsection{Sensitivity of the predictions to physical ingredients}  The calculation of sizes involves several components as outlined in \\S2.2. Given that this is the area in which, overall, the model predictions agree least well with the observations, it is instructive to vary some of the physical ingredients of the model to see if the agreement can be improved. In the tests which follow, we vary the strength of one ingredient at a time and assess the impact on the size-luminosity relation. We also show the effect of the parameter change on the form of the overall galaxy luminosity function and the mix of morphological types. These variant models are not intended to be viable or alternative models of galaxy formation, but instead allow us to gain some physical insight into how the model works.   \\begin{itemize} \\item[(i)]{\\it The strength of supernova feedback.}  SN feedback plays an important role in setting the sizes of disk galaxies, by influencing in which haloes gas can remain in the cold phase to form stars. Cole et~al. (2000) demonstrated that increasing the strength of SN feedback results in more gas cooling to form stars in more massive haloes, which leads to larger disks. Conversely, reducing the feedback allows gas to cool and form stars in smaller haloes resulting in smaller discs. The strength of SN feedback is parameterized using $V_{\\rm hot}$ and $\\alpha_{\\rm hot}$ as shown in Eq.~\\ref{eqsnfeedback}. The adopted values for these parameters are: $V_{\\rm hot} = 300\\,{\\rm km\\,s}^{-1}$ and $\\alpha_{\\rm hot} = 2$ in the Baugh2005 model and $V_{\\rm hot} = 485\\,{\\rm km\\,s}^{-1}$ and $\\alpha_{\\rm hot} = 3.2$ in the Bower2006 model. We perturb the models by increasing and reducing the value of $V_{\\rm hot}$ to its double and half the fiducial value in each model, and plot the results in Fig.~\\ref{Shenmedvhot}. The normalization of the size-luminosity relation for disk-dominated galaxies changes as expected on changing the strength of supernova feedback, moving to larger sizes on increasing $V_{\\rm hot}$ and smaller sizes on reducing $V_{\\rm hot}$. Reducing $V_{\\rm hot}$ in the Baugh2005 model leads to better agreement with the observed size-luminosity relation for disk-dominated galaxies, at the expense of producing slightly more faint galaxies. Similar trends are seen in the predictions for the size-luminosity relation of bulge-dominated galaxies. Note that there are very few bulge-dominated galaxies at faint magnitudes in the Bower2006 model, hence the noisy size-luminosity relation in this region. Changing the strength of feedback in this way has little impact on the slope of the size-luminosity relation.  \\item[(ii)] {\\it The threshold for disks to become unstable.}  The threshold for a disk to become unstable is set by the parameter $\\epsilon_{\\rm disk}$ (see Eq.~\\ref{eqinstab}). We show the result of varying this threshold in Fig.~\\ref{Shenmedinstab}. In the case of the Bower2006 model, we increase and reduce the threshold from the fiducial value of $\\epsilon_{\\rm disk}=0.8$; increasing the threshold means that more disks become unstable. The original Baugh2005 model does not test for the stability of disks, so in this case we switch the effect on and try two different values for the threshold. The result of turning on dynamical instabilities is straightforward to understand in this model.  For a given mass and rotation speed of disk, the stability criteria $\\epsilon \\propto \\sqrt{r_{\\rm disk}}$, and so disks with smaller values of $r_{\\rm disk}$ will preferentially be unstable. The removal of small disks raises the median disk size but reduces the fraction of galaxies that are disk-dominated. The impact of varying the stability threshold on the Bower2006 model is less easy to interpret: even though the fraction of faint disk-dominated galaxies fails dramatically on increasing the threshold, there is little change in the median size of the surviving disks. This change has a bigger impact on the size of bulge-dominated galaxies, with a sequence that is inverted compared with the Baugh2005 model. One result that is easily understood is the response of the luminosity function.  The burst of star-formation which can accompany the transformation of a dynamically unstable disk into a bulge is an important channel for generating black hole mass in the Bower2006, as discussed in Section 2.2.  If fewer disks become unstable, less mass is converted into black holes and AGN feedback has less impact on the cooling flows in massive haloes, leading to too many bright galaxies.   \\item[(iii)] {\\it The orbital energy of merging galaxies}  The parameter $f_{\\rm orbit}$ quantifies the orbital energy of two galaxies which are about to merge (Eq.~\\ref{Eorbit}). Our standard choice in both models is $f_{\\rm orbit}=1$, in which case Eq.~\\ref{Eorbit} is equal to the energy of two point masses in a circular orbit at a separation of $r_{1}+r_{2}$. We vary the value of $f_{\\rm orbit}$ trying $f_{\\rm orbit} = 0$, which corresponds to a parabolic orbit, and $f_{\\rm orbit}=2$. The results are plotted in Fig.~\\ref{Shenmedforbit}. As expected, the median size of disk-dominated galaxies is unaffected by varying $f_{\\rm orbit}$. Increasing the value of $f_{\\rm orbit}$ makes bulge-dominated galaxies smaller, with a larger effect seen for brighter galaxies. A smaller value of $f_{\\rm orbit}$ improves the shape of the size-luminosity relation of bulge-dominated galaxies in the Baugh2005 model; however, faint bulge-dominated galaxies are still too large after making this change.  \\item [(iv)]{\\it The contribution of the dark matter in galaxy mergers}  The parameter $f_{\\rm dark}$ controls the amount of dark matter associated with model galaxies during merger evenst (see Eq.~\\ref{DMcontrib}), which has an impact on the size of the merger remnant, through Eq.~\\ref{Eorbit} and Eq.~\\ref{rnew}.  We run the models using values of $f_{\\rm dark}=1$ and $f_{\\rm dark}=0$, the latter of which corresponds to the case of a galaxy without participating dark matter.  We can see that the reduction of $f_{\\rm dark}$ from the fiducial value leads to smaller sizes for the early type galaxies in both models.  The effect is particularly important at bright magnitudes. As expected, there is no change in the predicted sizes for late type galaxies. We do not find, either, a variation in the luminosity function, but there is an increase in the fraction of early type galaxies, particularly at intermediate magnitudes. We can see that we can improve the sizes of early type galaxies, matching with those inferred from SDSS observations for galaxies with magnitudes fainter than $L\\star$. However, this change is counterproductive at bright magnitudes, resulting in even smaller sizes.  \\end{itemize}   \\subsubsection{What drives the slope of the size-luminosity relation?}   We have seen that the prediction of the Bower2006 model for the size-luminosity relation of disk-dominated galaxies is much flatter than that of the Baugh2005 model (Fig.~\\ref{Shenmedvhot}). Moreover the Baugh2005 model is in better agreement with the observed relation. In the previous section we varied selected model parameters one at a time relative to the fiducial model, to show their impact on the model size-luminosity relation.  In this exercise, the most dramatic change in the Bower2006 predictions resulted from varying the strength of SNe feedback.  Reducing the value of the parameter $V_{\\rm hot}$, which sets the ``pivot'' velocity below which SNe feedback has a strong impact, leads to a shift in the size-luminosity relation to smaller sizes, with an improved match to the observed relation recovered for intermediate luminosity galaxies. In this section, we investigate the effect of varying several parameters together, essentially moving from the Bower2006 parameters for SNe feedback and the star formation timescale in disks, towards a set which more closely resembles that used in the Baugh2005 model.  The size-luminosity relations for disk- and bulge-dominated galaxies are plotted in Fig.~\\ref{flattosteep} for a sequence of models.  The starting point is the fiducial Bower2006 model. For each step in the sequence, one parameter is varied relative to the previous step, as shown in the key in Fig.~\\ref{flattosteep}. The first change made is to the value of $\\alpha_{\\rm hot}$, which controls the slope of the SNe feedback. Changing from the Bower2006 value of $\\alpha_{\\rm hot}=3.2$ to $\\alpha_{\\rm hot}=1$ gives a much improved match to the observed size-luminosity relation, particularly for intermediate luminosities. Faint disk-dominated galaxies are still somewhat too large, and bright galaxies in general are still too small. The next step is to retain the above change to $\\alpha_{\\rm hot}$, and also to change the value of $V_{\\rm hot}$ to that used in Baugh2005.  This results in a modest improvement in the size-luminosity relation for faint galaxies. Finally, the scaling of the quiescent star formation timescale with the disk dynamical time is switched off. The resulting size-luminosity relation is now in very good agreement with the observations for disk-dominated galaxies. However, bright bulge-dominated galaxies are still too small. Furthermore, the luminosity function and the predicted fraction of early types with luminosity are now much poorer matches to observations than in the fiducial Bower2006 model (lower two panels of Fig.~\\ref{flattosteep}).    \\begin{figure} \\includegraphics[width=8.5cm]{figs/Talkmedreal2008onlyBow.paperoneatatime.ps} \\caption{ The impact of cumulative parameter changes for a sequence of models, starting from the fiducial model of Bower2006 (blue). One parameter is changed in each step, but unlike the cases presented in Figs.~16, 17, 18, 19, the change made is relative to the previous model in the sequence.  The key lists the parameter change relative to the previous model in the sequence.  The upper two panels show the size-luminosity relation for bulge and disk-dominated galaxies respectively. The third panel shows the luminosity function and the bottom panel shows the morphological mix of galaxies as a function of luminosity.  } \\label{flattosteep} \\end{figure}  In summary, even though the sizes of disk-dominated galaxies can be brought into reasonable agreement with observations in a variant of the Bower2006 model with modified SNe feedback and disk star formation timescale, this is at the expense of agreement with the observed luminosity function and early-type fraction. Furthermore, neither the fiducial Baugh2005 model nor the fiducial or modified Bower2006 models are able to reproduce the observed sizes of early-type galaxies, in particular for bright galaxies. No single parameter change seemed able to solve the latter problem.  The most promising area to explore further appears to be the modelling of galaxy mergers; changing the amount of orbital energy brought in by merging galaxies in our prescription led to an increase in the sizes of the brightest bulge-dominated galaxies. ",
        "conclusions": "Observations of local galaxies have always played a central role in setting the parameters of galaxy formation models. However, the huge size of the SDSS sample combined with the quality and uniformity of the data allow much more precise and exacting tests of the physics of such models than was previously possible. To take full advantage of this opportunity, it is necessary for the model to be able to generate predictions which are as close as possible to the measurements made for real galaxies.  In this paper we use the {\\tt GALFORM} model, which predicts the size of both the disk and bulge components of galaxies. Hence we are able to take the model output for the luminosity and scale-lengths of a galaxy's disk and bulge and turn these into predictions for the quantities measured for SDSS galaxies: the Petrosian magnitude, the radius containing 50\\% of the Petrosian flux ($R_{50}$), the concentration parameter ($c$) and the S\\'{e}rsic index ($n$), the latter two quantities providing descriptions of the light profile of the galaxy.  The first major result of this work is to understand the correlation between different indicators of galaxy morphology. The concentration parameter, S\\'{e}rsic index and bulge-to-total ($B/T$) luminosity ratio have all been used to divide galaxies into morphological classes \\citep[e.g][]{Bershady00,Hog04,Ben07}.  The $B/T$ ratio is easy to compute theoretically, yet is perhaps the hardest of these quantities to measure observationally. Both the $c$-$B/T$ and $n$-$B/T$ planes show scatter. This can be traced to the ratio of the disk and bulge scale-lengths; galaxies with different values of this ratio occupy different loci in the $c$-$B/T$ and $n$-$B/T$ planes. The scatter is particularly large in the case of the S\\'{e}rsic index $n$ versus $B/T$. The scatter would be even larger if we simply generated galaxies by hand, taking the ratio of disk and bulge scale-lengths from a grid. The scatter we find is limited by the distribution of $r_{\\rm d}/r_{\\rm b}$ values predicted by {\\tt GALFORM}.  We compared the predictions of two different versions of the {\\tt GALFORM} model with SDSS data: that of \\citet{Bau05}, which has a top-heavy IMF in starbursts and feedback from superwinds, and \\citet{Bow06}, which has AGN feedback and a normal IMF in all modes of star formation. In the first stage of the comparison, none of the model parameters were adjusted to improve the fit to the data. The models gave reasonable matches to the total galaxy luminosity function, with the \\citeauthor{Bow06} model giving the best overall match to the shape. The match to the luminosity function of different colour subsamples is less impressive; both models overpredict the number of bright blue galaxies and fail to match the number of red galaxies.  The \\citeauthor{Bow06} model has a strongly bimodal colour distribution, whereas the \\citeauthor{Bau05} model shows only weak bimodality. The \\citeauthor{Bow06} model agrees better overall with the observed colours in SDSS, although the predicted bimodality appears somewhat too strong, and the positions of the peaks in the colour distribution do not agree exactly with what is observed.  Another clear success of the models is in predicting the correct trend of morphological type with galaxy luminosity. We used all three morphological indicators (concentration parameter, S\\'{e}rsic index and bulge-to-total luminosity ratio from disk+bulge fits) to separate galaxies into disk-dominated and bulge-dominated types. In the SDSS data, the fractions of these types are found to shift from being almost completely disk-dominated at low luminosity to almost entirely bulge-dominated at high luminosity, though with differences in the detailed behaviour depending on which morphological indicator is used. Both the \\citet{Bau05} and \\citet{Bow06} models successfully reproduce this general trend, although the \\citeauthor{Bow06} model agrees better in detail with the observed behaviour at intermediate luminosities. Both models qualitatively reproduce the observed correlation of colour with morphology (with bulge-dominated galaxies on average being redder than disk-dominated galaxies at every luminosity), although quantitatively the \\citeauthor{Bow06} model agrees better with the SDSS data, the \\citeauthor{Bau05} model predicting too many red disk-dominated galaxies.   Perhaps the most serious shortcoming of the models is in the predicted galaxy sizes. Whilst the \\citeauthor{Bau05} model gives a very good match to the luminosity-size relation for disks, the sizes of bulge-dominated galaxies do not match the observations. The slope of the size-luminosity relation for bulge-dominated galaxies in the Baugh et~al. model matches the observations for faint galaxies, but the normalization is too high. Brighter than $L_{*}$, the predicted relation flattens, with the consequence that the brightest bulge-dominated galaxies are around a factor of three too small (see also \\citealt{Al07}). The situation is worse for the sizes predicted by the \\citeauthor{Bow06} model; in this case the size - luminosity relation is flat for disk-dominated galaxies, while for bulge-dominated galaxies the predicted sizes at high luminosities fall even further below the observed relation than for the \\citeauthor{Bau05} model. We have demonstrated that a steeper slope for the size-luminosity relation for disk-dominated galaxies can be recovered in the \\citeauthor{Bow06} model if we set some physical processes to have the same parameters as used in the \\citeauthor{Bau05} model. The primary improvement in the model predictions is seen on reducing the strength of SNe feedback.  Also, by adjusting the star formation timescale in disks by switching off the dependence on the disk dynamical time, we can recover the observed slope of the size-luminosity relation even at high luminosities. However, this improvement in the size - luminosity relation comes at the expense of producing too many galaxies overall.   The differences between the predictions of the two models for the sizes of disk-dominated galaxies lie in the revised cooling model adopted by \\citet{Bow06}, the strength of supernova feedback, the inclusion of AGN feedback and the inclusion of dynamical instabilities for disks. In the \\citeauthor{Bow06} model, gas which is reheated by supernova feedback is reincorporated into the hot halo on a shorter timescale than in the \\citet{Bau05} model.  Neither model is able to match the observed size of bright bulge-dominated galaxies. We explore a range of processes in the models, varying the strength of supernova feedback, changing the threshold for disks to become unstable, and changing the prescription for computing the size of the stellar spheroid formed in a galaxy merger.  The latter seems the most promising solution, at least in the case of the \\citeauthor{Bau05} model. If we neglect the orbital energy of the galaxies which are about to merge (i.e. setting the parameter $f_{\\rm orbit}=0$), then the sizes of bright galaxies are in much better agreement with the observed sizes (though the faint bulge-dominated galaxies are still too large). A similar change in the predictions results from ignoring the adiabatic contraction of the dark matter halo in response to the gravity of the disk and bulge \\citep{Al07}.   The {\\tt GALFORM} model is one of the few able to make the range of predictions considered in this paper and hence to take full advantage of the constraining power of the SDSS. The model for disk sizes works well under certain conditions, albeit with too much scatter.  Our analysis suggests that the problems with disk sizes and colours are connected to the treatment of gas cooling and supernova feedback, while the problems with spheroid sizes are probably due to an overly-simplified treatment of the sizes of galaxy merger remnants. The prescription used to compute the size of spheroids is in need of improvement, which will require the results of numerical simulations of galaxy mergers.  This study highlights the need to make careful and detailed comparisons with observational data in order guide improvements in the treatment of physical processes in galaxy formation models."
    },
    "0812/0812.1743_arXiv.txt": {
        "abstract": "{We examine a sample of 30 edge-on spiral and S0 galaxies that have boxy and peanut-shaped bulges. We compute model stellar kinematics by solving the Jeans equations for axisymmetric mass distributions derived from $K$-band images. These simple models have only one free parameter: the dynamical mass-to-light ratio, which we assume is independent of radius. Given the simplicity of the modelling procedure, the model second velocity moments are strikingly good fits to the observed stellar kinematics within the extent of our kinematic data, which typically reach $\\approx$ 0.5--1\\,$R_{25}$ (where $R_{25}$ is the optical radius), or equivalently $\\approx$ 2--3\\,$R_\\mathrm{e}$ (where $R_\\mathrm{e}$ is the effective or half-light radius). We therefore find no evidence for a dominant dark matter component within the optical disk of spiral galaxies. This is equally true of the S0s in our sample, which significantly extends previous observational constraints on dark matter in these galaxies. The predicted kinematics do deviate slightly but systematically from the observations in the bulge region of most galaxies, but we argue that this is consistent with the claim that boxy and peanut-shaped bulges are bars viewed edge-on. ",
        "introduction": "Previous studies indicate with some certainty that dark matter does not make a significant contribution to the total mass within the optical radius, $R_{25}$, of disk galaxies \\citep[e.g.][]{Palunas:2000,Bell:2001,Kassin:2006}. Similar studies of ellipticals find little evidence for significant dark matter within the effective radius, $R_\\mathrm{e}$ \\citep[e.g.][]{Gerhard:2001,Rusin:2003,Romanowsky:2003,Cappellari:2006,Thomas:2007}. However, degeneracies in the modelling mean that the dark matter content of ellipticals is more uncertain. Moreover, $R_\\mathrm{e} \\ll R_{25}$.  Non-axisymmetric galaxies make it possible to lift degeneracies in the contributions from luminous and dark matter by ascribing non-circular motions to luminous matter only \\citep[e.g.][]{Englmaier:1999,Weiner:2001}. Fast observed pattern speeds also constrain dark matter in barred disks \\citep[e.g.][]{Debattista:2000,Aguerri:2003,Gerssen:2003}.  In this study we use the sample of thirty S0 and spiral galaxies with a boxy or peanut-shaped bulge of \\cite{Bureau:1999} to constrain their dark matter content. From previous numerical and observational work, such galaxies are thought to be barred disk galaxies viewed edge-on \\citep[e.g.][]{Kuijken:1995,Bureau:1999,Chung:2004,Mendez-Abreu:2008}. ",
        "conclusions": "Because our models are such good fits to the data, they provide no evidence for a dominant dark component within 0.5--1$ R_{25}$ or 2--4 $R_\\mathrm{e}$ in a sample of thirty S0 and spiral galaxies. This is consistent with previous studies of spiral galaxies. It also extends the radius out to which observational constraints may be placed on the dark matter content of S0s. In fact, since S0s have been demonstrated to be observationally equivalent to fast-rotating ellipticals \\citep{Emsellem:2007,Cappellari:2007}, this statement may also be tentatively applied to fast-rotating ellipticals.  The $K$-band dynamical mass-to-light ratios that we measure are typically in the range 1--2\\,$M_\\odot/L_{\\odot,K}$. This is consistent with those of the SAURON sample \\citep{Cappellari:2006}, but the mass-to-light ratios predicted by the single stellar population models of \\cite{Maraston:2005} with a Kroupa initial mass function do not exceed 1.4\\,$M_\\odot/L_{\\odot,K}$. This may thus hint at the presence of dynamically subdominant dark haloes. To place constraints on the amount, we have constructed Jeans models which explicitly include dark halos \\citep{Williams:2008}  We are limited by the small size of our non-boxy control sample, but the slight but systematic deviation of the models within the bulge regions is consistent with the claim that boxy bulges are bars viewed edge-on. The sense of the discrepancy (model $>$ data) is also consistent with the construction of the sample, which is dominated by bars viewed side-on (end-on bars appear spheroidal so were not included in this boxy sample). The radial extent of this discrepancy is smaller than the bars, which are probably about twice the length of the boxy bulge \\citep{Athanassoula:2005}, but simple models have shown that the axial ratio of the $x_1$ orbits supporting the bar decreases rapidly toward unity from its peak at around 20 percent of the length of the bar \\citep[see, e.g., Figure~2~(a) of][]{Bureau:1999a}.  We will follow up the present work by using the derived dynamical mass-to-light ratios in two ways. Firstly, we will compare them to those inferred from stellar population syntheis models based on absorption line strength indices, providing further constraints on the dark matter content and initial mass function of the galaxies. Secondly we will compute total dynamical masses and study the Tully-Fisher relation of these galaxies, which should reveal information about the possible evolutionary link between spirals and S0s.  \\begin{figure*} \\centering \\label{fig:results} \\includegraphics{williams1d} \\includegraphics{williams1e} \\includegraphics{williams1f} \\includegraphics{williams1g} \\includegraphics{williams1h} \\includegraphics{williams1i} \\addtocounter{figure}{-1} \\caption{ --- continued} \\end{figure*}"
    },
    "0812/0812.0603_arXiv.txt": {
        "abstract": "In this study we present simultaneous low-resolution longitudinal magnetic field measurements and high-resolution spectroscopic observations of the cool single giant FK~Com. The variation of the magnetic field over the rotational period of 2.4\\,days is compared with the starspot location obtained using Doppler imaging techniques, V-band photometry and V-I colours. The chromospheric activity is studied simultaneously with the photospheric activity using high resolution observations of the H$\\alpha$, H$\\beta$ and H$\\gamma$ line profiles. Both the maximum (272$\\pm$24\\,G) and minimum (60$\\pm$17\\,G) in the mean longitudinal magnetic field, $\\left<B_z\\right>$, are detected close to the phases where cool spots appear on the stellar surface. A possible explanation for such a behaviour is that the active regions at the two longitudes separated by 0.2 in phase have opposite polarities. ",
        "introduction": "The strongest magnetic fields in non-degenerate stars are measured in Ap stars (see, e.g., \\citealt{bab, land, hub1}). The development of spectropolarimetric techniques in the last years has enabled us to detect the weaker stellar magnetic fields seen in many cool stars, and to study the topologies of their magnetic fields in detail using Zeeman-Doppler imaging.  Zeeman-Doppler imaging \\citep{ZDI} is closely related to traditional Doppler imaging (see, e.g., \\citealt{vogt_DI, pisk90}) where high resolution, high signal-to-noise (S/N) spectra at different rotational phases of a star are used to measure the rotationally modulated distortions in the line-profiles. These distortions are produced by the inhomogeneous distribution of the observed characteristic, e.g., effective temperature or elemental abundance. Surface maps, so called  Doppler images, are constructed by combining all the observations from different rotational phases and comparing them with synthetic model line-profiles. In Zeeman-Doppler imaging similar techniques are used, but the observations consist of high resolution spectropolarimetric observations, either including circular polarisation (Stokes V) or preferably the full Stokes vector. Using this technique different types of stars have been mapped: young rapid rotators (e.g., \\citealt{3stars, marsden_V889Her}), more evolved stars (e.g., \\citealt{petit_HD199178, 3stars}), and M dwarfs (e.g., \\citealt{Donati_V374Peg}).  Zeeman-Doppler imaging provides detailed maps of the stellar surface magnetic fields, but observationally this technique is very demanding as high resolution, high S/N observations in polarised light are needed. On the other hand, information on the longitudinal magnetic field can be obtained also using low resolution spectropolarimetric observations in circular polarisation (see, e.g., \\citealt{hub_06}). With such observations we can retrieve the component of the magnetic field parallel to the line-of-sight, weighted by the local emergent spectral line intensity and averaged over the stellar hemisphere visible at the time of observation. Due to the low resolution, such spectropolarimetric observations, however, do not provide a detailed insight into the magnetic surface structure since the polarised line-profiles of temperature and magnetic field sensitive lines cannot be studied individually.  Here we present the first study of the mean longitudinal magnetic field of FK~Com, which is the prototype of the small class of FK~Comae stars \\citep{bopp_stencel}. These stars are magnetically very active late-type giants with photometric and spectroscopic characteristics similar to those of the very active RS~CVn stars. The main difference between the groups is that the RS~CVn stars are close binary systems in which the tidal effects produce synchronous rotation, and therefore also rapid rotation, whereas the FK~Comae stars do not show significant radial velocity variations caused by a companion, and are thus most likely single stars.  FK~Com itself, with its $v$\\,sin\\,$i$ of 160\\,km/s, is the fastest rotator of the FK~Comae stars, and hence also the most active among this group. The spectrum of FK~Com was first described by \\citet{merrill}. He noted a large projected rotational velocity, H$\\alpha$ and Ca II H\\&K emission and the variability of the H$\\alpha$ profile. Due to the very broad line profiles the radial velocity of FK~Com is difficult to measure very accurately, leaving still some room to speculate on its possible binarity. The radial velocity variations have been constrained to $\\pm 3$\\,km/s by \\citet{hue93} giving very strict limits to the mass of the possible companion. Small visual brightness variations in FK~Com of $\\sim$0.1\\,mag in the V-band with a period of 2.412 days were first reported by \\citet{chu66}. These variations are interpreted to be caused by large starspots on the surface. Several surface temperature maps of FK~Com have been obtained over the years using Doppler imaging technique (see, e.g., \\citealt{fkcom1, fkcom5}). These surface temperature maps mainly show high latitude spots with spot temperatures of typically 1000~K less than that of the unspotted surface.  The current work provides the first measurement of the magnetic field on FK~Com. The low resolution spectropolarimetric observations in circular polarisation are used to infer the behaviour of the mean longitudinal magnetic field at different rotational phases. The Doppler imaging technique is applied to the high resolution spectra to construct a simultaneous surface temperature map. The longitudinal magnetic field measurements are compared to the contemporaneous starspot locations obtained from the surface image and the chromospheric activity seen in the H$\\alpha$, H$\\beta$ and H$\\gamma$ lines. ",
        "conclusions": "The following conclusions can be drawn from the low resolution spectropolarimetry, high resolution spectroscopy and broad band photometry presented in this work:  \\begin{itemize} \\item The mean longitudinal magnetic field $\\left<B_z\\right>$, measured in April 2008, shows clear variations over the rotational phase of FK~Com, with a maximum $\\langle$$B_z$$\\rangle$\\,=\\,272$\\pm$24\\,G at the phase 0.08 and the minimum $\\langle$$B_z$$\\rangle$\\,=\\,60$\\pm$17\\,G at the phase 0.36. \\item The contemporaneous determination of the starspot locations by the Doppler imaging technique reveals three main spot groups on the surface of FK~Com in April 2008. Two groups are located at high latitudes close to 60$^{\\circ}$, and one is located close to the equator. The separation of the high latitude spot groups is 0.2 in phase. \\item Both the maximum and minimum value of $\\left<B_z\\right>$ occur within the phases of the broad photometric minimum. The maximum value coincides with the cool high latitude spot spanning the phases 0.00--0.14 and the minimum value with the high latitude and equatorial spots around the phases 0.21--0.35. This behaviour could be explained by the active region around the phase 0.1 having positive magnetic field polarity and the less cool spots around phase 0.3 a magnetic field of negative polarity. \\item The H$\\alpha$ profiles show enhanced emission at the phases 0.0--0.4, coinciding with the photospheric spots. The phases 0.6--0.9, when the photospheric spots are hardly seen, show least emission and the strongest central absorption core. Similar behaviour is also seen in the H$\\beta$ and H$\\gamma$ lines. \\item We would like to emphasise that obtaining high resolution spectropolarimetric observations of FK~Com, preferably before and after a flip-flop event, would provide valuable observational constraints for flip-flop Dynamos. \\end{itemize}"
    },
    "0812/0812.1814_arXiv.txt": {
        "abstract": "{ The periodicity of 5.5 years for some observational events occurring in $\\eta$~Carin\\ae\\/ manifests itself across a large wavelength range and has been associated with its binary nature. These events are supposed to occur when the binary components are close to periastron. To detect the previous periastron passage of $\\eta$~Car in 2003, we started an intensive, ground-based, optical, photometric observing campaign. } { We continued observing the object to monitor its photometric behavior and variability across the entire orbital cycle. } { Our observation program consisted of daily differential photometry from CCD images, which were acquired using a 0.8 m telescope and a standard $BVRI$ filter set at La Plata Observatory. The photometry includes the central object and the surrounding Homunculus nebula. } { We present up-to-date results of our observing program, including homogeneous photometric data collected between 2003 and 2008. Our observations demonstrated that $\\eta$~Car has continued increasing in brightness at a constant rate since 1998. In 2006, it reached its brightest magnitude ($V \\sim 4.7$) since about 1860s. The object then suddenly reverted its brightening trend, fading to $V = 5.0$ at the beginning of 2007, and has maintained a quite steady state since then. We continue the photometric monitoring of $\\eta$~Car in anticipation of the next ``periastron passage'', predicted to occur at the beginning of 2009. } ",
        "introduction": "$\\eta$~Carin\\ae\\ is the brightest ($V\\sim 5$) Luminous Blue Variable star in the sky. The central object is surrounded by a nebulosity, a product of the mass ejection that occurred during the {\\em``Great Eruption''} in the first half of the 19th century. It consists of a $18''$ long axis bipolar reflection nebula, which was called ``Homunculus'' by \\citet{Gaviola}.  The central star is suspected to be a binary system as first proposed by \\citet{D97}, from the evidence of the 5.5 years periodicity of the ``spectroscopic events''. These events consist of the attenuation or complete disappearance of high excitation lines \\citep{D96}.  The most recent ``spectroscopic event'' occurred in 2003.5 and, for that reason, $\\eta$~Car was the target of an intensive multiwavelength observational campaign. Clear dips were registered in X-ray luminosity \\citep{C05}, near-infrared $JHKL$ photometry \\citep{White04}, and 7-mm \\citep{Abraham05a} and 3-cm wave emissions \\citep{Duncan03}. In the optical range, we carried out a multicolor CCD monitoring of $\\eta$~Car. Our observations started in January 2003, and we detected a minimum in the four $BVRI$ bands \\citep[hereafter Paper I]{ibvs5477}, more than a week after detecting a minimun in X-ray luminosity. These sudden brightness fadings are occasionally referred to as ``eclipse-like'' events \\citep[e.g.][]{White04}, and are explained well by a close binary scenario, in which the hot secondary star moves in a highly eccentric orbit of period 2022.7 days \\citep{Damineli08a}. Many of these events are supposed to be produced during periastron passages \\citep{D97} and extend over only a few months due to the high orbital eccentricity. Nevertheless, some other features vary continuously throughout the complete cycle \\citep[e.g.][]{Duncan03,Damineli08b}   .  Following the 2003 observing campaign, we continued monitoring the photometric behavior during the complete orbital period, including the upcoming ``eclipse-like'' event, predicted for next January 2009. It is worth emphasizing the significant advantage of the daily availability of a telescope of the appropriate size, and retaining the same instrumental configuration (i.e. ``telescope + filter-set + detector''). This is a desirable factor when high-precision is sought in long-term ground-based photometric monitoring of $\\eta$~Car  \\citep[and references therein]{Sterken01}.  In this paper, we present the results of our long-term $BVRI$ CCD differential photometry of $\\eta$~Car, performed between 2003 and 2008, covering the entire orbital period. ",
        "conclusions": "In the following, we outline the main results derived from our photometry since January 2003: \\begin{itemize}  \\item The system maintained its tendency to brighten in the four photometric bands until 2006.35, when the maximum brightness in 150 years was reached.  \\item After this maximum in 2006.35, the brightness of $\\eta$~Car started to decline. Since 2007.2, the $BVRI$ light curves have remained at almost constant brightness, showing a slight brightness rise in the days of writing of this report. As we could expect considering the events of 2003, such an increase is probably associated with an upcoming periastron passage.  \\item Just after the last periastron passage in 2003.5, the photometry displayed a wide maximum, called an ``egress-maximum'' by \\citet{vG06}, which finished at a clearly evident minimum.  \\item The maximum peak was reached shortly after the orbital phase $\\phi \\sim 0.5$. This encourages us to relate it with some type of phenomena occurring during apastron passage. The light fluctuation detected about that maximum peak is quite symmetrical, covering 30\\% of the period. This would be consistent with the flux variations that occur after the periastron passage, reported by \\citet{Duncan03,Damineli08b}. Nevertheless, we cannot discount a possible coincidence that an S Dor phase just happened to occur at that particular orbital moment.  \\item The time at which a minimum occurred was estimated from the 2003.5 event, and by adopting a 2022.7 days period, the next optical photometric dip might be expected in January 28, 2009. The onset of the 2009-event might have occurred about one or two months before.  \\end{itemize}  We have provided dense and well defined multiband light curves of $\\eta$~Car + Homunculus, covering the entire orbital period of the proposed binary system. This consistent and homogeneous data set provides invaluable information for further modeling, analysis, and interpretation, which we expect to enhance our understanding of the physical nature of $\\eta$ Car."
    },
    "0812/0812.2210_arXiv.txt": {
        "abstract": "It is argued that cosmological models that feature a flow of energy from dark energy to dark matter may solve the coincidence problem of late acceleration (i.e., ``why the energy densities of both components are of the same order precisely today?\"). However, much refined and abundant observational data of the redshift evolution of the Hubble factor are needed to ascertain whether they can do the job. ",
        "introduction": "As is well known, the  popular vacuum-cold dark matter ($\\Lambda$CDM) model fits rather well most observational data, however it suffers from two main -and lasting- shortcomings, namely: the low value of the vacuum energy (about 123 magnitude orders below the quantum field theory estimation) and the coincidence problem. The latter refers to the fact that despite the vacuum energy does not vary with time and dark matter energy depends on redshift as $(1+z)^{3}$ both energy densities are nowadays of the same order. This is why very often the vacuum component is replaced by an evolving dark energy  field -see \\cite{rreviews,padm} for recent reviews. Nevertheless, this does not satisfactorily solve the problem since the initial value of the dark energy ought to be fine tuned up to obtain the desired result; it only substitutes a fine tuning by another \\cite{padm}.  An encouraging approach that is receiving increasing attention \\cite{list} rests on the assumption that  both components do not evolve separately but interact with each other, i.e., \\begin{equation} \\dot{\\rho}_{m} + 3 H \\rho_{m}  = Q \\, , \\qquad {\\rm and} \\qquad \\dot{\\rho}_{\\phi} + 3 H (\\rho_{\\phi} + P_{\\phi}) = -Q \\, . \\label{cons} \\end{equation} Here, $Q$ stands for the interaction term, and the subscripts $m$ and $\\phi$ denote non-relativistic (pressureless) dark matter (DM) and dark energy (DE), respectively. Clearly, this represents a step forward because for $Q > 0$ (i.e., for energy transfers from DE to DM) the ratio between both energies, $r \\equiv \\rho_{m}/\\rho_{\\phi}$, evolves more slowly than in the $\\Lambda$CDM model. Effectively,  in the latter, $\\dot{r} = -3H r$ while in the interacting case its time evolution of  governed by \\begin{equation} \\dot{r} = 3 H r\\left[w+\\frac{\\kappa^{2}\\,Q}{9\\,H^3}\\frac{(r+1)^2}{r} \\right] \\, , \\label{rr} \\end{equation} where $ w \\equiv P_{\\phi}/\\rho_{\\phi} < -1/3$ denotes the equation of state parameter of DE. Equation (\\ref{rr}) follows from (\\ref{cons}) and Friedmann's equation for spatially flat universes \\begin{equation} 3 H^2 = \\kappa^{2} (\\rho_{m} + \\rho_{\\phi}) \\quad \\qquad(\\kappa^{2} \\equiv 8\\pi\\, G)\\, . \\label{Fried1} \\end{equation} As we will see below, for reasonable and suitably chosen interaction terms, $r$ either tends to a constant or varies more slowly than the scale factor, $a(t)$, at late times. This certainly alleviates the coincidence problem. It is worthy of note that if $Q$ were negative (signifying an energy transfer from DM to DE) the problem would only worsen, moreover at early times the energy density of DE would have been negative, and, finally, the second law of thermodynamics could be violated \\cite{lima,db}. Except in chamaleon models \\cite{chamaleon}, baryons as well as radiation usually do not participate in the interaction. The former because of the tight constraints imposed by local gravity measurements \\cite{Peebles2,Hagiwara}, the latter because otherwise the photons would not longer follow geodesics which would affect precise measurements of deviations of radar signals grazing the sun.  It is not the aim of this article to solve the coincidence problem -we believe it cannot be done at present- but to devise a simple strategy to do it, within general relativity, as soon as sufficiently accurate and abundant data on the evolution of the Hubble factor becomes available. This paper extends and generalizes a previous work of the authors \\cite{rapid}.  Most studies of late acceleration implicitly assume that matter and dark energy interact gravitationally only. In the absence of an underlying symmetry that would suppress a matter-dark energy coupling (or interaction) there is no a priori reason to discard it. Further, the coupling is not only likely but inevitable \\cite{jerome} and its introduction is less arbitrary than otherwise. In fact, the accretion of DE by black holes \\cite{cai-wang,bean} implies a non-vanishing and positive $Q$. While its small value modifies only very little the dependence of the energy densities on redshift (and cannot substantially alleviate the coincidence problem), it sets a minimum value for the interaction.  The interaction may bring some further consequences. It can push the beginning of accelerated expansion era to higher redshifts. It can also  erroneously suggest an equation of state for the dark energy of phantom type (i.e., $w <-1$), see \\cite{amendola-gasperini} and references therein. Moreover,  the interaction may give rise to fluctuations in the count of galaxy clusters with redshift \\cite{mota}. Likewise, the coupling can alter the isothermal Maxwell-Boltzmann velocity distribution of weakly interacting massive particles in the galaxy halos \\cite{Tetradis1} whereby the average dark matter velocity is expected to increase significantly. Since the detection rates of the different experiments searching for dark matter strongly depend on this velocity it is crucial not to dismiss {\\em a priori} the coupling \\cite{Tetradis2}.   Ultimately, the existence or non-existence of the interaction is to be discerned observationally. Whereas the current empirical data are insufficient to settle this issue, a couple of studies seem to favor the former possibility: $(i)$ As it should  be expected, the interaction alters the time required for a self-gravitating, collapsing, structure to reach equilibrium as well as the equilibrium configuration itself. Therefore the Layzer-Irvine equation \\cite{layzer,Peebles1} is to be generalized to take into account the interaction. In this regard, from the analysis of the dynamics of 33 relaxed galaxy clusters -for which reliable x-ray, weak lensing and optical data are available-, it has been reported a small but not vanishing interaction \\cite{elcio,orfeu}. $(ii)$ Since the interaction modifies the evolution rate of the metric potentials the integrated Sachs-Wolfe component of the cosmic microwave background radiation (CMB) is enhanced. The cross-correlation of galaxy catalogs with CMB maps also hints at a small interaction \\cite{isw-olivares}.    This paper is organized as follows. Next section introduces three different expressions for $Q$ and show that all three lead to a late, everlasting, era of either constant or quasi-constant $r$ regardless of whether $w$ is lower or higher than $-1$. Section 3 shows the consistency of the three interacting models  with the recent observational bounds of Daly {\\it et al.} \\cite{daly1} based on data from supernovae type Ia, taken from different sources, and 30 radio-galaxies. The only cosmological assumption behind these bounds is that the  Universe can be considered homogeneous and isotropic at large scales. Section 4 devises a simple strategy to solve the coincidence problem. Finally, Section 5 summarizes and discusses our results.  We focus our attention on spatially flat Friedmann-Robertson-Walker models. As usual, a subscript zero means that the corresponding quantity is to be evaluated at present time. ",
        "conclusions": "Nowadays the leading candidate for dark energy is the popular $\\Lambda$CDM model. It is simple, it has just one free parameter -the vacuum energy density- and fits reasonably well most (if not all) observational data. However, it presents the puzzle (apparent or not) that we happen to live at a very special epoch: the transient epoch at which the ratio between the densities of matter and vacuum energy is of order unity. To solve this  one must look for alternative models in which the ratio $r$ stays constant or varies very slowly, around the present time, with respect to the Universe expansion. Interacting models (with $Q >0$, see Eq. (\\ref{rr})) possess this property. But it does not suffice. The successful model must also predict $r_{0} \\sim {\\cal O}(1)$.  In this paper we considered three interacting models that present the first characteristic and may also have the second one. To decide on the latter one must contrast the $r(z)$ function predicted by each model with observation. Regrettably, as far as we know, model-independent data about $r(z)$ do not exist. However, albeit with wide uncertainties, there exist model-independent data of the evolution of the Hubble factor up to $z = 1.8$ \\cite{daly1,simon}. Therefore, we have numerically determined $H(r)$ and $H(z)$ for each model and compared the latter  with these observational data. All three models show compatibility with the data but only in a  degree similar to the $\\Lambda$CDM model. Future, more accurate and extensive measurements will likely tell us whether the latter model fits the data better or worse than interacting models (the ones presented in this paper or maybe others). If, in the end, the $\\Lambda$CDM model makes the best fit, then we may conclude that the coincidence problem is not a problem any more; it is simply  a coincidence and therefore not a problem."
    },
    "0812/0812.0812_arXiv.txt": {
        "abstract": "We performed a series of high-resolution (up to $1024^3$) direct numerical simulations of hydro and MHD turbulence. Our simulations correspond to the \"strong\" MHD turbulence regime that cannot be treated perturbatively. We found that for simulations with normal viscosity the slopes for energy spectra of MHD are similar to ones in hydro, although slightly more shallower. However, for simulations with hyper viscosity the slopes were very different, for instance, the slopes for hydro simulations showed pronounced and well-defined bottleneck effect, while the MHD slopes were relatively much less affected.  We believe that this is indicative of MHD strong turbulence being less local than Kolmogorov turbulence. This calls for revision of MHD strong turbulence models that assume local ``as-in-hydro case'' cascading. Non-locality of MHD turbulence casts doubt on numerical determination of the slopes with currently available ($512^3$--$1024^3$) numerical resolutions, including simulations with normal viscosity.  We also measure various so-called alignment effects and discuss their influence on the turbulent cascade. ",
        "introduction": "Turbulence is ubiquitous in astrophysical fluids which are characterized by high Reynolds numbers. It affects key astrophysical processes, e.g. star formation (Elmegreen \\& Scalo 2004, McKee \\& Ostriker 2007). The observational signatures of turbulence are numerous and well documented. For instance, random fluctuations on all scales, which is a sign of turbulence, has been detected by a variety of observational techniques (see Crovisier \\& Dickey 1983, O'Dell \\& Castaneda, Armstrong et al. 1994, Lazarian 2009). Turbulence is universal, because the laminar flows with high Reynolds numbers are practical impossibility. It is driven by a variety of mechanisms such as supernova explosions starbursts, stellar winds, AGN jets, etc. In Keplerian flows turbulence is generated by magnetorotational instability (Velikhov 1959, Chadrasekhar 1960, Balbus \\& Hawley 1991). Galactic disks are subject to the cosmic ray-induced Parker's instability (Parker 1966).  Although, historically, hydrodynamic turbulence has been applied to astrophysics, now it is accepted that for almost all astrophysical fluid flows are coupled with magnetic fields, at least on large scales. This necessitates the use of the dynamic equations that include electric currents and magnetic fields. The simplest approach in this respect is the continuous non-relativistic one-fluid description, known as magnetohydrodynamics or MHD. This approach is broadly applicable to most of the astrophysical environments, such as Solar wind, Interstellar and Intercluster Medium, molecular clouds, stars interiors, and so on, although there are some exceptions, such as ultra-relativistic jets and shocks, where full relativistic equations should be used or small scales of Molecular Clouds, where, due to relatively low ionization rate, the two-fluid description of ions and neutrals is more appropriate.  The importance of astrophysical turbulence inspires much of theoretical and numerical work aimed at understanding its properties. We should clarify, however, that there are different types of MHD turbulence. In this paper we deal with strong MHD turbulence, which can not be treated perturbatively. The theory of weak Alfv\\'enic turbulence, which has limited applicability to astrophysics, is discussed elsewhere (Sridhar \\& Goldreich 1994, Ng \\& Bhattacharjee 1996, Lazarian \\& Vishniac 1999, Galtier et al. 2000, 2002). In order to study the basic properties of MHD cascade and to be able to directly compare to previous work, we restricted ourselves to so-called balanced MHD turbulence or turbulence with zero net cross-helicity. The properties of {\\it imbalanced turbulence} are studied in a companion paper Beresnyak \\& Lazarian (2009).  The issue of the spectral slopes of MHD turbulence has caused a substantial interest recently. A sizable number of papers attempting to measure true asymptotic spectral slope of high-Reynolds number MHD turbulence from direct numerical simulations appeared to date. For example, simulations of weakly compressible MHD turbulence performed in \\citet{haugen2003, haugen2004} used finite-difference code with numerical resolution of up to $1024^3$ with explicit viscous and resistive dissipation. The energy spectral slope for MHD was observed to be shallower than $-5/3$ which was interpreted as the influence of the bottleneck effect. Another example is the paper of \\citet{muller2005} who measured the spectral slopes of decaying MHD turbulence without mean field and driven MHD turbulence with strong mean field. The pseudospectral method with ordinary viscosity was ran at a numerical resolution of up to $1024^3$. The authors argued that the slope was close to $-3/2$ in the mean-field case.   The motivation behind these and many other papers was to understand the nature of the turbulent cascade. The turbulent energy transfer is, in a sense, a central issue of turbulence, be it hydrodynamic or MHD. And while hydrodynamic turbulence has its ``Standard Model'', the nature of MHD cascade is still debated.     An important first step in turbulence theory was made by Iroshnikov (1963) and Kraichnan (1965) who noticed that there is a local magnetic field which can not be excluded by a choice of reference frame, like the average velocity in hydrodynamics. Furthermore, they assumed that turbulence is weak, because perturbations are smaller than the mean field. This implicitly assumed local isotropic dynamics and happened to be a mistake, since the turbulent cascade preferred perpendicular direction, i.e. produced perturbations that are more and more anisotropic, this way increasing interaction and preventing turbulence from becoming weak.  The understanding of various aspects of MHD turbulence, including the role of turbulence anisotropy, compressibility, etc. resulted in a number of publications (see Dobrowolny, Mangeney, \\& Veltri 1980, Shebalin, Matthaeus \\& Montgomery 1983, Montgomery \\& Turner 1984, Higdon 1984). For the most part of the paper we will consider incompressible MHD turbulence, which properties are dominated by the Alfv\\'enic perturbations. Interestingly enough, some properties of Alfv\\'enic turbulence carry over not only to nearly incompressible low Mach number flows, but also for flows with Mach numbers larger than unity (Cho \\& Lazarian 2003, see also \\S 9.3).  It has been realized that interactions of Alfv\\'enic modes in MHD turbulence has a tendency of getting stronger as the cascades unfolds. Goldreich \\& Sridhar (1995, henceforth GS95) proposed a particular model of strong turbulence when the interaction strength is being controlled by two competing processes: a perpendicular cascade (a concept rigorously developed in a theory of weak Alfv\\'enic turbulence, e.g. Galtier et al 2000), which tends to increase the interaction, and a decorrelation due to cascading which tends to increase the frequency of perturbations, and thus decrease the interaction. GS95 concluded that the interaction has to be marginally strong (``critical balance''), and therefore, the cascade has to be of strong Kolmogorov type and has a spectral slope of around $-5/3$. Predictions of this model, such as scale-dependent anisotropy, were subsequently observed in three-dimensional simulations (Cho \\& Vishniac 2000, Maron \\& Goldreich 2001, etc).  Recently, however, a number of models, motivated by numerical spectral slopes shallower than $-5/3$, appeared (Boldyrev 2005, 2006, Gogoberidze 2007). The interest to this field has been heated by the numerical discovery of so-called scale-dependent polarization alignment (Beresnyak \\& Lazarian 2006), which was interpreted in a subsequent publication by Mason et al (2006) in favor of Boldyrev's (2006) modification of GS95 model.  This paper is bringing attention to serious difficulties that appear when one tries to measure true asymptotic slopes from direct three-dimensional numerical simulations which have rather moderate Reynolds numbers. By comparing spectra of hydrodynamic and MHD turbulence we found that MHD turbulence might be much more nonlocal that Kolmogorov turbulence, and, therefore, require much higher resolution to obtain spectral slope by brute force approach. The second part of this paper brings polarization alignment to more numerical scrutiny and compares the simulation results to theory. Our simulations allow to test some of the existing conjectures about properties of MHD turbulence. For instance, these results allow us to reject a conjecture in \\citet{boldyrev2006} that the alignment is limited by the magnitude of local field wandering.    In what follows in \\S2 we describe our approach based on comparing the properties of MHD and hydrodynamic turbulence, present our numerical setup in \\S3, present spectra in \\S4, discuss anisotropy in \\S5, discuss the interaction weakening and the alignment effects in \\S6 and \\S7, respectively, and provide some more hints of the non-locality of the MHD cascade in \\S8. In \\S9 we compare our numerical results with the existing theoretical predictions, as well as with the previous numerical work; we also discuss the applicability of findings obtained with incompressible simulations to the real-world compressible astrophysical turbulence. Our conclusions are summarized in \\S10. ",
        "conclusions": "We conclude that although the Kolmogorov-like Goldreich-Sridhar (GS95) model is appealing, simple and captures some essential physical properties of the strong MHD turbulence, such as scale-dependent anisotropy, it should be amended to explain cascade nonlocality and scale-dependent alignment effects. How to achieve this is the issue of the future research, as we demonstrate that the existing attempts to improve GS95 do not agree well with the presented numerical simulations. In addition, we issue a note of warning that the numerical measurements of the spectral slope that served as a motivation for many of theoretical studies are unlikely to represent the true theoretical slopes due to the non-locality of MHD cascade."
    },
    "0812/0812.0023_arXiv.txt": {
        "abstract": "When the Wide Field Camera 3 (WFC3) is installed on the {\\it Hubble Space Telescope (HST)}, the community will have access to powerful new capabilities for investigating resolved stellar populations.  The WFC3 Galactic Bulge Treasury program will obtain deep imaging in five photometric bands on four low-extinction fields. These data will have no proprietary period, and will enable a variety of science investigations not possible with previous data sets. To aid in planning for the use of these data and for future observing proposals, we provide an introduction to the Treasury program, its photometric system, and the associated calibration effort.  The observing strategy is based upon a new photometric system employing five WFC3 bands spanning the UV, optical, and near-infrared: F390W, F555W, F814W, F110W, and F160W (analogous but not identical to the ground-based filters Washington $C$, $V$, $I$, $J$, and $H$). With these bands, one can construct reddening-free indices of temperature and metallicity.  Using this photometric system, the program will target six fields in well-studied star clusters, spanning a wide range of metallicity, and four fields in low-extinction windows of the Galactic bulge.  The cluster data serve to calibrate the reddening-free indices, provide empirical population templates, and correct the transformation of theoretical isochrone libraries into the WFC3 photometric system.  The bulge data will shed light on the bulge formation history, and will also serve as empirical population templates for other studies. One of the fields includes 12 candidate hosts of extrasolar planets.  Color-magnitude diagrams (CMDs) are the most popular tool for analyzing resolved stellar populations. However, due to degeneracies among temperature, metallicity, and reddening in traditional CMDs, it can be difficult to draw robust conclusions from the data.  The five-band system used for the bulge Treasury observations will provide reddening-free indices that are roughly orthogonal in temperature and metallicity, and we argue that model fitting in an index-index diagram will make better use of the information than fitting separate CMDs.  We provide some results from simulations to show the expected quality of the data and their potential for differentiating between different star-formation histories. ",
        "introduction": "The Wide Field Camera 3 (WFC3) is scheduled to be installed on the {\\it Hubble Space Telescope (HST)} during the next servicing mission. Although the camera is intended as a general-purpose instrument, it offers particularly powerful new capabilities for studying resolved stellar populations.  The camera will provide wide-field high-resolution imaging with continuous spectral coverage from the ultraviolet into the near-infrared through a choice of 77 filters and 3 grisms.  Because it will be the first {\\it HST} instrument to provide wide-field imaging in the near-infrared, an obvious target for WFC3 is the Galactic bulge, given its high reddening and crowding.  We describe here an observing technique that employs five WFC3 bands to create reddening-free indices of temperature and metallicity that can be used to explore a highly reddened and complex stellar population, such as the bulge.  The tools and techniques needed for a bulge investigation should be useful to a wider range of resolved stellar population studies, and a large database of bulge stars with reliable parameters (temperature, metallicity, proper motion, etc.) can enable additional observing programs directed specifically at the bulge, from both ground and space.  These factors drove the creation of a bulge Treasury program (GO-11664), where all of the tools and products would be non-proprietary.  The primary purpose of this paper is to demonstrate the general utility of the program's population tools (\\S2), which include photometry of six star clusters that can be used as empirical population templates (\\S2.3).  In addition, this paper summarizes the scientific goals and data products specific to the Galactic bulge (\\S3).  That description of the bulge investigation also motivates the development of the tools, provides concrete examples of how they can be used in stellar population studies, and enables future bulge studies.  An early look at the program and its products allows them to be considered in the planning of observing proposals for upcoming {\\it HST} cycles. ",
        "conclusions": "With its broad wavelength coverage, high spatial resolution, rich filter set, and wide field of view, WFC3 brings powerful new stellar population tools to {\\it HST}.  Taking advantage of these capabilities, the new five-band photometric system described here will enable the construction of reddening-free indices of metallicity and temperature.  With these indices, observers can fit the star formation history in a population by modeling the distribution of stars in the $T_{\\rm eff}$-[Fe/H] plane, avoiding the reddening uncertainties arising in traditional CMD fitting.  Although this system was designed with the specific goal of understanding the formation of the Galactic bulge, it should be useful to a wide range of stellar populations work with WFC3.  Furthermore, because it is based upon analogs to ground-based bandpasses, it could in principle be used in other observatories (if properly calibrated to the exact bandpasses used). The Treasury program will provide population templates, isochrone transformations, and photometric indices in this system early in the WFC3 mission.  Because the bulge science described herein is part of the Treasury program, it will also provide accurate catalogs of metallicity, temperature, and proper motion for hundreds of thousands of bulge stars; these should be useful to other observational and theoretical efforts directed at the bulge.  We encourage observers to consider the availability of these products when planning stellar populations proposals with WFC3 in the upcoming {\\it HST} observing cycles."
    },
    "0812/0812.2356_arXiv.txt": {
        "abstract": "In massive stars, magnetic fields are thought to confine the outflowing radiatively-driven wind, resulting in X-ray emission that is harder, more variable and more efficient than that produced by instability-generated shocks in non-magnetic winds. Although magnetic confinement of stellar winds has been shown to strongly modify the mass-loss and X-ray characteristics of massive OB stars, we lack a detailed understanding of the complex processes responsible. The aim of this study is to examine the relationship between magnetism, stellar winds and X-ray emission of OB stars. In conjunction with a Chandra survey of the Orion Nebula Cluster, we carried out spectropolarimatric ESPaDOnS observations to determine the magnetic properties of massive OB stars of this cluster. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.2951_arXiv.txt": {
        "abstract": "We present infrared observations in search of a planet around the white dwarf, GD66. Time-series photometry of GD66 shows a variation in the arrival time of stellar pulsations consistent with the presence of a planet with mass $\\geq$2.4\\mj. Any such planet is too close to the star to be resolved, but the planet's light can be directly detected as an excess flux at 4.5\\um.  We observed GD66 with the two shorter wavelength channels of IRAC on Spitzer but did not find strong evidence of a companion, placing an upper limit of 5--7\\mj\\ on the mass of the companion, assuming an age of 1.2--1.7\\,Gyr. ",
        "introduction": "The advantage of searching for sub-stellar companions around white dwarf stars (WDs) has long been recognized. The improved contrast between a WD and a companion was first exploited by \\citet{Probst83} and the first L dwarf around a WD (or indeed any star) was discovered around GD165 by \\citet{Becklin88}. In recent years, improved instrumentation has enabled searches for lower mass objects, to the point where direct detection searches for planetary mass objects are possible from the ground \\citep[e.g.][Hogan et al. 2008, in press]{Burleigh02, Debes05cc2}. Searches utilizing the {\\it Spitzer} Space Telescope \\citep{Werner04} can obtain mass limits as low as $\\sim$10\\mj\\, \\citep[][]{Farihi08planet, Mullally07thesis, Debes07cc4}.  WDs are interesting targets for more reasons than just their luminosity. By looking at WDs we can sample a range of progenitor masses from about 1-8\\msolar, including higher mass stars not amenable to radial velocity surveys. Currently, the best main-sequence targets for direct detection of planetary companions are either transiting systems \\citep[e.g.]{Charbonneau05, Deming06} or very young stars \\citep{Marois08,Kalas08}. Detection of a planet around a WD would open up a older, more sedate region of parameter space where the predictions of planet models could be tested against observation for ages ranging from hundreds of millions to billions of years. WD planets are key to testing our models of planetary evolution.  WDs also lend insight into the ultimate fate of planetary systems. \\citet{Livio84} first estimated the conditions necessary for a gas giant planet to survive the red giant phase of its parent star. \\citet{Sackmann93} concluded that the Earth would survive the red giant phase of the Sun, but more recent calculations by \\citet{Schroder08} concluded the opposite. These calculations depend on stellar mass loss rates that are weakly constrained by observation due to the rapid evolution of evolved stars. Determining the distribution of planets around WDs will provide a key test for these models. The winds from evolved stars pollute the interstellar medium with the material that creates the next generation of terrestrial planets, so surveys for planets around dead stars will inform our understanding of planets around stars not yet born.  Early empirical success comes from the detection of a planet around a sub-dwarf star using pulsation timings \\citep{Silvotti07}, planets in an eclipsing sub-dwarf system \\citep{Lee08} and the detection of silicate rich debris disks around WDs \\citep[e.g.][]{Reach05b, Farihi07gd362}. Analysis of the composition of material raining down onto the surface of a WD from such a disk by \\citet{Zuckerman07} concluded that the pattern of elemental abundances is consistent with a species of asteroid. Asteroids are markers of planetary systems, and so the frequency of debris disk WDs may eventually provide an independent measure of the fraction of stars that bear (or bore) planetary systems.  \\citet[][hereafter M08]{Mullally08} published the results of a survey that sought to detect evidence of planets as variations in the arrival time of pulsations from variable WDs (in a manner analogous to sub-dwarf and pulsar planets). For one object, GD66, they discovered a variation consistent with a 2.1\\mj\\ planet in a 4.5\\,yr orbit at a distance of about 50\\,pc. Unfortunately, the span of their data was not long enough to cover an entire orbit and their detection remains provisional. In this paper we report on follow up observations with {\\it Spitzer} in an attempt to confirm this object by direct detection in the infrared. Although the candidate system is too distant to resolve into individual components, the low luminosity of the star means that the flux from the planet at 4.5\\um\\ rivals that of the WD. For wavelengths less than 10\\um  atmosphere models of planets are brightest at 4--5\\um, in a gap between absorbption bands of methane and water \\citep[e.g][]{Burrows03}. The flux from this bump increases with mass, and decreases with age, and provides a method of both detecting a planet and determining its mass. The goal of these observations is to measure an excess flux at 4.5\\um\\ over the bare photosphere of GD66 due to the light from the planet. We achieve this by measuring the flux ratio (i.e. the color) between channels 2 and 1 on IRAC \\citep[4.5 and 3.6\\um\\ respectively,][]{Fazio04} and comparing this ratio to a sample of other similar stars. This approach is capable of detecting companions independent of distance to the star, or the assumptions of atmosphere models.  These observations push the IRAC instrument in a new direction. While studies of transiting planets with IRAC have been able to detect changes in flux in a single band of order a few parts in $10^{-3}$, our goal was to measure an absolute flux excess relative to other wavelengths. With this approach, we established a sensitivity limit of $\\sim$0.5\\% with the two IRAC bands available during the warm mission. ",
        "conclusions": "We report on {\\it Spitzer} IRAC observations of the white dwarf GD66 in an attempt to directly detect a companion planet. We fail to detect an excess with any statistical significance. However, combined with our ground based timing observations we can now constrain the planet candidate's mass to between 2.4 and $\\approx$5--6\\mj."
    },
    "0812/0812.2483_arXiv.txt": {
        "abstract": "The emission mechanism and the origin and structure of magnetic fields in gamma-ray burst (GRB) jets are among the most important open questions concerning the nature of the central engine of GRBs. In spite of extensive observational efforts, these questions remain to be answered and are difficult or even impossible to infer with the spectral and lightcurve information currently collected. Polarization measurements will lead to unambiguous answers to several of these questions. Recent developments in X-ray and $\\gamma$-ray polarimetry techniques have demonstrated a significant increase in sensitivity enabling several new mission concepts, e.g. {\\it POET} (Polarimeters for Energetic Transients), providing wide field of view and broadband polarimetry measurements. If launched, missions of this kind would finally provide definitive measurements of GRB polarizations. We perform Monte Carlo simulations to derive the distribution of GRB polarizations in three emission models; the synchrotron model with a globally ordered magnetic field (SO model), the synchrotron model with a small-scale random magnetic field (SR model), and the Compton drag model (CD model). The results show that {\\it POET}, or other polarimeters with similar capabilities, can constrain the GRB emission models by using the statistical properties of GRB polarizations. In particular, the ratio of the number of GRBs for which the polarization degrees can be measured to the number of GRBs that are detected ($N_m/N_d$) and the distributions of the polarization degrees ($\\Pi$) can be used as the criteria. If $N_m/N_d > 30\\%$ and $\\Pi$ is clustered between 0.2 and 0.7, the SO model will be favored. If instead $N_m/N_d < 15\\%$, then the SR or CD model will be favored. If several events with $\\Pi > 0.8$ are observed, then the CD model will be favored. ",
        "introduction": "\\label{sec:intro}  Gamma-ray bursts (GRBs) are brief, intense flashes of $\\gamma$-rays originating at cosmological distances, and they are the most luminous objects in the universe. They also have broadband afterglows long-lasting after the $\\gamma$-ray radiation has ceased. It has been established that the bursts and afterglows are emitted from outflows moving towards us at highly relativistic speeds \\citep{taylor04}, and at least some GRBs are associated with the collapse of massive stars \\citep[e.g.,][]{hjorth03,stanek03}. Observations suggest that the burst is produced by internal dissipation within the relativistic jet that is launched from the center of the explosion, and the afterglow is the synchrotron emission of electrons accelerated in a collisionless shock driven by the interaction of the jet with the surrounding medium \\citep[for recent reviews,][]{piran05,meszaros06,zhang07}.  In spite of extensive observational and theoretical efforts, several key questions concerning the nature of the central engines of the relativistic jets and the jets themselves remain poorly understood. In fact, some of these questions are very difficult or even impossible to answer with the spectral and lightcurve information currently collected. On the other hand, polarization information, if retrieved, would lead to unambiguous answers to these questions. In particular, polarimetric observations of GRBs can address the following:  {\\it Magnetic composition of GRB jets --} It is highly speculated that strong magnetic fields are generated at the GRB central engine, and may play an essential role in the launch of the relativistic jets. However, it is unclear whether the burst emission region is penetrated by a globally structured, dynamically important magnetic field, and whether the burst is due to shock dissipation or magnetic reconnection \\citep[e.g.,][]{spruit01,zhang02,lyutikov03}.  {\\it Emission mechanisms of the bursts --} The leading model for the emission mechanism of the prompt burst emission is synchrotron emission from relativistic electrons in a globally ordered magnetic field carried from the central engine, or random magnetic fields generated in-situ in the shock dissipation region \\citep{rees94}. Other suggestions include Compton drag of ambient soft photons \\citep{shaviv95,eichler03,levinson04,lazzati04}, synchrotron self-Compton emission \\citep{panaitescu00}, and the combination of a thermal component from the photosphere and a non-thermal component (e.g., synchrotron) \\citep{ryde06,thompson07,ioka07}.  {\\it Geometric structure of GRB jets --} Although it is generally believed that GRB outflows are collimated, the distribution of the jet opening angles, the observer's viewing direction, and whether there are small-scale structures within the global jet are not well understood \\citep{zhang04,yamazaki04,toma05}.  To date, robust positive detections of GRB polarization have been made only in the optical band in the afterglow phase. Varying linear polarizations have been observed in several optical afterglows several hours after the burst trigger, with a level of $\\sim 1-3\\%$, which is consistent with the synchrotron emission mechanism of GRB afterglow \\citep[for reviews, see][]{covino04,lazzati06}. An upper limit $(< 8\\%)$ has been obtained for the early $(t\\sim200~{\\rm s})$ optical afterglow of GRB 060418 \\citep{mundell07}. Also for radio afterglows, we have several upper limits for the polarization degree \\citep{taylor05,granot05} \\citep[for some implications, see][]{toma08}. As for the prompt burst emission, strong linear polarization of the $\\gamma$-ray emission at a level of $\\Pi = 80 \\pm 20\\%$ was claimed for GRB 021206 based on an analysis of {\\it RHESSI} data \\citep{coburn03}, although this claim remains controversial because of large systematic uncertainties \\citep{rutledge04,wigger04}. Several other reports of high levels of polarization in the prompt burst emission are also statistically inconclusive \\citep{willis05,kalemci07,mcglynn07}.  Recently, more sensitive observational techniques for X-ray and $\\gamma$-ray polarimetry have been developed, and there are several polarimeter mission concepts. These include Polarimeters for Energetic Transients \\citep[{\\it POET},][]{hill08,bloser08}, Polarimeter of Gamma-ray Observer \\citep[{\\it PoGO},][]{mizuno05}, {\\it POLAR} \\citep{produit05}, Advanced Compton Telescope \\citep[{\\it ACT},][]{boggs06}, Gravity and Extreme Magnetism \\citep[{\\it GEMS},][]{jahoda07} {\\it XPOL} \\citep{costa07}, Gamma-Ray Burst Investigation via Polarimetry and Spectroscopy \\citep[{\\it GRIPS},][]{greiner08}, and so on.  Several of these missions, if launched, would provide definitive detections of the burst polarizations and enable us to discuss the statistical properties of the polarization degrees and polarization spectra. Although there are several polarimetry mission concepts described in the literature, {\\it POET} is the only one to date that incorporates a broadband capability for measuring the prompt emission from GRBs, and for this reason it provides a good case study for our simulations. {\\it POET} will make measurements with two different polarimeters, both with wide fields of view. The Gamma-Ray Polarimeter Experiment (GRAPE; 60-500~keV) and the Low Energy Polarimeter (LEP; 2-15~keV) provide a broad energy range for the observations. Suborbital versions of both {\\it POET} instruments are currently being prepared for flight within the next few years. GRAPE will fly on a sub-orbital balloon in 2011, and the Gamma-Ray Burst Polarimeter (GRBP, a smaller version of LEP) will fly on a sounding rocket.  Theoretically, it has been shown that similarly high levels of linear polarization can be obtained in several GRB prompt emission models; the synchrotron model with a globally ordered magnetic field, the synchrotron model with a small-scale random magnetic field \\citep{granot03,lyutikov03,nakar03}, and the Compton drag model \\citep{lazzati04,eichler03,levinson04,shaviv95}. Thus the detections of GRB prompt emission polarization would support these three models. In this paper, we show that these models can be distinguished by their statistical properties of observed polarizations. We performed detailed calculations of the distribution of polarization degrees by including realistic spectra of GRB prompt emission and assuming realistic distributions of the physical parameters of GRB jets, and show that {\\it POET}, or other polarimeters with similar capabilities, can constrain the GRB emission models. We use the limits of {\\it POET} for GRB detection and polarization measurements as realistic and fiducial limits. This paper is organized as follows. We first introduce the {\\it POET} mission concept in \\S~\\ref{sec:poet}. In \\S~\\ref{sec:theory}, we summarize the properties of the observed linear polarization from uniform jets within the three emission models. Based on these models, we perform Monte Carlo simulations of observed linear polarizations and show how the statistical properties of observed polarization may constrain GRB emission mechanisms in \\S~\\ref{sec:simulation}. A summary and discussion are given in \\S~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  Recently there has been an increasing interest in the measurement of X-ray and $\\gamma$-ray polarization, and the observational techniques can now achieve significant sensitivity in the relevant energy bands. Several polarimetry mission concepts, such as {\\it POET}, are being planned. The {\\it POET} concept has two polarimeters, GRAPE (60-500~keV) and LEP (2-15~keV) both of which have wide fields of view. If launched, missions of this type would provide the first definitive detection of the polarization of GRB prompt emission. This would enable the discussion of the statistical properties of the polarization degree and polarization spectra, which will give us diagnostic information on the emission mechanism of GRBs and the nature of the GRB jets that cannot be obtained from current spectra and lightcurve observations. We have performed Monte Carlo simulations of the linear polarization from GRB jets for three major emission models: synchrotron model with globally ordered magnetic field (SO model), synchrotron model with small-scale random magnetic field (SR model), and Compton drag model (CD model). We assumed that the physical quantities for the emission of the jets are uniform on the emitting surface and that the jets have sharp edges. Our jet angle distribution allows the detections of GRBs with very small opening angles (i.e., smaller than 1 degree) as suggested by several {\\it Swift} bursts \\citep{schady07,racusin08}. We have shown that the {\\it POET} mission or other polarimeters with similar capabilities, i.e., broadband spectral capabilities for the determination of $E_{p,{\\rm obs}}$ and sensitive broadband polarimetric capabilities to minimize MDP, can constrain the emission models of GRBs. Furthermore, these simulations indicate that an increase in the LEP effective area would be beneficial to compensate for the lower expected polarization at lower energies.  As shown in Figures~\\ref{fig:atorob}, \\ref{fig:arndb}, and \\ref{fig:acomp}, the SR and CD models require off-axis observations of the jets to achieve a high level of polarization, while the SO model does not. In this sense the SR and CD models are categorized as {\\it geometric} models, and the SO model as an {\\it intrinsic} model \\citep{waxman03,lazzati06}. The distribution of observed polarizations obtained by our simulations show that the geometric SR/CD models will be ruled out if the number ratio of the $\\Pi$-measurable bursts to detected bursts is larger than $30\\%$, and in this case the SO model will be favored if the measurable polarizations are clustered at $0.2<\\Pi<0.7$. If the number ratio is smaller than $15\\%$, the SO model may be ruled out, but we cannot distinguish between the SR and CD models with different distributions of $y_j=(\\gamma \\theta_j)^2$, $\\alpha$, and $\\beta$, where $\\gamma$ and $\\theta_j$ are the bulk Lorentz factor and the opening angle of the GRB jet, respectively, and $\\alpha$ and $\\beta$ are lower and higher indices of the energy spectrum. However, if several bursts with $\\Pi > 0.8$ are detected, the CD model which includes an adequate number of small $y_j$ bursts will be favored.  If the cumulative distribution of the measurable polarizations favors the SO model, the globally ordered magnetic field would be advected from the central engine. If we understand the strength of the magnetic field in the emitting region from the luminosity and the spectrum of the emission, we can constrain the strength of the field at the central engine. If the geometric SR/CD models are favored from the observations, it will be established, independently of the afterglow observations, that GRB outflows are not spherical but highly collimated. If the CD model is favored by the observations, we may constrain the distribution of the parameter $y_j = (\\gamma \\theta_j)^2$ of GRB jets. The CD model needs a dense optical/UV photon field interacting within the relativistic jets \\citep{lazzati00,eichler03}.  We have made some simplifications in our simulations, and there are some caveats. We have assumed that the jets are uniform on the emitting surfaces and have sharp edges. To compare the simulations and the observations further, more sophisticated modeling is required \\citep[e.g.,][]{zhang04,toma05}.  We have interpreted bursts as a simple combination of pulses, without taking account of the temporal variation of the Lorentz factor $\\gamma$ of the jet. If this is accounted for, each pulse may have different $y_j=(\\gamma \\theta_j)^2$ but the same $q=\\theta_v/\\theta_j$. We should then average the polarization with respect to fluence of each pulse having different $y_j$ \\citep{granot03,nakar03}. However, in the SO model, the cumulative distribution of measurable $\\Pi$ will not be changed significantly as long as $y_j > 10$, because $\\Pi$ is clustered into a small range for $q<1$ and $y_j>10$. To average the polarization in the case of $q>1$, the relation between the luminosity and the Lorentz factor for each pulse is required to predict the polarization distribution.  For the SR model we have assumed that the directions of the magnetic field are confined within the shock plane. They may be more isotropic in reality, in which case the polarization degree in the SR model will be reduced.  In the synchrotron model with a combination of the globally ordered magnetic field and the locally random field, $\\mathbf{B} = \\mathbf{B}_{\\rm ord} + \\mathbf{B}_{\\rm rnd}$, the linear polarization can be calculated by $\\Pi = (Q_{\\rm ord} + Q_{\\rm rnd})/(I_{\\rm ord} + I_{\\rm rnd}) \\approx (\\Pi_{\\rm ord}+\\eta \\Pi_{\\rm rnd})/(1+\\eta)$, where $\\{I,Q\\}_{\\rm ord}$ and $\\{I,Q\\}_{\\rm rnd}$ are the Stokes parameters from the ordered and random fields, respectively. $\\Pi_{\\rm ord}$ and $\\Pi_{\\rm rnd}$ are described by equations (\\ref{eq:atorob}) and (\\ref{eq:arndb}), and $\\eta \\equiv (B_{\\rm rnd}/B_{\\rm ord})^{\\alpha+1}$. This model will reduce the number ratio of $\\Pi$-measurable bursts to detected bursts to less than 30\\% and the clustering of measurable polarizations will be at $\\Pi < 0.7$."
    },
    "0812/0812.3103_arXiv.txt": {
        "abstract": "A reanalysis of the fluorine abundance in three Galactic AGB carbon stars (TX Psc, AQ Sgr and R Scl) has been performed from the molecular HF (1-0) R9 line at 2.3358 $\\mu$m. High-resolution (R$\\sim 50000$) and high signal to noise  spectra obtained with the CRIRES spectrograph and the VLT telescope or from the NOAO archive (for TX Psc) have been used. Our abundance analysis uses the latest generation of MARCS model atmospheres for cool carbon rich stars. Using spectral synthesis in LTE we derive for these stars fluorine abundances that are systematically lower by $\\sim 0.8$ dex in average with respect to the sole previous estimates by Jorissen, Smith \\& Lambert (1992). The possible reasons of this discrepancy are explored. We conclude that the difference may rely on the blending with C-bearing molecules (CN and C$_2$) that were not properly taken into account in the former study. The new F abundances are in better agreement with th prediction of full network stellar models of low mass AGB stars. These models also reproduce the $s$-process elements distribution in the sampled stars. This result, if confirmed in a larger sample of AGB stars, might alleviate the current difficulty to explain the largest [F/O] ratios found by Jorissen et al. In particular, it may not be necessary to search for alternative nuclear chains affecting the production of F in AGB stars. ",
        "introduction": "The origin  of the sole stable  isotope of fluorine,  $^{19}$F, is not well  known.  It  is easily destroyed by proton and alpha capture reactions in the stellar interiors so that, its abundance is the lowest among the light  elements  with atomic number $6\\leq Z\\leq  20$. Fluorine has  only a few accessible atomic  and molecular lines suitable  for abundance studies, detected whether in hot gaseous nebulae in the ultra-violet spectral region or in  a rather crowed  region at  $\\sim 2.3~\\mu$m  in cool  objects. Three  are the proposed sites of $^{19}$F   production:  neutrino  spallation   on   $^{20}$Ne  in gravitational  supernovae  (SNII)  (Woosley  \\& Haxton  1988), hydrostatic  nucleosynthesis in the He-burning core of heavily  mass-losing Wolf-Rayet (WR) stars (Meynet  \\&  Arnould  2000),   and  hydrostatic nucleosynthesis in the He-rich  intershell of thermally pulsing (TP) Asymptotic Giant Branch (AGB) stars  (Forestini et al.  1992). Up to date, the contribution of each source to the fluorine content in the Universe is still controversial. Renda el al. (2004)  concluded that the inclusion of all the three components is necessary to explain the  observed fluorine Galactic  evolution as inferred  from abundance determinations in a small  sample of field red giants (Cunha \\& Smith  2005; Cunha et al.  2008).  However,  recent measurements of fluorine in  the interstellar medium  (Federman et al. 2005)  yield no evidence of F  over-abundances caused by the neutrino process in SNII.  In addition, Palacios et al.   (2005) revisited the  F production in rotating WR star models concluding that the F yields from WR stars are significantly  lower than previously predicted  by Meynet  \\& Arnould  (2000), so that their contribution  to  the Galactic  F budget would be negligible.  This leaves  AGB  stars as  the only  significant producers.  The first  evidence of F production  in AGB stars was  provided by the study  of  Jorissen,  Smith  \\&  Lambert  (1992;  hereafter  JSL)  who determined  F abundances  from rotational  HF lines  in  extrinsic (binary) and intrinsic  stars of  near  solar metallicity  along  the AGB  spectral sequence   (M$\\rightarrow$MS$\\rightarrow$S$\\rightarrow$C).  JSL  found [F/Fe] ratios  up to  $\\sim 100$ times  solar and a  clear correlation between the F enhancement and the  C/O ratio. Because the C/O ratio is expected  to  increase  along  the  AGB evolution  due  to  the  third dredge-up (TDU), this occurrence has been  interpreted as an evidence of the F production in AGB stars. This result has been later supported by the large F enhancements found in  post-AGB stars (Werner et al. 2005) and planetary nebulae  (Zhang \\& liu 2005; Otsuka et al.  2008, and references  therein), the progeny of AGB stars. Nevertheless, up to now, nucleosynthetic models in  AGB  stars  (e.g. Forestini  et  al.  1992;  Lugaro et  al.  2004; Cristallo  et al.  2008)  have failed  to  quantitatively reproduce  the largest [F/Fe,O] ratios found in the  JSL's sample. Such a discrepancy has led to a deep revision  of the uncertainties associated with the major nuclear reactions affecting  the production/destruction of F in  AGB stars, to a search for alternative nuclear chains (e.g. Lugaro et al. 2004), or to invoke non-standard mixing processes capable to increase the F production below the convective envelope (e.g. the cool bottom process,  see Wasserburg, Boothroyd, \\& Sackmann 1995 and Busso et al.  2007).  However, no satisfactory   solution  has   been  found, making this large fluorine enhancement a major challenge for stellar nucleosynthesis.  Extant theoretical models show that the nuclear chain allowing the fluorine production in the He-rich and H-exhausted intershell zone of a TP-AGB star is $^{14}$N$(\\alpha,\\gamma)^{18}$F$(\\beta^+)^{18}$O$(p,\\alpha)^{15}$N$(\\alpha,\\gamma)^{19}$F. Note that the $^{15}$N production production requires the existence of a few protons, which can eventually be produced by the $^{14}$N$(n,p)^{14}$C reaction, but only when neutrons are released by the $^{13}$C$(\\alpha,n)^{16}$O reaction. The latter is the main neutron source in AGB stars, giving rise to the $s$-process nucleosynthesis (e.g. Gallino et al. 1998; Cristallo et al. 2008). Therefore, F and $s$-elements are expected to correlate in TP-AGB stars undergoing TDU. Nevertheless, some of the stars with the largest F abundances among those in the JSL sample are J-type carbon stars, which do not present $s$-element enhancements (Abia \\& Isern 2000).  As a part of a more extended work focused on the study of the F production in AGB stars and its  dependence with the stellar metallicity, we present here new F abundance determinations in three galactic AGB carbon stars included in the  JSL'sample. Based on the analysis of the R9 (1-0)  HF line at $\\lambda_{air}=2.33583~\\mu$m\\footnote{In the following, all the wavelengths will be given in air.}, we found F  abundances on average $\\sim  0.8$ dex  lower  than those  in  JSL. We  discuss the  possible reasons  of this discrepancy  and show  that the  new F  values nicely agree with the state of the art of AGB nucleosynthesis models. ",
        "conclusions": "Fluorine abundances  are derived in  three Galactic AGB carbon stars  of near solar metallicity using the R9 (1-0) HF molecular line and the state of the art of model atmospheres for cool C-rich objects. We show that this line is not very much affected by atomic and molecular blends and therefore is probably the best line to derive F abundances in cool carbon-rich objects. Our results show that the fluorine  abundance is systematically reduced in AGB carbon stars by $\\sim 0.8$ dex with respect to the sole previous analysis in this kind of stars by Jorissen, Smith \\& Lambert (1992). For the star TX Psc, where we find a significant F enhancement, the new F abundance agrees nicely with recent nucleosynthesis models of low mass AGB stars of near solar metallicity. We discuss the reasons of the discrepancy with previous measurements and we conclude that  most probably  blends with C-bearing molecular lines were not properly taken into account in the analysis by JSL. This systematic  reduction of  the F  abundance in AGB carbon stars, if confirmed in a  larger sample of stars, could  reduce the present difficulty to  explain the largest  [F/Fe,O] ratios previously found in these  stars. In  particular, it would  not be necessary to  search for other  nuclear chains and/or  non standard mixing mechanism  affecting the  production/destruction of F  in AGB stars. Our results however, do  not discard AGB stars as significant F producers, the only way to determine the real contribution of these stars to the  F content in the Universe  being the interplay between accurate observational data and theoretical modelling."
    },
    "0812/0812.3790_arXiv.txt": {
        "abstract": "The magnetorotational instability (MRI) plays an essential role in the formation of stars and black holes. By destabilizing hydrodynamically stable Keplerian flows, the MRI triggers turbulence and enables outward transport of angular momentum in accretion discs. We present the results of a liquid metal Taylor-Couette experiment under the influence of helical magnetic fields that show typical features of MRI at Reynolds numbers of the order 1000 and Hartmann numbers of the order 10. Particular focus is laid on an improved experiment in which split end caps are used to minimize the Ekman pumping. ",
        "introduction": "The basic part of PROMISE is a cylindrical containment vessel V made of copper (see Fig. 1). The inner wall of the vessel V is 10 mm thick, and extends in radius from 22 to 32 mm; the outer wall is 15 mm thick, extending from 80 to 95 mm.  This vessel is filled with the eutectic alloy Ga$^{67}$In$^{20.5}$Sn$^{12.5}$ which is liquid at room temperatures. The vessel V, which is completely made of copper in the first version of the experiment (henceforth called PROMISE 1), is fixed, via an aluminum spacer D, on a precision turntable T; the outer copper cylinder of the vessel represents the outer cylinder of the Taylor-Couette cell. The inner copper cylinder I of the Taylor-Couette flow is fixed to an upper turntable, and is immersed into the liquid metal from above. It has a thickness of 4 mm, extending in radius from 36 to 40 mm, thus leaving a gap of 4 mm between this immersed cylinder I and the inner wall of the containment vessel V.  The actual Taylor-Couette cell extends in radial direction over a cylindrical gap of width $d=r_{\\rm o}-r_{\\rm i}=40$ mm, and in axial direction over the liquid metal height $h=400$ mm, resulting in an aspect ratio of 10.   \\begin{figure}[t] \\begin{center} \\unitlength=\\textwidth \\includegraphics[width=0.95\\textwidth]{fig2.eps} \\caption{(a) The asymmetric end cap configuration in PROMISE 1 is replaced by a symmetric one in PROMISE 2. U - Upper end cap fixed to the frame, L - Lower copper end cap rotating with outer cylinder, I -  Inner rings, R - Outer rings, TF - Fixed ultrasonic transducer, TR - Rotating ultrasonic transducer. (b) and (c) Axial velocity perturbation for $f_{\\rm i}=0.06$ Hz (i.e. $Re=1779$), $f_{\\rm o}=0.0162$ (i.e. $\\mu=0.27$), $I_{\\rm coil}=50$ A (i.e. Ha=7.9). (b) PROMISE 1  with $I_{\\rm rod}=6000$ A (i.e. $\\beta=5.9$). (c) PROMISE 2 with $I_{\\rm rod}=7000$ A (i.e. $\\beta=6.9$).} \\end{center} \\end{figure}   In the PROMISE 1 configuration, the upper endplate is a non-rotating Plexiglass lid P fixed to the frame F. The bottom, however, is simply part of the copper vessel V, and rotates with the outer cylinder. With respect to both their rotation rates and electrical conductivities, there is thus a clear asymmetry in the end caps.  Figure 2a shows the changes made for PROMISE 2 to improve this situation: First, both end caps are made of insulating material in order to avoid short-circuiting of currents along the copper end cap at the bottom (those currents had been shown to be dangerous by possibly changing the rotation profile via azimuthal forces \\cite{JACEK1,JACEK3}). Second, both the upper and the lower end caps are split into two rings, the inner rotating with the inner cylinder and the outer rotating with the outer cylinder. In \\cite{JACEK2} it had been shown that this splitting yields a minimization of the Ekman pumping if the position of the splitting is at 0.4 of the gap width $d$. Third, the co-rotation of the two rings with one of the cylinders made it necessary to change the signal path of the ultrasonic transducers.  These two transducers provide full profiles of the  axial velocity along the beam-lines parallel to the axis of rotation. While in PROMISE 1 they are inserted in the upper Plexiglass end cap that is fixed to the laboratory frame, in PROMISE 2 they must be connected via a sliding contact to the signal processing computer (DOP 2000).  The configuration of magnetic fields is the same for PROMISE 1 and PROMISE 2. An axial magnetic field in the order of 10 mT is produced by a double-layer coil with 76 windings (C in Fig. 1). The omission of windings at two symmetric positions close to mid-height, as seen in Fig. 1a, was motivated by a coil optimization to maximize the homogeneity of the axial field  throughout the fluid volume. The coil is fed by currents up to 200 A, beyond which a significant heating of the coil sets in. The azimuthal field, also in the  order of 10 mT, is generated by a current through a water-cooled copper rod R of radius 15 mm. The special power supply (PS in Fig. 1c) for this axial current is capable of delivering up to 8000 A. ",
        "conclusions": "We have obtained experimental evidence for the existence of the MRI in current-free helical magnetic fields. The symmetrization of the axial boundary conditions and the use of split end caps in PROMISE 2 has led to a strong reduction of the Ekman pumping and hence to an avoidance of artefacts in the radial jet flow region. The MRI wave extends clearly beyond the Rayleigh line, and its behaviour is in good correspondence with both 2D simulations and 1D simulations for the absolute instability, but in stark contrast with 1D simulations for the convective instability. This indicates that the observed MRI wave is indeed a global instability and not only a noise triggered convective instability as claimed recently \\cite{LIU3}. Further dependencies of the MRI on parameters like $Re$, $\\beta$,  and $Ha$, as well as their comparison with numerical predictions will be published elsewhere.  \\Thanks{We are grateful to J\\=anis Priede for many fruitful discussions on MRI, and for sharing with us his results for the convective and absolute instability. This work was supported by the German Leibniz Gemeinschaft, within its ''Senatsausschuss Wettbewerb'' (SAW) programme.}"
    },
    "0812/0812.0793_arXiv.txt": {
        "abstract": "This is a report on the findings of the extragalactic science working group for the white paper on the status and future of TeV gamma-ray astronomy. The white paper was commissioned by the American Physical Society, and the full white paper can be found on astro-ph (arXiv:0810.0444). This detailed section discusses extragalactic science topics including active galactic nuclei, cosmic ray acceleration in galaxies, galaxy clusters and large scale structure formation shocks, and the study of the extragalactic infrared and optical background radiation. The scientific potential of ground based gamma-ray observations of Gamma-Ray Bursts and dark matter annihilation radiation is covered in other sections of the white paper. \\vspace*{1.5cm} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.0046_arXiv.txt": {
        "abstract": "The Orion star formation complex is the nearest region of on-going star formation that continues to produce both low and high mass stars. Orion is discussed in the larger context of star formation in the Solar vicinity over the last 100 Myr.  The Orion complex is located on the far side of the Gould's Belt system of clouds and young stars through which our Solar system is drifting.  A review is given of the overall structure and properties of the Orion star forming complex, the best studied OB association.  Over the last 12 Myr, Orion has given birth to at least ten thousand stars contained in a half dozen sub-groups and short-lived clusters.  The Orion OB association has been the source of several massive, high-velocity run-away stars, including $\\mu$ Columbae and AE Auriga.  Some of Orion's most massive members died in supernova explosions that created the 300 pc diameter Orion / Eridanus super-bubble whose near wall may be as close as 180 pc.  The combined effects of UV radiation, stellar winds, and supernovae have impacted surviving molecular clouds in Orion.  The large Orion A, IC~2118 molecular clouds and dozens of smaller clouds strewn throughout the interior of the superbubble have cometary shapes pointing back towards the center of the Orion OB association.  Most are forming stars in the compressed layers facing the bubble interior. ",
        "introduction": "Orion is the best studied region of star formation in the sky. Its young stars and gas provide important clues about the physics of star formation, the formation, evolution, and destruction of star forming clouds, the dynamics and energetics of the interstellar medium, and the role that OB associations and high mass stars play in the cycling of gas between the various phases of the ISM.  In this chapter, we start with an overview of the Solar vicinity that establishes the context in which nearby star forming regions must be understood.  Then we discuss the Orion complex of star forming regions as an example of star formation in an OB association. ",
        "conclusions": ""
    },
    "0812/0812.2872_arXiv.txt": {
        "abstract": "\\noindent We derive the redshift and the angular diameter distance in rotationless dust universes which are statistically homogeneous and isotropic, but have otherwise arbitrary geometry. The calculation from first principles shows that the Dyer-Roeder approximation does not correctly describe the effect of clumping. Instead, the redshift and the distance are determined by the average expansion rate, the matter density today and the null geodesic shear. In particular, the position of the CMB peaks is consistent with significant spatial curvature provided the expansion history is sufficiently close to the spatially flat \\LCDM model. ",
        "introduction": "\\label{sec:intro}  \\para{Observations and clumpiness.}  Observations of the universe are inconsistent with homogeneous and isotropic cosmological models with ordinary matter and standard gravity (meaning matter with non-negative pressure and gravity that is described by the four-dimensional Einstein-Hilbert action). Observational results are usually expressed in terms of their interpretation in the context of the homogeneous and isotropic Friedmann-Robertson-Walker (FRW) models, in which there is a one-to-one relationship between the expansion rate and the distance (given the spatial curvature). Analysed this way, observations imply that the expansion has accelerated in the past few billion years \\cite{FRWexp, Cattoen}. A model-independent statement would be that the observed distances at late times are a factor of about 2 higher than expected in FRW models with ordinary matter and gravity. This is usually taken as evidence for exotic matter with negative pressure or a modification of gravity. In particular, the homogeneous and isotropic, spatially flat model \\LCDM model fits observations of the distance scale and the expansion rate well by introducing vacuum energy. (We will refer to all homogeneous and isotropic models which contain only dust and vacuum energy as $\\Lambda$CDM, whatever the spatial curvature.) However, the universe is known to be far from exact homogeneity and isotropy at late times due to the formation of non-linear structures. Before concluding that new physics is needed, it is necessary to quantify the effect of clumpiness on the observations.  The influence of inhomogeneity and/or anisotropy on the average evolution was first mentioned in \\cite{Shirokov:1963} and was discussed in detail in \\cite{fitting} under the name ``fitting problem''. The effect on the expansion rate is also known as backreaction \\cite{Buchert:1995, Buchert:1999, Buchert:2001}; see \\cite{Ellis:2005, Rasanen:2006b, Buchert:2007} for reviews. The possibility that backreaction could lead to accelerated expansion and explain the observations without new physics was first considered in \\cite{Wetterich:2001, Schwarz} (and briefly mentioned in \\cite{Buchert:2000, Tatekawa:2001}). Both metric and matter perturbations were taken into account and the observables were correctly identified in \\cite{Rasanen}, where first order perturbation theory was expanded to second order, and this was extended to a consistent second order calculation in \\cite{Kolb:2004a}. In \\cite{Kai:2006, Rasanen:2006a, Rasanen:2006b} it was explained with toy models that the physical reason for average acceleration is that faster expanding regions come to occupy a larger fraction of the volume. Accelerated expansion has also been explicitly demonstrated with the exact Lema\\^{\\i}tre-Tolman-Bondi (LTB) solution \\cite{Kai:2006, Chuang:2005, Paranjape:2006a}. A semi-realistic model with an evolving distribution of non-linear structures was studied in \\cite{Rasanen:2008a, peakrevs}, and it was found that the expansion rate rises (relative to the homogeneous and isotropic case) by the right order of magnitude around the right time, some billions to tens of billions of years, though not rapidly enough to correspond to acceleration. The model involved several approximations, and a more careful treatment would be needed for detailed comparison with observations.  However, most observations, including those of the cosmic microwave background (CMB) \\cite{Komatsu:2008} and type Ia supernovae \\cite{SNobs}, do not directly measure the expansion rate, but rather cosmological distances and redshifts, which are defined in terms of light propagation. The few measurements that are sensitive to the expansion rate independent of the distance scale are those of the local Hubble parameter \\cite{Hubble, Jackson:2007}, the ages of passively evolving galaxies as a function of redshift \\cite{ages}, the Integrated Sachs-Wolfe (ISW) effect \\cite{Ho:2008, Giannantonio:2008} and the growth rate of matter fluctuations \\cite{lineargrowth}. Baryon acoustic oscillations depend on a mixture of the expansion rate and distance \\cite{BAO1, BAO2}.  \\para{The distance and the expansion rate.}  In addition to changing the expansion of the universe, non-linear structures affect the relationship between the expansion rate and light propagation. In a general spacetime, there is no direct relationship between the expansion rate and the distance scale, and it would in principle be possible to explain the observations without accelerated expansion. For example, in models where we are located near the center of an untypically large spherical void, the distances can be consistent with the observations, but generally there is no acceleration \\cite{Enqvist:2006, Bolejko:2008c}. (See \\cite{Enqvist:2007} for a review, \\cite{Rasanen:2006b} for more references, and \\cite{Caldwell:2007, Uzan:2008, GarciaBellido:2008a, Bolejko:2008b, Zibin:2008b, GarciaBellido:2008b} for observational constraints related to the inhomogeneity of these models.) Even if the large local void models are not realistic, studying them with the exact LTB solution has established unambiguously that inhomogeneities with density contrast of order unity and sizes smaller than the horizon can have a large impact on the distance scale. This is in contrast to perturbation theory arguments based on the amplitude of metric perturbations in the longitudinal gauge.  The speculative possibility of an unexpectedly large local void aside, the observed universe seems to be statistically homogeneous and isotropic on large scales, with a homogeneity scale of around 100 Mpc \\cite{Hogg:2004, Pietronero} (though see \\cite{morphology, SylosLabini}). (For discussion of statistically homogeneous and isotropic but locally clumpy dust universes, see \\cite{Rasanen:2006b, Rasanen:2008a}.) The relevant question is then what is the effect of such a distribution of non-linear structures on light propagation.  It was conjectured in \\cite{Rasanen:2008a} that in a statistically homogeneous and isotropic dust universe, light propagation can be treated in terms of the overall geometry (meaning the average expansion rate and average spatial curvature) if the structures are realistically small and the observer is not in a special location. Such a conjecture is in agreement with various studies of light propagation (see \\cite{Rasanen:2008a} for an overview and references), and it is also suggested by the fact that different observations are well explained in terms of the evolution of a single scale factor. However, until now there had been no proof of the conjecture, and it was not known whether additional conditions are necessary in addition to statistical homogeneity and isotropy.  In the present work, we clarify the relationship between the expansion rate and the distance scale in statistically homogeneous and isotropic dust universes which may contain non-linear structures. We relate the redshift and the distance to the dust geometry, and confirm that light propagation can be expressed in terms of average geometrical quantities, up to a term related to the null geodesic shear. In fact, the null shear aside, the average expansion rate (and the matter density today) is sufficient to determine the distance, and the spatial curvature enters only via its effect on the expansion rate. This implies that a clumpy model can be consistent with the observed position of the CMB acoustic peaks even when there is significant spatial curvature. The result also shows that the Dyer-Roeder prescription of multiplying the matter density by a smoothness factor does not correctly describe the effect of clumping.  In section 2 we set up the dust geometry, in section 3 we relate the redshift and the distance to the average geometry, and in section 4 we discuss and summarise our results. ",
        "conclusions": "\\label{sec:disc}  \\subsection{Observational and modelling issues}  \\para{Spatial curvature and the position of the CMB peaks.}  The position of the CMB acoustic peaks is often considered to be a measurement of spatial curvature. The reason is that the peak location in multipole space corresponds to the apparent size of the last scattering sound horizon, which provides a measure of the angular diameter distance to the last scattering surface at $z\\approx1100$ \\cite{Durrer:2008} (page 99). The observed position of the peaks is consistent with spatial flatness in a \\LCDM universe \\cite{Komatsu:2008}.  However, the peak position does not imply spatial flatness, even in a FRW universe. It is clear from the way $H$ and $K$ enter the distance \\re{DAclosed} that for any value of $K$, it is possible to adjust $H$ to compensate for the spatial curvature so as to keep the distance fixed. For example, the peak position is consistent with a FRW universe with large positive spatial curvature \\cite{Spergel:2006}. Such a model is not viable due to other constraints, such as the value of the Hubble parameter today. By replacing the vacuum energy with exotic matter with a time-dependent equation of state, it is possible to do a similar adjustment and allow spatial curvature without changing the expansion history as radically \\cite{Ichikawa, Clarkson:2007a}.  In the FRW case, it is true that given the expansion history $H(z)$, the peak position provides a measurement of the spatial curvature. The root of this argument is the relation \\re{DAclosed}, which expresses the distance in terms of the expansion rate and the spatial curvature. In a clumpy universe the distance is instead completely fixed by the expansion rate and the matter density (as well as the null shear), according to \\re{DAeqav}\\footnote{Note that we are comparing a clumpy universe with only dust to a FRW universe with arbitrary matter content. If we allowed other matter in the clumpy situation, the distance would also depend on the non-trivial $\\av\\rho+\\av{p}$. In this case, the Buchert equations would also be more complicated \\cite{Buchert:2001}.}. As \\re{DAeqavR} shows, the expression \\re{DAclosed} is inapplicable due to the non-trivial evolution of the spatial curvature as well as the fact that clumping contributes to the expansion rate via $\\sQ$. In terms of the density parameters defined by dividing the expansion law \\re{Ham} by $3 H^2$, the density of matter and the density of curvature do not sum to unity, because of the contribution of $\\sQ$ \\cite{Buchert:1999, Buchert:2003}. In a FRW model, any additional contribution that changes $H$ enters the distance also via $\\rho+p$ (which is related to the spatial curvature), but that is not the case here. In the FRW case with arbitrary matter, the spatial curvature is always proportional to $(1+z)^2$, and the evolution of $\\rho+p$ is not fixed, while in a clumpy dust universe, $\\rho+p$ is always proportional to $(1+z)^3$, but the evolution of the spatial curvature is complicated.  It is somewhat trivial that the spatial curvature enters the distance only via its effect on the expansion rate, since the spatial curvature can be written in terms of $H$, $\\Hdot$ and $\\av\\rho\\propto (1+\\av{z})^3$ using the Buchert equations \\re{Ray} and \\re{Ham}. However, it is not obvious that the equation for the average distance \\re{DAeqav} depends only on $H$ and the matter density, or that the dependence on $H$ is the same as in the FRW case, so that we recover the \\LCDM equation (because $\\rho+p$ is the same in both cases).  In summary, the CMB peak position can be consistent with large spatial curvature in a clumpy model. All that is required is that the expansion history and the matter density today is sufficiently close to that of the spatially flat \\LCDM model. (Since the Hubble rate enters via an integral, significant variation in $H(\\av{z})$ is still allowed \\cite{trans}.)  \\para{The effective equations of state.}  Observations of distances are typically analysed in terms of an effective equation of state in a FRW model. (Regarding the dependence of the results on the adopted parametrisation for the equation of state, see \\cite{trans}.) The equation of state determines the evolution history $H(z)$, which then gives the distance $D_A(z)$ via \\re{DAclosed}. If we want to express the evolution of a clumpy dust model this way, there are two different effective equations of state, because the relation between $H$ and $\\av{D_A}$ is different than in FRW models.  For the expansion rate, we can define an effective equation of state $w_H(z)$ such that a FRW model with this equation of state would have the same expansion history as the clumpy model. Formally this is done by writing  the spatial curvature and the backreaction variable $\\sQ$ as a single component in the Buchert equations \\re{Ray}--\\re{cons} \\cite{Buchert:2001, morphon}. For comparison to distance observations, we should introduce another equation of state $w_D(z)$, defined so that the resulting $H$ reproduces the real distance function $\\av{D_A}$ when plugged into the FRW distance formula \\re{DAclosed}. In both cases, we have to make a choice for the spatial curvature of the FRW fitting model. The simplest choice, which does not involve any loss of generality, is to take the FRW model to be spatially flat. In general, $w_H\\neq w_D$, so the effect of clumpiness cannot be treated just as an effective source in the FRW equations. This limits the usefulness of an effective description of backreaction in terms of a scalar field \\cite{morphon} (also, $w_H$ can violate the null energy condition, unlike the equation of state of a scalar field \\cite{Rasanen:2008a}).  Until calculations of the impact of structure formation are accurate to more than an order of magnitude \\cite{Rasanen:2008a, peakrevs}, it is not known how large the expected difference between $w_H$ and $w_D$ is. However, even without a theoretical prediction, it is possible to test the null hypothesis that the equations of state inferred separately from observations of the expansion rate and distance do not show any difference. This is the essence of the FRW consistency check proposed in \\cite{Clarkson:2007b}.  At the moment, while distances are measured relatively well, there are few observations probing the expansion rate as a function of redshift independent of the distance scale. The ages of passively evolving galaxies provide an interesting way to measure the expansion history, but at the moment the constraints are rather weak \\cite{ages}\\footnote{In a space which is not statistically homogeneous and isotropic, this data does not measure purely the expansion rate, since the shear also contributes to the redshift \\re{zt}. This was used to constrain local void models in \\cite{Bolejko:2008b}.}. Another measure is provided by baryon acoustic oscillations, which are sensitive to a combination of expansion rate and distance \\cite{BAO2}. (This was used to constrain local void models in \\cite{GarciaBellido:2008b}.) There does appear to be some discrepancy between the observations of the luminosity distance of type Ia supernovae and measurements of baryon acoustic oscillations, but only at the 2$\\sigma$ level \\cite{More:2008}, so it is not statistically significant. When analysing the data in the context of FRW models, this discrepancy would be interpreted as a violation of the reciprocity relation $D_L=(1+z)^2 D_A$ instead of the FRW relation \\re{DAclosed} between $H$ and $D_A$.  The expansion rate also has a role in the ISW effect, i.e. deviations of the redshift from the mean value $\\av{z}$ in different directions, as well as in the growth of density perturbations \\cite{lineargrowth}. The ISW signal is slightly higher than the spatially flat \\LCDM prediction, but the difference is not statistically significant (it is evaluated as 2$\\sigma$ in \\cite{Ho:2008} and 1$\\sigma$ in \\cite{Giannantonio:2008}). Neither the ISW signal nor the growth factor can be used at present to put accurate constraints on $H$ as a function of redshift.  \\para{Average observables.}  The equation \\re{DAeqav} determines the average angular diameter distance as a function of the average redshift, given the average expansion rate. The averages are taken on the hypersurface of constant proper time. However, we observe the redshift and the distance only in one fixed location. Nevertheless, for practical purposes, the averages $\\av{D_A}$ and $\\av{z}$ do correspond to directly observable quantities.  In order to model what could be observed in principle, we would have to know the details of the structures along each line of sight to calculate the relation between $D_A$ and $z$ for each direction. As noted earlier, the distance-redshift relation cannot in general be expressed as a function $D_A(z)$, because several values of $D_A$ can correspond to the same redshift. In practice differences between redshifts corresponding to the same distance are likely to be smaller than the observational resolution, except for nearby sources (or the CMB, for which the redshift is very accurately measured), because the regions which are collapsing (or have strong negative shear in the direction of the null geodesic) are small. Using the redshift integral \\re{zt}, a naive estimate of the blueshift due to a region one Mpc across which is collapsing with a rate of the present-day Hubble parameter is $10^{-3}$. As long as the variations in $D_A$ and $z$ are smaller than the observational errors, we can safely say that the averages correspond to the observed quantities. Observationally, the variation of the CMB peak position with direction is known to be small \\cite{Hansen:2004}. Note that standard CMB analysis also predicts only an ensemble average, and that in practice cosmological observations are often analysed using an average over the full sky, which by statistical homogeneity and isotropy corresponds to the average over the spatial hypersurface.  \\para{The Dyer-Roeder approximation.}  An approach where the effect of clumping on light propagation is modeled assuming that the light rays encounter only a fraction of the mass in the universe was introduced by Zel'dovich \\cite{Zeldovich:1964} and is known as the Dyer-Roeder approximation \\cite{Dyer}. In this prescription, the FRW equation \\re{DAeqFRW} for the distance is modified by multiplying the matter density by a constant $\\alpha$, which varies between 0 and 1, corresponding to a universe where the lines of sight are completely empty or completely filled with matter, respectively. The smoothness parameter $\\alpha$ was generalised to a function of redshift in \\cite{Linder} to account for the evolution of structures, and change of the expansion rate due to clumping was added to the equation in \\cite{Mattsson:2007b}.  According to our result \\re{DAeqav}, clumping is irrelevant for the contribution of the matter density, which is always proportional to $(1+\\av{z})^3$, with no extra prefactor. The reason is that mass is conserved, so if the line of sight goes through an underdense region somewhere, it must correspondingly go through overdensities elsewhere when considering distances of the homogeneity scale or larger, as discussed in connection with equations \\re{zHI} and \\re{muHI}. The Dyer-Roeder parameter $\\alpha$ is always unity, and clumping enters instead by changing the expansion rate. In addition, there is the null shear term; if interpreted as an effective, redshift-dependent $\\alpha$, it would correspond to $\\alpha>1$, contrary to the Dyer-Roeder case. In summary, the Dyer-Roeder prescription does not correctly describe the effect of clumping.  One possible caveat is that we have not taken into account the possibility that the observer or the sources could be in a special location, or that observations could be made along special lines of sight. For example, we would expect supernovae to be preferentially located in very overdense and thus highly untypical regions. The possibility that observations might be biased towards empty lines of sight has been brought up in \\cite{Mattsson:2007b, Weinberg:1976}.  The effect of the location of the observer is likely to be small, because the deviation from the mean is significant only over regions which are small compared to the overall distance travelled by the light, and the amplitude of the deviation is typically not correspondingly large (except in very special locations, such as near a black hole). This is another way of saying that the homogeneity scale is small. For the same reason, the location of the sources is not expected to have a large impact, though there could be a secular effect, as the degree to which the source locations are untypical could evolve with redshift. In any case, cosmological observations rely on various different sources, not only supernovae. Since different observations roughly agree, these kind of selection effects must be subdominant. This is also an argument against large effects due to special lines of sight used in observations. In particular, the CMB covers the full sky, so it is not subject to this kind of bias (apart from some uncertainty in the direction of the Galactic plane). Furthermore, since most cosmological observations (apart from nearby objects) are made over scales much larger than the homogeneity scale, the variation in the density along lines of sight should be small, and empty lines of sight should be very rare. Note that a clear line of sight is not necessarily empty, because most of the dust is dark matter, which is transparent. (See \\cite{More:2008} and references therein for more on cosmic transparency.)  \\subsection{Conclusion}  \\para{Summary.}  It was conjectured in \\cite{Rasanen:2008a} that light propagation in a statistically homogeneous and isotropic dust universe which may contain non-linear structures can be expressed in terms of average geometrical quantities (namely the expansion rate and the spatial curvature), assuming that the observer is not in a special location, and that structures have realistically small sizes. As reviewed in \\cite{Rasanen:2008a}, the literature on light propagation is mostly in agreement with this statement, but there had been no proof thus far.  We have now derived the equation for the angular diameter distance, assuming that the dust universe is statistically homogeneous and isotropic as well as rotationless, and that structures are small and their distribution evolves slowly compared to the time it takes for light to cross the homogeneity scale. It follows from these assumptions that the distance can be expressed in terms of the average dust geometry, apart from a term related to the null geodesic shear. Of the average geometry, the only term that is required is the average expansion rate, along with the present value of the matter density. In particular, the spatial curvature enters only via its effect on the expansion rate, unlike in Friedmann-Robertson-Walker (FRW) models with arbitrary matter content. Therefore, significant spatial curvature is not necessarily inconsistent with the position of the cosmic microwave background acoustic peaks. If the expansion history and present-day matter density of a clumpy model is close to that of the spatially flat \\LCDM FRW model, the peak position will also be near the \\LCDM case, and thus consistent with observations. This is important, because accelerated expansion due to structure formation involves large negative spatial curvature \\cite{Rasanen:2005, Rasanen:2006b, Rasanen:2008a}. The result also shows that the clumping is not correctly described by the Dyer-Roeder approximation, which changes the evolution of the matter density (which is in fact fixed) and misses the change of the expansion rate.  To complete the proof that the redshift and the distance can be expressed in terms of the average geometry alone, it would be necessary to show that the contribution of the null shear can be neglected. Observationally, the shear is known to be small \\cite{Munshi:2006}, but this should be shown to follow from statistical homogeneity and isotropy (or, if this is not the case, the required additional assumptions should be identified).  Since the distance-redshift relation is determined by the average expansion rate, we should calculate the average expansion rate in a realistic model to evaluate the effect of structure formation on cosmological observations. Using the Buchert equations \\cite{Buchert:1999}, this backreaction can be determined from purely statistical quantities (the variance of the expansion rate and the average dust shear). Evaluating these quantities in a realistic model is a challenging task, especially as the backreaction is a general relativistic effect related to spatial curvature and has no counterpart in Newtonian gravity \\cite{Buchert:1995, Buchert:1999, Buchert:2007, Rasanen:2008a}\\footnote{Regarding the differences between Newtonian gravity and the weak field limit of general relativity, see \\cite{Ellis:1971, Ellis, Senovilla:1997, Szekeres}.}. A first step was taken in \\cite{Rasanen:2008a, peakrevs}, where the average expansion rate was calculated in a semi-realistic model with an evolving distribution of structures. To compare with observations of distance and the expansion rate in detail, a more rigorous treatment is needed, with well-quantified errors. After that, the calculation of fluctuations around the average should follow, in order to evaluate the Integrated Sachs-Wolfe effect and the growth of density perturbations.  Before detailed analytical results on the effect of structure formation are worked out, it is possible to make general tests. The relation between the expansion rate and the distance scale in clumpy models is different than in FRW universes. Therefore, observations which probe the expansion rate and the distance separately can be used to test the null hypothesis that the two are related by the FRW consistency condition, as proposed in \\cite{Clarkson:2007b}. Comparing observations of type Ia supernovae, baryon acoustic oscillations and the ages of passively evolving galaxies seems promising in this respect.  \\ack  I thank Thomas Buchert for discussions and comments on the manuscript and Ruth Durrer, Teppo Mattsson and Subir Sarkar for helpful discussions.  \\\\  \\setcounter{secnumdepth}{0}"
    },
    "0812/0812.3209_arXiv.txt": {
        "abstract": "We use recently observed data:  the 192  ESSENCE type Ia supernovae (SNe Ia), the  182 Gold SNe Ia, the 3-year WMAP, the SDSS baryon acoustic peak, the X-ray gas mass fraction in clusters  and the observational $H(z)$ data to constrain models of the accelerating universe. Combining the 192 ESSENCE data with the observational $H(z)$ data to constrain a parameterized deceleration parameter, we obtain the best fit values of transition redshift and current deceleration parameter $z_{T}=0.632^{+0.256}_{-0.127}$, $q_{0}=-0.788^{+0.182}_{-0.182}$. Furthermore, using $\\Lambda$CDM model and two model-independent equation of state of dark energy, we find that the combined constraint from the 192 ESSENCE data and other four cosmological observations gives smaller values of $\\Omega_{0m}$ and $q_{0}$, but a larger value of $z_{T}$ than the combined constraint from the 182 Gold data with other four observations. Finally, according to the Akaike information criterion it is shown that the recently observed data equally supports three dark energy models: $\\Lambda$CDM, $w_{de}(z)=w_{0}$ and $w_{de}(z)=w_{0}+w_{1}\\ln(1+z)$. ",
        "introduction": "$}   {\\small {~~~}}~The type Ia supernovae (SNe Ia) investigations \\cite{1Riess}, the cosmic microwave background(CMB) results from WMAP \\cite{2Spergel} observations, and  surveys of galaxies \\cite{3Pope} all suggest that the expansion of present universe is speeding up rather than slowing down. If one considers that the evolution of universe complys with the standard cosmology, the accelerated expansion of the present universe is usually attributed to the fact that dark energy (DE) is an exotic component with negative pressure. Many kinds of DE models have already been constructed such as $\\Lambda$CDM \\cite{4Weinberg}, quintessence \\cite{5Ratra}, phantom \\cite{6Caldwell}\\cite{07Setarea}, generalized Chaplygin gas (GCG) \\cite{7Kamenshchik}\\cite{7zhu}, modified Chaplygin gas \\cite{8Benaoum}\\cite{9Lu}\\cite{9LuMCG2}, quintom \\cite{10Feng}, holographic dark energy \\cite{11Li}\\cite{12Zhang}, agegraphic dark energy\\cite{07Cai}\\cite{08Zhang},  and so forth.  On the other hand, to remove the dependence of special properties of extra energy components, a parameterized  equation of state (EOS) is assumed for DE. This is also commonly called the model-independent method. The parameterized EOS of dark energy $w_{de}$ which is popularly used in parameter best fit estimations, describes the possible evolution of DE. For example, $w_{de}(z)=w_{0}$= const \\cite{13Hannestad}, $w_{de}(z)=w_{0}+w_{1} \\ln(1+z)$ \\cite{14Gerke}. The parameters $w_{0}$, $w_{1}$ are obtained by the best fit estimations from cosmic observational datasets.  Recently, the 192 ESSENCE SNe Ia data \\cite{15Wu} was compiled by Ref. \\cite{16Davis} using the four sets of supernova (SN) data:  60 ESSENCE SNe \\cite{21Wood-Vasey},  57 Supernova Legacy Survey (SNLS) SNe \\cite{39Astier},  45 nearby SNe \\cite{1Riess}\\cite{40Hamuy}\\cite{41Jha},  and 30 new SNe at high redshift ($0.216 \\leq z \\leq 1.755$) recently discovered by the Hubble Space Telescope (HST) and classified as $^{\"}$Gold$^{\"}$ SNe by Ref. \\cite{17Riess}. The ESSENCE project \\cite{21Wood-Vasey} is a ground-based survey that design to detect about 200 SNe Ia in the range of $z = 0.2-0.8$ to measure the EOS of DE to better than 10 percent. The SNLS  and the nearby SNe data as the complementary cosmological probes have been refitted by \\cite{21Wood-Vasey} with the same lightcurve fitter used for the ESSENCE data. As regards the 30 HST SNe, it is necessary to perform a normalization\\cite{07042606Lazkoz}.  Ref. \\cite{16Davis} adopted the low redshift SNe that these samples had in common in order to normalize the luminosity distances of the samples, and the error in the normalization is included in the distance errors for the HST SNe \\cite{16Davis}\\cite{07042606Lazkoz}.   In Ref. \\cite{16Davis} the authors applied the 192 ESSENCE SNe Ia data, the 3-year WMAP CMB shift parameter \\cite{26Spergel}\\cite{18Wang}, the baryon acoustic oscillation (BAO) peak from Sloan Digital Sky Surver (SDSS) \\cite{19Eisenstein} to constrain the current values of EOS of DE $w_{0de}$ and dimensionless matter density $\\Omega_{0m}$ by using several model-independent EOS of DE. However, some other cosmological quantities  such as transition redshift $z_{T}$ and current deceleration parameter $q_{0}$ were not discussed. On the other hand, we know that the 182 Gold SNe Ia data \\cite{17Riess} is compiled from five distinct subsets defined by the group or instrument that discovered and analyzed the corresponding SNe data. These subsets are \\cite{28Nesseris}\\cite{30website182subset}: the High $z$ Supernova Search Team (HZSST) subset (41 SNe) \\cite{1Riess}\\cite{HZSST2}\\cite{HZSST3}\\cite{HZSST4}, the Supernova Cosmology Project (SCP) subset (26 SNe) \\cite{SCP}, the Low Redshift (LR) subset (38 SNe) \\cite{40Hamuy}\\cite{LR2}\\cite{28Jha}\\cite{LR4}, the HST subset (30 SNe) \\cite{17Riess} and the SNLS subset (47 SNe) \\cite{39Astier}. It can be found that there are 99 SNe that are in the 192 ESSENCE data but not in the 182 Gold data\\footnote{The two sets of SNe Ia data with their subsets are shown in the Appendix. From the Appendix it can be seen that there are 93 SNe Ia in common between the 192 ESSENCE data and the 182 Gold data(i.e., from number 81 to number 173 in  TABLE IV). They include 25 nearby SNe (or 25 LR SNe), 30 HST SNe and 38 SNLS SNe.}. Furthermore, relative to the  Gold sample  Ref. \\cite{16Davis}\\cite{21Wood-Vasey} applied an updated version of the MLCS2k2 method\\footnote{The basic framework for Multicolor Light Curve Shape method to measure the luminosity distances was laid out by Ref. \\cite{28Jha}  in 2002 (i.e., MLCS2k2 method) and it has already been applied to SN Ia cosmology such as the 157 Gold  SNe Ia data \\cite{22Riess} and the 182 Gold  SNe Ia data  \\cite{17Riess}. A  new version of the MLCS2k2 was developed with an expanded training set by Ref. \\cite{41Jha} and this light-curve fitting technique has also been applied to measure the luminosity distances to  ESSENCE, SNLS and nearby SNe Ia in Ref. \\cite{21Wood-Vasey}. Because the basic MLCS2k2 algorithms were designed  by Ref. \\cite{28Jha}, Refs. \\cite{21Wood-Vasey}\\cite{41Jha} continue to refer to this updated SN distance fitter as MLCS2k2, even though its implementation, applicability, and robustness have evolved substantially since then. For more details about MLCS2k2 please see Refs. \\cite{21Wood-Vasey}\\cite{41Jha}\\cite{17Riess}\\cite{28Jha}\\cite{22Riess}\\cite{MLCS2k2website}.} to measure distances to SNe Ia for the  ESSENCE sample, incorporating new procedures for $K$-correction and extinction corrections. So, the data points are also different even though  for the same SNe in the two SN  samples\\footnote{For the case of the 192 ESSENCE data, since the error in the normalization is included in the distance errors for the HST SNe,  the data points from the   30 HST SNe are also different  between the 192 ESSENCE data and the 182 Gold data.}. Therefore, we want to know what are the differences for the constraints on cosmological quantities from these two samples of SNe Ia respectively. In this paper, by using a parameterized deceleration parameter and model-independent EOS of DE, we apply the recent cosmic observations to constrain several cosmological quantities, such as $z_{T}$, $q_{0}$, and compare the differences for them when the constraints are obtained from  the 192 ESSENCE data and  the 182 Gold data, respectively. To avoid the degeneration of DE models and get the significant constraints on cosmological quantities, we combine other observational data with these two sets of SNe data, such as the 3-year WMAP CMB shift parameter, the BAO peak from SDSS, the X-ray gas mass fraction in clusters \\cite{20Allen} and the observational $H(z)$ data from the Gemini Deep Deep Survey (GDDS) \\cite{23Abraham} and archival data \\cite{24Treu}\\cite{25Treu}.   The paper is organized as follows. In section 2, we apply recent cosmic observations to constrain models of the accelerating universe by using a parameterized deceleration parameter $q$ and model-independent EOS of DE $w_{de}$. In section 3, we use the information criterion of model selection for DE models to estimate which model for an accelerating universe is distinguish by statistical analysis of observational datasets. Section 4 is the conclusion. ",
        "conclusions": "$} \\label{section3}  {~~~~In summary, we use the  192 ESSENCE SNe Ia data and the  182 Gold SNe Ia data combined with  other observed data such as the 3-year WMAP, the BAO peak from SDSS , the  X-ray gas mass fraction in clusters  and the observational $H(z)$ data from the GDDS and archival data,  to constrain models of the accelerating universe. Concretely,  using the  192 ESSENCE SNe Ia data and the 9 observational $H(z)$ data to constrain a parameterized deceleration parameter $q(z)=\\frac{1}{2}+\\frac{a+bz}{(1+z)^{2}}$ \\cite{35Gong}\\cite{15Gong}, we obtain the best fit model parameters, $a=-1.288^{+0.275}_{-0.276}$, $b=-0.068^{+1.010}_{-0.998}$. The best fit values of transition redshift $z_{T}$ and current deceleration parameter $q_{0}$ are given as $z_{T}=0.632^{+0.256}_{-0.127}$, $q_{0}=-0.788^{+0.182}_{-0.182}$. Replacing the 192 ESSENCE data with the 182 Gold data in combined constraint, it can be found that at 1$\\sigma$ error range, this result for $q_{0}$ is almost consistent with being the same for the two combined constraints. But the result for $z_{T}$  at 1$\\sigma$ error range tends to larger value than the case of the joint analysis involving the 182 Gold data and the observational $H(z)$ data. Furthermore, it is shown that the central value of $z_{T}$ (or $q_{0}$) for the former combined constraint is larger (or smaller) than latter one. For producing the differences of these cosmological quantities between these two combined constraints, the reason maybe is caused by the different way that the SNe magnitudes are calculated for the two samples of SNe Ia.  On the other hand, since some data points in the two sets of SN data are from the different subsets, some unknown system errors from the different instruments for SNe surveys are also possible to contribute to these differences for the cosmological quantities. We also expect the more advanced probers to explore SNe in  future.   Furthermore, for the cosmological quantities $\\Omega_{0m}$, $w_{0de}$, $z_{T}$ and $q_{0}$, we compare the differences for them between the combined constraint  from the 192 ESSENCE data  with  other four cosmic observations and the 182 Gold data with other four observations. Considering DE model $\\Lambda$CDM and two model-independent EOS of dark energy, $w_{de}(z)=w_{0}$, $w_{de}(z)=w_{0}+w_{1}\\ln(1+z)$, we plot the best fit forms of deceleration parameter $q(z)$ with $1\\sigma$ error by using two sets of SNe data  with other cosmological observations. It can be seen that the cosmic acceleration could have started between the redshift $z_{T}=0.706^{+0.063}_{-0.060}$ and $z_{T}=0.774^{+0.051}_{-0.050} $ for  two  combined constraints. By comparing the two combined constraints on the DE models: $\\Lambda$CDM, $w_{de}(z)=w_{0}$, and $w_{de}(z)=w_{0}+w_{1}\\ln(1+z)$, we find that the combined constraint from  192 ESSENCE+Hub+CMB+BAO+CBF tends to smaller current value of matter density $\\Omega_{0m}$ than the constraint from 182 Gold+Hub+CMB+BAO+CBF. And it can be seen that, though at $1\\sigma$ error range the differences between two combined constraints for the values of $z_{T}$ and $q_{0}$ are not very obvious, the central values of $z_{T}$ constrained from 192 ESSENCE+Hub+CMB+BAO+CBF are bigger than the cases  of 182 Gold+Hub+CMB+BAO+CBF, and the central values  of $q_{0}$ are smaller than the cases  of 182 Gold SNe Ia+Hub+CMB+BAO+CBF. At last, it is shown that for the case of $w_{de}(z)=w_{0}=$ const, the central value of $w_{0de}$ constrained from 192 ESSENCE+Hub+CMB+BAO+CBF is surprisingly close to $\\Lambda$CDM model ($w_{de}=-1$).  Since it is interesting to estimate which model for an accelerating universe is distinguish by statistical analysis of observational datasets over many models, by applying the recent observational data to the objective information criterion  of model selection, we compare with three DE models in TABLE \\ref{table2} to assess the strength of models. It is shown that, according to the AIC though the best model is $\\Lambda$CDM for using the combined datasets of 192 ESSENCE+Hub+CMB+BAO+CBF, other two models also have the same support with the best one because the values of $\\Delta$ AIC$_{i}$ for them are in the range 0-2. For the case of the BIC we find the best fit model is $\\Lambda$CDM for both combined constraints, and the more free parameters in DE model, the  weaker support from observational data for these three DE models.        \\textbf{\\ Acknowledgments } The research work is supported by NSF (10573003), NSF (10573004), NSF (10703001), and NSF (10647110), NBRP (2003CB716300) of P.R. China."
    },
    "0812/0812.3723_arXiv.txt": {
        "abstract": "Using dark matter halos traced by galaxy groups selected from the Sloan  Digital Sky Survey Data  Release 4, we  find that about 1/4  of the faint  galaxies  ($\\rmag >-17.05$,  hereafter  dwarfs)  that  are the  central galaxies  in their own  halo are  not blue  and star  forming, as  expected in standard models of  galaxy formation, but are red.  In contrast, this fraction is  about 1/2  for  dwarf satellite  galaxies.   Many red  dwarf galaxies  are physically associated with more massive halos.  In total, about $\\sim 45$\\% of red dwarf galaxies reside in  massive halos as satellites, while another $\\sim 25$\\% have a spatial distribution that is much more concentrated towards their nearest massive  haloes than  other dwarf galaxies.   We use mock  catalogs to show  that  the  reddest   population  of  non-satellite  dwarf  galaxies  are distributed within  about 3  times the virial  radii of their  nearest massive halos.   We suggest  that  this population  of  dwarf galaxies  are hosted  by low-mass halos that have passed  through their massive neighbors, and that the same environmental  effects that  cause satellite galaxies  to become  red are also responsible for the red colors of this population of galaxies.  We do not find any  significant radial  dependence of the  population of  dwarf galaxies with the  highest concentrations, suggesting that the  mechanisms operating on these galaxies affect color more  than structure.  However, over 30\\% of dwarf galaxies are red and isolated and their origin remains unknown. ",
        "introduction": "\\label{sec:intro}  In the current paradigm of galaxy formation, it is believed that virtually all galaxies initially form  as `disks' owing to the cooling  of gas with non-zero angular momentum in  virialized dark matter haloes. This  smooth gas accretion dominates  the galactic  gas supply  and hence  the fuel  for  star formation. Galaxies that  reside in the  centers of lower  mass halos, those  with masses less than  $M_{halo} \\lta  3\\times 10^{11} \\msunh$,  accrete gas  through very efficient  cold  mode accretion,  i.e.  gas that  is  never  heated (Keres  et al. 2005,  Keres et  al. 2008).   The central galaxies  that reside  in larger halos accrete their  gas through the classic, but less  efficient, hot mode of accretion where  the gas is shock  heated to near the  virial temperature near the virial  radius and then  must cool to  be accreted by the  central galaxy. Hence the  naive expectation would be  that dwarf galaxies  should be actively star forming and blue.  However, when a small halo is accreted by a larger halo, i.e.  when it becomes a  subhalo,  the central  galaxy  that  formed in  the  small  halo becomes  a satellite galaxy and may experience a number of environmental effects that may change its  properties.  For instance,  the diffuse gas  originally associated with  the  subhalo  may  be  stripped,  thus  removing  the  fuel  for  future star-formation  (e.g.   Larson,  Tinsley  \\& Caldwell  1980).   This  process, referred to as strangulation (Balogh \\&  Morris 2000), can result in a gradual decline of the  star formation rate in the satellite  galaxy, making it redder with  the passage of  time.  If  the external  pressure is  sufficiently high, ram-pressure  stripping  may  also be  able  to  remove  the entire  cold  gas reservoir of the satellite (e.g.  Gunn \\& Gott 1972), causing a fast quenching of its  star formation.  A satellite  galaxy is also subject  to tidal heating and stripping and  galaxy harassment (Moore et al 1996),  which may also cause the  satellite to  lose  its fuel  for  star formation.   These processes  are believed  to have  played  an important  role  in the  evolution of  satellite galaxies, and to  be responsible, to a large extent,  for the relation between galaxy  properties  and  their  environment. Indeed,  satellite  galaxies  are generally  found to  be redder  and  somewhat more  concentrated than  central galaxies  with similar  stellar masses  (e.g.  van  den Bosch  et  al.  2008a; Weinmann et al.  2008; Yang et al.  2008b; c; Guo et al. 2009).  Furthermore, recent investigations  based on cosmological $N$-body simulations have shown  that a significant fraction  of dark matter haloes  that are close to, but  beyond the  virial radius  of, a more  massive neighboring  halo, are physically connected  to their neighbor.  As  shown by Lin et  al.  (2003) and more  recently  by  Ludlow  et  al.   (2008), some  low-mass  halos  are  {\\it physically} associated  with more massive halos,  in the sense  that they were once subhalos  within the virial radii  of these more  massive progenitors and have subsequently been ejected.  This  population of halos was found to extend beyond three times the virial radii  of their host halos, and represents about $10\\%$ of the entire population of low-mass halos (Wang, Mo \\& Jing 2008).  If galaxies have managed to form in the progenitors of these ejected halos, it is likely that  the same environmental processes operating  on satellite galaxies may also  have affected the properties  of these galaxies.   In particular, we would  expect the  presence of  a red  population of  faint galaxies  that are closely associated with massive halos that once hosted them.  Galaxies are observed to be bimodal in the color-magnitude plane: red galaxies with  very little  star formation  (the red  sequence) and  blue  star forming galaxies that are typically disky (the blue cloud) (e.g. Kauffman et al. 2003, Baldry  et  al.  2004)  Extrapolating  the observed  division  line  (Yang  et al.  2008a) to  dwarf galaxies  we surprisingly  find that  for  central dwarf galaxies in the SDSS, with $r$-band magnitudes between -14.46 and -17.05, just over 1/4 are red.   In this paper we will investigate the  nature of these red dwarf galaxies.   Quantifying the spatial  distribution of this  population of galaxies is  clearly important, because it  allows us to  determine whether or not they  can be  explained as  a population of  satellite galaxies  that were ejected from larger halos.  In this  paper, we use the galaxy  group catalogue constructed by  Yang et al. (2007)  from  the  Sloan  Digital   Sky  Survey  Data  Release  4  (SDSS  DR4; Adelman-McCarthy  \\etal  2006) to  study  the  distribution  of central  dwarf galaxies around massive halos.  The structure of this paper is as follows.  In \\S\\ref{sec_data} we briefly describe the  criteria used to select galaxies and galaxy  groups.  In \\S\\ref{sec_analyze}  we study  the radial  distribution of dwarf  galaxies around  their nearest  neighbor  halos and  its dependence  on galaxy color  and concentration. Some  systematic effects that may  change our results are discussed in \\S\\ref{sec_systematics}.  In \\S\\ref{sec_mock}, we use mock catalogues to  test the reliability of our results  and to quantify their implications.   Finally, in  \\S\\ref{sec_discussion}, we  present  some further discussion regarding our results. ",
        "conclusions": "\\label{sec_discussion}   We set out to understand why  about 1/4 of {\\it central} dwarf galaxies, those with $r$-band magnitudes between -14.46 and -17.05, are red when one defines a dwarf  galaxy to be  red by  extrapolating the  division between  red sequence galaxies and the blue cloud, as determined by Yang et al. (2008a) for brighter galaxies, down to  dwarf magnitudes. Current models of  galaxy formation would naively  expect such  galaxies  to be  blue  since they  would be  efficiently accreting  gas through  cold mode  accretion  and rapidly  converting it  into stars.  In a recent study,  Ludlow et al.  (2008; see also Lin  et al.  2003) analysed the  properties of subhalos  in galaxy-sized  cold dark  matter halos  using a suite  of cosmological N-body  simulations. The  subhalos in  their definition refer to the whole population  of subhalos physically associated with the main system, including both subhalos that are found within the virial radius of the host halo  at the  present time, and  halos that  were once within  the virial radius of  the main  progenitor of  the host and  have survived  as self-bound entities until $z=0$.  They found that such populations can extend beyond {\\it three times}  the virial radius, and  contain objects on  extreme orbits, with some  approaching the nominal  escape speed  from the  system. On  average the subhalos identified  within the  virial radius represent  only about  {\\it one half} of all  associated subhalos, and many relatively  central halos may have actually been ejected  in the past from a more  massive system. Since galaxies are assumed  to form in  dark matter  halos, it is  interesting to see  if the results we obtain here can be  understood in terms of galaxy formation in this population of subhalos.  According to  the current theory  of galaxy formation, satellite  galaxies can experience various environmental effects  that can quench their star formation and make them  red (e.g., van den Bosch et al.  2008b and references therein). Because the galaxies in ejected subhalos have also been satellite galaxies, at least  for some period  of time,  they are  likely to  have been  subjected to similar environmental effects, and thus  to have experienced some quenching of their star  formation rates.  It is  thus likely that  the association  of red dwarf  galaxies  with  massive  halos   presented  here  is  produced  by  the association of ejected subhalos with their (former) hosts.  As shown in  Table~\\ref{tab1}, about 30\\% of the  dwarf galaxies are satellite galaxies. According  to the results obtained  by Ludlow et  al.  (2008), there should  thus  also  be a  significant  fraction  of  dwarf galaxies  that  are physically associated with nearby more  massive halos out to about three times the virial  radius.  If  these associated galaxies  (now outside  their hosts) have properties similar to the satellite galaxies, we would expect an enhanced fraction of  red dwarf  galaxies that are  distributed outside  massive halos. This is qualitatively consistent with our findings presented above and also as shown in Table~\\ref{tab1}.  However, since the observational data are obtained in redshift space and based on  galaxy groups that may contain interlopers and may  be incomplete,  a detailed  comparison between  the data  and  the models requires  the construction of  mock catalogues  that make  use of  the subhalo population and contain all the observational selection effects.  The fact that central dwarf  galaxies have concentrations that are independent of their distances  to the nearest massive halos  indicates that the processes that causes  them to become  red does not  have a significant impact  on their structure.  This  is similar to the results  obtained by van den  Bosch et al. (2008a) who  found that the transformation mechanisms  operating on satellites affect color more  than structure (see also Kauffmann et  al. 2004; Blanton et al. 2005a; Ball, Loveday \\& Brunner 2008; Weinmann et al.  2008).  Once again, this  similarity between red  dwarfs that  are satellites  and those  that are centrals suggests that both populations  may have experienced similar kinds of environmental effects.  However, as shown  in Table~\\ref{tab1}, in the S1+S2+S3  sample less than 42\\% of the  red dwarf central galaxies  have $\\rp/R_{180}\\le 3$.   In other words, more than  58\\% of the red  dwarf central galaxies  are not close enough  to a larger halo so that they could  have been preprocessed there, becoming red and then subsequently being ejected.  The origin of this population of central red dwarf galaxies,  which is almost 10\\%  of the combined  S1+S2+S3 dwarf sample, still  remains a  mystery within  the standard  paradigm of  galaxy formation. Furthermore,  if we had  used stellar  mass instead  of $r$-band  magnitude to define our dwarf  sample, the percentage of galaxies  in this population would likely increase.  This population of isolated red dwarfs are not merely a dust reddened star forming  population seen near edge-on because  their axis ratios are  consistent with  a randomly  oriented population.  Although as  Croton \\& Farrar  (2008)  probed  the origin  of  the  red  dwarfs  in voids  using  the semi-analytical  models, they only  found $\\sim  0.4\\%$ of  the dwarfs  in the total population  are red  centrals in  voids. While here  we find  that $\\sim 10\\%$ dwarfs  are red centrals  without close neighbours  with $\\rp/R_{180}\\le 3$.  An outstanding problem  for all galaxy formation models  concerns the low mass slope of  the galaxy mass  function.  CDM models  in general predict  too many low-mass dark matter haloes compared to  the number of low mass galaxies.  The mass function of dark matter haloes,  $n(M)$, scales with halo mass roughly as $n(M)\\propto M^{-2}$ at the low-mass end.  This is in strong contrast with the observed  luminosity function  of galaxies,  $\\Phi  (L)$, which  has a  rather shallow shape  at the faint end,  with $\\Phi(L) \\propto  L^{-1}$. To reconcile this difference  one usually  invokes some form  of feedback within  these low mass halos.  If the feedback mechanism were to prevent gas from entering these halos  at  late times,  such  galaxies would  appear  red.   For example,  the preheating mechanism of Mo et al. (2005), where gas is preheated by gas shocks within the  forming large scale structures  in which the low  mass dark matter halos themselves are forming, has this feature.  Therefore, this population of isolated, red, central dwarf galaxies  could represent the tail of the process that prevents  the vast majority  of low mass  dark matter halos  from forming galaxies  and  their further  study  could shed  new  light  on the  mechanism responsible."
    },
    "0812/0812.1984_arXiv.txt": {
        "abstract": "Exact analytical solutions are given for the three finite disks with surface density $\\Sigma_n=\\sigma_0 (1-R^2/\\alpha^2)^{n-1/2} \\textrm{ with } n=0, 1, 2$. Closed-form solutions in cylindrical co-ordinates are given using only elementary functions for the potential and for the gravitational field of each of the disks.  The  n=0 disk is the flattened homeoid for which $\\Sigma_{hom} = \\sigma_0/\\sqrt{1-R^2/\\alpha^2}$.  Improved results are presented for this disk. The n=1 disk is the Maclaurin disk for which $\\Sigma_{Mac} = \\sigma_0 \\sqrt{1-R^2/\\alpha^2}$.  The Maclaurin disk is a limiting case of the Maclaurin spheroid. The potential of the Maclaurin disk is found here by integrating the potential of the n=0 disk over $\\alpha$, exploiting the linearity of Poisson's equation. The n=2 disk has the surface density $\\Sigma_{D2}=\\sigma_0 \\left(1-R^2/\\alpha^2\\right)^{3/2}$.  The potential is found by integrating the potential of the n=1 disk. ",
        "introduction": "\\label{sec:intro}  Observations of edge-on galaxies now provide detailed structural and kinematic 3-D information. However attempts to use this data to generate mass distribution have not been completely successful.  One problem is that it is difficult to calculate the potential and force vectors for finite disks. There is a need for fully solved finite disks which can be used for theoretical studies, for benchmarking computer programs, and as basis functions to directly model 3-D observations.  Mass modeling commonly assumes that the disk mass in contained in an infinite disk.  These include the exponential disk \\citep{fre70}, the Mestel disk \\citep{mes63}, the Kuzmin-Toomre disk \\citep{too63,bin08,eva92,con00::1}, and the Rybicki disk \\citep{eva93}.  Only a few finite disks have been solved analytically for all $(R,z)$.  \\cite{las83} and \\cite{vok98} give a solution for a for all $(R,z)$ for a finite thin disk with constant density.  The gravitational attraction approaches infinity at the edge of this disk. The gravitational attraction of the finite Mestel disk \\citep{mes63,lyn78,hun84} and the truncated exponential disk \\citep{cas83} are finite at the edge of the disk but these disks have not been solved in closed form for points off the disk. \\cite{hur08,hur05::1} describe a method to approximate the potential of a power law disk for points both on and off the disk.  The family of finite disks with surface density $ \\Sigma_n(R;\\alpha) = \\left(1- {R^2}/{\\alpha^2}\\right)^{n-1/2}$ are related to the Maclaurin spheroid which has been studied since the time of Newton.  This family has recently been studied by \\cite{gon06} and \\cite{ped08}.  \\cite{gon06} use the method of \\cite{hun63} to obtain the general solution as a sum of Legendre polynomials in elliptical coordinates and give evaluated expressions for the potentials for the disks 1, 2, and 3. Here we derive complete closed form solutions in cylindrical co-ordinates for the potential and the gravitational fields of the n=0, 1, and 2 disks. We simplify the integration by moving to the imaginary domain in way which is similar to the complex-shift method introduced by Appell.  See \\cite{cio08,cio07} and references therein  The gravitational potential follows Poisson's equation $\\nabla^2\\Phi = 4\\pi\\G\\rho$. Poissons's equation is linear so that solutions may be added. That is, if $\\nabla^2\\Phi_1=4\\pi\\G\\rho_1$ and  $\\nabla^2\\Phi_2=4\\pi\\G\\rho_2$  then $\\nabla^2(\\Phi_1+\\Phi_2)=4\\pi\\G(\\rho_1+\\rho_2)$. Similarly, solutions may be differentiated with respect to a parameter to obtain a new solution. \\cite{too63} and \\cite{lyn89} exploited this linearity to find a family of velocity-density pairs by differentiating the \\cite{kuz56} model .   In the same way,  \\cite{sat80} derived new models from the \\cite{miy75} set of potential-density pairs by differentiating with respect to a parameter. Here we use a similar approach by integrating a known solution with respect to a parameter. If the integral is tractable the result is a new potential-density pair.  Maple 11 was used in this work. Maple was invaluable in simplifying the unwieldy intermediate results but needed a good deal of coaxing to handle the messier equations. The Maple results were cross checked in several ways.  A common notation is used throughout.   Cylindrical co-ordinates $(R,z)$ are used; $\\Sigma$ is the disk surface density; $\\sigma_0$  is the surface density at $R=0$; and $\\alpha$ is the disk radius. $\\Phi$ is the potential where $\\Phi$ is always negative and $\\Phi(\\infty) =0$;  $F_R$ and $F_z$ are the gravitational field vectors with the standard sign convention, ie: $\\vec{F} = - \\vec{\\nabla} \\Phi  $~.  The principal values of elementary functions are used so that, for instance, for $z=x+\\I y=r \\cos(\\theta) + \\I  r  \\sin(\\theta)$; ~~$\\sqrt{z}=\\sqrt{r}\\cos(\\theta/2) +\\I \\sqrt{r}\\sin(\\theta/2)$, valid for $-\\pi<\\theta<\\pi$.  In this way the function $\\sqrt{z}$ is unambiguous with continuous derivatives except on the negative real axis where it is discontinuous. ",
        "conclusions": "We have presented new solutions for a family of finite disks.  Closed form expressions in cylindrical coordinates using elementary functions are given for the potential and gravitational force for the disks with surface density $\\Sigma_n=\\sigma_0 (1-R^2/\\alpha^2)^{n-1/2} \\textrm{ with } n=0, 1, 2$.  Expressions are also given for the limiting cases of $R=0$ and $z=0$.  These solutions fill a need and should make it easier to model 3-D gravitational phenomenon involving disk galaxies.  This is particularly important due to the recent availability of detailed kinematic data above the plane of the disk."
    },
    "0812/0812.4874_arXiv.txt": {
        "abstract": "\\input abstract.tex ",
        "introduction": "\\input intro.tex ",
        "conclusions": "We have presented a combined analytic and numerical analysis of the effects on a circumbinary accretion disk of central mass loss caused by prompt gravitational wave emission from a binary black hole inspiral and merger.  Our significant results are:  1) The primary effect of abrupt mass loss caused by the merger of supermassive black holes will be a rapid drop in the mass accretion rate across the inner edge of the disk.  If the mass accretion rate is close to Eddington and the inner edge can extend without significant disruption close to the innermost stable circular orbit in the last stages of binary inspiral, then this drop could result in large changes in the X-ray spectrum due to, e.g., dramatically reduced Comptonization from hot optically thin gas inside the ISCO.  The flux changes from this effect would be detectable with current instruments if the source were located precisely enough, but a full survey of the likely error ellipses would require an instrument with a field of view of at least several tenths of a square degree and sensitivity in a $\\sim 10^4$ second exposure to fluxes below $\\sim 10^{-15}$~erg~cm$^{-2}$~s$^{-1}$.  2) The expected enhancement of disk emission following a binary merger due to the circularization of elliptical orbits is not likely to be observed. While shocks are identifiable in our simulations, they contribute little to the bremsstrahlung emissivity of the disk. Moreover, such shocks are separated by rarefactions that act to reduce the brightness of the disk.  3) In addition to serving as probes of cosmology and general relativity, observations of EM signatures from merging binary black holes have the potential to tell us much about the inner structure of accretion disks. If the observed emission from such systems is tied to the accretion rate in particular, the duration of the observed flux deficit can be used to estimate the viscous timescale of the disk."
    },
    "0812/0812.0450_arXiv.txt": {
        "abstract": "We present new mid-infrared ($5 - 35\\mu$m) and ultraviolet (1539 -- 2316 \\AA) observations of the interacting galaxy system Arp 143 (NGC 2444/2445) from the Spitzer Space Telescope and GALEX.  In this system, the central nucleus of NGC 2445 is surrounded by knots of massive star-formation in a ring-like structure. We find unusually strong emission from warm H$_2$ associated with an expanding shock wave between the nucleus and the western knots.  At this ridge, the flux ratio between H$_2$ and PAH emission is nearly ten times higher than in the nucleus. Arp 143 is one of the most extreme cases known in that regard. From our multi-wavelength data we derive a narrow age range of the star-forming knots between 2 Myr and 7.5 Myr, suggesting that the ring of knots was formed almost simultaneously in response to the shock wave traced by the H$_2$ emission.  However,  the knots can be further subdivided in two age groups: those with an age of 2--4 Myr (knots A, C, E, and F), which are associated with $8\\mu$m emission from PAHs, and those with an age of 7-8 Myr (knots D and G), which show little or no $8\\mu$m emission shells surrounding them.  We attribute this finding to an ageing effect of the massive clusters which, after about 6 Myr, no longer excite the PAHs surrounding the knots. ",
        "introduction": "\\label{intro}  \\begin{figure} \\plotone{f1.eps} \\caption{Overlay of the SL (sparse) and LL (complete) maps on a V-band image of Arp 143. North is up.} \\label{slits} \\end{figure}  The direct study of warm molecular hydrogen emission from galaxies has been improving considerably through the advent of more sensitive space-bourne spectrometers on the {\\it Infrared Space Observatory} (e. g. \\citet{rigop02,lutz03}) and more recently 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. Support for this work was provided by NASA through a GO2 award issued by JPL/Caltech.}. With 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{houck04} on {\\it Spitzer} \\citep{werner04} it has been possible to improve upon systematic studies of the mid-IR pure-rotational lines of molecular hydrogen in a variety of environments, from nearby normal galaxies \\citep{roussel07}, to galaxies in more extreme environments, like groups and clusters \\citep{ogle07, egami06, ferland08} and ULIRGS \\citep{armus06, higdon05} (see also a review on molecules by \\citet{omont07}).  Recently unusually strong mid-IR H$_2$ lines have been discovered in the giant shock-wave structure in Stephan's Quintet \\citep{apple06}, and in over a dozen low-luminosity radio galaxies \\citep{ogle07,ogle08}, where H$_2$ are often the strongest in mid-infrared wavelengths. Strong H$_2$ equivalent widths were also found near the ``overlap'' region in the Antenna interacting galaxy \\citep{haas05}. These systems have extremely large H$_2$ equivalent widths and line luminosities (L$_{H_2}> 10^{41-42}$ ergs/s), and relatively low star formation rates. Another defining characteristic is the large ratio of H$_2$ luminosity to PAH line strength, and the large (0.1 to 30\\%) fraction of H$_2$ luminosity to the total bolometric luminosity of the objects. Even larger H$_2$ line luminosities have been found associated with several massive galaxies in X-ray clusters (e. g. \\citet{egami06}). This new class of powerful H$_2$ emitting galaxy appears to be shock excited, and models of Stephan's Quintet, where a large-scale shock is strongly implicated suggest that a significant fraction of the bulk kinetic energy in the shock must be funnelled into the H$_2$ line in order to explain the results \\citep{boula08,guillard08}. In this context, it is of considerable interest to find other examples of the same kind of emission in nearby systems. This paper describes strong H$_2$ emission from the strongly interacting galaxy pair Arp~143.  Arp~143 is an interacting pair of galaxies (NGC 2444/2445) with many of the necessary ingredients for a study of shocks generated by galaxy collisions. The southern gas-rich component of the pair, NGC 2445, is notable for the ring of stellar clusters and HII regions, whereas the SO companion is devoid of activity. It is possible that NGC 2445 shares some similarities with collisional ring galaxies (like the Cartwheel ring), because the powerful star-forming knots lie along a crescent-shaped wave of HI emission \\citep{apple92}, and there is kinematic evidence that the structure is expanding through the disk \\citep{higdon97}. Early unpublished spectroscopy by \\citep{jeske86} and optical and near-IR photometry of the knots \\citep{apple92} suggest the star clusters are quite young $<$30-60Myrs, much younger than the dynamical age of the expanding HI wave--thus providing evidence that they are triggered by a collision with NGC 2444. \\citet{jeske86} estimated sub-solar (LMC-like) metallicity in the disk knots in the range 12+log[O/H]~=~8.56-8.77, similar to other known collisional ring galaxies (see Bransford et al. 1996). The nucleus has mildly super-solar values of 12+log[O/H]~=~9.18--similar to nuclear starbursts.  However, this simple picture is complicated by the unsolved mystery of the large 150kpc-long HI plume \\citep{apple87} which extends just to the north of NGC 2444. The existence of such a long plume implies that the two galaxies involved in the collision have had a previous major tidal encounter in the past \\citep{hibbard99,hibbard00}. Perhaps the simplest interpretation of the system is that NGC 2444 may have experienced an initial close passage with NGC 2445 $\\sim$100~Myrs ago, stripping HI into the plume from NGC 2444. After an initial non-pentrating encounter, the main stellar body of NGC 2444 has returned to collide recently with the disk of NGC 2445 in a manner similar to collisional ring galaxies \\citep{apple96}.    In this paper we will present new multi-wavelength imaging and spectroscopic data of Arp~143, another example of a galaxy with very strong Mid-IR H$_2$ emission. Our goal is to connect the recent dynamical events in Arp~143 with the shock and star formation history of the galaxy.  The paper will use {\\it Spitzer}, {\\it GALEX\\footnote{GALEX (Galaxy Evolution Explorer) is a NASA small explorer launched in 2003 April. We gratefully acknowledge NASA's support under Guest Investigator program no. 45.}}, and new ground based optical and near-IR images as well as mid-IR spectroscopy to explore the physical conditions in Arp~143. This will include the study of the dust and HII region properties, ages and luminosities of star formation regions and the excitation conditions of the warm H$_2$ gas.  After a description of the images and spectra in \\S 2, we will present in \\S 3 the multiwavelength photometry and the spectral analysis. In \\S 4, the discussion of the results will follow two main directions: the analysis of the shock region using mainly results from the spectral analysis, and the star formation history of the knots based on an interpretation of the UV-infrared SEDs of each knot.  We will assume a distance to Arp~143 of 56.7 Mpc based on its corrected Virgo-centric velocity of 4142 km/s and an assumed H$_0$=73 km/s/Mpc. At this distance, 1 arcsecs corresponds to 275pc. ",
        "conclusions": "Mid-infrared observations of Arp 143 have been presented in this paper, along with ancillary data including GALEX UV images. Multiwavelength images from UV to mid-infrared show a bright dusty nucleus surrounded by young star forming knots. The four main conclusions of this study are the following:  \\begin{itemize} \\item Spectral line maps of the H$_2$ rotational lines show a ridge of warm H$_2$ emission that curves between the nucleus and the western knots. The H$_2$ line flux related to PAH flux is nearly 10 times higher in the ridge than in the nucleus. The flux ratios between the sum of H$_2$ S(0) and S(2) lines over the sum of PAH 7.7$\\mu$m and $8.6\\mu$m lines reveal that this H$_2$ ridge observed in Fig.~\\ref{h2fig} arises from shocked gas behind the wave that provoked the onset of the ring of star forming knots. This is one of the few cases were this kind of feature is seen.  \\item With the use of photometry and fitting the SEDs, we improved greatly the age determination of the knots. The knots in the ring are all very young, varying from 2-7.5 Myr old and the nucleus is about 500 Myr old. The ring of knots are a product of a shock wave that has been expanding from the nucleus since $\\sim85$ Myr ago, and were formed simultaneously in situ. Behind the shock the molecular gas condenses, and it is traced by the H$_2$ emission ridge. The improvement of age determination is achieved in this study especially due to the GALEX FUV and NUV bands, which are crucial to break the age-extinction-metallicity degeneracy.  \\item The distribution of ages of the knots correlate with the presence of PAH emission. Younger 2-4 Myr old knots are associated with PAH emission shells, whereas older 7-8 Myr knots contain little or no PAH emission. The most plausible explanation for this would be an effect of the ageing of the massive stars created in a single instantaneous burst all over the ring, within a timescale of $\\sim6$ Myr, after which the UV radiation is no longer energetic enough to excite the PAHs.  \\item Given the current reservoir of molecular gas, and assuming that the current star formation rate maintains itself constant, the star formation in Arp 143 will last for about 2 Gyr. However, as it is very likely that the SFR will increase dramatically, this can only be an upper limit for the duration of the starburst. \\end{itemize}"
    },
    "0812/0812.0720_arXiv.txt": {
        "abstract": "The strong gravitational field of neutron stars in the brany universe could be described by spherically symmetric solutions with a metric in the exterior to the brany stars being of the Reissner--Nordstr\\\"{o}m type containing a brany tidal charge representing the tidal effect of the bulk spacetime onto the star structure. We investigate the role of the tidal charge in orbital models of high-frequency quasiperiodic oscillations (QPOs) observed in neutron star binary systems. We focus on the relativistic precession model. We give the radial profiles of frequencies of the Keplerian (vertical) and radial epicyclic oscillations. We show how the standard relativistic precession model modified by the tidal charge fits the observational data, giving estimates of the allowed values of the tidal charge and the brane tension based on the processes going in the vicinity of neutron stars. We compare the strong field regime restrictions with those given in the weak-field limit of solar system experiments. ",
        "introduction": "The string and M-theories describing gravity as higher-dimensional interaction appearing to be effectively 4D at low energies inspired braneworld models with the observable universe being a 3-brane (domain wall) to which the standard non-gravitational matter fields are confined; on the other hand, gravity enters the extra spatial dimensions allowed to be much larger than $l_\\mathrm{P}\\sim 10^{-33}~\\mathrm{cm}$ \\cite{Ark-Dim-Dva:1998:}. Gravity can be localized near the brane even with an infinite size extra dimension with the warped spacetime satisfying the 5D Einstein equations with a negative cosmological constant, as shown by Randall and Sundrum \\cite{Ran-Sun:1999:}. Consequently, an arbitrary energy-momentum tensor could be allowed on the brane with effective 4D Einstein equations \\cite{Shi-Mae-Sas:2000:}.  The Randall--Sundrum model gives standard 4D Einstein equations in the low-energy limit, but significant deviations occur at very high energies, e.g., in the vicinity of compact objects, such as black holes \\cite{Dad-etal:2000:} or neutron stars \\cite{Ger-Maa:2001:}. There are high-energy effects of local origin influencing pressure of matter and non-local effects of ``back-reaction'' origin arising from the influence of the Weyl curvature of the bulk space on the brane.  The combination of the high-energy (local) and bulk stress (non-local) effects alters significantly the matching problems on the brane in comparison with the standard 4D gravity \\cite{Maa:2004:}, and the stresses induced by the bulk gravity imply that the matching conditions do not have unique solution on the brane \\cite{Maa:2004:}. The 5D Weyl tensor is needed as a minimum condition for uniqueness, but no exact 5D solution is known recently \\cite{Ger-Maa:2001:}.  Assuming a spherically symmetric metric induced on the brane, solutions of the effective 4D Einstein equations could be found relatively easily. The vacuum (in both the bulk and brane) solutions represent black holes (naked singularities) of the Reissner--Nordstr\\\"{o}m (RN) type endowed with a brany tidal charge parameter $b$, describing the bulk tidal stress effect on the black hole structure \\cite{Dad-etal:2000:}. In the brany RN metric, the tidal charge $b$ stands instead of the charge parameter $Q^2$ of the standard RN spacetimes~\\cite{Mis-Tho-Whe:1973:Gra:}; $b$ can be both positive and negative, but there are indications that $b<0$ should properly represent the effects of the bulk space Weyl tensor \\cite{Dad-etal:2000:,Sas-Shi-Mae:2000:,Maa:2004:}.\\footnote{The stationary and axially symmetric vacuum solutions describing rotating black holes (naked singularities) are found to be of the Kerr--Newman type endowed again with the tidal charge reflecting the effect of the bulk stresses \\cite{Ali-Gum:2005:}.}  In the case of spherically symmetric neutron stars (or quark stars) simplified by the uniform energy density profile, the local high-energy and non-local bulk graviton stress effects make the matching conditions unambitious, since we do not know the 5D Weyl tensor. Note that the uniform energy density distribution model of neutron stars could, in principle, be used in order to roughly determine the neutron star spacetime structure, since in realistic models with a variety of realistic equations of state, the energy density profile happens to be nearly uniform in most of the neutron star interior \\cite{Web-Gle:1992:ASTRJ2:upr:,Gle:1997:CompactStars:}.  Two different exact exterior solutions related to the uniform density interior are known, both having asymptotically Schwarzschild character \\cite{Ger-Maa:2001:}. The first one is of the known RN type with the tidal charge and mass parameters being influenced by the energy density of the star and the brane tension. The second one is more complicated, with metric components being explicitly dependent on the energy density and brane tension \\cite{Ger-Maa:2001:}. Quite recently, the solution of the RN type was used in the weak-field limit in order to obtain restrictions on the brany parameter and the brane tension from the standard tests on the Mercury perihelion precession, deflection of light near the edge of Sun and the radar echo experiments \\cite{Boh-etal:2008:SolarSys:}. Here, we also restrict our investigations to the metric of the RN type in order to obtain an idea of the tidal charge relevance in a strong gravitational field near neutron stars (or black holes) and to find restrictions on the allowed values of the brany parameter. The choice of the RN type of the uniform star brany solutions is supported by a recent study of static exteriors of non-static brany stars showing that spherical non-rotating star solutions admit only Schwarzschild or RN type exterior \\cite{PdeLeon:2008:}. Since the junction conditions of the internal and external neutron star spacetimes relate the tidal charge to the brane tension and the energy density (mass and radius) of the neutron star, we can, in principle, also put restrictions on the allowed values of the brane tension \\cite{Ger-Maa:2001:}.  There is a variety of papers related to both weak-field \\cite{Kee-Pet:2006:brane-lensing:} and strong-field \\cite{Boz:2005:,Eir:2005:,Whi:2005:} optical lensing phenomena in the braneworld black hole spacetimes. The optical phenomena connected to the accretion disc radiation could also reflect in a relevant way the influence of the tidal charge \\cite{Sch-Stu:2007:RAGtime8and9CrossRef:SpBranKBH, Sch-Stu:2007:RAGtime8and9CrossRef:OEBraK}. Here we focus our attention on the strong-field phenomena closely related to the perihelion precession effect discussed in the weak-field limit \\cite{Boh-etal:2008:SolarSys:}.  High-frequency quasiperiodic oscillations (QPOs) observed in Galactic low-mass binary systems, namely, in both the black hole \\cite{Rem-McCli:2006:ARASTRA:,Kli:2004:,Kli:2000:ARASTRA:,Ter-Abr-Klu:2005:ASTRA:QPOresmodel} and neutron star binary systems \\cite{Bel-Men-Hom:2005:ASTRA:,Bel-Men-Hom:2007:astro-ph/0702157:,Bar-Oli-Mil:2005:MONNR:,Tor:2007:ASTRA:RevTwPk6}, or in extragalactic black hole sources \\cite{Str:2007:astro-ph/0701390:,Lac-Cze-Abr:2006:astro-ph0607594:,Asc-etal:2004:ASTRA:,Ter:2005:astro-ph0412500:}, % represent such an interesting phenomenon in the strong field regime enabling estimates of the tidal charge on the basis of orbital models. It is widely accepted that the high-frequency QPOs are related to the orbital motion in the inner part of an accretion disc around a compact object. There are strong indications that some resonant phenomena could be relevant in both black hole systems where twin peak QPOs with fixed frequencies and their ratios are observed \\cite{Ter-Abr-Klu:2005:ASTRA:QPOresmodel}, and neutron star systems with wide spread of frequencies and their ratios where the twin peak QPOs are cumulated around small number ratios \\cite{Tor-etal:2008:AA:,Tor-etal:2008:AA2:}. The hypothetical resonances of the orbital epicyclic oscillations representing plausible models of QPOs were recently studied for brany Kerr black holes \\cite{SK:2008:PhysRevD:}. Here, we study these models as related to the brany neutron stars described approximately by the RN spacetimes, i.e., neglecting the effects related to rotation of neutron stars. We focus our attention to the relativistic precession model \\cite{Ste-Vie:1998:ASTRJ2L:} that fits qualitatively well in the wide spread data given by observations of a variety of neutron star binary systems \\cite{Bel-Men-Hom:2005:ASTRA:,Bel-Men-Hom:2007:astro-ph/0702157:,Bar-Oli-Mil:2005:MONNR:,Tor:2007:ASTRA:RevTwPk6}.  % However, recently it has been discussed in the case of the atoll source 4U~1636$-$53 \\cite{Tor-etal:2008:AA:,Tor-etal:2008:AA2:,Tor-Bak-Stu-Cec:2007::prep} and some other atoll sources \\cite{Tor-etal:2007:RAGtime8and9CrossRef:MutRelkHzQPO} that the resonant phenomena should be relevant even in the framework of the relativistic precession (RP) model. We, thus, present a first study of strong-field phenomena for brany neutron stars that could be compared  with the weak-field results (perihelion precession, radar echo delay and light deflection) based on the RN solutions used to describe approximately the spacetime in the vicinity of Sun \\cite{Boh-etal:2008:SolarSys:}. In order to obtain rough restrictions on both the brany tidal charge and brane tension, we use an ensemble of typical neutron star atoll sources and the special cases of the Z-sources Sco X-1 and Circinus X-1 that demonstrate a wide variety of QPOs.  In~\\sref{sec:dve}, the effective 4D Einstein equations are presented, and the braneworld neutron star spacetime of the RN type (approximated by the uniform energy density distribution \\cite{Ger-Maa:2001:}) is briefly described. In~\\sref{sec:tri}, the circular geodesics are discussed. In~\\sref{sec:ctyri}, frequencies of the orbital epicyclic oscillations are given; their radial profiles are discussed, and the resonance conditions are briefly described for some frequency ratio functions modeling twin peak QPOs. In~\\sref{sec:pet}, the RP model is used to fit data observed in the representative ensemble of neutron star X-ray binaries exhibiting QPOs, and the restrictions on the tidal charge and brane tension are given. In~\\sref{sec:sest}, concluding remarks are presented. ",
        "conclusions": ""
    },
    "0812/0812.5111_arXiv.txt": {
        "abstract": "We investigate the biases and uncertainties in estimates of physical parameters of high-redshift Lyman break galaxies (LBGs), such as stellar mass, mean stellar population age, and star formation rate (SFR), obtained from broad-band photometry. These biases arise from the simplifying assumptions often used in fitting the spectral energy distributions (SEDs).  By combining $\\rm{\\Lambda}$CDM hierarchical structure formation theory, semi-analytic treatments of baryonic physics, and stellar population synthesis models, we construct model galaxy catalogs from which we select LBGs at redshifts $z$ $\\sim$ 3.4, 4.0, and 5.0.  The broad-band photometric SEDs of these model LBGs are then analysed by fitting galaxy template SEDs derived from stellar population synthesis models with smoothly declining SFRs. We compare the statistical properties of LBGs' physical parameters -- such as stellar mass, SFR, and stellar population age -- as derived from the best-fit galaxy templates with the intrinsic values from the semi-analytic model.  We find some trends in these distributions: first, when the redshift is known, SED-fitting methods reproduce the input distributions of LBGs' stellar masses relatively well, with a minor tendency to underestimate the masses overall, but with substantial scatter.  Second, there are large systematic biases in the distributions of best-fit SFRs and mean ages, in the sense that single-component SED-fitting methods underestimate SFRs and overestimate ages. We attribute these trends to the different star formation histories predicted by the semi-analytic models and assumed in the galaxy templates used in SED-fitting procedure, and to the fact that light from the current generation of star-formation can hide older generations of stars. These biases, which arise from the SED-fitting procedure, can significantly affect inferences about galaxy evolution from broadband photometry. ",
        "introduction": "Over the last decade, astronomy has experienced a boom of panchromatic surveys of high redshift galaxies -- such as the Great Observatories Origins Deep Survey \\rm{\\citep[GOODS,][]{gia04a}} and the Cosmic Evolution Survey \\rm{\\citep[COSMOS,][]{sco07}} -- thanks to the developments of space telescopes, like \\it{Hubble Space Telescope} \\rm{(HST)} and \\it{Spitzer Space Telescope} \\rm{($Spitzer$)}, and also of ground-based facilities, including Keck observatory, Very Large Telescope (VLT), and Subaru telescope.  Spectral energy distributions (SEDs) constructed from the photometric data of wide wavelength coverage have been used to constrain galaxy parameters, such as stellar masses, SFRs, and stellar population ages, of high redshift galaxies via comparison with simple stellar population synthesis models \\citep[e.g.][]{pap01,sha01,sha05,shi07,sta07,ver07}.  Accompanied byselection techniques to separate high-redshift galaxy samples efficiently via their broad-band colors and magnitudes, these surveys have provided increasing information about the nature of high-redshift galaxies.  However, despite recent advances in our knowledge of the properties of high redshift galaxies, more accurate estimation of physical parameters is necessary for addressing several important issues in galaxy evolution and cosmology.  For example, it is as yet uncertain whether the Lyman break galaxies \\rm{\\citep[LBGs,][]{ste96,gia02}} -- high-redshift, star-forming galaxies selected according to their rest-frame ultraviolet (UV) colors -- will evolve into large ellipticals \\citep{ade98,ste98}, or into smaller spheroids, such as small ellipticals/spiral bulges or subgalactic structures, which will later be merged to form larger galaxies at $z \\sim 0$ \\citep{low97,saw98,som01}. Accurate knowledge of physical parameters, such as stellar masses and SFRs, serves as one of the important elements in discriminating among the possible evolutionary descendants of LBGs.  Also, accurately constraining LBG's stellar-population ages and stellar masses is necessary for determining how much the LBG populations contributed to cosmic reionization.  Better estimation of these parameters is also crucial in comparing different galaxy populations at different redshifts.  The stellar populations of LBGs as well as other populations of high-redshift galaxies have been studied by several authors using SED-fitting methods to compare the photometric SEDs of observed galaxies with various galaxy spectra from stellar population synthesis models. \\citet{pap01} investigated 33 spectroscopically-confirmed LBGs with a redshift range, $2.0 \\leq z \\leq 3.5$. They found that the mean lower limit of stellar masses of these LBGs is $\\sim 6 \\times 10^{9}$ $M_{\\sun}$ with upper limits of $\\sim$ 3-8 times larger and that the mean age is $\\sim$ 120 Myr (assuming solar metallicity and \\citet{sal55} IMF) with broad range between 30 Myr and 1 Gyr.  More importantly, they concluded that the most robust parameter constrained through SED-fitting is stellar mass, while stellar population age and especially star formation rate (SFR) are only poorly constrained. Also, they speculated that SED-fitting methods (single-component fitting) can only give lower limits to galaxies' stellar masses because the results from the SED-fitting methods are largely driven by the light from the most massive, most recently formed stars. \\citet{sha01} studied the physical parameters of 74, $z\\sim$ 3 LBGs with spectroscopic redshifts. They confirmed that among various physical parameters, stellar mass is the most tightly constrained, and that constraints on other parameters such as dust extinction and age are weak. With the assumption of solar metallicity and a Salpeter IMF, the median age (defined as time since the onset of current star formation) was found to be $t_{sf} \\sim$ 320 Myr with a large spread from several Myr to more than 1 Gyr. The median stellar mass was 1.2$\\times 10^{10} ~h^{-2} ~M_{\\sun}$, which is higher than the \\citet{pap01} value. However, the \\citet{sha01} LBGs were generally brighter than LBGs in the \\citet{pap01} sample. The two studies found similar stellar masses for LBGs with similar rest-frame UV luminosities. \\citet{sha01} also found that about 20 $\\%$ of the LBGs have best-fit ages older than 1 Gyr, stellar masses larger than $10^{10} ~M_{\\sun}$, and SFR $\\sim$ 30 $M_{\\sun}$ $yr^{-1}$, and they interpreted these as galaxies which have formed their stars over a relatively long period in a quiescent manner. \\citet{sha01} also noted that more luminous galaxies are dustier than less luminous ones and that younger galaxies are dustier and have higher SFR.  More recently, \\citet{ver07} studied 21 $z\\sim$ 5 LBGs, six of which have confirmed spectroscopic redshifts and the remaining 15 of which have photometric redshifts.  These LBGs were found to be moderately massive with median stellar mass $\\sim$ $2 \\times 10^{9}$ $M_{\\sun}$, and to have high SFRs with a median of $\\sim$ 40 $M_{\\sun} \\rm{yr^{-1}}$.  They also found that the stellar mass estimates are the most robust of all derived properties.  Best-fit ages have a large spread with a median value $\\sim$ 25 Myr, assuming a metallicity of one-fifth solar (0.2 $Z_{\\sun}$) and a Salpeter IMF. \\citet{ver07} also compared their $z \\sim$ 5 LBGs with \\citet{sha01}'s $z \\sim$ 3 LBGs, assuming the same IMF and metallicity (i.e. solar metallicity and Salpeter IMF), and concluded that these two samples of LBGs with similar rest-frame UV luminosity are clearly different. Specifically, $z\\sim$ 5 LBGs are much younger ($\\lesssim$ 100 Myr) and have lower stellar masses ($\\sim ~10^{9} ~M_{\\sun}$) than $z\\sim$ 3 LBGs.  The fraction of young (age $<$ 100 Myr) galaxies is $\\sim$ $70 \\%$ for $z \\sim$ 5 LBGs while it is $\\sim$ $30 \\%$ for $z \\sim$ 3 LBGs.  They also concluded that these $z\\sim$ 5 LBGs are the likely progenitors of the spheroidal components of present-day massive galaxies, based on their high stellar mass surface densities. \\citet{sta07} analysed 72 $z\\sim$ 5 galaxies (not just LBGs) with photometric redshifts. They performed an SED-fitting analysis on a spectroscopically confirmed subset of 14 $z\\sim$ 5 galaxies to derive best-fit stellar masses ranging between $3 \\times 10^{8}$ and $2 \\times 10^{11}$ $M_{\\sun}$, and ages from 1 Myr to 1.1 Gyr. Three out of these 14 galaxies have stellar masses in excess of $10^{11} ~M_{\\sun}$.  Using stellar mass estimates from both spectroscopically confirmed and candidate galaxies, they calculated a stellar mass density at $z\\sim$ 5 of $6 ~\\times ~10^{6} ~M_{\\sun} ~\\rm{Mpc^{-3}}$, which is much larger than the integration of the star formation rate density (SFRD) from $z\\sim$ 10 to 5. They attributed this discrepancy either to significant dust extinction or to undetected low-luminosity star-forming galaxies at z $\\gtrsim$ 5. \\citet{shi07} studied the properties of 1088 massive LBGs whose best-fit stellar masses are larger than $10^{11} ~M_{\\sun}$ at $z\\sim$ 3. They derived stellar masses of these massive LBGs through SED-fitting with the assumption of a Salpeter IMF, and noted that LBGs which are detected in mid-infrared wavelength are the ones with large stellar mass and severe dust extinction among LBG population.  The SED-fitting method is also used to constrain physical parameters for high-z Lyman-$\\alpha$ emitting galaxies \\citep[LAEs,][]{fink07} and for submillimetre galaxies \\citep{dye08}. \\citet{fink07} analysed $z \\sim 4.5$ LAEs through SED fitting and found their ages range from 1 Myr to 200 Myr (assuming a constant star formation history), and stellar masses from $2 \\times 10^{7} - 2 \\times 10^{9}$ $M_{\\sun}$.  \\citet{dye08} used SED fitting to study physical parameters of SCUBA (Submillimetre Common-User Bolometer Array) sources, finding an average stellar mass of $\\sim 10^{11.8 \\pm 0.1} M_{\\sun}$.  However, the estimated parameter values and implications derived from these values are prone to several sources of error including errors inherent in SED-fitting methods.  It is unclear whether or not SED-fitting methods deliver biased estimates of parameters such as stellar mass, age, and SFR.  It is also unclear how much worse the parameter estimates become if spectroscopic redshifts are unavailable and the photometric data must be used to constrain not only the star formation history, but also the redshift.  The critical question is how far we can trust the physical parameters derived using mainly photometric data under the assumption of simple star-formation histories (SFHs), when real galaxies are expected to have more complex formation/evolution histories with several possible episodes of star-formation.  In this paper, we address this issue by comparing the statistical distributions of intrinsic physical parameters of model Lyman break galaxies (LBGs) from semi-analytic models (SAMs) of galaxy formation, with the estimates of those same parameters derived from stellar population synthesis models through the commonly used SED-fitting method.  There have been several studies which tried to constrain the stellar populations of LBGs or of high-redshift galaxies by comparing them with predictions from theoretical models, such as semi-analytic models \\citep[e.g.][]{som01,idz04} or cosmological hydrodynamic simulations \\citep[e.g.][]{nag05,nig06,fin07}.  \\citet{som01} compared $z \\sim 3$ and $z \\sim 4$ observed LBGs with semi-analytic models to investigate possible scenarios for LBGs, and \\citet{idz04} put constraints on the properties of $z \\sim 4$ LBGs through comparison of colors of observed LBGs and SAMs model LBGs.  \\citet{nag05} and \\citet{nig06} used cosmological simulations to place constraints on properties of UV-selected, $z \\sim 2$ galaxies and $z \\sim 4-6$ LBGs.  \\citet{fin07} constructed a set of galaxy templates from hydrodynamical simulations and used them to place constraints on properties of 6 $z \\leq 5.5$ observed galaxies. Our approach is distinct from these studies. Here, instead of comparing model galaxies with observed ones, we perform SED fitting on model galaxies, derive best-fit parameters, and compare them with the intrinsic values given in the model, trying to understand quantitatively as well as qualitatively the biases and uncertainties of SED-fitting methods in constraining the physical parameters of LBGs.  To accomplish this, we construct galaxy formation histories by combining $\\rm{\\Lambda}$CDM hierarchical structure formation theory with semi-analytic treatments of gas cooling, star formation, supernova feedback, and galaxy mergers. Then, using the stellar population synthesis models of Bruzual $\\&$ Charlot \\citep[][hereafter BC03]{bru03}, and a simple model for dust extinction, we derive the photometric properties of these model galaxies. After selecting LBGs through appropriate color criteria, we compare the photometric spectral energy distributions (SEDs) in observed-frame optical (HST/ACS) through mid-infrared ($Spitzer$/IRAC) passbands of these model LBGs with those of SED templates from stellar population synthesis model of BC03 using a $\\chi^{2}$-minimization method. In this SED-fitting procedure, a large multi-parameter space is explored to minimize any prior. Each parameter -- such as star-formation time scale, age, dust-extinction, and metallicity -- in this parameter space spans a broad as well as physically realistic range. The distributions of physical parameters of these model LBGs derived from this SED-fitting method are compared with the input distributions of these parameters from the semi-analytic model.  Various parameters in the semi-analytic model are calibrated to reproduce the colors of observed LBGs in the southern field of the GOODS (GOODS-S).  The GOODS \\citep{gia04a} is a deep, multiwavelength survey which covers a combined area of $\\sim 320$ $\\rm{arcmin^{2}}$ in two fields -- GOODS-N centered on the Hubble Deep Field-North (HDF-N), and GOODS-S centered on the Chandra Deep Field-South (CDF-S).  The extensive wavelength coverage of the photometric data in the GOODS fields including HST/ACS (${\\it Advanced~ Camera~ for~ Surveys}$) and $Spitzer$/IRAC (${\\it Infrared~ Array~ Camera}$) is beneficial for the derivation of various physical parameters of galaxies including stellar mass, SFR, and age.  Deep observations (reaching $z_{850} \\sim 26.7$) in the GOODS fields can probe the high redshift universe in a comprehensive and statistically meaningful manner.  This has enabled many authors to investigate high-redshift LBG populations in the GOODS fields \\citep[e.g.][]{gia04b,idz04,pap04,lee06,rav06,yan06,sta07,ver07}. In this work, we use bandpasses for the HST/ACS filters F435W ($B_{435}$), F606W ($V_{606}$), F775W ($i_{775}$), F850LP ($z_{850}$), VLT/ISAAC (Infrared Spectrometer And Array Camera) $J$, $H$, $Ks$ bands, and $Spitzer$/IRAC 3.6, 4.5, 5.8, 8.0 $\\mu m$ channels for SED-fitting.  The CTIO (Cerro Tololo Inter-American Observatory) MOSAIC U band is used for selecting U-dropouts.  A description of the semi-analytic model of galaxy formation used in this work is given in \\S~2, with the Lyman break galaxy sample selection and the SED-fitting procedures explained in \\S~3.  The statistical properties of the derived physical parameters for model LBGs are shown in \\S~4. We analyse, in detail, the effects of various factors on the biases in SED-fitting in \\S~5, and the bias in estimating SFR from rest-frame UV magnitude in \\S~6. We discuss the effects of these biases on the galaxy evolution studies in \\S~7 and a summary and our conclusions are given in \\S~8. Throughout the paper, we adopt a flat $\\rm{\\Lambda}$CDM cosmology, with ($\\Omega_{m}, \\Omega_{\\Lambda}$) = (0.3,0.7), and $H_{0}$ = 100$h$ $km$ $s^{-1}$ $Mpc^{-1}$, where $h$=0.7, and all magnitudes are expressed in the AB magnitude system \\citep{oke74}. ",
        "conclusions": "In this paper, we examine how well the widely-used SED-fitting method can recover the intrinsic distributions of physical parameters -- stellar mass, SFR, and mean age -- of high-redshift, star-forming galaxies. To this end, we construct model high-redshift galaxies from the semi-analytic models of galaxy formation, make a photometric catalog via the BC03 synthesis model, and select LBGs through the appropriate color selection criteria based on their broadband colors. Then, we perform SED-fitting analysis, comparing the photometric SEDs of these model galaxies with various galaxy spectral templates from the BC03 stellar population model to derive the distributions of best-fit stellar masses, SFRs, and mean ages. We use this test to explore (rather exhaustively) the errors and biases that arise in such SED fitting and the underlying causes of these errors and biases.  Here are the summary of the main results of this work.  1. When we fix the redshift to the given value in the SAM catalog and use $ACS$/$ISAAC$/$IRAC$ passbands, the SED-fitting method reproduces relatively well the input distributions of stellar masses with a minor tendency to underestimate the stellar masses and with substantial scatter for individual galaxies. The mean stellar masses are underestimated by about 19$\\sim$25 $\\%$, which is due to the fact that the old generations of stars can be hidden by the current generation of star formation. The distributions of SFRs and mean ages show larger offsets than the stellar mass distributions. The SFRs are systematically underestimated and the mean ages are systematically overestimated, and these trends mainly reflect the difference in the SFHs predicted by the semi-analytic models and assumed in the simple galaxy templates used in the SED fitting. The well-known `age-extinction degeneracy' plays an important role in biasing the derived SFRs.  2. When we use redshift as an additional free parameter, the discrepancy between the intrinsic- and SED-derived stellar mass distributions increase (i.e., the overall stellar mass underestimation becomes worse), while the bimodalities which appear in the SFR $\\&$ mean age distributions become more significant. The distributions of offsets of individual galaxies indicate that there exist sub-population(s) of LBGs whose behaviors are distinct from the majority of LBGs in the SED-fitting. The SED-fitting generally underestimates the redshift slightly.  3. The age overestimates are clearly related to the intrinsic age and specific star formation rate of each galaxy. The overestimation of mean ages is worse for galaxies with younger ages and higher SSFRs. Inspection of the SFHs of individual SAM galaxy confirms that the main origin of the bias in the age estimation is the difference of assumed SFHs in SAM galaxies and the simple galaxy templates used in the SED fitting. This bias in the age-estimation propagates into the stellar mass and SFR estimations, in the sense that the age-overestimation leads to the mass-overestimation and SFR-underestimation. The SFR-underestimation is further enhanced by the `age-extinction degeneracy'. Another source of biases is the dominance of the current generation of star formation over the old generation(s) of star formation in the SED-fitting. This causes both of the stellar masses and SFRs to be underestimated.  4. We perform two types of two-component SED-fitting: (1) adding a young, bursty component to an old component with long-lasting SFH, and (2) combining an old, bursty component with a younger, long-lasting component. The changes of behaviors in these two types of two-component fitting depend on galaxies' SFHs. Generally, compared with the best-fit values in the single-component fitting, the best-fit stellar masses are generally smaller in the two-component fitting with a young, bursty component embedded in an older, long-lasting component, while they become larger in the two-component fitting with an old burst combined with a younger, long-lasting component. For the majority of galaxies (mainly, with type-1/type-2 SFHs), the best-fit ages become younger and the best-fit SFRs become higher in both types of the two-component fitting. The behavior of galaxies with type-3 SFH is opposite: the best-fit ages are older and the best-fit SFRs are smaller than the values derived in the single-component fitting.  5. We perform the SED-fitting with different combinations of passbands -- by omitting $IRAC$ data or $ISAAC$ data. If we fit the galaxy SEDs with $ACS$ and $ISAAC$ data only omitting $IRAC$ data, the derived distributions of stellar masses, SFRs and mean ages are significantly affected. The detailed behaviors of change in the SED-fittings with and without IRAC data strongly depend on galaxy's redshift and SFH. Alternatively, if we fit the SEDs of LBGs with $ACS$ and $IRAC$ data only, the derived distributions of stellar masses, SFRs, and means age do not show significant changes except for the enhanced bimodalities in the SFR/age distributions for B-/V-dropouts. This indicates that the significant changes occur only for a small fraction of galaxies, while the effects of $ISAAC$ data in the SED-fitting are insignificant for the majority of LBGs. These experiments demonstrate the usefulness of the observed-frame MIR data from $IRAC$ in constraining physical properties of the high-$z$, star-forming galaxies.  6. When the allowed range of $\\tau$ (star formation $e$-folding timescale) is limited to be insufficiently short or long, biases in the SED-fitting increase, in general, compared with the case when sufficiently broad range of $\\tau$ is allowed (from 0.2 Gyr to 15.0 Gyr in this study). The mean values of the best-fit stellar mass and age are smaller than the values derived with the full range of $\\tau$ used. The age shifts are larger when we limit $\\tau$ to be small ($\\leq$ 1.0 Gyr) than when we use very large value of $\\tau$ (= 15.0 Gyr). Detailed behaviors of change for individual galaxy depend on the SFH.  7. The experiments with the SED templates constructed from the BC03 model (\\S~\\ref{test}) isolate the effect of limiting the allowed range of $\\tau$. If $\\tau$ is restricted to be $\\tau \\leq 1$ Gyr, both the stellar mass and mean age are underestimated while the SFR is overestimated. If only very long values of $\\tau$ are used, the stellar mass is overestimated for the majority of model galaxies (from BC03 model) as a result of the mis-assignment of light to older stars (with consequently higher mass-to-light ratio). The biases in the SFR and age estimation depend on the age of galaxy. If the age (more exactly, value of $t$, i.e. time since the onset of star formation) is long, compatible to the age of the universe, the mean age is underestimated and $E(B-V)$ is greatly overestimated leading to the severe overestimation of SFR. Otherwise, the mean age is overestimated and SFR is underestimated. In the case when one tries the single-component SED-fitting for the galaxies with two clearly distinguished generations of star formation (resembling repeated-burst models), the stellar mass, SFR, and mean age are all underestimated.  8. Star-formation rates estimated from the UV luminosity alone may be less biased than those estimated from SED-fitting, provided one has a reasonable estimate of $E(B-V)$. This is mainly due to the fact that the results are less subject to the degeneracy between age and dust-extinction when we use only rest-frame UV photometry. However, the bias depends on galaxies' redshift, and more significantly on the estimation of mean dust-extinction, which is challenging. A relatively small change in $E(B-V)$ would result in large bias in UV-derived SFRs.  9. We show that biases arising in the SED-fitting procedure can affect studies of the high-redshift, star-forming galaxies. The different directions and amounts of biases depending on galaxy's SFH can produce the artificial bimodalities in the age or SFR distributions. This can affect the interpretation of the properties and nature of these galaxies. Also, the stellar mass underestimation for massive LBGs, combined with the SFR underestimation, can affect the interpretation of possible evolutionary paths for these massive LBGs.  In conclusion, we show that single-component SED-fitting generally slightly underestimates the LBGs' stellar mass distributions, while the SFR distributions are significantly underestimated and the age distributions are significantly overestimated. The main causes of these biases are: (1) the difference of assumed SFHs between in the SAM galaxies and simple templates used in the SED-fitting, and (2) the effects of the current generation of star formation masking the previous generation(s) of stellar population. The well-known `age-extinction degeneracy' (or `age-extinction-redshift degeneracy' in the case when redshifts are allowed to vary freely as an additional free parameter during the SED-fitting procedure) plays an additional role, mainly in the estimation of SFR distributions. Consequently, the directions and amounts of the biases in the SED-fitting strongly depend on galaxy's star formation history (SFH). If we change various inputs or fitting parameters in SED-fitting, such as the range of $\\tau$ -- the e-folding timescale of star formation, combinations of passbands used, or assumed SFHs, the derived distributions of best-fit stellar masses, SFRs and ages can change dramatically. Due to the compensating causes of biases, the best-fit stellar mass distributions are more stable against these changes than the SFR/age distributions. Moreover, the behaviors of individual galaxy in various settings of the SED-fitting strongly depend on galaxy's SFH.  These biases arising in the SED-fitting can have significant effects in the context of the galaxy formation/evolution studies as well as cosmological studies. Besides the effects discussed in \\S~\\ref{discus}, biases in the estimation of stellar mass and SFR (which are dependent on the various input settings in the SED-fitting, on available passbands and more importantly on individual galaxy's SFH) can affect, for example, the estimation of the stellar mass function as well as the global density of the stellar mass and SFR, which are important probes in galaxy formation/evolution. The blind comparison among various works -- which were done with different input settings in the SED-fitting, with the different sets of available photometric data (e.g. works done before or after $Spitzer$-era), or with differently-selected galaxy samples -- can also misleads us. Therefore, appropriate caution should be applied to the estimates of physical parameters of high-redshift, star-forming galaxies through the SED-fitting, and also to the interpretation in the context of galaxy formation/evolution based on (the comparisons of) the derived results."
    },
    "0812/0812.2731_arXiv.txt": {
        "abstract": "In this paper, we consider logarithmic radiative corrections and higher order terms to the supersymmetric hilltop F- and D-term hybrid inflation models. Conventional F- and D-term hybrid inflation only predicts $n_s \\gae 0.98$. We show that via a positive quadratic and a negative quartic correction the spectral index can be reduced to $n_s=0.96$ suggested from latest WMAP result and also cosmic string problem appeared in SUSY hybrid inflation can be solved with mild tuning of the parameters if $\\kappa \\lae 0.01$ for F-term inflation and $g \\lae 0.05$ for D-term inflation. ",
        "introduction": "\\label{1} Inflation \\cite{Starobinsky:1980te, Sato:1980yn, Guth:1980zm} (for review, \\cite{Lyth:1998xn, Lyth:2007qh, Linde:2007fr}) is an vacuum dominated epoch in the early universe when the scale factor grew exponentially. This scenario is used to set the initial condition for the hot big bang and provided primordial density perturbation as the seed of structure formation. In the framework of slow-roll inflation, the slow-roll parameters are defined by \\begin{equation} \\epsilon \\equiv \\frac{M_P^2}{2} \\left(\\frac{V^\\prime}{V}\\right)^2, \\end{equation} \\begin{equation} \\eta \\equiv M_P^2\\frac{V^{\\prime\\prime}}{V}, \\end{equation} where $M_P=2.4\\times 10^{18} \\mbox{ GeV}$ is the reduced Planck mass. The spectral index can be expressed in terms of the slow-roll parameters as \\begin{equation} n_s=1+2\\eta-6\\epsilon. \\label{n} \\end{equation} The spectrum is given by \\begin{equation} P_R=\\frac{1}{12\\pi^2M_P^6}\\frac{V^3}{V'^2} \\end{equation} With the slow-roll approximation the value of the inflaton field $\\phi$, in order to achieve $N$ e-folds inflation, is \\begin{equation} N=M^{-2}_P\\int^{\\phi(N)}_{\\phi_{end}}\\frac{V}{V'}d\\phi. \\label{efolds} \\end{equation} From observation \\cite{Komatsu:2008hk}, $P^{1/2}_R\\simeq 5\\times 10^{-5}$ at $N \\simeq 60$.  Suppose we are going to build a small field inflation model, i.e. $\\phi \\lae M_P$. For a wide range of the potentials, for example, polynomial and logarithmic forms, we have \\begin{equation} V^\\prime \\simeq \\phi V^{\\prime\\prime}. \\end{equation} Therefore, since $\\phi \\lae M_P$ \\begin{equation} |\\epsilon|=\\frac{\\phi^2}{2M_P^2}|\\eta|^2 \\ll |\\eta|. \\end{equation} From Eq. (\\ref{n}) \\begin{equation} n_s \\simeq 1+2\\eta. \\end{equation} The latest WMAP 5-year result prefers the spectral index around $n_s=0.96$ \\cite{Komatsu:2008hk} which implies $\\eta=-0.02$. This suggests the inflation took place near a maximum of the potential which is concave-downward. This kind of models is called ``hilltop inflation'' \\cite{Boubekeur:2005zm, Kohri:2007gq}.  In hybrid inflation \\cite{Linde:1993cn}, a false vacuum is provided by the waterfall field. By adding the false vacuum, even use a quadratic potential (like the one in chaotic inflation \\cite{Linde:1983gd}), the model can still be a small field model \\cite{Copeland:1994vg}. % In the framework of supersymmetry, there are two standard types of hybrid inflation models: F-term \\cite{Dvali:1994ms, Copeland:1994vg} and D-term  \\cite{Binetruy:1996xj, Halyo:1996pp, Riotto:1997wy, Lyth:1997pf}. Both types of models predict the spectral index $n_s \\gae 0.98$. In conventional forms of models, cosmic strings are produced after inflation. However, there is a $10\\%$ upper bound for the contribution from cosmic strings to the CMB angular power spectrum \\cite{Pogosian:2003mz, Pogosian:2004ny, Wyman:2005tu}. This puts a very strong constraint to the parameters in both F- and D-term hybrid inflation. In \\cite{Lin:2006xta}, it was shown that the cosmic string problem for D-term inflation can be solved by using a negative quadratic correction to the scalar potential, but a spectral index $n_s<0.96$ is required. In D-term inflation, if we reduce the superpotential coupling to $\\lambda \\lae O(10^{-4}-10^{-5})$ and the $U(1)_F$ gauge coupling $g \\lae 2\\times 10^{-2}$ \\cite{Endo:2003fr, Rocher:2004my}, then the cosmic string contribution to the CMB can be reduced to less than $10\\%$. This also makes $n_s \\simeq 1$. For $10\\%$ cosmic string contribution, $n_s \\simeq 1$ is actually what we would like to have \\cite{Bevis:2007gh, Battye:2006pk}. However if the cosmic string contribution is further reduced, for example, less than $5\\%$, $n_s \\simeq 1$ is not favored. On the other hand, in the F-term inflation, the constraint for the superpotential coupling from cosmic string study is $\\kappa \\lae 7\\times 10^{-7}$ \\cite{Rocher:2004my, Rocher:2004et}. The above parameters will be defined in the subsequent sections.  It is well-known that there is the $\\eta$-problem \\cite{Copeland:1994vg, Dine:1995uk} in the F-term hybrid inflation. But if we tune the coupling mildly, it can be used as '$\\eta$-correction'. If the correction is negative \\cite{BasteroGil:2006cm}, $n_s=0.96$ can be achieved. In this paper, we show that we can also have $n_s=0.96$ in the case the quadratic correction is positive if we include a negative quartic term. Similar models can be realized in D-term hybrid inflation, where the higher order correction terms can come from non-minimal gauge kinetic function.  It would be interesting if we can have \\emph{both} $n_s=0.96$ and solving the cosmic string problem. In this paper, we show that by considering some generic higher order terms, the spectral index can be reduced to fit the data and the cosmic string problem can also be solved.  This paper is organized as follows. In section \\ref{2}, we introduce a simple parameterization of our model and show the analytic solutions of it. In section \\ref{cosmicstring}, the formation of cosmic string and its contribution to the CMB anisotropy is explained. In section \\ref{3}, we discuss the hilltop version of D-term hybrid inflation. In section \\ref{4}, we discuss the F-term hybrid inflation case. Finally, we present our conclusions. ",
        "conclusions": "\\label{5} In this paper we shown that the higher order corrections to the SUSY hybrid inflation models can have very interesting results if we include quartic terms in addition to the quadratic term. If the quadratic correction is positive and the quartic correction is negative, a hilltop form for the inflaton potential is possible, so that the potential of such a hilltop form makes it possible to have the spectral index reduced to $n_s=0.96$ after some mild tuning of the coupling of the higher order terms. For D-term inflation, if $g \\lae 0.05$, cosmic string problem can be solved with $\\alpha=O(10^{-2})$ and $\\beta=O(1)$. For F-term inflation, if $\\kappa < 0.01$, cosmic string problem can be solved with $\\alpha=10^{-1}-10^{-3}$ and $\\beta=10-10^3$. Which needs more tuning than the case of D-term inflation.  Lower bounds for $g$ and $\\kappa$ are not crucial in this work. The F- and D-term hybrid inflation works even for much lower $g$ or $\\kappa$, which means a lower scale $V_0$ and smaller $\\epsilon$ (a flatter potential). However, this more or less violates the spirit of hybrid inflation which was introduced to prevent small couplings to occur in the successful new inflation \\cite{Linde:1981mu} or chaotic inflation.  The model proposed in this paper can be applied not only restricted to SUSY hybrid inflation. For example, the D3/D7 model \\cite{Dasgupta:2002ew} has an effective description as a D-term inflation model and suffers the problem of producing too much cosmic string contribution to the CMB. If higher order terms are generated in this model, cosmic string problem can be evaded."
    },
    "0812/0812.0358_arXiv.txt": {
        "abstract": "We study the utility of a large sample of type Ia supernovae that might be observed in an imaging survey that rapidly scans a large fraction of the sky for constraining dark energy.  We consider both the information contained in the traditional luminosity distance test as well as the spread in Ia supernova fluxes at fixed redshift induced by gravitational lensing.  As would be required from an imaging survey, we include a treatment of photometric redshift uncertainties in our analysis.  Our primary result is that the information contained in the mean distance moduli of supernovae Ia and the dispersion of supernova Ia distance moduli complement each other, breaking a degeneracy between the present dark energy equation of state and its time variation without the need for a high-redshift ($z \\gsim 0.8$) supernova sample. Including lensing information also allows for some internal calibration of photometric redshifts. To address photometric redshift uncertainties, we present dark energy constraints as a function of the size of an external set of spectroscopically-observed supernovae that may be used for redshift calibration, $\\nspec$. Depending upon the details of potentially-available, external supernova data sets, we find that an imaging survey can constrain the dark energy equation of state at the epoch where it is best constrained $\\wpiv$, with a 1$\\sigma$ error of $\\sigma(\\wpiv) \\approx 0.03-0.09$. In addition, the marginal improvement in the error $\\sigma(\\wpiv)$ from an increase in the spectroscopic calibration sample drops once $\\nspec \\sim \\mathrm{a}\\ \\mathrm{few} \\times 10^{3}$.  This result is important because it is of the order of the size of calibration samples likely to be compiled in the coming decade and because, for samples of this size, the spectroscopic and imaging surveys individually place comparable constraints on the dark energy equation of state.  In all cases, it is best to calibrate photometric redshifts with a set of spectroscopically-observed supernovae with relatively more objects at high redshift ($z \\gsim 0.5$) than the parent sample of imaging supernovae. ",
        "introduction": "\\label{section:introduction}      Type Ia Supernovae (SNeIa) can be used as well-calibrated standard candles with relatively little dispersion in their intrinsic luminosities \\citep[e.g.,][]{phillips93,hamuy_etal95,riess_etal96,prieto_etal06,guy_etal07,jha_etal07}. This has allowed for recent measurements of the relation between cosmological distance and redshift, which provide the strongest contemporary evidence for an accelerating cosmological expansion \\citep[see][for recent applications]{astier_etal06,riess_etal07,wood-vasey_etal07}. Several studies have examined the spread in observed fluxes (rather than their average) due to the magnification of SNeIa as a source of bias in cosmological parameter extraction from the distance--redshift test \\citep{sasaki87,linder_etal88,kantowski_etal95,wambsganss_etal97, holz98,wang99,barber00,holz_linder05}, a cross-check of lensing shear maps \\citep[][and references therein]{cooray_etal06}, or as an interesting signal in its own right \\citep{metcalf99,dodelson_vallinotto06,cooray_etal06,albrecht_etal06}.    The present paper is inspired by these studies and other efforts to explore how SNeIa might be utilized as part of a large, photometric survey \\citep[e.g.,][]{albrecht_etal06,zhan_etal08,hannestad_etal08,zhang_chen08}. The aforementioned studies considered the lensing of a relatively small ($ \\lsim 10^3$) spectroscopic sample of SNeIa, with redshifts that are known precisely \\citep[with errors $\\lsim 10^{-4}$, as prescribed by][]{huterer_etal04}. We examine the utility of a significantly larger ($\\sim 10^6$) photometric sample of supernovae that may be collected by a future, wide, fast, imaging survey such as the proposed Large Synoptic Survey Telescope (LSST)\\footnote{{\\tt http://www.lsst.org}} in which the vast majority of SNeIa will not have accompanying spectroscopic observations. It is thought that the identification and determination of photometric redshifts for SNeIa can be done relatively reliably \\citep{pinto_etal04,wang_etal07}.  Consequently, the dispersion in SNeIa fluxes at fixed {\\em photometric} redshift has three sources: the intrinsic dispersion of SNeIa, including errors in calibrating the standard candles and applying K-corrections and extinction corrections; photometric redshift errors; and any dispersion induced by gravitational lensing. Moreover, each of these factors are thought to contribute to the total dispersion at fixed photometric redshift at comparable levels (see \\S~\\ref{section:methods}). The implication is that a very large sample of supernovae from an imaging survey can be used to overwhelm the intrinsic dispersion, thus enabling the measurement of SNeIa dispersion induced by photometric redshift errors and gravitational lensing which should shed light on both SNeIa photometric redshifts and cosmological parameters at relatively low redshift.    The modest goal of this paper is to make preliminary estimates of the utility of such a photometric SNeIa sample to learn about cosmology from the dispersion among SNeIa fluxes in conjunction with the traditional distance--redshift test. We find that the spread in observed SNeIa fluxes at fixed redshift provides constraints that are not, by themselves, competitive with other probes (such as weak lensing shear), but are yet interesting, subject to unique systematic errors, and obtained without significant additional observational effort if such a sample is to be used to perform the traditional distance--redshift test or to map baryon acoustic oscillations \\citep[as studied in][]{albrecht_etal06,zhan_etal08}. In fact, because such a measurement will be most constraining on dark energy at low redshift ($z \\sim 0.2-0.4$) it will have some degree of complementarity with other measures, such as cosmic shear which best constrains the dark energy equation of state near $z \\sim 0.7$. More importantly, the information contained in the traditional distance--redshift test (most sensitive to the dark energy equation of state near $z \\sim 0.1$) and the dispersion of SNeIa brightnesses complement {\\em each other} and help to make cosmology with photometric SNeIa more robust to uncertainties in photometric redshifts.    Of course, the lack of spectroscopy is a {\\em major} complication in using SNeIa from a photometric survey.  As the use of photometric SNeIa samples for cosmology requires some treatment of photometric redshift uncertainties, we also address the dependence of cosmological constraints upon the calibration of SNeIa photometric redshifts.  In particular, we find that utilizing the dispersion in SNeIa distance moduli allows for mild self-calibration in the uncertainties of photometric redshift distributions.   Perhaps more interestingly, we present results for dark energy constraints as a function of the size of a spectroscopic sample of SNeIa that can be used to calibrate photometric redshifts for two different assumptions about the redshift distribution of this spectroscopic calibration set.  This is a convenient way to express cosmological constraints as a function of prior knowledge of the photometric redshift distribution of SNeIa. Interestingly, we find that if the goal is to seek the best constraint on the equation of state parameter at the epoch where it is best constrained by the data (the so-called ``pivot'' equation of state, see \\S~\\ref{section:methods}, results are similar for a constant equation of state model), then there is decreasing marginal value in expanding the size of the spectroscopic calibration set beyond a few thousand SNeIa, a number of SNeIa that should be achievable in the coming decade.  However, this statement does depend upon model details and upon the assumption that it is the pivot equation of state that is of most interest.   There are several important caveat to our analysis.  We neglect the, perhaps considerable, additional complications in supernova type identification and calibration arising from the lack of spectroscopy.  Aside from neglect of several potentially-important observational realities, our calculation is also an idealized one.  We have treated the influence of gravitational lensing on observed fluxes in a simplified manner, in large part because current analytic models for the proposed lensing signal are not adequate for future surveys and a full treatment of this effect is computationally demanding.  However, part of the purpose of this paper is to demonstrate that such a signal does contain valuable information and to emphasize that as we yet lack the theoretical tools needed to perform a robust analysis of future data, developing such tools should be a high priority in the run-up to forthcoming photometric surveys.    The present paper is organized as follows.  In \\S~\\ref{section:methods} we describe the our treatment of SNeIa observables and photometric redshifts as well as our methods for forecasting cosmological parameter constraints.  We present our primary results in \\S~\\ref{section:results}.  We discuss our results, including implications and important caveats in \\S~\\ref{section:discussion}. ",
        "conclusions": "We have studied simultaneous constraints on dark energy coming from both the evolution of the mean luminosity distance of SNeIa as a function of redshift and the dispersion among SNeIa fluxes in a SNeIa data set that may arise from a large, wide, and fast photometric survey such as that proposed for the LSST.  Sources of dispersion among the measured fluxes of SNeIa in such a survey include intrinsic disperision among SNeIa (including standard candle calibration), uncertainty in photometric redshifts, and magnification due to gravitational lensing.  We have shown that the additional dispersion information complements the traditional luminosity distance test in several ways.   First, the dispersion information breaks a degeneracy between the contemporary dark energy equation of state $\\wzero$, and dark energy equation of state evolution, parameterized here by the common, benchmark parameter $\\wa$.  Using the luminosity distance (mean flux) test alone, this degeneracy can be broken by observing high-redshift SNeIa to increase the lever arm of the data in redshift. The dispersion information leads to constraints that break the degeneracy between $\\wzero$ and $\\wa$ without the nead for a high-redshift ($z \\gsim 0.8$) SNeIa sample (see Fig.~\\ref{fig:w0wa}).   Second, photometric surveys require some calibration of photometric redshifts. Including dispersion information may allow for mild internal self calibration of uncertainties in the true redshift distribution of observed SNeIa given their photometric redshifts.  In the context of our models, this results in constraints on photometric redshift model parameters (see \\S~\\ref{section:methods}) that are a factor of $\\sim 2-5$ more constraining than the priors on these parameters from a spectroscopic sample of size $\\nspec=3000$. The realized level of improved calibration depends upon the details of both the model and the additional SNeIa calibration data that are available, but in all cases the luminosity distance test alone does not serve to calibrate the photometric redshift distribution.   Lastly, it is entirely possible that there will be evolution in the distribution of SNeIa intrinsic luminosities that significantly degrade the forecasts for dark energy constraints that we present (see \\S~\\ref{section:results}). All approaches to SNeIa cosmology, even those that exploit the traditional luminosity distance test alone, rely upon some knowledge and/or assumptions about the intrinsic distribution of SNeIa luminosities to obtain cosmological constraints.  In cases where the evolution of the spread in luminosities at fixed redshift is significant over the redshift range $0.2 \\lsim z \\lsim 0.8$, a photometric survey cannot use lensing to constrain dark energy directly.  However, we have shown that such a survey can measure the evolution of the intrinsic SNeIa luminosity distribution at levels comparable to, or better than, a spectroscopic survey containing a few thousand SNeIa.  This should aid the luminosity distance test by providing a better understanding of possible SNeIa evolution and demographics and, therefore, a better understanding of potential parameter biases and uncertainties.    As we have mentioned in the preceding paragraph, in examining a photometric SNeIa sample, it is necessary to consider uncertainties in the photometric redshift distribution.  We have assumed that some calibration will be done with a spectroscopic sample of an unspecified size, $\\nspec$. Figure~\\ref{fig:nspecfom} gives the expected area of the 95\\% ellipse in the $\\wzero$-$\\wa$ plane, $\\fom$, as a function of $\\nspec$.  $\\fom$ decreases rapidly with $\\nspec$ in all cases because increasing $\\nspec$ allows access to information over a broader range of observed redshifts and helps to break degeneracies in the $\\wzero$-$\\wa$ plane.  Figure~\\ref{fig:nspecwpiv} shows the constraints on the equation of state parameter at that redshift where the data best constrain it, the so-called pivot equation of state $\\wpiv$, as a function of $\\nspec$.  This figure shows that the marginal improvement in $\\sigma(\\wpiv)$ with increasing $\\nspec$ decreases beyond $\\nspec \\sim \\mathrm{a}\\ \\mathrm{few}\\ \\times 10^3$.  This result is interesting because the upper limit on the achievable size of a spectroscopic sample is near $\\nspec \\sim 10^4$ and, moreover, with $\\nspec \\gsim 10^4$ constraints from the spectroscopic sample alone (which has miniscule redshift errors, $\\sigma_z \\sim 10^{-4}$) would begin to dominate over that achievable with a large, photometric sample even in the ideal limit where statistics, and not systematics, dominate the error budget.  Both Fig.~\\ref{fig:nspecfom} and Fig.~\\ref{fig:nspecwpiv} illustrate that at fixed $\\nspec$ it is more fruiful to employ a spectroscopic calibration set weighted more toward high-redshift ($z \\gsim 0.5$) SNeIa than the distribution of SNeIa from the imaging survey.      The previous paragraph closed with an allusion to systematic error, and indeed sources of systematic error are a major caveat to the present study.  The main advantage of a photometric sample is its size and the additional statistical leverage that come with a large sample size.  However, once systematic errors dominate the error budget, it is no longer possible to take advantage of this benefit. What is more, a photometric survey places emphasis on many potential sources of systematic error.  We have accounted for some of these complications.  For example, we have allowed for consistent determination of the photometric redshift distribution and for variation of the average SNeIa luminosity and intrinsic dispersion, including mild evolution.  This is indicative that our results will be somewhat robust to evolution in the intrinsic SNeIa population \\citep[e.g.][]{hamuy_etal95,howell01,sullivan_etal06,howell_etal07,sarkar_etal08b}. Other sources of systematic error are likely to be important.  In fact, the SNAP collaboration, which aims to observe and exploit a spectroscopically-observed SNeIa sample from space, has assumed a fundamental limit in the ability to calibrate supernovae to within $\\delta \\mu \\sim 0.02$~mag over a redshift difference of $\\delta z \\sim 0.1$ to specify their survey \\citep[e.g.][]{kim_etal04}. Yet it remains to be seen whether such a systematic limit will be realized, and it is possible that it may be higher or lower. Further, without spectroscopy, not only are redshifts uncertain, but supernovae type identification, extinction and K-corrections, and standard candle calibration all become more difficult \\citep[e.g.][]{pinto_etal04,prieto_etal06}.  We have not explicitly accounted for any of these, largely because systematic levels are uncertain and modeling them is both complex and dependent upon survey strategy.  These complications place a detailed treatment of systematics beyond the scope of our study (the studies by the DETF, as well as those of \\citealt{zhan_etal08} and \\citealt{hannestad_etal08} neglected a detailed treatment of systematics as well). If the assumed SNAP systematics floor is realized, our forecasts for dark energy equation of state constraints increase by a factor of $\\sim 3$, significantly reducing the additional cosmological information that may be extracted from the photometric sample of SNeIa.    \\begin{figure*}[t!] \\includegraphics[width=8.0cm]{pmu.eps} \\includegraphics[width=8.0cm]{phistrel.eps} \\caption{ The non-Gaussianity of the lensing distribution.   Both panels show the distribution of the shift in distance modulus $\\delta \\mu$ of SNeIa.  The {\\em left} panel is illustrative.  The {\\em solid, smooth} curve shows the distribution of distance moduli for sources at redshift $z=0.8$ due only to lensing.  This result was computed using the fitting formula of \\citet{wang_etal02}.  We display this quantity at relatively high redshifts only because the formula of \\citet{wang_etal02} is no longer internally self-consistent below $z \\sim 0.6$. The {\\em dotted} curve shows a Gaussian distribution with the same mean and dispersion as the lensing distribution.  The Gaussian distribution is truncated at the minimum magnification, corresponding to a maximum shift in distance modulus $\\delta \\mu_{\\mathrm{max}}$.  The histogram shows the distribution of distance modulus shifts from a random realization of $19,000$ SNeIa at $z=0.8$. The {\\em right} panel shows something closer to what would be seen in any observed set of SNeIa.  In this panel, we show the distribution of distance moduli (relative to the mean) for a random sample of SNeIa in a bin of width $\\delta z = 0.03$ at $z=0.8$.  The sample consists of $19,000$ SNeIa, which is roughly the number of SNeIa expected in such a bin according to our fiducial model described in \\S~\\ref{section:methods}.  In this case, the distribution includes the intrinsic dispersion, the dispersion caused by poorly-controlled photometric redshifts, and the lensing distribution.  For clarity, the probability distribution is plotted relative to the same distribution under the assumption of the lensing contribution is a Gaussian, $P_{\\mathrm{Gaussian\\ Lens}}(\\delta \\mu)$. The error bars reflect the statistical error on $P(\\delta \\mu)$ from a sample of this size and portray of the detectability of the unique shape induced by lensing in the absence of significant errors. } \\label{fig:plens} \\end{figure*}     As described in \\S~\\ref{section:methods}, the calculations we have performed to estimate the additional information contained in the variety of SNeIa distance moduli at a fixed redshift are greatly simplified. Our idealization may lead to an over-estimate of the constraining power of SNeIa.  However, we have also neglected some potentially-available information.  In particular, the distribution of lensing magnifications is highly non-Gaussian \\citep[][]{schneider_wagoner87,sasaki87,wambsganss_etal97,munshi_jain00,wang_etal02} and we have neglected any information beyond the dispersion because including the full information is computationally-intensive and current theoretical estimates of the magnification distribution are inadequate for our application. The left panel of Figure~\\ref{fig:plens} shows the non-Gaussianity of the lensing distribution at a fixed redshift $z=0.8$ using the fitting formula of \\citet{wang_etal02}.  The right panel of Figure~\\ref{fig:plens} shows the distribution of distance moduli in a bin of width $\\delta z = 0.03$ centered at redshift $z=0.8$.  In our fiducial model, the bin contains roughly $n \\approx 19,000$ SNeIa and Fig.~\\ref{fig:plens} demonstrates that the non-Gaussianity of the lensing distribution would be detectable even by examining the distribution of distance moduli in this single bin. For the realization shown, the non-Gaussianity of the lensing distribution is detectable at slightly more than $\\sim 2\\sigma$ from this bin alone. This strongly suggests that the lensing deviation from Gaussianity may be detectable in a complete analysis of a future photometric survey, where the full shape of the magnification distribution would be utilized and where all supernovae could be considered (which requires better predictions).  This is important for two reasons. First, the unique distribution of fluxes induced by lensing may aid in distinguishing the gravitational lensing effect from other additional sources of dispersion and from evolution in the intrinsic properties of SNeIa.  Second, to the degree that this non-Gaussianity can be detected, it will bring additional information to bear and improve dark energy constraints.     To conclude, a future imaging survey that scans a large fraction of the sky rapidly will enable the discovery of $\\sim 10^5 - 10^6$ SNeIa.  The survey of the LSST is the canonical example of such an endeavor.  The number of SNeIa with light curves of sufficient quality to be used for cosmology depend upon specific survey strategies; however, such a SNeIa survey may be able to constrain dark energy properties using both the traditional luminosity distance test and the spread of supernovae apparent brightnesses as a function of photometric redshift.  Unfortunately, current theoretical treatments of gravitational lensing, in particular estimates of the magnification distribution, are not reliable enough to analyze any such data set.  In the coming years, it will be necessary to refine these estimates, including potential theoretical uncertainties \\citep{huterer_takada05,zhan_knox04,white04,rudd_etal08}. In the end, the utility of such a sample to constrain dark energy is subject to limitations of systematic error and SNeIa evolution.  However, if observational systematics or SNeIa evolution prove to be limiting, such a large photometric survey will provide useful information about these issues.  In either case, an LSST-like imaging survey will revolutionize SNeIa cosmology."
    },
    "0812/0812.3392_arXiv.txt": {
        "abstract": "As supernova remnants expand, their shock waves are freezing in and compressing the magnetic field lines they encounter; consequently we can use supernova remnants as magnifying glasses for their ambient magnetic fields. We will describe a simple model to determine emission, polarization, and rotation measure characteristics of adiabatically expanding supernova remnants and how we can exploit this model to gain information about the large scale magnetic field in our Galaxy. We will give two examples: The SNR DA530, which is located high above the Galactic plane, reveals information about the magnetic field in the halo of our Galaxy. The SNR G182.4+4.3 is located close to the anti-centre of our Galaxy and reveals the most probable direction where the large-scale magnetic field is perpendicular to the line of sight. This may help to decide on the large-scale magnetic field configuration of our Galaxy. ",
        "introduction": "Recently, there have been several studies of the Milky Way's magnetic field by the means of observing the rotation measure of compact polarized objects like extra-galactic point sources (e.g. \\cite{brown07}) or pulsars (see talks by A. Noutsos and J.-L. Han). However, we still do not know the large-scale magnetic field configuration or even the number of field reversals within our Galaxy. One problem is that through Faraday rotation studies of extra-galactic point sources we only derive the average magnetic field parallel to the line of sight $B_\\parallel$ through our Galaxy weighted by the electron density $n_e$, because the rotation measure $RM$ is given by: \\begin{equation} RM = 0.81 \\int_l B_\\parallel n_e dl, \\label{RM} \\end{equation} here $RM$ is given in rad/m$^2$, $B$ in $\\mu$G, $n_e$ in cm$^{-3}$, and the pathlength $l$ in pc. In addition extragalactic sources may suffer from intrinsic Faraday rotation of unknown magnitude. Faraday rotation studies of pulsars average $B_\\parallel$ between us and the pulsar weighted by $n_e$. In addition pulsars distances are usually quite uncertain and pulsars probe only a very small area in space since they are only about 15\\,km in diameter. One big problem of the averaging procedure is that there could be numerous field reversals along the line of sight, which would be averaged out. This ambiguity could be solved if we had anchor points for the magnetic field within our Galaxy. We propose to determine these anchor points with polarization and Faraday rotation studies of supernova remnants (SNRs), since these can be used as magnifying glasses of their ambient magnetic field. ",
        "conclusions": ""
    },
    "0812/0812.3995_arXiv.txt": {
        "abstract": "We study global non-axisymmetric oscillation modes trapped near the inner boundary of an accretion disc. Observations indicate that some of the quasi-periodic oscillations (QPOs) observed in the luminosities of accreting compact objects (neutron stars, black holes and white dwarfs) are produced in the inner-most regions of accretion discs or boundary layers. Two simple models are considered in this paper: The magnetosphere-disc model consists of a thin Keplerian disc in contact with a uniformly rotating magnetosphere with and low plasma density, while the star-disc model involves a Keplerian disc terminated at the stellar atomosphere with high density and small density scale height. We find that the interface modes at the magnetosphere-disc boundary are generally unstable due to Rayleigh-Taylor and/or Kelvin-Helmholtz instabilities. However, differential rotation of the disc tends to suppress Rayleigh-Taylor instability and a sufficiently high disc sound speed (or temperature) is needed to overcome this suppression and to attain net mode growth. On the other hand, Kelvin-Helmholtz instability may be active at low disc sound speeds.  We also find that the interface modes trapped at the boundary between a thin disc and an unmagnetized star do not suffer Rayleigh-Taylor or Kelvin-Helmholtz instability, but can become unstable due to wave leakage to large disc radii and, for sufficiently steep disc density distributions, due to wave absorption at the corotation resonance in the disc. The non-axisymmetric interface modes studied in this paper may be relevant to the high-frequency QPOs observed in some X-ray binaries and in cataclysmic variables. ",
        "introduction": "Quasi-periodic variabilities have been observed in the timing data of various types of accreting objects. Several types of quasi-periodic oscillations (QPOs) are observed in X-ray binaries with accreting black holes (BHs) or neutron stars (NSs) (e.g., Remillard \\& McClintock 2006; Van der Klis 2006). Oscillations are also seen in the outbursts of accreting white dwarf (WD) systems (e.g., Patterson 1981; see Warner 2004 for a review).  In accreting NS and BH X-ray binaries the observed QPO frequencies ($40-450$ Hz for the high-frequency QPOs in the BH systems and $\\gtrsim 300$ Hz  for kHz QPOs in the NS systems) imply a source close to the central compact object where the Keplerian orbital frequencies are high. Since the BH systems lack a hard surface where oscillations may occur, it is likely that the source of the variability is in the inner regions of the disc itself or in some interface regions between the disc and the plunging flow. Gilfanov et al (2003), however,  found that, based on spectral analysis of the disc emission components, the quasi-periodic variability in Low Mass NS X-ray Binary systems are most likely caused by variations in the disc boundary layer, rather than the disc itself.  In Cataclysmic Variables (CVs) the Dwarf Nova Oscillations (DNOs) seen during during outbursts  have frequencies roughly corresponding to the Keplerian rotation rate at the WD surface (e.g., Patterson, 1981; Warner, 2004; Knigge et al, 1998), which imply an origin at or near the inner disk boundary.  Several models involving accretion disk boundary dynamics have been proposed in different contexts. Popham (1999) studied the effect of a non-axysymmetric bulge at the optically thick to optically thin transition radius as a model for DNOs. Piro \\& Bildsten (2004) examined the surface wave oscillations that would occur within the thin equatorial belt around a non-magnetized WD formed by the accretion spreading layer, while Warner \\& Woudt (2002) considered accretion onto a slipping belt. In the context of accreting magnetic (neutron) stars, Arons \\& Lea (1976) and Elsner \\& Lamb (1977) considered the interchange instability at the magnetosphere boundary. Spruit \\& Taam (1990) and Spruit, Stehle \\& Papaloizou (1995) investigated the stability of thin rotating magnetized discs. Of particular relevance to the present paper is the work of Li \\& Narayan (2004), who examined a simplified cylindrical model of the Rayleigh-Taylor and Kelvin-Helmholtz instabilities at the boundary between a magnetosphere and an incompressible rotating flow. There have also been a number of numerical simulations of the interface at the magnetosphere-disc boundary (see Romanova et al., 2008 and Kulkarni \\& Romanova, 2008 and references therein).  In this paper we study global non-axisymmetric oscillation modes confined near inner boundary of the accretion disc (interface modes). We consider two simple models. The first model involves the magnetosphere-disc boundary similar to the model of Li \\& Narayan (2004): we consider an uniformly rotating incompressible magnetosphere with low gas density (where magnetic pressure dominates), which truncates a thin barotropic accretion disc (where gas pressure dominates). This situation may arise from magnetic field build up due to accretion (e.g., Bisnovatyi-Kogan \\& Ruzmaikin, 1974, 1976; Igumenshchev et al., 2003; Rothstein \\& Lovelace, 2008) or by the magnetosphere of a central (neutron) star. Unlike Li \\& Narayan (2004), who restricted their model to incompressible fluid, our discs are compressible and we show that because of the differential rotation of the disc, finite disc sound speed plays an important role in the development of the instability of the interface modes.  In our second model we examine the interface modes for accretion onto a non-magnetic stellar surface. Though the structure of the boundary layer is non-trivial and may affect boundary modes (see, e.g., Carroll et al., 1985; Collins et al., 2000), we consider the instabilities for a thin disc truncated by a sharp transition to a dense uniformly rotating stellar atmosphere. This simplified model may provide insight for modes with characteristic radial length scale much greater than the radial length scale of the boundary layer.  In Section 2 we describe the basic setup  for the magnetospheric boundary model, and in Section 3 we discuss the resulting interface mode instabilities. We describe the star-disc boundary and analyze its possible instabilities in Section 4. We then conclude in Section 5 with a discussion of possible applications of our findings. ",
        "conclusions": "We have studied the non-radial oscillation modes at the interface between an accretion disc and a magnetosphere or stellar surface. Although the models explored in this paper are perhaps too simplified compared to realistic situations, they offer some insight into the behavior of the interface modes in various astrophysical contexts.  Our study of the interface modes at the magnetosphere-disc boundary extended the work by Li \\& Narayan (2004), who considered incompressible disc flow (and therefore could not treat real discs). The model can have very strongly unstable modes due to Rayleigh-Taylor and Kelvin-Helmholtz instabilities. In systems where the magnetosphere has developed from advection of frozen magnetic flux, the magnetosphere is expected to be roughly rotating with at the Keplerian rate. Since there is no shear at the interface, only the Rayleigh-Taylor instability may occur. However the disc vorticty (due to differential rotation) acts to suppress the instability, leading to a cutoff below a critical sound speed. Thus a sufficiently hot disc is required to generate unstable low-$m$ modes. For magnetospheres rotating with the central star, shearing is expected between the magnetosphere and disc, and the Kelvin-Helmholtz instability becomes active. This can help to drive the instability for low-$m$ modes to overcome the vorticity in low sound-speed discs.  In discs that terminate near the ISCO in a general relativistic potential, the vorticity approaches zero at the inner disc radius, and unstable interface modes can be found for any sound speed.  We can expect a strong dependence of the interface mode growth rates on the sound speed, and thus accretion rate, while the real frequencies of these oscillations remain close to $(1-1.3) m\\Omega_{\\rm in}$, and do not depend strongly on the sound speed. It is worth noting that the same boundary condition that gives rise to the interface modes also gives rise to intertial-acoustic modes (or p-modes) in the disc.  Although higher-$m$ interface modes are more unstable in our model, these are less likely to be observed due to the averaging out of the luminosity variation over the observable emitting region. In addition, if the effect of viscous damping is considered (e.g., Wang \\& Robertson 1985), small wavelength or high-$m$, perturbations are suppressed.  These results are for perturbations with no vertical structure, and are applicable mainly to the midplane of accretion discs interacting with a magnetosphere. Global 3D numerical studies of Rayleigh-Taylor instability induced accretion onto magnetized stars have been performed by Romanova, Kulkarni \\& Lovelace (2008) and Kulkarni \\& Romanova (2008), and show such small $m$ instabilities in the disk midplane. The low-$m$ oscillations at the magnetosphere-disc interface may be relevent to the high-frequency QPOs observed in some NS and BH X-ray binary systems (Li \\& Narayan, 2004; see Section 1 of Lai \\& Tsang for a critical review of various theoretical models), although to obtain the correct QPO frequencies for the BH systems, the disc inner radius must lie outside the inner-most stable circular orbit.  For the star-disc boundary model considered in in this paper, the interface mode growth rates are much smaller than for the magnetospheric case, since the effective gravity now acts to stabilize the system, and the Rayleigh-Taylor instability is inactive. The modes discussed here are unstable due primarily to propagation through the corotation. With sufficiently steep disc density profile ($\\Sigma \\propto r^{-p}$ with $p > 3/2$), corotation absorption can also help to drive these modes, as studed previously by Tsang \\& Lai (2008) and Lai \\& Tsang (2008). Such modes may be responsible for the high-frequency (of order the Keplerian frequency at the stellar surface) dwarf nova oscillations observed in CVs, although oscillations with longer periods would require a different explanation."
    },
    "0812/0812.4860_arXiv.txt": {
        "abstract": "The evolutionary scenarios which are commonly accepted for PG\\,1159 stars are mainly based on numerical simulations. These stellar evolution scenarios have to be tested and calibrated with real objects with known stellar parameters. One of the most crucial but also quite uncertain parameters is the stellar mass. PG\\,1159 stars have masses between 0.5 and 0.8 solar masses, as derived from asteroseismic and spectroscopic determinations. Such mass determinations are, however, themselves model-dependent. Moreover, asteroseismologically and spectroscopically determined masses deviate systematically for PG\\,1159 stars by up to 10\\,\\%. SDSS\\,J212531.92-010745.9 is the first known PG\\,1159 star in a close binary with a late main sequence companion allowing a dynamical mass determination. The system shows flux variations with a peak-to-peak amplitude of about 0.7 mag and a period of about 6.96\\,h. In August 2007, 13 spectra of SDSS\\,J212531.92-010745.9 covering the full orbital phase range were taken at the TWIN 3.5\\,m telescope at the Calar Alto Observatory (Alm\\'{e}ria, Spain). These confirm the typical PG\\,1159 features seen in the SDSS discovery spectrum, together with the Balmer series of hydrogen in emission (plus other emission lines), interpreted as signature of the companion's irradiated side. A radial velocity curve was obtained for both components. Using co-added radial-velocity-corrected spectra, the spectral analysis of the PG\\,1159 star is being refined. The system's lightcurve, obtained during three seasons of photometry with the G\\\"ottingen 50cm and T\\\"ubingen 80cm telescopes, was fitted with both the \\texttt{NIGHTFALL} and \\texttt{PHOEBE} binary simulation programs. An accurate mass determination of the PG\\,1159 component from the radial velocity measurements requires to first derive the inclination, which requires light curve modelling and yields further constraints on radii, effective temperature and separation of the system's components. From the analysis of all data available so far, we present the possible mass range for the PG\\,1159 component of SDSS\\,J212531.92-010745.9. ",
        "introduction": "% \\subsection{Evolutionary scenarios and masses for PG\\,1159 stars} PG\\,1159 stars are hot hydrogen-deficient (pre-)white dwarfs with effective temperatures between 75\\,000 and 200\\,000\\,K, and $\\log{(g/\\textrm{cm}\\,\\textrm{s}^{-2})}$\\,=\\,5.5--8.0 (Werner 2001). They are in the transition between the asymptotic giant branch (AGB) and cooling white dwarfs. Spectra of PG\\,1159 stars are dominated by absorption lines of He\\,{\\sc ii}, C\\,{\\sc iv} and O\\,{\\sc vi}. Roughly half of them have a planetary nebula. Current theory suggests (e.\\,g.~Werner 2001) that they are the outcome of a late or very late helium-shell flash, a phenomenon that drives the currently observed fast evolutionary rates of three well-known objects (FG~Sge, Sakurai's object, V605 Aql). Flash-induced envelope mixing produces a H-deficient stellar surface. The photospheric composition then essentially reflects that of the region between the H- and He-burning shells in the precursor AGB star. The He-shell flash forces the star back onto the AGB. The subsequent, second post-AGB evolution explains the existence of Wolf-Rayet central stars of planetary nebulae and their successors, the PG\\,1159 stars. These in turn are the progenitors of the helium rich WD sequence. Currently, 40 PG\\,1159 stars are known, Figure\\,\\ref{fig:tracks} shows their positions in a log\\,$T_{\\rm eff}$-$\\log g$-diagram. \\par Recent evolutionary models for PG\\,1159 stars (Miller Bertolami \\& Althaus 2006b, Werner \\& Herwig 2006) lead to spectroscopic masses which still suffer from serious uncertainties concerning the efficiency of the overshooting process in the He-flash driven convection zone during the thermal pulse phase on the AGB. This efficiency is represented in the models with an a priori unknown free parameter. The other ingredient for the determination of spectroscopic masses, the stellar atmosphere models, have to deal with uncertainties in the line broadening theory, so lacking reliable physical input there makes the determination of $\\log{g}$ particularly difficult. \\par The overall resulting uncertainties in model masses manifest themselves as discrepancies between current asteroseismological masses and spectroscopic masses, as is evident from Table~3 in Werner \\& Herwig (2006; an updated version has been adopted as Table~1 in Schuh \\& Nagel 2007). The determination of a dynamical mass as an alternative method and test case for the models is currently restricted to only one suitable object: SDSS\\,J212531.92$-$010745.9. \\begin{figure} \\begin{center} \\includegraphics[width=0.85\\textwidth]{mbtracks.eps} \\end{center} \\caption{\\label{fig:tracks}Positions of PG\\,1159 stars with known $T_{\\rm eff}$ and $\\log{g}$ in the log $T_{\\rm eff}$-$\\log{g}$-diagram. The current state-of-the-art post-AGB evolutionary tracks, first presented by Miller Bertolami et al.\\ (2006a), as well as the positions of known PG1159 stars, are adopted from Figure~3 in Miller Bertolami \\& Althaus (2006b, dotted lines). The masses correspond to evolutionary sequences for 0.87, 0.664, 0.609, 0.584, 0.564, 0.542, 0.53, and 0.512~M$_\\odot$ from left to right. We also give lines of constant radii for the theoretical tracks (dashed lines with annotations in R$_\\odot$). The position of SDSS\\,J212531.92$-$010745.9 found by Nagel et al.\\ (2006) is shown as a circle; the position we find from our new spectral analysis (see section~\\ref{sec:spectralanalysis}) is indicated by the error ellipse. } \\end{figure} \\begin{figure} \\begin{center} \\includegraphics[width=0.9\\textwidth]{vorhandenemodelle_bw.eps} \\end{center} \\caption{\\label{fig:modelgrid}The model atmospheres were calculated with \\texttt{NGRT} (Werner et al.\\ 2003). Shown are the models calculated by end of October 2008 in the $\\log{g} - T_{\\textrm{eff}}$-plane for six different \\textsc{C}-abundances represented by different plot symbols (squares: \\textsc{C/He}=0.01, triangles: \\textsc{C/He}=0.03, diamonds: \\textsc{C/He}=0.05, crosses: \\textsc{C/He}=0.07, asterisks: \\textsc{C/He}=0.10, plus signs: \\textsc{C/He}=0.13). Other elemental abundances (all by number) are fixed at \\textsc{N/He}=0.01 and \\textsc{O/He}=0.01. The grid is still being completed: The whole range from $T_{\\textrm{eff}}=60\\ldots100\\,\\textrm{kK}$ in steps of 5 kK and from $\\log{(g/\\textrm{cm}\\,\\textrm{s}^{-2})}=5.6\\ldots8.0$ in steps of 0.2\\,dex will be filled with models in all six \\textsc{C}-abundances by December 2008. Model spectra were calculated from the atmosphere models using \\texttt{profile} and were convolved with a gauss profile to match the spectral resolution of $1.2\\,\\textrm{\\AA}$ of the Calar Alto TWIN Spectograph in the configuration used. } \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=0.9\\textwidth]{optimum_bw_wl2.eps} \\end{center} \\caption{\\label{fig:spectralfit} Cross-correlating specific spectral features of the Calar Alto spectra radial velocities were deduced. The three panels show three distinctive PG1159 features in the median spectrum (solid black line) along with the two model spectra (grey lines), one of which (dotted line) belongs to the model atmosphere with $\\log{(g/\\textrm{cm}\\,\\textrm{s}^{-2})} = 7.60$ and $T_{\\textrm{eff}}=90000\\,\\textrm{K}$ and $\\textsc{C/He}=0.05$ and was found as first preliminary result by Nagel et al.\\ (2006). The solid grey line is the model spectrum with the highest cross-correlation value found by the quantitative spectral analysis with the extended model grid ($\\log{(g/\\textrm{cm}\\,\\textrm{s}^{-2})} = 7.1$, $T_{\\textrm{eff}}=72500\\,\\textrm{K}$, $\\textsc{C/He}=0.07$). The lowermost panel shows the narrow \\textsc{Civ} lines used for the radial velocity determination of the PG1159 star. } \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=0.86\\textwidth]{RVplot_relativ_gewichtet_bw.eps} \\end{center} \\caption{\\label{fig:radialvelocities}Radial velocity curves for the PG\\,1159 star (triangles, light gray) and the companion (diamonds, black) determined from spectra obtained with TWIN at Calar Alto. The best fit sine curves yield semi-amplitudes of $k_1 = 94.3 \\pm 15.0\\, \\textrm{km}\\,\\textrm{s}^{-1}$ for the PG\\,1159 (light gray sine curve) and $k_2 = 113.0 \\pm 3.0\\, \\textrm{km}\\,\\textrm{s}^{-1}$ for the the companion (black sine curve). The zero points of the two velocity curves are shown as dashed lines, the dotted lines indicate their uncertainties. The shift of the zero points relative to each other by about 11\\,km\\,s$^{-1}$ could be explained by gravitational redshift of the PG\\,1159 star but is not significant. } \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=0.86\\textwidth]{lc_phoebe.ps} \\end{center} \\caption{\\label{fig:lightcurve}The observed light curve of all nights from 2005 to 2007 folded onto the orbital period and the simulated light curve of a binary system, consisting of a PG\\,1159 star and a heated M dwarf, calculated with \\texttt{PHOEBE} for $m_1$=0.55\\,M$_\\odot$, $\\log{(g/\\textrm{cm}\\,\\textrm{s}^{-2})}$=7.1, $R_1$=0.035\\,R$_\\odot$(black line). } \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=0.9\\textwidth]{msin3i_constraints.eps} \\end{center} \\caption{\\label{fig:msin3iconstraints}Using the photometric period and the mass ratio from the radial velocity amplitudes, the masses can be obtained as a function of the inclination. For the radial velocity curve solution in Figure\\,\\ref{fig:radialvelocities}, the ($m_1$,$m_2$) pairs lie on the line (error range: dashed) shown above, with $m_{1,\\textrm{min}}$=0.15\\,M$_{\\odot}$ and $m_{2,\\textrm{min}}$=0.12\\,M$_{\\odot}$. The dark areas mark forbidden ranges of the inclination: $m_1<1.44$\\,M$_{\\odot}$ requires $i>20^{\\circ}$, and the fact that the system is not eclipsing sets a limit of $i< 80^{\\circ}$, implying $m_1>0.15$\\,M$_{\\odot}$; more stringent constraints are clearly required. A comparison with the evolutionary tracks in Figure\\,\\ref{fig:tracks} (although these may not be appropriate for a presumed PCEB system) favours a higher inclination (lower masses) than assumed in Figure\\,\\ref{fig:lightcurve}. } \\end{figure}  \\subsection{Discovery of one PG\\,1159 close binary among many SDSS white dwarf binaries} The important role of binaries containing white dwarfs, for dynamical mass and other parameter determinations, is evident from the number of publications devoted to the topic in these proceedings. The contributions by Heller et al.\\ (2009), Rebassa-Mansergas et al.\\ (2009), Schreiber et al.\\ (2009), and further contributions on specific individual systems, reflect the remarkable increase in number of known white dwarfs with companions that the Sloan Digital Sky Survey (SDSS) has brought about. Among the PG\\,1159 stars, however, SDSS\\,J212531.92$-$010745.9 is the only known close binary system (Nagel et al.\\,2006). Eisenstein et al.\\ (2006) also list it as a spectroscopically confirmed white dwarf from the SDSS. \\par The spectrum of SDSS\\,J212531.92$-$010745.9 shows significant features that are typical for PG\\,1159 stars, for example the strong C\\,{\\sc iv} absorption lines at $4650-4700$\\AA\\ and He\\,{\\sc ii} at $4686$\\AA. Furthermore, the spectrum shows the Balmer series of hydrogen in emission. We have shown from our extensive follow-up photometry and modelling that this is due to a close cool companion which is heated up by irradiation from the hydrogen-deficient PG\\,1159 star. Thus, the radial velocities for both stars can be measured in this double-lined system. \\par Nagel et al.\\ (2006) first reported a photometric period of 6.95573(5)\\,h with an amplitude of 0.354(3)\\,mag which represents the orbital period of the binary system.  From an initial comparison of the SDSS archive spectrum (Data Release~4) with NLTE model spectra they derived, as preliminary results, an effective temperature of 90\\,000\\,K, $\\log{(g/\\textrm{cm}\\,\\textrm{s}^{-2})}\\,\\sim$\\,7.60 and the abundance ratio \\textsc{C/He}\\,$\\sim$\\,0.05 for the PG\\,1159 component. They simulated the light curve of the binary system using \\texttt{NIGHTFALL}, showing that in in addition to the constant contribution from the PG\\,1159, the irradiated side of the companion leads to a periodic brightening due to the reflection effect. For the inclination they obtained $(70\\pm5)\\,^{\\circ}$.  The system is not eclipsing and has not been found to be pulsating. While their solution for a mean radius of $(0.4\\pm0.1)\\,\\rm R_{\\odot}$, a mass of \\,$(0.4\\pm0.1)\\,\\rm M_{\\odot}$ and a temperature of the irradiated surface of about 8\\,200\\,K for the companion was in good agreement with the observed light curve, a manifold of solutions remains equally plausible as long as we cannot reliably and indepedently fix the PG\\,1159 mass and radius. \\par We note that in the SDSS Data Release~6 catalogue Adelman-McCarthy et al.\\ (2008) quote a proper motion of +5(3)\\,mas\\,yr$^{-1}$ in RA and +4(3)\\,mas\\,yr$^{-1}$ in DEC, a value for the redshift corresponding to 138(60)\\,km\\,s$^{-1}$, and Sloan colours of u'=17.144(0.009)\\,mag, g'=17.538(0.005)\\,mag, r'=17.757(0.006)\\,mag, i'=17.794(0.008)\\,mag, z'=17.828(0.020)\\,mag. \\subsection{Orbit ephemeris from photometric observations} From time-series photometry of SDSS\\,J212531.92$-$010745.9 performed in 2005, 2006 and 2007 with the T\\\"ubingen 80\\,cm and the G\\\"ottingen 50\\,cm telescopes, Schuh et al.\\ (2008b) find the ephemeris of the predicted \\emph{maxima} times to be ~ \\mbox{$ {\\rm HJD} = 2454055\\hbox{$.\\!\\!^{\\rm d}$} 2134(4)  +    0\\hbox{$.\\!\\!^{\\rm d}$} 289822(2) \\cdot E. $} Further multi-colour observations are available, which show that the photometric amplitude in the optical is roughly 35\\% smaller in the blue than in the red. ",
        "conclusions": ""
    },
    "0812/0812.3737_arXiv.txt": {
        "abstract": "We perform a time-resolved spectral analysis of bright, long Gamma-ray burst GRB 061007 using Suzaku/WAM and Swift/BAT. Thanks to the large effective area of the WAM, we can investigate the time evolution of the spectral peak energy, $E^{\\rm t}_{ \\rm peak}$ and the luminosity $L^{\\rm t}_{\\rm iso}$ with 1-sec time resolution, and we find that the time-resolved pulses also satisfy the $E_{\\rm peak}-L_{\\rm iso}$ relation, which was found for the time-averaged spectra of other bursts, suggesting the same physical conditions in each pulse. Furthermore, the initial rising phase of each pulse could be an outlier of this relation with higher $E^{\\rm t}_{\\rm peak}$ value by about factor 2. This difference could suggest that the fireball radius expands by a factor of $2-4$ and/or bulk Lorentz factor of the fireball is decelerated by a factor of $\\sim$ 4 during the initial phase, providing a new probe of the fireball dynamics in real time. ",
        "introduction": "Many characteristics of Gamma-ray Bursts (GRBs) have been revealed by the previous observations, and now it is widely believed that GRBs are the one of the most powerful explosion in the Universe, originating at cosmological distances. The association with the energetic supernovae was found for some long duration GRBs. This provides a strong evidence that the progenitor of the long GRBs is the core collapse of a massive star. The observed prompt gamma-ray spectra are often fitted well by a smoothly connected broken power-law called as the Band function (\\cite{Band 1993}). However, we are still not able to uniquely associate the measured properties to the radiation mechanism of the prompt gamma-ray emission.  Although, the observed behavior of GRBs, such as the intensity and the peak energy of the $\\nu F _\\nu$ spectrum ($E_{\\rm peak}$) are dramatically different from burst to burst, some clear correlations between these parameters were reported, such as the correlation between isotropic equivalent radiation energy $E_{\\rm iso}$ and the peak energy in  the rest frame of the burst ($E_{\\rm peak,src}$) (\\cite{Amati 2002}; \\cite{Amati 2006}), or the isotropic equivalent peak luminosity $L_{peak,\\rm iso}$ and the $E_{\\rm peak,src}$ (\\cite{Yonetoku 2004}), and the total radiation energy corrected for the jet opening angle and the $E_{\\rm peak,src}$ (\\cite{Ghirlanda 2004b}). These relations are often used to estimate the inferred redshift value, and Yonetoku et al. (2003) estimated redshift for large BATSE sample and found that some of them have large redshift with z $>$ 10. Such high-z GRBs are important to constrain the cosmological parameters. Furthermore, the same kind of correlation have been reported not only for each burst but also for individual pulses of a burst internally. Several authors studied this hardness-intensity correlation (HIC) for time-resolved spectral analysis. Golenetski et al. 1983 found a power-law correlation between observed energy flux and the spectral characteristic energy E$_0$, and the power-law correlation index was found to be a typical value of 1.5-1.7. Kargatis et al. (1994) confirmed this correlation but the power-law index showed a larger spread. They obtained the power-law index of 2.2$\\pm$1.0. Strohmayer et al. (1998) also found that this power-law correlation is valid using Ginga data. Borgonovo \\& Ryde (2001) also studied this correlation by strong BATSE sample. They found the specific feature that there is the sharp transition from one power-law track to the second track with the same power-law index in the decay phase of the pulse. They explained this track-jump behavior as the transition to a new pulse. The fact that the HIC holds for both each pulse and within the pulse implies that the same physical condition is realized in each pulse. Therefore, to investigate this relation it is very important to consider the radiation mechanism of the GRB prompt emission. Recently, Liang et al. (2004) analyzed the $L_{\\rm iso}$ and $E_{\\rm peak,src}$ of the 2408 time-resolved spectra for 91 BATSE GRBs, assuming that the burst rate as a function of redshift is proportional to the star formation rate. Yoshida et al. (2007) measured the $L_{\\rm iso}$ and $E_{\\rm peak,src}$ value for time-resolved spectra of HETE-2 bursts. Both of them found that $L_{\\rm iso} \\propto E_{\\rm \\rm peak,src}^2$, and Yoshida et al. (2007) found that this relation is tighter for each spike-averaged spectra rather than that of simply 5.2 s-divided spectra. Therefore, this relation could be different for finer time-resolved spectra. However, most of previous analysis was done using data for each pulse or for the decay phase of such pulses, while the detailed properties of time-resolved spectra, including initial rising phase for each pulse have not been considered and thus are still ambiguous. In this paper, we report on the time-resolved spectroscopy of the bright long GRB 061007 observed by the Suzaku/WAM and Swift/BAT. This burst was also observed by many ground-based telescopes and the redshift was measured to be 1.261  (\\cite{Osip 2006}; \\cite{Jakobsson 2006}), and Schady et al. (2007) suggested that this burst has an early jet break time which is within 80 s of the prompt emission, and a highly collimated outflow is likely case for this burst. Using this bright burst with a measured redshift, we can investigate the relation between $L_{\\rm iso}$ and $E_{\\rm peak,src}$ within each pulse, including both the pulse-rise and the decay phase.  This is possible since Suzaku/WAM has the largest effective area from 300 keV to 5 MeV than any previous missions. Furthermore, since Swift/BAT covers an energy range from 15 keV to 150 keV, the joint analysis using Suzaku/WAM and Swift/BAT gives us further constraints on the spectral parameters. Therefore, this joint, broad-band Suzaku/WAM and Swift/BAT analysis allows the excellent time-resolved spectral analysis with both fine time resolution and high sensitivity.  We also investigate the time-averaged properties of this burst such as $E_{\\rm iso}$ -- $E_{\\rm peak,src}$  (Amati) relation and $L_{\\rm iso}$ -- $E_{\\rm peak,src}$ (Yonetoku) relation in order to compare with other bursts. Thus, we denote our ``time-resolved''  peak energy and luminosity relation as  $E^{\\rm t}_{ \\rm peak,src}$ -- $L^{\\rm t}_{ \\rm iso}$  relation in this paper to distinguish from those time-averaged relationships. ",
        "conclusions": "\\subsection{$E^{\\rm t}_{ \\rm peak}$--$L^{\\rm t}_{ \\rm iso}$ relation for time-resolved spectra}  From time-resolved spectral analysis, we find that the time-resolved $E_{\\rm peak}$; $E^{\\rm t}_{ \\rm peak}$ changes are correlated with the burst intensity. In order to confirm this, we perform a correlation analysis between $E^{\\rm t}_{ \\rm peak}$ and time resolved luminosity; $L_{\\rm iso}$ for time-resolved spectra. The $L^{\\rm t}_{ \\rm iso}$ value is measured as the 1-sec time-averaged value for each of the time-resolved spectra in 30 keV to 10000 keV energy range. Figure \\ref{EpLiso} shows the relation between $E^{\\rm t}_{ \\rm peak}$ and $L^{\\rm t}_{ \\rm iso}$. We can see a clear positive correlation between these two parameters. To quantify this, we calculate the Spearman rank-order coefficients ($r_s$), and we obtain the correlation coefficient of 0.78. This coefficient corresponds to the significance probability that there is no correlation of 5.4 $\\times$ 10$^{-13}$. Therefore, we confirm that the correlation surely exists at $> 3 \\sigma$ confidence level. We then fit this correlation by the simple power-law model like the Hardness-Intensity correlation, the Amati relation, and the Yonetoku relation. The fit gives a best fit parameter as, \\begin{equation} \\Biggl(\\frac{E^{\\rm t}_{ \\rm peak}}{1 ~\\rm keV} \\Biggr) = 492(\\pm 24) \\times \\Biggl(\\frac{L^{\\rm t}_{ \\rm iso}}{10^{52} ~\\rm {erg~ s^{-1}}} \\Biggr)^{0.43(\\pm 0.03)}. \\label{epliso_alldata} \\end{equation}  The power-law index slightly flatter than the value obtained by Golenetskii et al. 1983 but roughly agrees with the previous analysis by Yonetoku 2004; $E_{\\rm peak} \\propto L_{\\rm iso}^{0.5}$. Furthermore, our data set can constrain both the power-law index and the normalization factor as  well as the previous analysis using only one burst data set. The $\\chi^2$ value for this fit is 83/56 and we still see a certain deviation in this relation around $L^{\\rm t}_{ \\rm iso}/10^{52} ~{\\rm erg ~s^{-1}} = 2 - 5$ (see Figure 8), and the data point which have large dispersion from this relation tend to have a higher $E^{\\rm t}_{ \\rm peak}$ value by a factor of 2 compared with many other data. We calculate the $E^{\\rm t}_{ \\rm peak}/ (L^{\\rm t}_{ \\rm iso})^{0.43}$ for each time region in order to separete these data point from other data which follow the equation(\\ref{epliso_alldata}). If the correlation follows $E^{\\rm t}_{ \\rm peak} \\propto L^{\\rm t~0.43}_{\\rm iso}$ for any time region, this parameter should retain constant. Figure \\ref{EpoverLiso045} shows the $E^{\\rm t}_{ \\rm peak}/ (L^{\\rm t}_{ \\rm iso})^{0.43}$  parameter as a function of time. We also show the best-fit constant value with 3 $\\sigma$ confidence region. The 8 data points (T=4,5,6,7,31,32,48,50) exceed 3 $\\sigma$ limit of the constant value, and thus we consider these data points as the outliers on the  $E^{\\rm t}_{ \\rm peak}-L^{\\rm t}_{ \\rm iso}$ plane. We confirm that these outliers with large dispersion from the main population are not caused by the systematic uncertainties in the analysis procedures such as the uncertainties of detector response matrix, systematic effect to the obtained spectral parameters, and the effect of the intense spectral evolution during the initial rising phase. The details of this verification processes are shown in the Appendix.   We then divide the derived parameter pairs into the main population, where the $E^{\\rm t}_{ \\rm peak}/ (L^{\\rm t}_{ \\rm iso})^{0.43}$ correlation is acceptable, and the outlier population, where that relation is broken at 3 $\\sigma$ significance. We calculate the correlation coefficient and perform the power-law fit for each population separately, and obtain the coefficients of 0.90 ($P_{rs}$ = 6.9 $\\times 10^{-18}$) for the main population, and 0.62 for the outliers. The correlation coefficient of the main population is marginally improved when we exclude the outliers. The outliers also have a certain degree of the correlation. The power-law best-fit parameters for those two populations are found as follows,   \\begin{eqnarray} \\Biggl(\\frac{E^{\\rm t}_{ \\rm peak}}{1 ~\\rm keV}\\Biggr) &=& 456(\\pm 25) \\times \\Biggl(\\frac{L^{\\rm t}_{ \\rm iso}}{10^{52} ~\\rm {erg~ s^{-1}}} \\Biggr)^{0.46(\\pm 0.03)} ~~~~~~~~~~ {\\rm for~~ main~~ population} \\nonumber \\\\ &=&  909(\\pm 147) \\times \\Biggl(\\frac{L^{\\rm t}_{ \\rm iso}}{10^{52} ~\\rm {erg~ s^{-1}}} \\Biggr)^{0.25(\\pm 0.14)}  ~~~~~~~~~~ {\\rm for~~ outliers.} \\label{EpLisorelation} \\end{eqnarray}  For the main population, the best-fit parameter of the power-law fit does not change significantly from that we derived from the entire data. We find that the power-law model is still valid for this relation. Furthermore, this sparation also led to a very interesting result. From figure \\ref{EpoverLiso045}, we can see that all of these outliers correspond to the rising phase of both initial weak peak and the second and third intense pulses. This indicates that if we exclude the initial rising phase of each pulse, we obtain much tighter $E^{\\rm t}_{ \\rm peak,src}$--$L^{\\rm t}_{ \\rm iso}$ correlation for all GRB. Therefore, it is probably more appropriate to exclude initial rising phase of each pulse when we use this relation as the redshift indicator like time-averaged $E_{\\rm peak,src}$ -- $L_{\\rm iso}$ (Yonetoku) relation. Moreover, it is likely that these outliers, in other words initial rising phase of each pulse have a different correlation with higher normalization and possibly smaller power-law index. However, we cannot give a strong conclusion for this because this power-law index varies, strongly depending on the criteria between the main population and outliers. Furthermore, we cannot constrain the power-law index of outliers due to the small number of data. Therefore, the robust conclusion is that the initial rising phase of burst have a different $E^{\\rm t}_{ \\rm peak}-L^{\\rm t}_{ \\rm iso}$ correlation with higher $E^{\\rm t}_{ \\rm peak}$ value from that of decay phase of burst. Note that this argumentation do not form a circular logic because the main point of this study is not to improve the fit of  $E^{\\rm t}_{ \\rm peak}-L^{\\rm t}_{ \\rm iso}$ correlation by removing the outliers but to just find that these outliers correspond to the initial phase of each pulse. This kind of trend that the initial phase of long GRBs has harder spectrum than that of decay phase is already reported by \\cite{Ghirlanda 2004}. However, the WAM data revealed that this harder-rising phase of long GRBs also follow the another power-law relation with that of softer-decay phase.   This difference between these different phases affects mainly the normalization of the relation. This fact can be useful to constrain the emission mechanism or dynamics of the emission site of the prompt emission of GRBs.  \\subsection{Implication for the emission radius and the bulk Lorentz factor}  Let us consider theoretical implications of our results that the spectral peak energy $E_{\\rm peak}$ evolves from hard to soft in the rising phase of each pulse. Although the emission mechanism of prompt GRBs is still enigmatic, the leading models are (1) the synchrotron shock model \\citep{zm02} and (2) the photosphere model \\citep{mesree00, th07,ioka07}.  In the synchrotron shock model, the peak energy is identified with the typical synchrotron energy $E_{\\rm peak}=\\Gamma \\hbar \\gamma_m^2 e B/m_e c$ of electrons that are shocked in the relativistic outflow with a Lorentz factor $\\Gamma$. Assuming that a fraction $\\epsilon_e$ of total energy goes into the electron acceleration and a fraction $\\epsilon_B$ goes into the magnetic amplification, we have the typical Lorentz factor of electrons as $\\gamma_m \\sim \\epsilon_e m_p/m_e$ and the magnetic luminosity as $\\epsilon_B L_{tot} = 4 \\pi r^2 c \\Gamma^2 B^2/8\\pi$ where $r$ is the shock radius and $L_{tot}$ is the total luminosity. Since almost all electron energy is radiated, the photon luminosity is given by $L_{\\gamma}=\\epsilon_e L_{tot}$. Combining these relations, we have the peak energy as a function of $r$, $\\Gamma$ and $L_{\\gamma}$ as \\begin{equation} E_{\\rm peak} \\sim 3~  \\epsilon_B^{1/2} \\epsilon_e^{3/2} L_{\\gamma,52}^{1/2} r_{13}^{-1} \\ {\\rm MeV}, \\label{eq:Epsyn} \\end{equation}  \\noindent where $Q_x=Q/10^x$ in the cgs unit. In this model, the low energy photon index should not exceed $-2/3$, and even $-3/2$ if electrons cool fast as expected in GRBs \\citep{ghi00}. However, many previous observations found flatter low energy index than $-2/3$ (e.g., \\cite{Preece 1998}; \\cite{Ghirlanda 2002}). Since our results for GRB 061007 exhibit flatter low energy photon index than $-2/3$ for the spectrum of one time interval, the photosphere model should also be considered.   In the photosphere model, the peak energy is identified with the thermal peak $E_{\\rm peak} \\sim \\Gamma T'$ of the photosphere under which the outflow energy is internally dissipated and thermalized. Most opacity of the photosphere is provided by $e^{\\pm}$ pairs or electrons associated with baryons (e.g., \\cite{mi08}). This model has several advantages that the peak energy is stabilized even for high radiative efficiency and that the low energy photon index can be as hard as the thermal one, $\\sim 1$ \\citep{th07,ioka07}. With the Stefan-Boltzmann formula, we have the comoving energy density of photons as $a T'^4= L_\\gamma/4\\pi r^2 c \\Gamma^2$, so that \\begin{equation} E_{\\rm peak} \\sim  1~ \\Gamma_{3}^{1/2} r_{10}^{-1/2} L_{\\gamma,52}^{1/4} \\ {\\rm MeV}. \\label{eq:Epphoto} \\end{equation}   From equations (\\ref{eq:Epsyn}) and (\\ref{eq:Epphoto}), the high peak energy above the Yonetoku relation $E_{\\rm peak} \\propto L_{\\gamma}^{1/2}$ suggests that the shock radius $r$ is small and/or the Lorentz factor $\\Gamma$ is large. This means that the fireball is expanding and/or decelerating in the initial rising of pulses, and we could be directly observing the fireball dynamics through the evolution of the Yonetoku relation. Since the shift of the peak energy is about a factor of two, the radius increases by a factor of $2$-$4$ and/or the Lorentz factor decreases by a factor of $\\sim 4$. Because the time dependence could be different between the synchrotron shock model $E_{\\rm peak} \\propto r^{-1} \\propto t^{-1}$ and the photosphere model $E_{\\rm peak} \\propto \\Gamma^{1/2} r^{-1/2} \\propto \\Gamma^{1/2} t^{-1/2}$, we could potentially discriminate models with similar but more detail spectral observations."
    },
    "0812/0812.0114_arXiv.txt": {
        "abstract": "Formulae for the energies of degenerate non-relativistic and ultra-relativistic Fermi gases play multiple roles in simple arguments related to the collapse of a stellar core to a neutron star.  These formulae, deployed in conjunction with the virial theorem and a few other basic physical principles, provide surprisingly good estimates of the temperature, mass, and radius (and therefore also density and entropy) of the core at the onset of collapse; the final radius and composition of the cold compact remnant; and the total energy lost to neutrino emission during collapse. ",
        "introduction": "Not long after `supernovae' came to be distinguished from `common novae' by virtue of their greatly higher intrinsic luminosity,\\cite{Baade1934On-Super-Novae} and shortly after the discovery of the neutron in the early 1930s, Baade and Zwicky declared: ``With all reserve we advance the view that supernovae represent the transitions from ordinary stars to {\\em neutron stars,} which in their final stages consist of extremely closely packed neutrons.''\\cite{Baade1934Supernovae-and-,Baade1934Cosmic-Rays-fro} For the `core-collapse' varieties of supernovae (Types Ib/Ic/II) this basic picture prevails today with good theoretical and observational support.\\cite{Arnett1996Supernovae-and-,Woosley2002The-evolution-a,Filippenko1997Optical-Spectra,Lattimer2004The-Physics-of-}  The virial theorem and formulae for the energies of degenerate Fermi gases are especially important to this paper's simple arguments and calculations related to the collapse of the of a massive star's core to a neutron star. The virial theorem states that a self-gravitating body consisting of particles that are either non-relativistic or ultra-relativistic satisfies \\begin{equation} 2E_{N} + E_{R} = -E_G, \\label{virial} \\end{equation} where $E_{N}$ and $E_{R}$ are the contributions to the body's internal energy in non-relativistic and ultra-relativistic particles respectively, and $E_G$ is the gravitational energy.\\cite{Collins-II1978The-Virial-Theo} While hydrostatic equilibrium requires that pressure and density decrease with increasing distance from the center of a spherically symmetric star, an assumption of uniform density allows for simple and surprisingly good estimates. The gravitational energy of a uniform-density sphere of mass $M$ and radius $R$ is \\begin{equation} E_{G} = -\\frac{3}{5}\\frac{G M^2}{R}, \\label{gravity} \\end{equation} where $G$ is the gravitational constant. As will be seen below, degenerate gases dominate the internal energies of both the pre-collapse stellar core and the resulting neutron star; hence formulae for the internal energies $(E_F)_{N}$ and $(E_F)_{R}$ of non-relativistic and ultra-relativistic uniform-density degenerate Fermi gases confined to a spherical volume of radius $R$ will play a prominent role in use of Eq. (\\ref{virial}). These energies are \\begin{eqnarray} (E_F)_{N} &=& \\frac{3}{5} N_N \\,(\\epsilon_F)_N, \\label{internalEnergyN}\\\\ (E_F)_{R} &=& \\frac{3}{4} N_R \\,(\\epsilon_F)_R, \\label{internalEnergyR} \\end{eqnarray} where $N_N$ and $N_R$ are the numbers of particles in non-relativistic and ultra-relativistic gases composed of a given particle type, and $(\\epsilon_F)_N$ and $(\\epsilon_F)_R$ are the respective Fermi energies: \\begin{eqnarray} (\\epsilon_F)_N &=& \\left(\\frac{9 \\pi N_N}{2 g}\\right)^{2/3} \\frac{\\hbar^2}{2m R^2}, \\label{fermiEnergyN}\\\\ (\\epsilon_F)_R &=& \\left(\\frac{9 \\pi N_R}{2 g}\\right)^{1/3} \\frac{\\hbar c}{R}. \\label{fermiEnergyR} \\end{eqnarray} Here $g$ is the spin degeneracy of the particle type, $m$ is the particle mass (relevant to the non-relativistic case), $\\hbar$ is the reduced Planck constant $h/2\\pi$, and $c$ is the speed of light (relevant to the ultra-relativistic case). In addition to their utility in representing the internal energies of self-gravitating bodies, Eqs. (\\ref{internalEnergyN})-(\\ref{fermiEnergyR}) will also prove useful in a simple model of the nucleus used in arguments about the process of electron capture on nuclei that signals the onset of core collapse.  After the existence of a degenerate iron core at the center of a massive star is argued for in Sec. \\ref{sec:ironCore}, subsequent sections utilize Eqs. (\\ref{internalEnergyN})-(\\ref{fermiEnergyR}) in various ways to determine the properties of this pre-collapse core (Sections \\ref{sec:mass} and \\ref{sec:radius}) and of the cold neutron star that is the final result of its collapse (Section \\ref{sec:remnant}). The energy lost to neutrino emission---which allows collapse to proceed---is addressed in Section \\ref{sec:energyLoss}. Section \\ref{sec:conclusion} contains concluding remarks. ",
        "conclusions": "\\label{sec:conclusion}  The observed supernova phenomenon, which led Baade and Zwicky to speculate on the formation of neutron stars, exhibits explosion energies of about $10^{44}$~J---a tiny fraction of the $10^{46}$~J of gravitational potential energy released in stellar collapse. Because the explosion is a comparatively minor detail in energetic terms, it is not surprising that sophisticated simulations are required to understand it.\\cite{Mezzacappa2005ASCERTAINING-TH,Woosley2005The-physics-of-}  Nevertheless, some features of the stellar core collapse itself can be understood very simply, using estimates that make use of the virial theorem and formulae for the energies of degenerate gases. (More detailed but less self-contained studies of the nuclear physics, weak interactions, and hydrodynamics related to the state of the pre-collapse core and the process of collapse, which retain an analytic or semi-analytic flavor, are available in the literature.\\cite{Bethe1979Equation-of-sta,Fuller1982Stellar-weak-in,Fuller1985Stellar-weak-in,Yahil1983Self-similar-st}) The virial theorem, which relates internal energy to gravitational energy, is a concise summary of the hydrostatic equilibrium that determines stellar structure; an assumption of uniform density then provides simple approximate relations between the mass and radius of both the pre-collapse stellar core and the post-collapse neutron star. The pre-collapse core is supported by electron degeneracy pressure, and when these electrons are relativistic an estimate of the core mass just before collapse is obtained. The radius (and therefore density) of the pre-collapse core can be obtained by estimating conditions for the onset of electron capture, through which energy (and pressure support) are lost through neutrino emission. This calculation involves another use of the concept of degeneracy, namely to crudely model the nucleus as consisting of degenerate proton and neutron ideal gases. Taking the cold post-collapse neutron star to be supported by degenerate non-relativistic nucleons, the virial theorem and the condition of `chemical' equilibrium provide an estimate of the neutron star radius and demonstrate why the name `neutron star' is well-deserved. Finally, the difference in total energy of the pre-collapse core and the post-collapse neutron star implied by virial theorem constitutes an estimate of the energy loss to neutrino emission that allows collapse to occur."
    },
    "0812/0812.0322_arXiv.txt": {
        "abstract": "The KASCADE experiment measures extensive air showers induced by cosmic rays in the energy range around the so-called knee. The data of KASCADE have been used in a composition analysis showing the knee at 3--5~PeV to be caused by a steepening in the light-element spectra~\\cite{Ant05}. Since the applied unfolding analysis depends crucially on simulations of air showers, different high energy hadronic interaction models (QGSJet and SIBYLL) were used. The results have shown a strong dependence of the relative abundance of the individual mass groups on the underlying model. In this update of the analysis we apply the unfolding method with a different low energy interaction model (FLUKA instead of GHEISHA) in the simulations. While the resulting individual mass group spectra do not change significantly, the overall description of the measured data improves by using the FLUKA model. In addition data in a larger range of zenith angle are analysed. The new results are completely consistent, i.e. there is no hint to any severe problem in applying the unfolding analysis method to KASCADE data. ",
        "introduction": "Due to the rapidly falling intensity with increasing energy, cosmic rays of energies above $10^{15}\\,$eV can be studied only indirectly by observations of extensive air showers (EAS) which are produced by the interactions of cosmic particles with nuclei of the Earth's atmosphere. The observation of a change of the power law slope~\\cite{knee} of the size spectrum of EAS and consequently of the all-particle energy spectrum at $\\sim 3\\cdot 10^{15}\\,$eV 50 years ago has not yet been convincingly explained~\\cite{Hau03}. Several theories for the origin of the knee predict different knee positions for particles of different primary mass. Therefore the energy spectra of single elements or at least mass groups are of considerable interest. Indeed, recent analyses \\cite{Agl04,Ant05} find a steepening in the energy spectra of the light components in the knee region. Whereas the measurement method of detecting air showers alleviates statistical problems, one has to rely on the results of simulations and the description of hadronic interactions while reconstructing the properties of the primary particles. Since the required energy and important kinematic regions of these interactions are beyond the range of collider or fixed target experiments, the interaction models used are uncertain and differ in their predictions. On the other hand, a thorough analysis of EAS data offers the opportunity of testing~\\cite{Mil06} and improving the validity of these hadronic interaction models.  The KASCADE-Grande experiment~\\cite{Nav04}, located on site of the Forschungszentrum Karlsruhe (Germany), is designed to measure EAS in the energy range between 0.5~PeV and 1~EeV. The installation consists of the original KASCADE~\\cite{Ant03} experiment and an extension by the Grande array, covering an effective area of 0.5~km$^{2}$.  The data analysis pursued by the KASCADE collaboration invokes an unfolding procedure of the two-dimensional shower size spectrum (total electron number vs.~muon number) into energy spectra of five individual mass groups~\\cite{Ant05}. Despite the success of this method for the reconstruction of the shape of the spectral forms, a strong dependence of the result for the elemental abundances on the interaction model underlying the analysis was found. Also, an insufficient description of the measured data by the employed simulations has been demonstrated. While in~\\cite{Ant05} the results for the analysis based on two different high energy interaction models were compared, we give in the present paper an update of the composition analysis concentrating on the influence of the low energy interaction model on the result. Furthermore, the analysis is repeated for measured data of different zenith angle intervals, thus testing the consistency of the analysis. ",
        "conclusions": "The analysis in terms of energy spectra of individual mass groups as described in~\\cite{Ant05} has been applied to KASCADE data based on a different low energy hadronic interaction model and to data sets in different intervals of the zenith angle.  Using the FLUKA model instead of the GHEISHA model has no significant effect on the overall picture of the solution, but a slightly better description of the data is achieved, i.e.~i.e.~the correlation of electron to muon number in the simulations with FLUKA is in better agreement with the data which results in a lower $\\chi^2_{dof}$ distribution of the solution. The results of the unfolding analyses of the two-dimensional shower size spectrum for different zenith angular ranges show no strong or unexplainable systematic differences. Thus, the results give no hint to any severe problem in the simulation or the analysis, and reaffirm the conclusions~\\cite{Ant05} drawn from the analysis of the nearly vertical shower set: The knee is observed at an energy around $\\approx 5\\,$PeV with a change of the index $\\Delta \\gamma \\approx 0.4\\,$. Considering the results of the mass group spectra, in all analyses an appearance of knee-like features in the spectra of the light elements is ascertained. In all solutions the positions of the knees in these spectra is shifted to higher energy with increasing element number.  By applying the analysis to different data sets and based on different interaction models, it has been demonstrated that unfolding methods are capable to reconstruct energy spectra of individual mass groups from air shower data, in addition to the all-particle spectrum. But still, the limiting factor of the analysis are the properties of the hadronic interaction models used and not the quality or the understanding of the KASCADE data. Furthermore, the procedure of the KASCADE data analysis, and in future also the analysis of KASCADE-Grande data measuring higher primary energies and muons at larger distances~\\cite{Zab05}, gives valuable hints for the improvement of hadronic interaction models. The data can be confidently used when improved interaction models, based on more and extended accelerator experiments, become available.   \\begin{ack} This work is supported by the BMBF of Germany, the Ministry of Science and Higher Education of Poland, and the Romanian Ministry of Education and Research (grant CEEX 05-D11-79/2005). \\end{ack}"
    },
    "0812/0812.0787_arXiv.txt": {
        "abstract": "A simple formula is introduced which indicates the amount by which projections of dark matter direct detection experiments are expected to be degraded due to backgrounds. ",
        "introduction": "Experiments sensitive to weakly interacting dark matter particles are placing ever tighter constraints on dark matter-nucleon cross sections (e.g., \\cite{cdms,xenon,warp}). The most recent upper limits rule out a substantial chunk of parameter space for well-motivated theories (see, e.g., \\cite{Jungman:1995df}). Even more exciting are the plans for future experiments which have excellent discovery potential~\\cite{dmsag}.  As these experiments get larger and more expensive, it becomes important for the community to agree on uniform standards for projecting the sensitivity of future experiments. Projections are often carried out simply by scaling exposures, defined as the product of fiducial detector mass $M$ times integrated detector live time $T$. That is, the projected improvement in the upper limit on the cross section is taken to be \\be \\frac{\\sigma^{\\rm UL}}{\\sigma^{\\rm UL,curr}} \\rightarrow \\frac{1}{\\lambda}  \\equiv \\frac{(MT)^{\\rm curr}}{MT} .\\ee  Neglected in these projections is the effect of backgrounds. The purpose of this note is to provide a simple formalism that accounts for the degradation of the future experiment's reach due to backgrounds. A given experiment's projected constraint on the cross section should be multiplied by the {\\it background penalty factor} introduced here to get a baseline estimate of that experiment's constraining power. This baseline estimate assumes that the future incarnation has the same background rate per volume as the current version of the experiment. Experimentalists could plausibly argue that they expect to reduce their background rates by some factor (here defined as $r$) and therefore they project a tighter constraint than the baseline estimate. The framework introduced here can also incorporate these arguments by computing the background penalty factor as a function of $r$. ",
        "conclusions": "Projections of the future reach of dark matter experiments should account for backgrounds. The Background Penalty Factor introduced here is a way to go beyond the standard projection which assumes zero backgrounds. This penalty factor is a function of the expected background rate, so projections using BPF, as illustrated in Fig.~\\rf{rlimit1_signb}, can quantitatively inform experimenters (and reviewers) the extent to which background rates need to be reduced.  {\\bf Acknowledgments:} This work was supported by the Fermi Research Alliance, LLC under Contract No. DE-AC02-07CH11359 with the US Department of Energy. I thank Dan Bauer, Richard Gaitskell, Bernard Sadoulet, Richard Schnee, and Steve Yellin for helpful comments and discussions. Fortran code to compute the BPF is available at {\\tt http://home.fnal.gov/\\~{}dodelson/dm.html}."
    },
    "0812/0812.1781_arXiv.txt": {
        "abstract": "Magnetic fields have been observed in galaxies, clusters of galaxies and probably in superclusters. While mechanisms exist to generate these in the late universe, it is possible that magnetic fields have existed since very early times. A field existing before the formation of the cosmic microwave background will leave imprints from its impact on plasma physics that might soon be observable. This thesis is concerned with investigating methods to predict the form of such imprints.  In chapter \\ref{Chapter-PerturbationTheory} we review in detail a standard, linearised cosmology based on a Robertson-Walker metric and a universe filled with photons, massless neutrinos, cold dark matter, a cosmological constant, and baryons. We work in a synchronous gauge for ease of implementation in a Boltzmann code, and keep our formalism general. We then consider the statistics of the cosmic microwave background radiation, assuming that only scalar (density) perturbations cause significant impact.  Chapter \\ref{Chapter-MagnetisedCosmology} introduces an electromagnetic field of arbitrary size and presents the equations governing the magnetised cosmology, and the structure of the electromagnetic fields, in greater detail than has hitherto been shown. We then invoke a hierarchy of approximations, treating the conductivity of the universe as infinite and thus removing the electric field, and considering the energy density of the magnetic field to be small. We present the resulting system in a computationally useful form and end by reviewing previous studies into the damping scales induced by photon viscosity.  Chapter \\ref{Chapter-SourceStats} considers the intrinsic statistics of the magnetic stresses. We approach this issue in two ways. Analytical methods are exact but reliant on an underlying Gaussian distribution function for the magnetic field, and simulating fields on a finite grid allows a great freedom in the form of the field but imposes an undesirable granularity and an unphysical infra-red cut-off. We construct the two-point moments in Fourier space, extending and improving the analytical results, some of which we present for the first time. There is excellent agreement between the results from analysis and those from the simulated fields. At the one- and three-point level we find significant intrinsic non-Gaussianities.  In chapter \\ref{Chapter-CMB} we turn to the observable impacts a primordial magnetic field. We briefly review previous studies into constraints from the epoch of nucleosynthesis before turning to consider the cosmic microwave background. We briefly consider the potential impact of an evolving damping scale, which may cause decoherence in the sources. Assuming coherence, which is likely to be accurate for very infra-red fields, the statistics of the source can be mapped onto the cosmic microwave background in an extensible manner, modelling the source statistics and the photon evolution entirely seperately. We demonstrate that our approach is valid by reproducing the signals for Gaussian power law fields on the microwave sky. After outlining how we will improve our predictions by employing a Boltzmann code, we show that it is a purely technical matter to extend the method to the three-point function. With detector sensitivity increasing, the non-Gaussianity of the cosmic microwave background will be well constrained in the near future and our work allows us to employ a new probe into the nature of an early-universe magnetic field. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.3321_arXiv.txt": {
        "abstract": "In the framework of a two-moment photo-hydrodynamic modelling of radiation transport, we introduce a concept for the determination of effective radiation transport coefficients based on the  minimization of the local entropy production rate of radiation and matter. The method provides the nonequilibrium photon distribution from which the effective absorption coefficients and the variable Eddington factor (VEF) can be calculated.  The photon distribution depends on the frequency dependence of the absorption coefficient, in contrast to the distribution obtained by methods based on entropy maximization. The calculated mean absorption coefficients are not only correct in the limit of optically thick and thin media, but even provide a reasonable interpolation in the cross-over regime between these limits, notably without introducing any fit parameter.  The method is illustrated and discussed for grey matter and for a simple example of non-grey matter with a two-band absorption spectrum. The method is also briefly compared with the maximum entropy concept. ",
        "introduction": "\\label{Intro} Excessive effort is required for modelling and simulation of radiation heat transfer in media with complex optical absorption spectra, like hot gases or plasma \\cite{Tien1968}. If the radiation model is part of a larger model for hot, compressible mixtures of various chemically reacting species consisting of complex ions, electrons, neutral molecules etc., and subject to transonic and turbulent flow, an exact treatment is not nearly possible. Applications range from arc physics in welding or electrical switching \\cite{Jones1980,Eby1998,Nordborg2008}, atomic explosions, up to % astrophysics \\cite{Goupil1984}. A frequently used approximate radiation model is photo-hydrodynamics, which is based on a two-moment expansion of the radiative transfer equation and a variable Eddington factor (VEF) closure \\cite{Levermore1984}. This concept can even be realized in a multiband framework, where the relevant quantities are decomposed according to their spectral properties \\cite{Ripoll2008}.\\\\ \\indent The main problem of photo-hydrodynamic models is the optimal choice of the effective transport parameters, i.e., the effective absorption coefficients and the Eddington factor, which generally depend on the hydrodynamic and thermodynamic variables. The most prominent examples are the Planck mean absorption coefficient for optically thin media, the corresponding Rosseland mean for optically thick media \\cite{SiegelHowell1992}, and the constant Eddington factor of $1/3$ for isotropic radiation.\\\\ \\indent Besides these special limit cases, the optimal definition of effective transport parameters  is not straightforward \\cite{Ripoll2001,Ripoll2004}. Accurate treatment of the general case is particularly important when radiation in the cross-over range between optically thin and thick limit dominates the physical behavior of a system, or if a medium is simultaneously transparent and opaque for different relevant wavelength bands. A simple approximate approach to solve that problem can consist in the construction of fitting expressions that interpolate between the limit cases. This has been done, for instance, by Sampson \\cite{Sampson1965} for the absorption coefficient and by Kershaw \\cite{Kershaw1976} for the variable Eddington factor. Another suggestion by Patch \\cite{Patch1967} generalizes heuristically averages, which are exact for special cases. A very common approach is based on entropy maximization \\cite{Levermore1984,Minerbo1977,Anile1991}. However, Struchtrup \\cite{StruchtrupUP} has pointed out a weakness of this procedure due to the neglect of the of radiation-matter coupling for the determination of the nonequilibrium photon distribution function, which plays indeed a main role for equilibration \\cite{Planck1906}. One of the consequences is that already the Rosseland mean aborption coefficient in the near-equilibrium case is not correctly reproduced by a two-moment photo-hydrodynamics with the maximum entropy closure. In this paper, we propose to use an entropy {\\em production rate} principle, and we will show that the limit cases are correctly reproduced and the cross-over between them is provided in a natural way, i.e., without further model parameters.\\\\ \\indent  Maximization and minimization of the entropy production rate $\\dot S$ have turned out to be powerful approaches for modelling many complex nonequilibrium systems \\cite{Martyushev2006}. We mention that whether the optimum of $\\dot S$ is a maximum or a minimum depends on the type of constraints \\cite{Christen2006a} and emphasize that entropy production optimization is in general not an exact physical law except near equilibrium, i.e., in linear deviation from equilibrium. However, it often provides useful approximate results even far away from equilibrium, provided a predominance of strongly irreversible equilibration processes. It is also important for the discussion below that the method is not restricted to systems in (partial) local equilibrium, i.e., where the notion of (probably several) local temperatures can be introduced (e.g., electron and ion temperatures in a plasma, light pencil temperatures for radiation \\cite{Oxenius1966,Kroll1967}, etc.).  Entropy production optimization has been shown to be applicable to local nonequilibrium systems, for instance to Knudsen layers in material ablation or evaporation processes \\cite{FordLee2001,Christen2007a}. In such cases, the notion of entropy can still be defined \\cite{LandauLifshitz}.\\\\ \\indent Entropy production of radiation has been discussed by Oxenius \\cite{Oxenius1966} and Kr\\\"oll \\cite{Kroll1967}. Various results related to entropy production principles in radiation have been reported. Essex \\cite{Essex1984} has shown that the entropy production rate is minimum in a grey atmosphere in local radiative equilibrium. Also this author has pointed out that a consideration of the interaction between radiation and matter is crucial because it contains the appropriate equilibration, i.e. entropy production, mechanism. Later on, W\\\"urfel and Ruppel \\cite{WurfelRuppel1985,Kabelac1994} discussed entropy production rate maximization by introducing an effective chemical potential of the photons, related to their interaction with matter. Finally, we mention Santillan et al.\\ \\cite{Santillan1998} who showed that for a constraint of fixed radiation power, the black bodies are those which maximize the entropy production rate.\\\\ \\indent  This paper is organized as follows. In order to fix notation and to introduce the relevant quantities, in Sect. \\ref{Photohydrodynamics} the photo-hydrodynamic equations are recalled. We will consider a system that is characterized by similar assumptions as in \\cite{Patch1967}. Scattering as well as photon time of flight effects are assumed to be negligible, the matter is non-relativistic and in thermal equilibrium, and the ordinary index of refraction is unity. Section \\ref{dSdt} provides the basic result, i.e., Eq. (\\ref{maximgen1}) determining the non-equilibrium photon distribution function. From this distribution function the transport coefficients can be calculated.  In Sect. \\ref{Photo} we show that the approach provides the Rosseland % mean in the corresponding limit case. In Sect.\\ \\ref{general} cases far from equilibrium are discussed. In particular, the emission limit, grey matter, and a simple artificial but illustrative example for non-grey matter are investigated. Our results are compared with the method of entropy maximization in Sect. \\ref{Comparison_ME}. ",
        "conclusions": "\\label{conc} The description of radiative heat transfer with photo-hydrodynamic equations requires a closure procedure to restrict the number of coupled equations. If radiative energy density $e$ and momentum density $\\vec p$ are to be considered, effective absorption coefficients and a variable Eddington factor (VEF) have to be determined as a function of $e$ and $\\vec p$. Using the minimization of the (local) entropy production rate of radiation {\\em and} matter, a closure procedure is introduced that yields the photon distribution function, which depends on the spectral absorption coefficient, and which allows for the calculation of transport coefficients in general.\\\\ \\indent It has been shown that the derived expressions are exact near equilibrium, in the emission limit, and correctly describe directed radiation with a VEF $\\chi = 1$. Whereas the first fact is not very surprising as minimum entropy production is known to be exact near equilibrium, the other correct limiting behavior demonstrates that entropy production optimization can provide sound and useful results also far from equilibrium.\\\\ \\indent For the general case, effective absorption coefficients and a variable Eddington factor are found, which reasonably interpolate between the limiting cases, notably without introduction of any fit parameter. The variable Eddington factor for grey matter is close to the one proposed by Kershaw \\cite{Kershaw1976}, and not far from the Levermore-Eddington factor, which can be derived from a maximum entropy approach.\\\\ \\indent In contrast to our approach, the distribution function obtained from the maximum entropy (ME) method does not contain the absorption spectrum. Hence, ME cannot describe frequency resolved absorption effects on the distribution function. An example of such an effect, which is taken into account by our method, is the enhanced photon equilibration in wave number bands with larger absorption as discussed in Sect. \\ref{twoband}. Another difference is that in the two-moment photo-hydrodynamic framework, ME is unable to reproduce the Rosseland mean absorption coefficient in the equilibrium limit.\\\\ \\indent Roughly speaking, the maximum entropy approach is a zero order approximation in the sense that radiation has equilibrated to a conditional maximum entropy state. Our approach can be understood as a first order approximation by assuming that the radiation is away from equilibrium and equilibrates according the minimum entropy production method. To summarize, we conjecture that this method may serve as a useful approach in various radiation heat transfer problems. \\\\ \\indent A benchmarking and critical discussion of specific application examples will be necessary in the future. Furthermore, the feasibility of improvements in the framework of multiband extensions \\cite{Ripoll2008} or by taking into account higher order moments should be investigated. \\newpage"
    },
    "0812/0812.1438_arXiv.txt": {
        "abstract": "Wide-band {\\it Suzaku} data on the merging cluster Abell 3667 were examined for hard X-ray emission in excess to the known thermal component. {\\it Suzaku} detected X-ray signals in the wide energy band from 0.5 to 40 keV. The hard X-ray ($> 10$ keV) flux observed by the HXD around the cluster center cannot be explained by a simple extension of the thermal emission with average temperature of $\\sim 7$~keV. The emission is most likely an emission from a very hot ($kT > 13.2$ keV) % thermal component around the cluster center, produced via a strong heating process in the merger. In the north-west radio relic, no signature of non-thermal emission was observed. Using the HXD, the overall upper-limit flux within a $34'\\times34'$ field-of-view around the relic is derived to be 5.3$\\times10^{-12}$ % erg s$^{-1}$ cm$^{-2}$ in the 10--40 keV band, after subtracting the ICM contribution estimated using the XIS or the {\\it XMM-Newton} spectra. Directly on the relic region, the upper limit is further tightened by the XIS data to be less than $7.3\\times10^{-13}$ erg s$^{-1}$ cm$^{-2}$, % when converted into the 10--40 keV band. The latter value suggest that the average magnetic field within the relic is higher than $1.6~\\mu$G. % The non-thermal pressure due to magnetic fields and relativistic electrons may be as large as $\\sim 20$\\% of the thermal pressure in the region. ",
        "introduction": "Enigmatic extended radio sources in rich galaxy clusters have been known for over 30 years (Willson 1970; for recent review, see Feretti 2005). Sources located at the cluster center are referred to as ``radio halos'', while those on the cluster periphery are called ``relics''. Steep spectra and spectral cutoffs at a few GHz detected in several halos/relics indicate that relativistic electrons with typically GeV energy are undergoing cooling due to synchrotron and inverse Compton (IC) emission. In every case, the radio halos/relics are found in irregular clusters which are apparently undergoing mergers. This suggests that the radio emitting electrons are accelerated by shocks or turbulence generated by energetic cluster merging.  The same relativistic electrons scatter Cosmic Microwave Background (CMB) photons up to the hard X-ray band (IC emission). By comparing the radio and the hard X-ray emissions, not only the electron energy density but also the intra-cluster magnetic field intensity can be estimated (e.g. Sarazin 1988). Detecting the IC emission is difficult because of the strong ICM component dominant below $\\sim 20$ keV and the low expected flux of the emission in hard X-rays above this energy. Currently, there are only a few reports claiming its detection, mostly from the {\\it Beppo-SAX} PDS in the energy band around 40 keV or higher (e.g. Fusco-Femiano et al.~1999; Nevalainen et al.~2004), though a few of the detections are still controversial (Rossetti \\& Molendi 2004; Fusco-Femiano et al.~2007). There are reports from RXTE (e.g. Petrosian et al.~2006; Rephaeli et al. 2006) as well. Another case includes the {\\it ASCA} GIS results around 4 keV on galaxy groups (Fukazawa et al.~2001; Nakazawa et al.~2006). These results would require a magnetic field as low as 0.1 $\\mu$G (e.g. Fusco-Femiano et al.~1999), which at the first glance is inconsistent with the radio rotation measure estimation of $2\\sim10$ $\\mu$G (e.g. Clarke et al.~2001; Carilli \\& Taylor 2002). Further observations by independent instruments of IC and other hard X-ray emission processes (e.g. Inoue et al.~2005), is thus desirable.  Significant heating of the ICM itself must also be taking place in merging clusters (e.g. Takizawa 2000; Ricker \\& Sarazin 2001). X-ray hardness ratio images sometimes show temperatures exceeding 10 keV (e.g. Watanabe et al.~1999; Briel et al.~2004). Since the emission can be multi-phase within the line of sight, the actual highest temperature can be much higher, though it is not easy to identify using the contemporary imaging X-ray detectors working below $< 12$~keV.  In many ways, Abell 3667 is the ideal cluster to study mergers, radio relics and hard X-ray emission. It is a very bright X-ray cluster at a low redshift of $z = 0.0556$ (Struble \\& Rood 1999). The {\\it ROSAT} and {\\it ASCA} observations showed that it is a spectacular merger with shock heated gas (Markevitch et al.~1999). {\\it Chandra} and {\\it XMM-Newton} observations have shown much evidence for an ongoing merger, such as a cold front (Vikhlinin et al.~2001a,b) and highly inhomogeneous temperature structure with temperature ranging from 4 to $> 10$ keV (Mazzotta et al.~2002; Briel et al.~2004). The optical galaxy distribution shows elongation towards the north-west south-east axis (Sodre et al.~1992; Johnston-Hollitt et al.~2008), supporting the binary merger scenario as well.  The cluster is famous for its pair of radio relics (e.g. Roettgering et al. 1997). The relic to the north-west is the brightest and largest among the diffuse radio sources associated with cluster of galaxies, having a flux of 3.7 Jy (Johnston-Hollit 2004) and a width of $\\sim 20'$ at 1.4 GHz. It is located about $30'$ or 2 Mpc from the cluster X-ray centroid. The radio signal is detected from 85 MHz to 2.3 GHz, and its average photon index is $\\Gamma = 2.1$ (Roettgering et al. 1997). Here the photon index $\\Gamma$ is defined as $N_{\\rm photon}(E) = N_0 \\times E^{-\\Gamma}$. The radio relics have very sharp outer edges and the radio spectra steepen with distance from the edge. Thus, the radio relics are considered to reflect the position of the particle acceleration in a merger shock. The cluster is also a good candidate for detection of IC emission. The {\\it Beppo-SAX} PDS provided a rather marginal evidence of hard excess emission from the cluster (Fusco-Femiano et al. 2001; Nevalainen et al. 2004). However, the large field of view (FOV) of the PDS makes the data strongly affected by contamination from thermal emission from the entire cluster, mainly from its central region, and possible AGNs.   The hard X-ray detector (HXD; Takahashi et al. 2006; Kokubun et al. 2006) onboard {\\it Suzaku} (Mitsuda et al. 2006) is characterized by its low detector background in the 10--40 keV band and its narrow FOV of $34'$ full width at half maximum (FWHM). Using the HXD, we can spatially distinguish hard X-ray components from the cluster center and the north-west relic, with the highest sensitivity in the 10--40 keV band. In addition, the X-ray CCD cameras (XISs; Koyama et al. 2006) onboard {\\it Suzaku} are characterized by their large effective areas and low and stable detector backgrounds, which make them very powerful devices to observe extended hard emission.  In May 2006, we observed Abell 3667 cluster using {\\it Suzaku} with three pointings running from the north-west relic to the center, each separated by $\\sim 17'$. Here we report the results obtained by these observations. Observation logs and data reduction are described in the next section. Section \\ref{chap:spec} is devoted to the spectral and imaging analysis, followed by the discussion in section \\ref{chap:discussion} and conclusions in section \\ref{chap:conclusion}. In this paper, we utilized canonical cosmology value of $H_{0} = 70$ Mpc km$^{-1}$s. At the redshift of $z = 0.0556$, $1'$ corresponds to 62 kpc. For a mean cluster temperature of 7.2 keV, the virial radius should be $r_{200} \\sim 2.6$ Mpc $= 42'$ (e.g. Neumann \\& Arnaud 1999). Unless otherwise noted, all errors are at the 90\\% confidence limit. ",
        "conclusions": "\\label{chap:conclusion}  From {\\it Suzaku} mapping observations of Abell~3667, upper-limits on the non-thermal (hard) X-ray emission were obtained. On the north-west radio relic itself, the XIS limit of 7.3$\\times10^{-13}$ erg s$^{-1}$ cm$^{-2}$ extrapolated to the 10--40 keV band provides a lower-limit magnetic field of $> 1.6~\\mu$G.  % The non-thermal energy density there is estimated to be higher than 7\\%, and likely near 20\\%.  By combining the XIS and the {\\it XMM} mapping data to the HXD-PIN spectra, the upper limit on the 10--40 keV excess hard emission is derived as 1.0--5.3$\\times10^{-12}$ erg s$^{-1}$ cm$^{-2}$, depending on the modeling of the ICM. Since the north-west relic has a magnetic field of $> 1$ $\\mu$G and hence must be almost an order of magnitude darker than this hard X-ray limit, the relativistic electrons must be distributed in regions other than the relic with low magnetic field (of $B < 0.3~\\mu$G) if IC hard X-ray emission near the flux limit exists.  The PIN data from the pointing on the cluster center shows hard X-rays in excess to the averaged ICM emission, which is thought to have a temperature of around 7~keV. Although other interpretations cannot be rejected, its soft spectral shape indicates the emission to be possibly from a very hot thermal component with a temperature of $\\sim 20$ keV, generated temporary in the cluster merger phase.  In the near future, X-ray observatories with hard X-ray imaging optics operating up to 40 keV will be launched, such as the ASTRO-H (former NeXT; e.g. Takahashi et al 2006), NuSTAR and Simbol-X (e.g. Ferrando et al. 2006). These detectors will drastically improve the sensitivity to both point sources and relatively clumpy extended hard X-rays, such as the local non-thermal emission associated with the radio halos/relics and the very hot component. In addition, the Soft Gamma-ray Detector onboard ASTRO-H, (a collimated instrument) features the highest sensitivity at around 100 keV using the novel narrow-field-of-view Compton camera concept. It will open a new window to the widely distributed non-thermal emission, which is difficult to separate from the ICM thermal emission below $\\sim 50$~keV. \\\\  K.N. and K.M. are supported in part by a Grant-in-Aid from the Ministry of Education, Science, Sports, and Culture of Japan (18104004), and M.T. (16740105,19740096) and S.I (19047004, 19540283), as well. C.L.S. and D.R.W. were supported by NASA {\\it Suzaku} grants NNX06AI44G and NNX06AI37G and {\\it XMM-Newton} grant NNX06AE76G. D.R.W. was also supported by a Dupont Fellowship and a Virginia Space Grant Consortium Fellowship.  A.F. was partially supported by NASA grant NNG05GM5OG to UMBC. Basic research in radio astronomy of T.C. at the NRL is supported by 6.1 Base funding.   \\appendix"
    },
    "0812/0812.2710_arXiv.txt": {
        "abstract": "We present results of our search for X-ray line emission associated with the radiative decay of the sterile neutrino, a well-motivated dark matter candidate, in {\\it Suzaku} Observatory spectra of the Ursa Minor dwarf spheroidal galaxy. These data represent the first deep observation of one of these extreme mass-to-light systems and the first dedicated dark matter search using an X-ray telescope. No such emission line is positively detected, and we place new constraints on the combination of the sterile neutrino mass, $m_{\\rm st}$, and the active-sterile neutrino oscillation mixing angle, $\\theta$. Line flux upper limits are derived using a maximum-likelihood-based approach that, along with the lack of intrinsic X-ray emission, enables us to minimize systematics and account for those that remain. The limits we derive match or approach the best previous results over the entire 1--20 keV mass range from a single {\\it Suzaku} observation. These are used to place constraints on the existence of sterile neutrinos with given parameters in the general case and in the case where they are assumed to constitute all of the dark matter. The range allowed implies that sterile neutrinos remain a viable candidate to make up some -- or all -- of the dark matter and also explain pulsar kicks and various other astrophysical phenomena. ",
        "introduction": "The observational evidence that nonbaryonic dark matter comprises most of the mass in the universe is extremely strong, but the nature of dark-matter particles remains a mystery.  No particle included in the Standard Model of particle physics has the characteristics required to explain the dark matter.  In its original formulation, the Standard Model described all known particles, including neutrinos, which were assumed to be massless.  The discovery that neutrinos have mass has forced one to go beyond this minimal model: a modest extension of the Standard Model by gauge singlet fermions can accommodate the neutrino masses.  One or more of these gauge-singlet fermions can have Majorana masses below the electroweak scale, in which case they appear as sterile neutrinos in the low-energy theory.  If one of the sterile neutrinos has mass in the 1--20~keV range and has small mixing angles with the active neutrinos (as expected), such a particle is a plausible candidate for dark matter~\\citep{dw94}.  The same particle could be produced in a supernova explosion, and its emission from a cooling neutron star could explain the pulsar kicks \\citep{ks97,Fuller03,k04,kmm}, could facilitate core collapse supernova explosions~\\citep{Fryer,hf07} and can affect the formation of the first stars~\\citep{bk06,s07} and black holes~\\citep{mb05,mb06}. Therefore, there is a strong motivation to search for signatures of sterile neutrinos in this mass range.  The most promising way to discover (or rule out) relic sterile neutrinos is with the use of X-ray telescopes (XRTs), such as {\\em Suzaku}. The sterile neutrinos can decay radiatively~\\citep{pw,barger} and produce a lighter neutrino and a photon amenable to X-ray observation~\\citep{afp,aft,dh}.  Since this is a two-body decay, a line is expected in the X-ray spectrum.  The dwarf spheroidal galaxies are ideal targets for these observations because of their proximity, high dark matter density \\citep{s08a}, and absence of additional obscuring X-ray sources (see below). The Ursa Minor and Draco dwarf spheroidals are the optimal targets in this class for large field-of-view X-ray spectroscopic investigation, based on their average dark matter surface densities. From {\\it XMM-Newton} observations of the Ursa Minor dwarf spheroidal, sterile neutrino line flux limits comparable to the best previous constraints were derived -- even though strong background flaring limited the amount of good observing time to 7 ks \\citep{bnr}.  The {\\em Suzaku} Observatory \\citep{m07} provides the most sensitive instruments for current searches for weak sterile neutrino radiative decay lines in the $\\sim 0.5-10$ keV bandpass because of its low and stable background \\citep{ta08}, and the relatively sharp spectral resolution of its CCD spectrometers \\citep{k07,n08}.  In this paper, we discuss our analysis of {\\it Suzaku} observations of the Ursa Minor system, considerably improving on the {\\it XMM-Newton} constraints; a companion article on the Draco system is in preparation. These represent the first substantial X-ray data of such extremely dark-matter-dominated systems.  \\subsection{Context}  Sterile neutrinos may be produced through non-resonant oscillations, as proposed by Dodelson and Widrow \\citep{dw94}, and via various other mechanisms~\\citep{d02,sf99,st,k06,kadota,pk08}.  A generic prediction is that relic sterile neutrinos must decay into a lighter neutrino and a photon. Since this is a two-body decay of a particle with mass $m_{\\rm st}$, decay photons produce an intrinsically narrow line with energy $E_\\gamma = m_{\\rm st}/2$ that can be detected using X-ray spectroscopy \\citep{aft}. While the decay timescale, $\\Gamma_{\\rm st}^{-1}$, greatly exceeds the Hubble time, the emissivity may reach detectable levels in regions of large dark matter concentration. Upper limits on $\\Gamma_{\\rm st}$ from X-ray observations of the cosmic background \\citep{ak06,b06a,a07}, galaxy clusters \\citep{b06b,r07,brm}, and M31 \\citep{w06,r07,birs} may be complemented and improved by observing dwarf spheroidal galaxies \\citep{b06c}.  There is no consensus as to whether dark matter is cold or warm based on observations of small-scale structure. Cold and warm dark matter reproduce large-scale structure equally well, but differ in their predictions on scales comparable to the free-streaming length. Ly$\\alpha$ forest clouds provide a tracer population of the perturbation spectrum on the appropriate (non-linear) scale that may be used to place an upper limit on the free-streaming length \\citep{n00}, and have been interpreted as implying that the mass of the dark matter sterile neutrino, $m_{\\rm st}$, must be greater than $\\sim 8~$keV \\citep{v08,blrv} for production exclusively by the Dodelson-Widrow (DW) non-resonant mixing mechanism. However, $N$-body simulations and observations of dwarf spheroidal galaxies imply a lower mass \\citep{g07,s07a}. In scenarios with alternative production mechanisms that do not involve oscillations, sterile neutrinos may effectively be colder, allowing masses as low as 3~keV to be consistent with the Ly$\\alpha$ observations~\\citep{k06,pk08,pet08}. In this case, the photometric and kinematic data from dwarf spheroidal galaxies~\\citep{g07,s07a} can be satisfied by dark matter in the form of sterile neutrinos with mass between 0.5~keV and 1.3~keV~\\citep{b08}.  Constraints on sterile neutrino parameters may be further relaxed in mixed dark matter models \\citep{p07,blrv} and in models where sterile neutrinos are produced in resonant oscillations in the presence of significant lepton asymmetries \\citep{sf99,ls08}. Since the sterile neutrino distribution function may substantially deviate from thermal equilibrium \\citep{a06a,a06b,pet08,b08}, simple scaling arguments and alterations to the power spectrum (e.g., imposing a cutoff) do not suffice. New large-scale structure formation simulations with the appropriate revisions are required to re-evaluate the constraints from Ly$\\alpha$ data \\citep{ls08,b08}.  Moreover, their interpretation relies on a thorough understanding of the uncertain thermal and ionization evolution of the intergalactic medium, which may be affected by the properties and interactions of dark-matter particles~\\citep{bk06,gao,reviewLya}. Hence, we present limits solely derived from our X-ray data.  The sterile neutrino mass and the mixing angle $\\theta$ determine the rate of radiative decay to active neutrinos and X-ray photons: $\\nu_s\\rightarrow\\gamma \\nu_a$.  The radiative decay width is equal to~\\citep{pw} \\begin{eqnarray} \\Gamma_{\\nu_s\\rightarrow\\gamma \\nu_a} &=& \\frac{9}{256\\pi^4}\\, \\alpha_{_{\\rm EM}}\\, {\\rm G}_{_{\\rm F}}^2 \\, \\sin^2 \\theta \\, m_{\\rm s}^5 \\nonumber \\\\ & = & \\frac{1}{1.8\\times 10^{31}\\, {\\rm s}}\\ \\left( \\frac{\\sin^2 \\theta}{10^{-10}} \\right) \\ \\left( \\frac{m_{\\rm s}}{\\rm keV}\\right)^5, \\end{eqnarray} where $ \\alpha_{_{\\rm EM}}=1/137.0$ and $G_{_{\\rm F}}=1.2\\times 10^{-5} {\\rm GeV}^{-2}$.  For a given production mechanism, one may be able to eliminate one of these two parameters by requiring that the total density match that of dark matter in the universe as inferred from observations.  However, given the various possible mechanisms for producing sterile neutrinos, and the possibility that sterile neutrinos may account for some -- but not all -- dark matter \\citep{p07}, one should keep an open mind and search all the available parameter space in the $m_{\\rm st}-\\theta$ plane. ",
        "conclusions": "With a single {\\it Suzaku} observation of the Ursa Minor dwarf spheroidal, we have derived limits on the mass and mixing angle of the sterile neutrino dark matter candidate that are comparable to or better than previous constraints. These limits are illustrated in Figure~\\ref{figure:Suzaku_limits} for the case where sterile neutrinos comprise all of the dark matter regardless of how they are produced in the early universe and, in the more general case, when their abundance is not assumed, but is calculated under the assumptions that minimize it. The best estimate of the total projected dark matter mass within 400 pc of $6\\times 10^7$ $M_{\\odot}$ is adopted here. If solely produced by the DW mechanism, sterile neutrinos cannot constitute 100\\% of the dark matter in Ursa Minor unless they are less massive than 2.5 keV.  Because of its orbit, {\\it Suzaku} has a lower and more stable background than either the {\\it Chandra} or {\\it XMM-Newton} Observatories, and its spectrometers have sharper energy resolution. Our approach in deriving limits utilizes the C-statistic to evaluate a very general model fit to the data with minimal pre-processing and no background subtraction. Thus, there are no a priori assumptions about the accuracy of the NXB estimate that we only use to set line energies and shapes, or whether an extra line feature lies above or below the total continuum. Our final limits are those that correspond to a change in the C-statistic that occurred in 1\\% of our Monte Carlo simulations. These factors may account for why our constraints represent only modest improvements over previously published limits despite the more optimal characteristics of {\\it Suzaku}. As the NXB characterization improves, it may be possible to tighten upper limits in the $m_{\\rm st}>2$ keV regime; at lower masses, the ubiquity of the GXB will remain a barrier until significantly higher energy resolution is attained -- e.g., with {\\em Astro-H}\\footnote{http://www.isas.ac.jp/e/enterp/missions/astro-h/\\\\index.shtml}.  Since there is some uncertainty as to the dynamical state of dwarf spheroidals on an individual basis, X-ray constraints obtained from additional systems are warranted (e.g., our paper -- in preparation -- on the {\\it Suzaku} observation of the Draco system). The Sloan Digital Sky Survey (SDSS) has approximately doubled the number of galaxies identified with the local group population and revealed a new sub-population of faint Milky Way satellites \\citep{bel07}, including some with very high mass-to-light ratio \\citep{s08a,s08b} that -- with the right observational setup -- may provide the best targets for investigations of this type.  Finally, predictions of the Ly$\\alpha$ forest power spectrum in fully self-consistent large-scale structure simulations with sterile neutrinos can provide complementary constraints in the form of a lower limit on the particle mass.  Sterile neutrinos remain a viable dark matter candidate. For example, a sterile neutrino with a mass of $m_{\\rm st}\\sim 4$ keV and $\\sin^2\\theta\\sim 4\\times 10^{-10}$ could account for all of the dark matter (or $\\sim 15$\\%$-75$\\% if produced solely by the DW mechanism) and explain pulsar kicks \\citep{kmm} without violating the line flux upper limits."
    },
    "0812/0812.0009_arXiv.txt": {
        "abstract": "We investigate the physical mechanisms of tidal heating and satellite disruption in cold dark matter host haloes using N-body simulations based on cosmological initial conditions.  We show the importance of resonant shocks and resonant torques with the host halo to satellite heating.  A resonant shock (torque) couples the radial (tangential) motion of a satellite in its orbit to its phase space.  For a satellite on a circular orbit, an ILR-like resonance dominates the heating and this heating results in continuous satellite mass loss.  We estimate the requirements for simulations to achieve these dynamics using perturbation theory.  Both resonant shocks and resonant torques affect satellites on eccentric orbits. We demonstrate that satellite mass loss is an outside-in process in \\emph{energy space}; a satellite's stars and gas are thus protected by their own halo against tidal stripping.  We simulate the evolution of a halo similar to the Large Magellanic Cloud (LMC) in our Galactic dark matter halo and conclude that the LMC stars have not yet been stripped.  Finally, we present a simple algorithm for estimating the evolution of satellite mass that includes both shock heating and resonant torques. ",
        "introduction": "\\label{sec:intro}  Physical processes affecting satellite galaxy evolution in their host haloes are an important component of galaxy formation in the cold dark matter (CDM) cosmogony as galaxies are built up from the assembly of small structures.  This assembly includes the process of satellite galaxies merging with their host galaxies.  Moreover, recent CDM cosmological simulations predict the existence of a large number of \\emph{subhaloes} \\citep{Ghigna98, Klypin99}. Consequently, understanding the detailed physical processes affecting satellite evolution are key ingredients to understanding galaxy formation in the CDM cosmogony.  Several basic questions about satellite halo evolution remain.  First, how is the satellite stripped?  In other words, what parts of the initial mass distribution might persist to the present day?  Second, what is the rate of satellite halo disruption?  Third, how does a satellite's internal structure evolve?  In a satellite-galaxy merger, the stars and gas of the satellite galaxy can be stripped and become halo stars and halo gas \\citep[e.g.][]{QBB00,BKW01}.  The interaction between the satellite galaxy and its host during the course of the merger results in evolution of the satellite galaxy \\citep[e.g.][]{MKL96}.  The remaining components of the satellite galaxy merge with the host galaxy and this merging causes the host galaxy to gain mass \\citep[e.g.][]{Murali02}.  Current cosmological simulations can only provide statistical properties of subhaloes since even in the highest resolution cosmological simulations \\citep{Ghigna00, DLucia04, Diemand04, Gao04, OL04, Diemand.etal:07, Springel.etal:08} the detailed physical processes of individual subhalo evolution have not been accurately studied, owing to limited resolution.  To investigate the physical processes in detail, we perform high resolution idealised simulations with cosmologically motivated initial conditions instead of using cosmological simulations.  In these idealised simulations, a live satellite orbits in a static host halo.  Although too simplified to reproduce a satellite's evolution in realistic detail, several authors have used similar non-cosmological simulations to study satellite disruption with alternative simulation methods \\citep[e.g.][]{Hayashi.etal:03,Kazantzidis.etal:04,Read.etal:06, Boylan-Kolchin.Ma:07,Penarrubia.Navarro.McConnachie:08}.  These studies have demonstrated important features of the satellite evolution such as mass loss history and density profile evolution.  In this paper we use the higher resolution simulation to investigate satellite disruption particularly focusing on the detailed physical processes affecting satellite evolution such as resonant dynamics.  When a time-dependent force acts on a bound system such as a galaxy or a dark matter halo, resonant interactions play an important role in the system evolution.  Recent studies of resonant dynamics in galaxy evolution claim that high resolution simulations are required to accurately reproduce these resonant effects \\citep{WK07a,WK07b} and \\citet{WK07a} provide a procedure to determine minimum particle number guidelines.  Since our idealised, high resolution simulations are designed to satisfy these particle number guidelines, they allow us to investigate the role of resonant dynamics in satellite disruptions.  Resonant interactions couple a time-dependent perturbing force with orbits in the system.  The frequency spectrum of the time-dependent perturbing force characterises the interaction.  For a satellite orbiting in its host halo, the satellite's orbital frequencies, and possibly the rate of orbital decay, determine the time-dependence of the external force.  A general, eccentric satellite orbit in a spherical halo has both a radial and azimuthal frequency, making the resonant coupling for an eccentric orbit complex. Empirically, we may characterise the overall effects of the interaction as a \\emph{resonant shock} and a \\emph{resonant torque}. A resonant shock represents coupling with the radial orbital frequency and a resonant torque represents coupling with the azimuthal orbital frequency.  A resonant shock is a generalisation of the standard impulsive, gravitational shock.  During a resonant shock, some orbits within the satellite gain energy through resonant coupling even though they are not in the impulsive limit, i.e. the time scale of the perturbation near pericentre is much longer than the internal orbital time scale. A resonant torque couples the rotation in the external potential to orbits in the satellite.  For a resonant torque, the magnitude of the external potential does not have to change; a change in the position angle of the satellite frame relative to the centre of the halo is sufficient to produce a torque.  Resonant shocks have been previously considered and included in the impulsive approximation as an adiabatic correction \\citep[e.g.][]{Spitzer87,Weinberg94a,Weinberg94b,GO99}. However, resonant torques have not been similarly considered in satellite evolution studies although they have been extensively investigated in the dynamics of barred galaxies.  In this study, we will carefully investigate these resonant effects on satellite evolution using this distinction.  Globular cluster evolution in a host galaxy is well established \\citep[e.g.][]{Spitzer87,Chernoff.Weinberg90}.  A globular cluster experiences both tidal truncation and heating by both compressive gravitational shocks and tidal shocks.  Because satellite halo evolution in a host halo is similar to globular cluster evolution in a host galaxy, many galaxy formation studies employ simple analytic formulae taken from these globular cluster evolution studies to estimate satellite galaxy evolution.  However, unlike globular clusters, the satellite--host mass ratio is \\emph{not} vanishingly small.  This breaks the spatial symmetry in mass loss, as shown in \\citet{Choi.etal:07}, and changes the relative importance of resonant coupling.  In this paper, we present numerical simulation results of satellite galaxy disruption.  In \\S\\ref{sec:method}, we present an overview of our numerical techniques: the N-body simulation code, the generation of initial conditions, and the relevant perturbation theory.  In \\S\\ref{sec:circ}, we present the results of a circular orbit simulation.  We show that resonant torque effects result in significant satellite mass loss.  In \\S\\ref{sec:ecc}, we present the results of eccentric orbit simulations.  We show that satellite heating by gravitational shocks at pericentre, which also includes internal structure evolution, is the dominant process responsible for disrupting the satellite.  In \\S\\ref{sec:stripping}, we show that the process of satellite stripping is an outside-in process in satellite \\emph{energy} space.  Using this finding, we suggest an explanation for the `missing' LMC stellar stream.  We also discuss the evolution of the satellite density profile.  In \\S\\ref{sec:massloss}, we provide an improved analytic estimate for satellite mass loss, and we summarise in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  Using high resolution simulations with cosmologically motivated initial conditions, we investigate the physical processes responsible for the evolution of satellite galaxies in their host halo.  We identified and explored the following important physical mechanisms that result in satellite galaxy disruption.  Our main results are as follows: \\begin{enumerate} \\item Resonant mechanisms of two types play a key role in satellite disruption: the resonant shock and the resonant torque. \\item We studied satellites on circular orbits to isolate the effects of resonant torques.  The ILR-like resonance $l_1:l_2:l_3=-1:2:2$ (see equation (\\ref{eq:res})) dominates the torque resonance because all the $l=1$ resonances and the 0:2:2 (corotation-like) and the 1:2:2 (OLR-like) resonances are located outside the satellite's tidal radius. \\item Some important resonances require more particles than are typically used in cosmological simulations.  For example, to accurately reproduce the heating from the -1:2:2 resonance, a satellite simulation needs more than $10^{5}$ particles within the virial radius for our expansion code and possibly more for codes affected by small-scale noise, e.g. a tree code.  Too few particles results in less mass loss and it suggests that the lifetime of dark matter subhaloes in current cosmological simulations could be \\emph{overestimated}. \\item For satellites on eccentric orbits, gravitational shocks dominate the heating but the heating rate is significantly underestimated by the impulse approximation for several reasons. First, strong low-order resonances couple to the phase space in the `adiabatic' regime and the work done on the satellite drives structural evolution.  Second, the ongoing tidal truncation affects the dynamics of the escaping material and the satellite's size.  The interplay of these mechanisms leads to continuous mass loss and evolution. \\item The number of resonances in a satellite increases as the satellite's orbital frequency increases; the strength of each resonant interaction also increases.  This affects satellites on low-energy and low-angular momentum orbits which have small pericentres and high satellite azimuthal frequencies. \\item Satellite mass stripping is an outside-in process in energy space.  We discuss the morphological consequences for the stripping signatures of stellar and gaseous streams. \\item We present an improved algorithm to estimate satellite mass loss including both resonant shocks and torques.  We show that the mass-loss history computed using this algorithm reproduces the general features seen in the simulations. \\end{enumerate}  The satellite evolution using N-body simulation has been studied by a number of authors recently \\citep[e.g.][]{Hayashi.etal:03,Kazantzidis.etal:04, Boylan-Kolchin.Ma:07}. Their approach differs ours in several ways. We use a highly idealised simulation configuration to better discern the details of the dynamical mechanisms, focussing on the detailed physical processes affecting satellite evolution such as resonant dynamics that have not been rigorously addressed elsewhere.  For example, we equilibrated the initially tidally truncated satellite to reduce the artificial perturbation from sudden introduction of the host halo potential.  We also explored a circular orbit simulation to clearly demonstrate the effects of resonant torque.  We performed a linear perturbation calculation to estimate the required resolution for resonant effect, and made the simulations satisfy this requirement. In addition, we use an expansion code to reduce the force fluctuations on small scales. Owing to these efforts, we were able to demonstrate the importance of the resonant torque in satellite disruption.  \\citet{Hayashi.etal:03} used an NFW halo as a satellite initial conditions and a Tree code for a potential solver. They found that the simple tidal-limit approximation underestimates the mass loss, as we do, and found structural evolution, also similar to our findings. \\citet{Kazantzidis.etal:04} and \\citet{Boylan-Kolchin.Ma:07} also found that similar internal structure evolution using tree-code simulations. This consensus suggests that the inner cusp of satellite halo is not strongly affected by the tides from the host halo. However, owing to their rather complicated configuration, they could not discern the effect of the resonant torque, although some authors \\citep[e.g.][] {Hayashi.etal:03} has noticed that analytic formulae underestimate the mass loss.  Although we have improved our understanding of the detailed physical processes responsible for satellite disruption, some issues remain.  By separating the heating mechanisms into two distinct regimes, we achieved an improved understanding of resonant dynamics for an eccentric orbit.  However, we have not compared this approximation with a comprehensive perturbation theory calculation. The resonant heating of satellites on eccentric orbits is similar to heating by other subhaloes.  This interaction is an important source of satellite evolution in addition to the interaction with the smooth host halo.  In addition, the initially spherical satellite is deformed during disruption owing to the host halo's tidal field and we have not yet accounted for this deformation.  A comprehensive treatment of this deformation might be necessary to understand satellite disruption in detail.  Lastly, we also need a better understanding of the non-linear processes that occur during satellite evolution.  In this study, we characterised the linear processes; understanding the detailed consequences of the non-linear processes is a daunting future task. We should then finally be able to fully constrain the satellite disruption mechanism, which is an essential ingredient of galaxy formation and evolution."
    },
    "0812/0812.0379_arXiv.txt": {
        "abstract": "Using high resolution SPH simulations in a fully cosmological $\\Lambda$CDM context we study the formation of a bright disk dominated galaxy that originates from a ``wet'' major merger at z$=$0.8. The progenitors of the disk galaxy are themselves disk galaxies that formed from early major mergers between galaxies with blue colors. A substantial thin  stellar disk grows rapidly following the last major merger and the present day properties of the final remnant are typical of early type spiral galaxies, with an $i$ band B/D $\\sim$ 0.65, a disk scale length of 7.2 kpc, $g-r$ = 0.5 mag, an HI line width ($W_{20}$/2) of 238 km/sec and total magnitude $i$ = -22.4.  The key ingredients for the formation of a dominant stellar disk component after a major merger are:  i) substantial and rapid accretion of gas through cold flows followed at late times by cooling of gas from the hot phase, ii) supernova feedback that is able to partially suppress star formation during mergers and iii) relative fading of the spheroidal component. The gas fraction of the progenitors' disks does not exceed 25\\% at z$<$3, emphasizing that the continuous supply of gas from the local environment plays a major role in the regrowth of disks and in keeping the galaxies blue.  The results of this simulation alleviate the problem posed for the existence of disk galaxies by the high likelihood of interactions and mergers for galaxy sized halos at relatively low z. ",
        "introduction": "Within the $\\Lambda$CDM framework the build up of galaxies and their parent dark matter (DM) halos occurs through a series of mergers and accretion of more diffuse matter \\citep{frenk85, millennium05}. The classic models of galaxy formation \\citep{whiterees78,fall83,whitefrenk91} and subsequent work \\citep{dalcanton97,mo98,silk01,bower06,zheng07,somerville08} assume that gas infalls inside the parent dark matter halo and subsequently cools to temperatures low enough to fragment and form stars. The morphology of galaxies and the fraction of stars in their disk and spheroidal components are set by a combination of competing physical processes, as numerical simulations have shown that violent relaxation in mergers between similar size galaxies can destroy or dynamically heat their stellar disks and turn them into spheroids \\citep{barneshernquist96}. On the other hand, remaining and subsequently accreted gas can regrow a disk \\citep{baugh96, steinmetznavarro02} if given enough time. It is then a prediction of hierarchical models that the morphology of a galaxy is not assigned ``ab initio'' but it might change often over cosmic times, as the spheroidal to disk light ratio changes back and forth.  Observed merger remnants have disky isophotes \\citep{rothberg04} or evidence for young disk components \\citep{mcdermid06}, hinting at disk regrowth and supporting the above scenario.    Observationally, major mergers are observed to be common in the redshift range 0 -- 1. Their number density has been estimated to be of the order of 10$^{-(3-4)}$h$^{-1}$Mpc$^{-3}$ Gyr$^{-1}$, \\citep{lin08,jogee08}, possibly increasing with redshift. As numerical work suggests that a consistent supply of gas is important to regrow a disk component \\citep{robertson06}, it is relevant that observational evidence for mergers between actively star forming galaxies has also emerged, showing that at redshift $\\sim$1, 70\\% of the merging pairs are between blue  (or ``wet'', defined as rest frame g-r $<$0.65) pairs; however the fraction of blue and  likely gas rich mergers decreases  by the present time \\citep{lin08}.  Numerical work has also highlighted a possible problem for the existence of disk galaxies at low redshift, showing that interactions, mergers and accretion events are common at every redshift for L$\\star$ galaxy sized halos in $\\Lambda$CDM models \\citep{maller06,stewart08} and potentially destructive for disks \\citep{toth92,stelios07,bullock08,purcell08}.  This makes it potentially difficult to reconcile the observed present day population of disk dominated galaxies with $\\Lambda$CDM, if a quiet merging history is essential for their existence. But do disk galaxies really require a quiet merging history to be observed as such at the present time? What are the time scales of disk growth and destruction?  How quickly do disk galaxies regrow after a merger?  Theoretical models also highlight the difficulties in predicting the outcomes of galaxy mergers and the subsequent regrowth of  stellar  disks. These difficulties arise from several factors. Given the decoupling of the baryonic cores from the parent halos it is difficult to predict robust merging rates of galaxies (as compared to their DM halos) unless cosmological simulations with hydrodynamics are used \\citep{maller06}. Also, numerical simulations have shown how the bulge to disk ratio (B/D) resulting from a major merger might depend on the orbital parameters, the internal spin of the progenitors, the gas fraction of the parent disks and the efficiency of SN feedback \\citep{barneshernquist96,cox06,scannapieco08}.   Recent observational and theoretical work has pointed out ways to alleviate the possible problem of a high merger rate for the survival of disk galaxies. \\cite{hopkins08} showed that angular momentum loss of the gas component is not necessarily catastrophic even in 1:1 mergers. Disks can survive or rapidly regrow, provided that the gas fraction in the disks of the progenitors is high \\citep{robertson06,robertson08,bullock08}.  If feedback is able to suppress star formation during the merger event \\citep{brook04,robertson06,zavala08,g07} the existing cold gas can settle on a new disk plane and start regrowing a stellar disk.  Hydrodynamical simulations indeed show that cold flows (cold gas that flows rapidly to the center of galaxies from filamentary structures around halos \\citep{keres05,dekel08}) play a major role in the build up of disks in galaxies \\citep{brooks08}. Gas accreted through cold flows arrives to the central  stellar disk on a time scale a few Gyrs shorter than gas that is first shocked to the virial temperature of the host halo and then cools onto the disk, leading not only to early disk star formation, but the creation of a large reservoir of cold gas.  No matter its origin, late infalling gas would likely have a higher angular momentum content than material accreted at earlier times \\citep{quinn92} and gas in the merging disks would acquire a coherent spin set by the orbital parameters of the binary system and the internal spins of the parent galaxies.  A feedback$+$cold flows model is particularly attractive as the mechanism for the survival/regrowth of  gas (and then stellar) disks as it provides a natural explanation to the fact that more massive galaxies tend to have large B/D ratios \\citep{benson07,graham08}, while smaller galaxies are likely disk dominated.  At large galaxy masses energy feedback from supernovae (SNe) becomes ineffective at suppressing star formation while cold flows become inefficient at carrying a supply of fresh gas necessary to regrow a stellar disk \\citep{dekel06}. Combined with the relative higher frequency of major mergers for massive galaxies \\citep{guo08} this framework leads to the build up of a larger stellar spheroid and disfavors the quick rebuilding of stellar disks in massive systems.     \\begin{figure*} \\includegraphics[width=14cm]{fig1.eps} \\caption{  {\\bf SDSS rest frame, unreddened $i$-band surface brightness images of the merging system at three different stages}. Left (panel a): z=1.0 close to the first pericentric passage. Center (panel b): the merger remnant at z$=$0.6, 1Gyr after the stellar cores of the two galaxies have coalesced. Right (panel c): the system observed at the present day: the halo has faded considerably while the galaxy has grown an extended thin disk. The $i$-band disk scale length R$_d$ is 7.2 kpc and the light ratio of the kinematically identified Bulge and Disk components (B/D) is 0.49, while a GALFIT 2D decomposition gives a B/D = 0.65.  } \\label{fig1} \\end{figure*}   Only recently have numerical simulations been able to directly compare the observable quantities of the outputs with real data \\citep{jonsson06,chakrabarti07,lotz08,covington08,rocha08}, rather than simply predicting the mass distribution of the simulated stellar systems.  This is a crucial point as the stellar spheroids will fade drastically after a major merger, while reforming disks will contain a large fraction of younger and brighter stars. Different mass to light (M/L) ratios for the two components will skew the observed photometric light ratios compared to the underlying stellar masses.  Despite all of this progress, the effect of mergers on the existing disk components, and of feedback and cold flows on the regrowth of stellar disks have not been studied in detail in fully cosmological simulations of major mergers. Here we present results from a fully cosmological, smooth particle hydrodynamic (SPH) simulation where late gas accretion plays a major role in the rapid regrowth of a dominant stellar disk in an L$\\star$ galaxy after a z $=$0.8 major merger. In this work we focus on the physical processes that drive the regrowth of the disk as identified by its kinematical properties, but we also measure the structural properties of the simulated galaxy based on the light distribution, i.e. in a way closely comparable to observations. This result helps solve the apparent contradiction that strong interactions at relatively low redshifts are common even for DM halos that are likely to host bright disk galaxies. The paper is organized as follow: \\S 2 discusses the simulations, \\S 3 describes the evolution of the system, \\S 4 describes the observational properties of the merger remnant and \\S 5 the assembly of the disk component. The results are then discussed in \\S 6.   \\section[]{Description of the Simulation}   The simulation described in this paper is part of a campaign of high resolution simulations aimed at studying the formation of field galaxies in a WMAP3 cosmology ($\\Omega_0$=0.24, $\\Lambda$=0.76, h=0.73, $\\sigma_8$=0.77, {\\bf $\\Omega_b$=0.042}).  Our sample of halos has been selected from low resolution volumes of size 50 and 25 Mpc (the latter for dwarf size halos) using the ``zoom-in'' tecnique to ensure a proper treatment of tidal torques \\citep{katz93}.  The sample covers two magnitudes in total mass (from 2$\\times$10$^{10}$ to 2$\\times$10$^{12}$ M$_\\odot$) sampling a representative range in halo spins and epochs of last major merger. The galaxy described in this paper has a total mass of 7$\\times$10$^{11} $M$\\odot$, spin $\\lambda$ = 0.04 (as defined in \\cite{bullock01}) and a last major merger at z$=$0.8. The environment density at z$=$0 is typical of field halos with an over density $\\delta \\rho$/$\\rho$ = 0.1 (0.2) measured over a sphere of radius 3 (8) h$^{-1}$ Mpc. The final galaxy has 650k, 320k and 1.8m DM, gas and star particles within the virial radius of the galaxy at the present time, with particle masses: 1.01$\\times$10$^6$, 2.13$\\times$10$^5$ and 6.4$\\times$10$^4$ M$_{\\odot}$ for each DM, gas and star particle at the moment of formation.  The force spline softening is 0.3 kpc. The minimum smoothing length for gas particles is 0.1 times the force softening.  The simulations were performed with the N-body SPH code GASOLINE \\citep{wadsley04} with a force accuracy criterion of $\\theta$ = 0.725, a time step accuracy of $\\eta$ = 0.195 and a Courant condition of $\\eta_C$ = 0.4.  The adopted star formation and SN schemes have been described in detail in \\cite{stinson06} and \\cite{g07}.  Our ``blastwave'' feedback scheme is implemented by releasing energy from SN into gas surrounding young star particles.  The affected gas has its cooling shut off for a time scale associated with the Sedov solution of the blastwave equation, which is set by the local density and temperature of the gas and the amount of energy involved. At the resolution of this study this translates into regions of $\\sim$ 0.2-0.4 kpc in radius being heated by SN feedback and having their cooling shut off for 10-30 million years. The effect is to regulate star formation in the disks of massive galaxies and to lower the star formation efficiency in galaxies with circular velocity in the 50 $<$ V$_c$ $<$150 km/sec range \\citep{brooks07}.  At even smaller halo masses (V$_c$ $<$ 20-40 km/sec, with V$_c$ = sqrt(G $\\times$ M/r)) the collapse of baryons is partially suppressed by the cosmic UV field \\citep{hoeft06,g07}, here modeled following \\cite{haardtmadau96}.  The simulation applied a correction to the UV flux for self shielding of dense gas as introduced in \\citet{pontzen08} and low temperature metal cooling \\citep{wadsley08}. It is important to note that the only two free parameters in the SN feedback scheme (the star formation efficiency and the fraction of SN energy coupled to the ISM) have been fixed to reproduce the properties of present day galaxies  (star formation rates, Schmidt law, cold gas turbulence, disk thickness) over a range of masses \\citep{g07}.  Without further adjustments this scheme has been proven to reproduce the relation between metallicity and stellar mass \\citep{brooks07,maiolino08} and the abundance of Damped Lyman $\\alpha$ (DLA) systems at z=3 \\citep{pontzen08}. However, even at this resolution the central regions (r$<$1-2 kpc) of galaxies remain still partially unresolved and likely form bulges that are too concentrated \\citep{g08,mayer08}.  In this respect the mass of the bulge component of our simulated galaxy has to be considered an upper limit imposed by current resolution limits, especially at high z, where the number of resolution elements per galaxy is less.  To properly compare the outputs from the simulation to real galaxies and make accurate estimates of the {\\it observable} properties of galaxies, we used the Monte Carlo radiation transfer code {\\it{SUNRISE}} \\citep{jonsson06} to generate artificial optical images (see Figure 1) and spectral energy distributions (SEDs) of the outputs of our run. {\\it{SUNRISE}} allows us to measure the dust reprocessed SED of every resolution element of our simulated galaxies, from the far UV to the far IR, with a fully 3D treatment of radiative transfer. Filters mimicking those of the SDSS survey \\citep{adelman06} are used to create mock observations.  To measure the rotational velocity of the remnant and other galaxies in the sample in a way comparable with observations we used the spatial and kinematic distribution of cold gas in the disk of our simulated galaxy. Using the HI fraction calculated by GASOLINE, we determined the HI line width, ($W_{20}$), by finding the width of the HI velocity distribution at 20\\% of the peak.  The value of $W_{20}$/2 is then used as a measure of the galaxy rotation velocity at different redshifts. This measurements reflects the mass weighted position-velocity distribution of cold gas inside the central region of the galaxy and in  observed bright galaxies it is usually associated with the peak rotational velocity.  Because the mass inside the central few kpc is likely overestimated owing to resolution effects \\citep{mayer08}, $W_{20}$/2 provides a slightly larger measurement of the rotational velocity than that obtained from the rotational speed of young stars at 2.2 or 3.5 disk scale lengths (typically 10--20kpc for bright galaxies), and thus is a useful upper limit for a comparison with real data.   \\begin{table}{} \\centering \\caption{Merger remnant properties at z$=$0.8 (shortly after the merger) and at the present time.  $R_d$ is the disk scale length and B/D is the bulge to disk light ratio measured for the kinematically identified components.  (B/D ratios in parentheses are measured using a 2D photometric decomposition with GALFIT of the face on projection of the light distribution).  Total magnitudes and colors have been measured in SDSS filters, including the effects of dust (disk inclination 45deg) and in the AB system.  } \\begin{tabular}{@{}llr@{}} \\hline -  & z=0. & z=0.8 \\\\ \\hline Tot Mag $i$ band  &  -22.3   & -22.7 \\\\ R$_d$ i band (kpc)& 7.2 &  4 \\\\ B/D $i$ band & 0.49 (0.65) & 1.1 (1.4) \\\\ B/D (stellar mass) & 0.87 & 1.16   \\\\ $g-r$  & 0.5  & 0.4 \\\\ SFR  & 2.2 M$_{\\odot}$/yr & 6 M$_{\\odot}$/yr\\\\ disk stellar mass       & 3.24$\\times$10$^{10}$ & 2.07$\\times$10$^{10}$ \\\\    \\hline \\end{tabular} \\end{table}    \\begin{figure} \\includegraphics[width=8cm]{fig2.ps} \\caption{SFR vs time for the last major merger progenitors (red dotted and blue dashed) and for the whole galaxy (solid) . The peaks of the dotted and dashed curves correspond to the two major  mergers between the original four  progenitors.} \\label{fig2a.ps} \\end{figure}   \\begin{figure} \\includegraphics[width=9cm]{fig3.ps} \\caption{Lower Panel: Time evolution of the cold disk gas (T$<$4$\\times$10$^4$ K) scale length and the stellar disk scale length (measured in the unreddened $i$ band) of the final merger remnant.  Upper Panel: The B/D ratio for the final remnant as a function of redshift: dashed: stellar masses. Dotted: B/D light ratio (unreddened $i$ band) with a GALFIT decomposition based on the 2D, face on light profile. Continuous: B/D light ratio (unreddened $i$ band) of the kinematically defined disk and bulge components.} \\label{fig3b} \\end{figure} ",
        "conclusions": "We have analyzed an SPH, fully cosmological simulation of the evolution of a galaxy that grows an extended thin stellar disk after a major (1.6:1) merger at z$=$0.8. The disk dominates the light distribution 3.5 Gyrs after the merger. By the present time the galaxy shares several properties with the observed population of disk dominated L$\\star$ galaxies.    This result is particularly relevant as the assembly history of the galaxy's parent DM halo is different from many previously published studies of cosmological simulations of disk galaxies, as it undergoes a late major merger, whereas previous works tended to select galaxies with a relatively quiet merger history. The two progenitors underwent major mergers themselves, at z $\\sim 2$.  The end result of this simulation strongly contradicts the notion that disk formation requires a ``quiet'' halo merging history.  In fact, the merging history of the halo picked for this study has often been considered hostile to the formation of extended disks at the present time. Our study provides strong support to the notion that galaxy disks can reform in a few Gyrs after a gas rich major merger, while a fraction of the pre-existing stellar disk can survive, even if faded by age and thickened by the strong interaction.    The results of our analysis are made particularly robust by measuring the properties of the remnant light distribution rather than just that of the stellar mass.  This approach is crucial, as it highlights the effects of fading of the light from older stellar populations and plays a major role in quantifying the predominance of the newly formed (and hence bluer and brighter) stellar disk. The main structural parameters of the galaxy at z$=$0 suggest that it has an Sa or Sb morphology: reddened $i$ band B/D $\\sim$ 0.65, Rd = 7.2 kpc, $g-r$ = 0.5, a $W_{20}$/2 HI velocity width of 238 km/sec.  Moreover, being able to measure the reddened colors and brightness evolution of the progenitors allows us a comparison with high-z galaxies, showing that they are representative of the population of blue, gas rich and moderately star forming galaxies observed at z$>$1. Verifying that the progenitors have some of the observed properties of high redshift galaxies is important, as it makes the present day properties of the merger remnant more relevant, having been built from realistic progenitors.    As our simulation includes a full treatment of the cosmological environment it includes a realistic treatment of a number of hydrodynamical processes that are necessary for the regrowth of stellar disks and that have to be simultaneously included: continuous inflow of gas from the cosmic web, stellar feedback, and cooling from hot halo gas. This explains the difference between our results and those of the collisionless simulations of \\cite{purcell08}, which showed significant disk heating due to interactions and accretion events. Lack of gas infall prevented the disks in their simulations to regenerate a new thin component (gas resupply was in fact advocated as a possible solution). In our simulation the effect of interactions with infalling satellites, both luminous and dark, is naturally included. However, interactions and minor mergers are not strong enough to significantly disrupt or thicken the disk as it reforms from new gaseous material.       Idealized hydrodynamic simulations of binary galaxy mergers \\citep{robertson06,hopkins08} did not include the continued infall of gas from cold flows and the hot halo. These works pointed out that without any subsequent accretion/growth onto the disk after the merger, a cold gas fraction in excess of 50\\% would be required in the disks of the progenitors in order to re-build a disk dominated galaxy after the merger.  Here we have shown that such high disk gas fractions are not necessary, as cold flows and cooling from a hot halo make disk regrowth possible even at low redshifts, when the gas fraction in the progenitors' disks and the remnant is only $\\sim$20\\% at any given time. The analysis of the build up of the disk in our simulated galaxy highlights the dominant role of cold flows at high redshift. Cold flows funnel gas to the central disk on a time scale a few Gyrs shorter than gas that is first shocked to the virial temperature of the host halo and then cools onto the disk, leading not only to early disk star formation, but the creation of a large reservoir of cold gas.  Thus, if a substantial fraction of this cold gas reservoir survives a merging event, it can provide a faster accretion rate onto the reforming disk than that available from cooling the hot gas in galaxy halos alone. Rapid disk reformation will also be aided if cold gas accretion onto the galaxy halo is still occurring, while gas cooling from the hot phase forms just 30\\% of the post merger stellar disk. However, this young component is more evident when the system is observed in bluer bands. It is important that these processes of disk formation and destruction be carefully implemented in analytical models that study the properties of large sample of galaxies.   More work is also needed to make detailed quantitative predictions about the morphology of galaxies formed in cosmological simulations and to make the results of our study more general.  The maximum halo mass and lowest redshift at which mergers can regrow disks are obviously a function of the feedback efficiency \\citep{brook04b,scannapieco08} and could therefore provide useful qualitative tests of models of SF and feedback.  In dense environments, such as groups or clusters, the gas reservoir associated with cold flows and cooling halo gas will likely be disrupted by tidal forces and ram pressure stripping. Hence the disk re-growth process should be much less efficient, as expected by the observed correlation between galaxy morphology and environment density.  Also, higher resolution cosmological simulations will have to address the role of secular processes on the detailed structure of the bulge and thin disk components and their role in setting the B/D ratio of galaxies \\citep{debattista04,genzel08aph, weinzirl08aph}. Still, results of the work presented here greatly alleviate the problem posed for the existence of disk galaxies by the high likelihood of interactions and mergers for galaxy sized halos at relatively low z."
    },
    "0812/0812.0186_arXiv.txt": {
        "abstract": "We present the first application of mid-infrared Period-Luminosity relations to the determination of a Cepheid distance beyond the Magellanic Clouds.  Using archival IRAC imaging data on NGC~6822 from Spitzer we were able to measure single-epoch magnitudes for sixteen long-period (10 to 100-day) Cepheids at 3.6$\\mu$m, fourteen at 4.5$\\mu$m, ten at 5.8$\\mu$m and four at 8.0$\\mu$m. The measured slopes and the observed scatter both conform to the relations previously measured for the Large Magellanic Cloud Cepheids, and fitting to those relations gives apparent distance moduli of $\\mu_{3.6}$ = 23.57$\\pm$0.06, $\\mu_{4.5}$ = 23.55$\\pm$0.07, $\\mu_{5.8}$ = 23.60$\\pm$0.09 and $\\mu_{8.0}$ = 23.51$\\pm$0.08~mag. A multi-wavelength fit to the new IRAC moduli, and previously published BVRIJHK moduli, allows for a final correction for interstellar reddening and gives a true distance modulus of 23.49 $\\pm$ 0.03~mag with E(B-V) = 0.26~mag, corresponding to a metric distance of 500$\\pm$8~kpc. ",
        "introduction": "We have embarked upon a program to recalibrate extragalactic distance scale using Cepheid variables observed at mid-infafred wavelengths. Two calibration papers (Freedman et al.  2008 and Madore et al. 2009) gave the first-epoch and then the dual-epoch calibrations of the four mid-infared period-luminosity relations at 3.6, 4.5, 5.8 and 8.0$\\mu$m based upon 76 Cepheids in the Large Magellanic Cloud as originally measured by the SAGE project (Meixner et al. 2006). Single-phase observations at these wavelengths have scatter contributed in almost equal measure by the natural width of the Cepheid instability strip and by the randomly sampled amplitudes of the individual Cepheids themselves.  In the calibration papers we found that  single-phase scatter of $\\pm$0.14~mag is indicative of all four of the mid-infrared PL relations. In addition to the small scatter, the main advantage of observing Cepheids at these long wavelengths is the overall insensitivity of these magnitudes to interstellar extinction. For instance, compared to the optical B band, the extinction suffered by stars measured at 4.5$\\mu$m (say) is reduced by a factor of 60 according to the extinction curve of Rieke \\& Lebovsky (1996, their Table 3). Thus any star with a line-of-sight B-band extinction 0.60~mag or less will have mid-IR extinction corrections of less than 1\\%. For the majority of extragalactic Cepheids reddenings are not all that high and this effectively means that mid-IR distance moduli are effectively true distance moduli.   Hubble (1925) was the first discover Cepheids in NGC~6822. In Hubble's own words NGC~6822 is ``the first object definitely assigned to a region outside the galactic system'' based on ``familiar relations such as those connecting periods and luminosities of Cepheids ...''. Hubble listed 11 Cepheids with periods ranging from 12 to 64 days.  Nearly half a century later, Kayser's (1967) comprehensive examination of the resolved stars in NGC~6822 brought the number of Cepheids up to 13 and modernized the study of this galaxy by introducing both B and V observations of the main body of the galaxy. Kayser also selectively used UBV data of bright field stars along the line of sight to NGC~6822 to show that there was appreciable reddening (mostly foreground) between us and the Cepheids in NGC~6822. Kayser derived a reddening of E(B-V) = 0.27~mag; subsequent determinations ranged from 0.19 to 0.42~mag (see Gallart et al. 1996). Such a large range of reddening values suggests that a true distance modulus based, say, on V-band data, explicitly corrected for reddening, would immediately carry a systematic uncertainty whose range would amount to 0.7~mag. Indeed, most of the variance in subsequently published distances to NGC~6822 are strongly correlated with the adopted and/or derived reddenings.   More recent applications of the Cepheids in determining the distance to NGC~6822, following Kayser's (1967) study, include the multi-wavelength BVRI CCD study of six Cepheids by Gallart et al. (1996), and the near-infrared (H-band) study of 9 Cepheids by McAlary et al. (1983). More recently Gieren et al. (2006) obtained J and K photometry for 56 Cepheids in NGC~6822 based on a massive monitoring program undertaken by Pietrzynski et al. (2004) which resulted in the cataloging of 116 variables with periods ranging from 17 to 124 days. ",
        "conclusions": ""
    },
    "0812/0812.4888_arXiv.txt": {
        "abstract": "{Like other starburst galaxies, M\\,82 hosts compact, massive ($>5\\times 10^5$\\,M$_{\\odot}$) young star clusters that are interesting both in their own right and as benchmarks for population synthesis models. } {In addition to assessing or reassessing the properties of some of the brightest near-IR sources in M\\,82, this paper addresses the following questions. Can population synthesis models at $\\lambda/\\delta \\lambda \\simeq 750$ adequately reproduce the near-IR spectral features and the energy distribution of these clusters between 0.8 and 2.4\\,$\\mu$m? How do the derived cluster properties compare with previous results from optical studies? } { We analyse the spectra of 5 massive clusters in M\\,82, using data acquired with the spectrograph SpeX on the InfraRed Telescope Facility (NASA/IRTF) and a new population synthesis tool with a highly improved near-IR extension, based on a recent collection of empirical and theoretical spectra of red supergiant stars. } {We obtain excellent fits across the near-IR with models at quasi-solar metallicity and a solar neighbourhood extinction law. Spectroscopy breaks a strong degeneracy between age and extinction in the near-IR colours in the red supergiant-dominated phase of evolution. The estimated near-IR ages cluster between 9 and 30\\,Myr, i.e. the ages at which the molecular bands due to luminous red supergiants are strongest in the current models. They do not always agree with optical spectroscopic ages. Adding optical data sometimes leads to the rejection of the solar neighbourhood extinction law. This is not surprising considering small-scale structure around the clusters, but it has no significant effect on the near-IR based spectroscopic ages. } { The observed IR-bright clusters are part of the most recent episode of extended star formation in M\\,82. The near-IR study of clusters that are too faint for optical observation adds important elements to the age distribution of massive clusters in dusty starbursts. Further joint optical and near-IR spectroscopic studies will provide strong constraints on the uncertain physics of massive stars on which population synthesis models rest. } ",
        "introduction": "Starburst galaxies host large populations of star clusters, and it has become increasingly clear that global star formation processes and cluster formation processes are intimately linked (Meurer et al. 1995, Larsen 2006, Elmegreen 2006). This is an important motivation for detailed studies of young star clusters, their distributions in terms of mass and age, their mass-to-light ratios and their stellar initial mass functions.  Because star formation originates in molecular clouds, starburst galaxies tend to be dusty, and young star clusters can be severely obscured. As beautifully illustrated by near-IR and optical Hubble Space Telescope (HST) images of galaxies such as M\\,82 or the Antennae (de Grijs et al. 2001, Alonso-Herrero et al. 2003, McCrady et al. 2003, Rossa et al. 2007), the brightest near-IR star clusters can go undetected at optical wavelengths, and vice-versa. Optical studies tend to focus on lines of sight through holes in the dust, and the derived distributions of cluster properties are subject to biases. In order to obtain a more complete picture of the cluster populations, it seems crucial to combine optical and near-IR data. This can be done with multiband photometry or with spectroscopy.  The importance of near-IR spectroscopy in studies of dusty star forming galaxies has long been recognised (Rieke \\& Lebofsky 1979), and spectra in the H or K windows have been used to constrain star formation histories (e.g. Rieke et al. 1993, Oliva et al. 1995, Lan\\c{c}on \\& Rocca-Volmerange 1996, Satyapal et al. 1997, F\\\"orster Schreiber et al. 2001).  However, near-IR studies of young stellar populations remain difficult. Among the relevant problems are the following (see main text for details and references).\\\\ (i) Evolutionary tracks for massive stars are uncertain, in particular in the red supergiant phase which is the predominant source of near-IR light at  ages of order 10$^7$\\,yr. \\\\ (ii) Model atmospheres for luminous red supergiants are not yet reliable, and this affects the mapping of isochrones into colour diagrams.  \\\\ (iii) The stellar population of a cluster is stochastic, and the intrinsically small number of bright stars leads to potentially large spreads and systematics in the expected cluster properties. \\\\ (iv) The subtraction of very irregular starburst galaxy backgrounds can be hazardous. \\\\ (v) It matters to photometric studies that the shape of the dust attenuation law depends on the unknown spatial distribution of the dust as well as on the nature of the dust; there are degeneracies in colours, for instance between age and extinction. \\\\ As a result of some or all of the above, discrepancies between photometric studies that do or do not include near-IR passbands can occur. The consistency between results obtained from near-IR and from optical spectroscopic data needs to be tested.  With {\\em extended} spectroscopy of massive star clusters, by which we mean spectroscopy that extends through optical and near-IR wavelengths, it becomes possible to isolate some of the above aspects.  By using stellar absorption features exclusively, dust attenuation issues are mostly eliminated. By focusing on very massive star clusters, stochasticity effects can be kept below acceptable levels. Contamination by other stellar populations on neighbouring lines of sight remains an issue in crowded areas, but can be addressed with multicomponent models, more safely with spectroscopy than with photometry. In the end, the spectroscopic analysis is mostly sensitive to the ingredients of the population synthesis models, such as the stellar evolution tracks and the stellar spectra assigned to each point along theses tracks. One may expect two types of results. If the best-fit models based on the near-IR data are not consistent with the optical data, this points to inadequacies of the input physics of the models, i.e. likely systematic errors in both optical and near-IR cluster ages. Conversely, a good fit to the optical and near-IR data set will provide significantly more robust cluster properties than a fit to a single spectral window or to photometric data.  In this paper, we analyse extended spectra of five massive clusters in the nearby starburst galaxy \\object{M82}, using a population synthesis tool with a highly improved near-IR spectroscopic extension.  We show that it is possible to obtain a good representation of all the near-IR features (0.8-2.4\\,$\\mu$m, $\\lambda/\\delta \\lambda \\sim 750$). We investigate whether the near-IR results are consistent with the optical constraints, and what we can learn about extinction and the possible origins of difficulties faced in photometric age-dating procedures.  Section \\ref{models.sec} introduces the new population synthesis tool. In Sections \\ref{clusters.sec} and \\ref{analysis.sec} we describe and analyse the spectra of star clusters in M\\,82. Selected aspects of the analysis as well as the issue of stochastic fluctuations are discussed in Sect.\\,\\ref{discussion.sec}. Conclusions are summarised in Sect.\\,\\ref{conclusion.sec}. ",
        "conclusions": "\\label{conclusion.sec}  Using new synthetic spectra of single stellar populations that extend to 2.5\\,$\\mu$m and include red supergiant spectra at resolution $\\lambda/\\delta \\lambda \\simeq 750$, we have analysed the near-IR spectra of a few of the most massive star clusters in the starburst galaxy M\\,82.  We demonstrate that very good fits to all the near-IR photospheric features seen at this resolution can now be obtained. In particular, the new synthetic spectra can be used for a precise subtraction of the stellar background in emission line studies. The models also significantly improve predictions of the low resolution energy distribution around 1\\,$\\mu$m over those based on previous stellar libraries. Further improvements around 1\\,$\\mu$m are to be expected from the inclusion of stars warmer than 5000\\,K in the library of near-IR spectra used here.  Between ages of 10 and 60\\,Myr, we have found that the degeneracy between age and extinction is essentially perfect in near-IR broad band photometry (Z,Y,J,H,K), when the extinction law of Cardelli et al. (1989) is used. A simultaneous fit of the near-IR spectra {\\em and} the optical\\,+\\,near-IR energy distribution was obtained for cluster M82-L, and this required a modified extinction law (e.g. the law of Cardelli et al., 1989, with R$_V<3.1$). This new evidence for non-standard extinction laws towards star clusters in starburst galaxies, which is not surprising considering the very inhomogeneous spatial distribution of the dust in these objects, is an additional difficulty in any attempt to derive ages from photometry alone. The use of extended spectra allows to constrain both the ages and the shape of the extinction law. It is worth recalling, however, that a given shape of the extinction law can correspond to a variety of values of the total amount of obscuration (e.g. Witt \\& Gordon 2000).  A table of red supergiant numbers and flux contributions has been provided. We argue that the stochastic nature of the stellar mass function must be kept in mind at all cluster masses, but that with masses above 10$^5$\\,M$_{\\odot}$ the analysis of individual extended spectra nevertheless allows us to test {\\em selected} parameters of the population synthesis models. For instance, we expect it will be possible to exploit the largest differences between various sets of current stellar evolution tracks (e.g. tracks for rotating and for non-rotationg stars). But we discourage attempts to test the shape of the upper IMF unless cluster mass approaches 10$^6$\\,M$_{\\odot}$.  The absolute ages derived from the near-IR spectra depend on model parameters that are still highly uncertain, to a large part because the physics of red supergiants (evolution, spectra) are particularly complex. More work on the stellar models and more confrontations with star cluster data will be needed. Care needs to be taken in matching the apertures of multi-wavelengths observations. It is promising that data quality now allows us to exploit ``details\" that were neglected until now, such as the differences between spectra of supergiants of class Ia, Iab and Ib/II.  The near-IR ages found with the current model assumptions for the observed IR-bright clusters in M\\,82 are concentrated between 9 and 35\\,Myr. Indeed, their spectra display deep bands of CN and CO, and therefore favour the model ages at which the contributions of luminous red supergiants are strongest. Cluster F, with weaker bands, is the oldest cluster of our sample. In most cases, the near-IR molecular bands of the models at the ages derived from {\\em optical} studies are marginally acceptable or too weak. Changing the adopted evolutionary tracks or the parameters assigned to the spectra of the input stellar library can result in deeper near-IR bands over a wider range of ages. Work is in progress to perform these tests, which should in time allow us to select the theoretical model inputs most appropriate for the analysis of star clusters in starburst galaxies."
    },
    "0812/0812.1463_arXiv.txt": {
        "abstract": "Most transiting planets orbit very close to their parent star, causing strong tidal forces between the two bodies. Tidal interaction can modify the dynamics of the system through orbital alignment, circularisation,  synchronisation, and orbital decay by exchange of angular moment. Evidence for tidal circularisation in close-in giant planet is well-known. Here we review the evidence for excess rotation of the parent stars due to the pull of tidal forces towards spin-orbit synchronisation. We find suggestive empirical evidence for such a process in the present sample of transiting planetary systems. The corresponding angular momentum exchange would imply that some planets have spiralled towards their star by substantial amounts since the dissipation of the protoplanetary disc. We suggest that this could quantitatively account for the observed mass-period relation of close-in gas giants. We discuss how this scenario can be further tested and point out some consequences for theoretical studies of tidal interactions and for the detection and confirmation of transiting planets from radial-velocity and photometric surveys. ",
        "introduction": "%  Tidal interaction in a binary system leads to dissipation and echanges of momentum  that tend, over time, to align the rotation axes of the two components, synchronise their rotation and orbital periods, and circularize the orbit. Evidence of these three effects are clearly observed in close binaries \\citep[for a recent review see ][]{maz08}.  \\citet{zah77} and \\citet{hut81} have estimated the timescales of these processes for pairs of stars. Although the general features of tidal evolution in binary systems are well established, many of the details are still sketchy. Even the main process through which tides affect the orbital elements is not firmly determined, and the key intrinsic parameter $Q$ (tide quality factor)  is expected to vary by orders of magnitude for different objects and different orbital configurations. As a result, both the parametrisation and the absolute scale of the alignement, synchronisation and circularisation timescale are still very uncertain.  In the context of tidal interaction,  systems with a close-in gas giant planet can simply be considered as binary systems with an extremely large mass ratio. Theoretical timescales for tidal evolution from \\citet{zah77} and \\citet{hut81} are generally used to justify the assumption of a circular orbit when studying close-in transiting planets, or to identify anomalous values of eccentricity or stellar rotation in systems with a close-in planet.    A key difference with stellar binary systems is that, since the mass ratio is so extreme, the synchronisation of the stellar rotation will require a very large inward spiralling of the planetary orbit, through exchange of angular momentum.  Tidal timescales have been calculated specifically for planetary systems in several studies \\citep[e.g.][]{ras96,mar02,dob04}, but these papers also stress that timescales in no way represent the whole story, and detailed integration of the dynamical evolution of these systems is necessary. \\citet{jac08} have recently carried out such integrations for several known transiting planets, and shown that the tidal evolution varies case by case, and could result in dramatic changes of the orbital parameters over time. \\citet{hut80} have established the limit below which an asymptotic equilibrium state towards which the system can evolve no longer exists, and \\citet{lev09} show that most known transiting planets lie below this limit. Even systems which satisfy the equilibrium condition may not be stable, because the parent star is not in a steady state. Its rotation is slowed down by magnetic drag, which can counteract the tidal forcing and evacuates angular momentum away from the system. As discussed by \\citet{dob04} for instance, if the magnetic drag is large enough, synchronisation is not possible.  In the absence of an aymptotic equilibrium, the cirularisation and synchronisation timescales lose their conventional meaning, and full numerical simulations are required to determine the evolution of the system. These integrations, though, require a good knowledge of the tidal quality factor $Q$ both for the planet and the star. \\citet{mat08} recently reviewed the knowledge of the tidal quality factor $Q$ for close star-planet systems, and concluded that the observed circular orbits would indicate values of $Q$ different by several orders of magnitudes from the values adopted for Solar-System giant planets. Theoretical tidal dissipation models indicate that the $Q$ factor could even vary by orders of magnitudes depending on the frequency of forcing (G. Lesur, priv. comm.).  For these reasons, we choose in this study to consider the present data on transiting planets for signs of tidal evolution without a priori assumptions from theory, beyond the simple fact that the strength of tidal effects scales with the mass ratio of the system and a steep function of the orbital distance. Besides the well-known evidence for circularisation of planetary orbits, we look for signs of evolution of the stellar rotation towards synchronisation. Spin-orbit alignement is another tidal procees accessible to observations in transiting systems, but the present data is not sufficient for statistical analysis \\citep{win08,heb08}. The fourth observable tidal effect -- tidal locking of the planet's rotation -- is out of reach of observations for extra-solar planets.  Obviously, these processes are only relevant if the orbits are not initially circular, aligned and synchronised. In the Solar System for instance , planetary orbits are nearly circular and aligned, not because of tidal effects, but as a relic of initial conditions after formation. The residual eccentricities are sustained by secular interactions with other planets. Formation scenarios of close-in planets by migration in a disc predict initially aligned orbits, with some possible eccentricity injected by interaction with the protoplanetary disc. Scenarios invoking violent dynamical interaction between planets after the dissipation of the disc predict a wider distribution of eccentricities \\citep{for99,cha08,nag08}. The present sample of known planets shows that after their formation planets occupy orbits with a very wide range of eccentricities, circular orbits being rare in the absence of tidal effects.  The present status of empirical evidence for tidal effects in close-in planetary system was summarized by \\citet{maz08}. Evidence for {\\it tidal circulatisation} of planetary orbits is obvious in the period-eccentricity relation of the non-transiting exoplanets: planets with periods of a few days predominantly have circular orbits, while at longer periods planetary orbits cover a wide range of eccentricities. This is coherent with timescale calculations, given the leeway offered by the absence of independent constraint on the $Q$ factor of hot Jupiters \\citep[e.g.][]{hal05}. {\\it Synchronisation} of the stellar rotation with the planetary orbit can be excluded for all close-in planetary systems from the observed rotation rate of the host star, with a handful of exceptions located near the high-mass end of the planet range, including $\\tau Boo$, and HD 162020 \\citep{but02,udr02}. This is also compatible with timescale calculations. The star synchronisation timescale for these two system is of the order of the Hubble time, because of the proximity and high mass of the planets. It is much higher than the Hubble time for other systems.  In Section 2, we re-consider the empirical evidence for tidal effects provided by the rapidly growing sample of transiting extra-solar planets. We find tentative but fairly convincing evidence for more indications of tidal evolution. In Section~\\ref{sec3} we outline a possible coherent global picture that underlines the importance of tides to understand the orbital characteristics of close-in exoplanets. ",
        "conclusions": "We have examined the set of known transiting extrasolar planets as a whole for signs of tidal evolution. We find evidence for well-defined limits in parameter space beyond which the orbits of planets are circularized, and beyond which the host star shows sign of excess spin due to tidal effects, roughly following the dependence expected from simple scale arguments. We find that all presently observed eccentric orbits and fast-rotating host stars for transiting planets are compatible with tidal evolution -- and that alternative explanations in terms of 3-body mechanisms or young stellar ages are not required. The data for non-transiting close-in planets, although less precise, confirm this global picture.  These tentative conclusions would imply that below a certain distance ($P\\sim 3-5 $ days), the orbital properties of planets are controlled by tidal effects. Not only the orbital eccentricity, but also the orbital distance and stellar spin, can be very different from the ones set after planetary formation. In order of increasing mass, we would expect close-in planets to be unaffected ($M<<M_J$), circularized ($M\\sim M_J$), spiralling in ($M \\sim 1-2 M_J$), destroyed ($M \\sim 2-3 M_J$), synchronised ($M > \\sim 3 M_J$). The observed mass-period relation of hot Jupiters \\citep{maz05} would be a footprint of this mass-dependent \"sheperding\" of close-in planets.  One important practical implication is that tidal spin-up of the host star will produce a strong detection bias against massive, close-in planets, especially in Doppler surveys."
    },
    "0812/0812.1668.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {We introduce our imaging survey of possible young massive globular clusters in M31 performed with the Wide Field and Planetary Camera 2 (WFPC2) on the  Hubble Space Telescope (HST).  We obtained shallow (to B$\\sim 25$) photometry of individual stars in 20 candidate clusters.  We  present here details of the data reduction pipeline that is being applied to all the survey data and describe its application to the brightest among our targets, van den Bergh~0 (VdB0), taken as a test case.} %{ methods heading (mandatory) {Point spread function fitting photometry of individual stars was obtained for all the WFPC2 images of VdB0 and the completeness of the final samples was estimated using an extensive set of artificial stars experiments. The reddening, the age and the metallicity of the cluster were estimated by comparing the observed color magnitude diagram (CMD) with theoretical isochrones.  Structural parameters were obtained from model-fitting to the intensity profiles measured within circular apertures on the WFPC2 images.} %results heading (mandatory) {Under the most conservative assumptions, the stellar mass of VdB0 is $M> 2.4\\times 10^4 ~M_{\\sun}$,  but our best estimates lie in the range $\\simeq 4-9\\times 10^4 ~M_{\\sun}$.  The CMD of VdB0 is best reproduced by models having solar metallicity and age $\\simeq 25$ Myr. Ages less than $\\simeq 12$ Myr and greater than $\\simeq 60$ Myr are clearly ruled out by the available data. The cluster has a remarkable number of red super giants ($\\ga 18$) and a CMD very similar to Large Magellanic Cloud clusters usually classified as young globulars such as NGC~1850, for example.} % conclusions heading (optional), leave it empty if necessary {VdB0 is significantly brighter ($\\ga 1$ mag) than Galactic open clusters of similar age. Its present-day mass  and half-light radius ($r_h=7.4$ pc) are more typical of faint globular clusters than of open clusters. However, given its position within the disk of M31, it is expected to be destroyed by dynamical effects, in particular by encounters with giant molecular clouds, within the next $\\sim 4$ Gyr.} ",
        "introduction": "Much of the star formation in the Milky Way is thought to have occurred within star clusters  (Lada et al. \\cite{lada91};  Carpenter et al. \\cite{carp}).   Therefore, understanding the formation and evolution of star clusters is an important piece of the galaxy formation puzzle. Our understanding of the star cluster systems of spiral galaxies has  largely come from studies of the Milky Way. Star clusters in our Galaxy have traditionally been separated into two varieties, open and globular clusters  (OCs and GCs hereafter).  OCs are conventionally regarded as young ($<10^{10}$ Gyr), low-mass  ($<10^4 M_{\\sun}$) and metal-rich systems that reside in the Galactic disk.  In contrast, GCs are characterized as old, massive systems. In the Milky Way,  GCs can be broadly separated into two components: a metal-rich disk/bulge  subpopulation, and a spatially extended, metal-poor halo subsystem  (Kinman \\cite{tom}, Zinn \\cite{zinn};  see also Brodie \\& Strader \\cite{brodie};  Harris \\cite{h01}, for general reviews  of GCs).  However, the distinction between OCs and GCs has become increasingly blurred. For example, some OCs are sufficiently luminous and old to be confused with GCs  (e.g., Phelps \\& Schick \\cite{phelps}). Similarly, some GCs are very low-luminosity systems (e.g., Koposov et al. \\cite{kopos}) and at least one has an age that is consistent  with the OC age distribution (Palomar 1; Sarajedini et al. \\cite{sara}). Moreover, a third category of star cluster,  ``young massive clusters'' (YMCs) are observed to exist in both merging  (e.g., Whitmore \\& Schweizer \\cite{wischw}) and quiescent galaxies  (Larsen \\& Richtler \\cite{lari}), Indeed, YMCs have been known to exist in the Large Magellanic Cloud  for over half a century (Hodge \\cite{ho61}).  These objects are significantly more luminous than  OCs (M$_V\\la-8$ up to M$_V\\sim -15$), making them promising candidate young GCs. Once thought  to be absent in the Milky Way, recent observations suggest that their census may be quite incomplete, as some prominent cases have been found  recently in the Galaxy as well (Clark et al. \\cite{clark}; Figer \\cite{figer}).  Thus, a picture has emerged that, rather than representing distinct entities, OCs, YMCs and GCs may represent regions within a continuum of cluster properties dependent upon local galaxy conditions (Larsen \\cite{lars}). The lifetime of a star cluster is  dependent upon its mass and environment. Most low-mass star clusters in disks are rapidly disrupted via interactions with giant molecular clouds (Lamers \\& Gieles \\cite{lamers}; Gieles et al. \\cite{gieles}). These disrupted  star clusters are thought to be the origin of much of the present field star populations (Lada \\& Lada \\cite{lala}). Surviving disk clusters may then be regarded as OCs or YMCs, depending upon their  mass. Star clusters in the halo may survive longer since they are subjected to  the more gradual dynamical processes of two-body relaxation and evaporation.  The clusters which survive for an Hubble time -- more likely to occur away from the  disk -- are termed GCs (see also Krienke \\& Hodge \\cite{krie1}). To date, no known {\\it thin} disk GCs have been identified in the Milky Way.  After the Milky Way, M31 is the prime target for expanding our knowledge of cluster systems in spirals. However, our present state of knowledge about the M31 cluster system is far from complete.  Similar to the Milky Way, M31 appears to have at least two GC  subpopulations; a metal-rich, spatially concentrated subpopulation of GCs  and a more metal-poor, spatially extended GC subpopulation (Huchra et al. \\cite{huchra};  Barmby et al. \\cite{barm00}). Also, again similar to the Milky Way GCs, the metal-rich GCs in M31 rotate and  show \"bulge-like\" kinematics (Perrett et al. \\cite{perrett}). However, unlike the case in  the Milky Way, the metal-poor GCs also show significant rotation (Huchra et al. \\cite{huchra}; Perrett et al. \\cite{perrett}, Lee et al.~\\cite{lee}). Using the Perrett et al. (\\cite{perrett}) data, Morrison et al. (\\cite{morri}) identified what  appeared to be a {\\it thin} disk population of GCs, constituting some  27\\% of the Perrett et al. (\\cite{perrett}) sample. Subsequently, it has been shown that at least a subset of these objects  are in fact young ($\\leq$ 1 Gyr), metal-rich star clusters rather than old ``classical''  GCs (Beasley et al. \\cite{beas}; Burstein et al. \\cite{burst}; Fusi Pecci et al.~\\cite{ffp}; Puzia et al. \\cite{puzia}).  Fusi Pecci et al. (\\cite{ffp}; hereafter F05) presented a comprehensive study of bright young disk clusters in M31, selected from the Revised Bologna Catalogue\\footnote{\\tt www.bo.astro.it/M31} (RBC, Galleti et al.~\\cite{rbc}) by color [$(B-V)_0\\le 0.45$] or by the strength of the $H\\beta$ line in their spectra ($H\\beta\\ge 3.5\\AA$). While these clusters have been noted since Vetesnik (\\cite{vetes}) and have been studied by various authors, a systematic study was lacking. F05 found that these clusters, that they termed -- to add to the growing menagerie of star cluster species -- ``Blue Luminous Compact Clusters'' (BLCCs), are fairly numerous in M31 (15\\% of the whole GC sample), they have positions and kinematics typical of {\\em thin disk} objects, and their colors and spectra strongly suggest that they have ages (significantly) less than 2 Gyr.  Since they are quite bright ($-6.5\\la M_V\\la -10.0$) and -- at least in some cases -- morphologically similar to old GCs (see Williams \\& Hodge \\cite{will}, hereafter WH01), BLCCs could be regarded as YMCs, that is to say, candidate  young globular clusters. In particular, F05 concluded that if most of the BLCCs have an age $\\ga 50-100$ Myr they are likely brighter than Galactic Open Clusters (OC) of similar ages, thus they should belong to a class of objects that is not present, in large numbers, in our own Galaxy. Unfortunately, the accuracy in the age estimates obtained from the integrated properties of the clusters is not sufficient to determine their actual nature on an individual basis, i.e., to compare their total luminosity with the luminosity distribution of OCs of similar age (see Bellazzini et al. \\cite{cefa} and references therein).  In addition to the question of the masses and ages of these BLCCs, it has become clear that the BLCC photometric and spectroscopic samples in M31 may suffer from significant contamination. Cohen, Matthews \\& Cameron (\\cite{coh}, hereafter C06) presented NIRC2@KeckII Laser Guide Star Adaptive Optics (LGSAO) images of six candidate BLCCs. Their $K\\arcmin$ very-high spatial resolution images revealed that in the fields of four of the candidates there was no apparent cluster. This lead C06 to the conclusion that some/many of the claimed BLCC may in fact be just {\\em asterisms}, i.e. chance groupings of stars in the dense disk of M31. While the use of the near infrared $K\\arcmin$ band (required by the LGSAO technique) may be largely insensitive to very young clusters that are dominated by relatively few hot stars, which emit most of the light in the blue region of the spectrum, the inference is that the true number of massive young clusters of M31 may have been severely overestimated.  Therefore, in order to ascertain the real nature of these BLCCs we have performed  an HST survey to image 20 BLCCs in the disk of M31 (program GO-10818, P.I.: J. Cohen).  The key aims of the survey are:  \\begin{enumerate}  \\item to check if the imaged targets are real clusters or asterisms, and to determine the fraction of contamination of BLCCs by asterisms;  \\item to obtain an estimate of the age of each  cluster in order to verify whether it is brighter than Galactic OCs of similar age. Ultimately the survey aims to provide firm conclusions on the existence of BLCCs (YMCs) in M31 as a distinct class of object with respect to OCs (see Krienke \\& Hodge \\cite{krie1,krie2}, and references therein).  \\end{enumerate}  In the present contribution we describe the data reduction and analysis strategies that we will apply to our cluster sample to estimate their ages and metallicities. The overall procedure is described using the brightest among the observed clusters, VdB0, as a specific case. We conclude this section with a brief presentation of the cluster VdB0, below.  The present paper is organized as follows. The observations and the data reduction procedure are described in detail in Sect.~2; the principal assumptions that will be adopted in the whole survey are also reported in this section. Sect.~3 is devoted to the analysis of the surface brightness profile and of the Color Magnitude Diagram of VdB0, including total luminosity, age and metallicity estimates. In Sect.~4 our main results are briefly summarized and discussed.   %______________________________________________________________ \\begin{table*} \\caption{Positional and Photometric parameters for VdB0 from the RBC$^a$} \\label{table:1} \\centering \\begin{tabular}{c c c c c c c c c c c c c} \\hline\\hline NAME & alt NAME & RA$_{J2000}$ & Dec$_{J2000}$ &X &Y & U & B & V & R & J & H & K\\\\ \\hline VdB0 & B195D$^b$  & 00:40:29.3 & +40:36:14.7 & -47.2$\\arcmin$ &-4.3$\\arcmin$& 14.97 & 15.31 & 15.06 & 14.92& 13.77 & 13.14 & 12.99\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$]  X and Y are projected coordinates in the direction along (increasing Eastward) and perpendicular to the major axis of M31 (increasing Northward) respectively, in arcmin, see Galleti et al. 2004, and references therein. \\item[$^{\\mathrm{b}}$] see Sect.~2.5. \\end{list} \\end{table*} %_____________________________________________________________  %                                                One column figure %----------------------------------------------------------- FIG 0 \\begin{figure} \\centering %\\includegraphics[width=9cm]{mosaic.ps} \\includegraphics[width=9cm]{mosaic_lr.ps} \\caption{F450W mosaic of the whole field sampled by our WFPC2 observations. The cluster VdB0 is at the center of the PC camera.} \\label{imatut} \\end{figure} %   \\subsection{The cluster van den Bergh~0 (VdB0)}  VdB0 was indicated as an open cluster by Hubble (\\cite{hub36}) in the image on the frontispiece of his book {\\em The Realm of the  Nebulae}\\footnote{S. van den Bergh kindly drove our attention to this curious occurrence.}. van den Bergh (\\cite{vdb69}) presents VdB0 as the brightest open cluster of M31, reporting an integrated spectral type A0. He also notes that the cluster contains the Cepheid variable V40 (Hubble \\cite{hubble}). A check of Hubble's (\\cite{hubble}) finding charts revealed that two sources are labeled \\# 40  in his plate VII: one of them seems indeed associated  with the cluster, while the other is $\\sim 8\\arcmin$  away from  VdB0, near the association  OB78 = NGC 206 (van den Bergh \\cite{vdbOB}; see also Hodge  \\cite{open}). The cluster was re-discovered by Hodge \\cite{open}, who classified it as an open cluster (C107, see also Hodge \\cite{atlas}). Finally, Battistini et al. (\\cite{bat87}) listed the cluster as their class D candidate globular cluster number 195 (B195D in the RBC). The failure to identify B195D with VdB0 was due to the fact that the coordinates provided by van den Bergh (\\cite{vdb69}) were in error by $\\simeq 17\\arcsec$. For this reason VdB0 and B195D survived as independent entries in M31 GC catalogues until the present day. In our survey we imaged both the clusters and the WFPC2 images revealed unequivocally that the two targets are in fact the same cluster. In particular the images intended to observe B195D have the cluster in the center of the PC camera while in the VdB0 images the cluster lie in the corner of the PC opposite to the WF cameras, such that part of the cluster is out of the image. In the following (and in the future) we will refer to the cluster as VdB0. The dataset analysed here is the one with the cluster centered on the PC images, hence the actual label in the header of the fits files is B195D.  %                                                One column figure %----------------------------------------------------------- FIG 1 \\begin{figure} \\centering %\\includegraphics[width=9cm]{imaB.ps} %\\includegraphics[width=9cm]{imaI.ps} \\includegraphics[width=9cm]{imaB_lr.ps} \\includegraphics[width=9cm]{imaI_lr.ps} \\caption{F450W (upper panel) and F814W (lower panel) images of the whole PC camera, with VdB0 at the center. The superposed circles have radius r=160, 205, 260, 288 and 330 pixels, from inside out, and mark the edges of the annuli whose CMDs are shown in Fig.~\\ref{clus}, below. The light stripes associated with stars in the F450W image are due to the effect of CTE that is particularly strong in this shallow low-background image.} \\label{ima} \\end{figure} %  VdB0 is located at a projected distance of $R_p= 10.8$ kpc from the center of M31 to the South-West, just $\\sim 4\\arcmin$ from the major axis of the galaxy (see Tab. ~\\ref{table:1}), near the edge of one of the most prominent substructures of the M31 disk, the so called {\\em 10 kpc ring} (see Hodge \\cite{hodge} and Barmby et al. \\cite{spitzer}, and references therein) and within a the large OB association OB80 (van den Bergh \\cite{vdbOB}, A80 in Hodge (\\cite{atlas}) atlas). Its radial velocity ($V_r=-567$ km/s, Perrett et al. \\cite{perrett}) is in full agreement with the rotation curve of the HI disk of M31 (Carignan et al. \\cite{rotcur}), thus confirming the physical association with the thin disk of the parent galaxy (F05). The strong value of the $H_{\\beta}$ index supports the idea that the cluster is younger than 1 Gyr ($H_{\\beta}=4.3$ \\AA,  Perrett et al.~\\cite{perrett}\\footnote{Note that Perret's et al. measures refers to B195D, i.e. the ``alter ego'' of VdB0 whose available coordinates were the most appropriate for the cluster. In this context, it is interesting to note that, adopting a calibration based on old  GCs, Perrett et al. found [Fe/H]=-1.64 for VdB0, from integrated spectral indices (see F05).}). The existing estimates of both $V_r$ and  $H_{\\beta}$ are nicely confirmed by recent high signal-to-noise spectra acquired at the Italian Telescopio Nazionale Galileo (S. Galleti, private communication).  With the assumed reddening and distance, the integrated V magnitude reported in the RBC (see Tab.~\\ref{table:1}) gives an absolute magnitude $M_V=-10.03$, much brighter than any Galactic open cluster older than 10 Myr (see Bellazzini et al. \\cite{cefa}, and below); it appears quite extended and irregular in shape even in ground based images. In these ways VdB0 stands out among the members of our candidate BLCC sample that are, in general, fainter and more compact than it. ",
        "conclusions": "We have outlined the data reduction and scientific analysis strategy that we adopt for our HST-WFPC2 survey of M31 candidate YMCs, whose complete results will be presented in future contributions. As an exemplary case, we have described the study of the cluster VdB0. We have found that VdB0 is a very bright and extended cluster of approximately solar metallicity and of age $\\sim 25$ Myr, with a rich population of blue and red supergiants. %The cluster appear to be %significantly younger than the average age of the surrounding thin disk %field population ($\\sim 100$ Myr, see above).  %_____________________________________________________________ %                 A figure as large as the width of the column %------------------------------------------------------------- \\begin{figure} \\centering %\\includegraphics[width=9cm]{vdb0_gs.ps} \\includegraphics[width=9cm]{vdb0.ps} \\caption{Integrated V mag and total mass as a function of age for various clusters. Galactic Open Clusters (OC, from the WEBDA database) are plotted filled circles, Galactic Globular Clusters (GC, $M_V$ from the most recent version of the Harris (1996) catalogue, i.e. that of February 2003; the ages have been arbitrarily assumed to be 12.0~Gyr for all the clusters) are plotted as $\\times$ symbols. VdB0 is represented as a crossed square at $M_V=-10.42$, from Tab.~\\ref{table:3}. The continuous lines are fixed-stellar-mass models from the set by Maraston (1998, 2005) for SSPs of solar metallicity, with a Salpeter's Initial Mass Function (IMF) and intermediate Horizontal Branch morphology. Note that in this plane, the dependence of the models from the assumed IMF, metallicity and HB morphology is quite small (see B08). The outlier OC at log Age$\\simeq 9.0$ is Tombaugh~1. The long dashed line is the VDB0 evolutionary track including the mass loss by dynamical effects according to the formulas by LG06. The cluster is expected to dissolve within $<4$~Gyr from the present epoch.} \\label{vdb0} \\end{figure} % %_____________________________________________________________   Having clearly ascertained that VdB0 is a real cluster, it remains to be established if it is more similar to ordinary open clusters of the Milky Way than to to the Young Massive Clusters that may be considered as possible precursors of ``disk globulars''. The similarity with LMC objects  typically classified as ``Young Globular Clusters''  such as NGC1850 (see Sect.~3., above)  is quite remarkable and it suggests that VdB0 is not an ordinary OC (but see also point 1, below).  A more general way to compare clusters of different ages, taking into account the fading of the luminosity of SSPs as they age, it is to plot them into a diagram comparing age to some indicator of the stellar mass of the cluster (see, for example,  Whitmore, Chandar \\& Fall \\cite{whitmore}, Gieles, Lamers \\& Portegies-Zwart \\cite{gieles}, and de Grijs, Goodwin \\& Kouwenhoven \\cite{degris}, for recent applications and references). Here we adopt log(Age) vs. absolute integrated magnitude as in B08.   %_____________________________________________________________ %                 A figure as large as the width of the column %------------------------------------------------------------- \\begin{figure} \\centering \\includegraphics[width=9cm]{mir.ps} \\caption{The same as Fig.~\\ref{vdb0} but for near infrared colors. Integrated magnitudes of GCs are taken from Cohen et al. (\\cite{cohir}); the IR magnitudes for VdB0 are taken from Tab.~4. The dotted lines are $M=10^4 M_{\\sun}$ and $M=10^5 M_{\\sun}$ iso-mass models assuming a Kroupa \\cite{kroupa} IMF instead of a Salpeter (\\cite{salp}) IMF, plotted here to illustrate the weak effect of assumptions on IMFs.} \\label{mir} \\end{figure} % %_____________________________________________________________  In Fig.~\\ref{vdb0} VdB0 is compared with Galactic Open Clusters (data taken from the WEBDA database\\footnote{\\tt http://www.univie.ac.at/webda/integre.html}), with  Galactic Globular Clusters (from the latest version of Harris \\cite{h96}  assuming a uniform age of 12 Gyr, a reasonable approximation for our purpose), and with a grid of SSP models with solar metallicity and Salpeter's IMF from the set by Maraston\\footnote{\\tt http://www-astro.physics.ox.ac.uk/~maraston/} (\\cite{mara1,mara2}). As a SSP ages massive stars die while the mass of the most luminous stars decreases (passive evolution). Keeping the total mass fixed, the luminosity of the population fades and, as a consequence, the stellar mass-to-light (M/L) ratio increases. The continuous lines plotted in Fig.~\\ref{vdb0} describe the passive evolution of SSPs of various (stellar) masses: under the adopted assumptions the mass of a cluster of given age and $M_V$ can be read from the grid of iso-mass tracks.  The path of the track passing through the cluster shows what  its luminosity will be in the future if the cluster did not lose stars through dynamical processes (evaporation, tides, ecc.). The latter is clearly not the case in general, and in particular for VdB0. In addition to the relatively mild evaporation driven by two body encounters, it will suffer from the strain of the M31 tidal field and from  encounters with Giant Molecular Clouds (GMC), as the cluster is embedded in the dense thin disk of M31 (Lamers \\& Gieles \\cite{lamers}, hereafter LG06, and references therein). To take these effects into account we used the analytical approach presented by LG06 to produce an evolutionary track including the cluster mass loss by stellar evolution, galactic tidal field, spiral arm shocking, and encounters with giant molecular clouds, plotted in Fig.~\\ref{vdb0} as a long-dashed curve. The LG06 formulas describe the evolution of a cluster located within the Milky Way (thin) disk at the Solar circle. They should provide a reasonable approximation for VdB0 which lies in the disk of M31, at a similar distance from the center of a similarly massive spiral galaxy (van den Bergh \\cite{vdbbook}). The required inputs are the cluster mass, for which we adopted the value that can be read from the SSP grid of Fig.~\\ref{vdb0} (see below), and the half-light radius, which we obtained in Sect.~3.3, above (see Tab.~\\ref{table:3}). The initial expulsion of gas not used in star formation may lead young clusters (age $<50$ Myr) to lose their virial equilibrium and it may represent an additional relevant factor driving toward the destruction of clusters like VdB0 that is not included in the LG06 approach (Bastian \\& Goodwin \\cite{bastian6}; Goodwin \\& Bastian \\cite{goodwin}; Bastian et al. \\cite{bastian8}).  Fig.~\\ref{vdb0} is worth of some detailed considerations:  \\begin{enumerate}  \\item{} Independently of the exact value of $M_V$ adopted, VdB0 is significantly brighter ($\\ga 1$ mag) than Galactic OCs of similar ages, actually it is brighter than Galactic OCs of {\\em any} age. The same is true also if all other known M31 OCs are considered (Hodge \\cite{open}; Krienke \\& Hodge \\cite{krie1,krie2}). However it should be noted that the population of disk clusters in M31 may be so huge ($\\sim 80000$ clusters, according to Krienke \\& Hodge \\cite{krie1}) that even the extreme tails of the luminosity distribution may be populated. (This should not be the case for the LMC, for example, as it is orders of magnitude less massive than M31). Hence it is premature to draw a conclusion from an individual cluster; when the whole sample is analyzed we will get a deeper insight on the actual nature of VdB0.  \\item{} Assuming the RBC value for the integrated V magnitude, E(B-V)=0.0 instead of E(B-V)=0.2 and a grid of iso-mass tracks adopting a Kroupa IMF, we can obtain an extremely conservative strong {\\em lower limit} to the stellar mass of VdB0, $M=2.4\\times 10^4 M_{\\sun}$. Under the same assumptions but adopting the best-fit value  E(B-V)=0.2 we obtain $M=6.5\\times 10^4 M_{\\sun}$ with a Salpeter IMF and $M=4.2\\times 10^4 M_{\\sun}$ with a Kroupa IMF. These are at the threshold between the OC and GC mass distributions (see van den Bergh \\& Lafontaine \\cite{vdblf} and B08) and also at the upper end of the mass distribution of Galactic YMC (see Figer \\cite{figer} and Fig.~\\ref{mvrhN}, below). The conclusion that VdB0 is much more massive than MW clusters of similar ages seems inescapable, unless extreme IMFs are considered (i.e. IMF truncated at low masses, see Sternberg \\cite{stern}).  \\item{} If $M_V=-10.42$ is adopted, as obtained from large aperture ground-based V photometry in Sect.~3.4, the total stellar mass is  $M=9.5\\times 10^4 M_{\\sun}$ with a Salpeter IMF and $M=6.0\\times 10^4 M_{\\sun}$ with a Kroupa IMF.  \\item{} The evolutionary tracks including the LG06 treatment of mass-loss by dynamical effects show that, independent of the actual mass (within the range outlined above), it is unlikely that the cluster VdB0 would survive  for an Hubble time.  Hence it is very probable that it will never have the opportunity to evolve into a classical (faint) GC. The disruption timescale is dominated by encounters with GMCs; considering this effect alone (Eq.~7 of LG06) the cluster is predicted to dissolve within $\\simeq 3.6$ Gyr if its mass is $M=9.5\\times 10^4 M_{\\sun}$, as obtained from our best estimate of the integrated V magnitude and assuming a Salpeter's IMF.  \\item {} In the same grid of Fig.~\\ref{vdb0} and under the same assumptions the masses of the BLCCs observed by WH01 - adopting their age estimates -  range  from $8.0\\times 10^3 M_{\\sun}$, (G293) in the realm of OCs, to $\\simeq 2 \\times 10^4 M_{\\sun}$ (G44 and G94) and $8\\times 10^4 M_{\\sun}$ (G38), very similar to that of  VdB0 and significantly larger than OCs of similar ages.  \\end{enumerate}  To obtain independent and more robust estimates of the present-day stellar mass of VdB0 we used the Near Infrared (NIR) version of the log~Age vs. absolute integrated magnitude plane. In Fig.~\\ref{mir}, J,H and K absolute magnitudes of VdB0 extracted from the Extended Sources Catalogue (XSC) of 2MASS are compared with Maraston's SSP models of solar metallicity and Salpeter's (continuous lines) or Kroupa's (dotted lines) IMFs and with Galactic GCs (from Cohen et al. \\cite{cohir}, ages assumed as above)\\footnote{For J,H,K colors we adopt  $A_J=0.871E(B-V)$, $A_H=0.540E(B-V)$, and $A_K=0.346E(B-V)$, from Rieke \\& Lebofsky \\cite{rieke}. The J,H,K absolute magnitudes of the Sun are taken from Holmberg, Flynn \\& Portinari \\cite{holm}.}. NIR integrated magnitudes for significant samples of OCs are not available, at present. To account for the whole extent of the cluster we extracted $r=15\\arcsec$ aperture photometry, that is provided in the XSC, instead of the $r=5\\arcsec$ adopted in the RBC, see Tab.~\\ref{table:1} and Tab.~\\ref{table:3}).  NIR magnitudes are more reliable mass tracers than visual magnitudes as NIR M/L ratios are smaller and have smaller variations with age, compared to optical M/L ratios. For example, according to  Maraston (\\cite{mara1,mara2}) models, a solar metallicity Salpeter-IMF SSP at Age = 10 Gyr has (M/L)$_V$=5.5, while  (M/L)$_K$=1.4; the same SSP has $\\frac{d(M/L)_V}{dt}\\simeq 0.55$ while  $\\frac{d(M/L)_K}{dt}\\simeq 0.13$. The independent estimates of the stellar mass from J,H, and K magnitudes are essentially identical, ranging from 6 to 9 $\\times 10^4 M_{\\sun}$, assuming a Salpeter IMF,  and  from 4 to 5.5 $\\times 10^4 M_{\\sun}$, assuming a Kroupa IMF.  These estimates are in fair agreement with those obtained from the integrated V photometry.   %_____________________________________________________________ %                 A figure as large as the width of the column %------------------------------------------------------------- \\begin{figure} \\centering \\includegraphics[width=9cm]{massrhN.ps} \\caption{VdB0 (crossed square) is compared to other clusters in the {\\em logarithm of the mass} vs. {\\em logarithm of the half-light-radius} plane. Filled circles are Galactic GCs from Mackey \\& van den Bergh (\\cite{macksyd}). Arrows are Galactic OCs: we plot the radii where a break in the surface brightness profile occurs, taken from Kharchenko et al. (\\cite{kar}, their ``core radii''). These should be considered as upper limits for actual $r_h$, which are not available for most OCs. The masses of the OCs have been computed using the grid of SSP models shown in Fig.~\\ref{vdb0}, while for GCs we adopted age=12.0 Gyr and a grid of SSP models having $[Z/H]=-1.35$. Open pentagons are the clusters studied by WH01. Open triangles are the massive young MW clusters listed by Figer (\\cite{figer}); masses and radii are taken from his Table~1. Note that the radii reported by Figer for these clusters are not half-light-radii, however they should be a reasonable proxy. The good match between the two quantities has been verified in the case of RSG1, for which Figer report $r = 1.3$~pc, and Davies et al. (\\cite{davies}) obtain $r_h=1.5\\pm 0.3$ pc.} \\label{mvrhN} \\end{figure} % %_____________________________________________________________  %_____________________________________________________________ %                                             Simple A&A Table %_____________________________________________________________ % \\begin{table} \\caption{Newly derived coordinates, half-light radius, integrated magnitudes, reddening, age and metallicity for the cluster VdB0. The origin of each parameter is described in the last column.} \\label{table:3} \\centering \\begin{tabular}{c c c} \\hline\\hline par & value & note \\\\ \\hline &&\\\\ $\\alpha_{J2000}$& $00^h$ $40^m$ $29.4^s$               & from 2MASS-XSC\\\\ $\\delta_{J2000}$& $+40\\degr$ $36\\arcmin$ $15.2\\arcsec$ & from 2MASS-XSC \\\\ \\hline &&\\\\ $r_h$       & 1\\farcs93 $\\pm$ 0\\farcs66  & from intensity profile (i.p.) fit \\\\ %$r_t$       & $\\simeq 15\\arcsec $  & K66 best-fit model of the F814Wi.p.\\\\ \\hline &&\\\\ B     &  14.94$\\pm$ 0.09  & r=14\\farcs4 ap. phot. on M06 images\\\\ V     &  14.67$\\pm$ 0.05  & r=14\\farcs4 ap. phot. on M06 images\\\\ R     &  14.45$\\pm$ 0.11  & r=14\\farcs4 ap. phot. on M06 images\\\\ I     &  14.01$\\pm$ 0.11  & r=14\\farcs4 ap. phot. on M06 images\\\\ J     &  13.26$\\pm$ 0.07  & r=15\\farcs0 ap. phot. from 2MASS-XSC\\\\ H     &  12.76$\\pm$ 0.12  & r=15\\farcs0 ap. phot. from 2MASS-XSC\\\\ K     &  12.77$\\pm$ 0.15  & r=15\\farcs0 ap. phot. from 2MASS-XSC\\\\ \\hline \\hline &&\\\\ age & 25 Myr               & value of adopted best-fit isochrone\\\\ Z   & 0.019                & value of adopted best-fit isochrone\\\\ E(B-V)   & 0.20            & adopted best-fit value\\\\ \\hline \\end{tabular} \\end{table} %   Finally, in Fig.~\\ref{mvrhN} we compare VdB0 with Galactic OCs, GCs and YMC, plus the BLCCs studied by WH01, in the log of the {\\em stellar mass} versus log of the half-light radius plane (similar to Mackey \\& van den Bergh \\cite{macksyd} and  Federici et al. \\cite{luckyrt}). The radii at which the break in the profile (core/corona transition) of Galactic OCs (from Kharchenko et al. \\cite{kar}) occurs is taken as a strong upper limit for their $r_h$. VdB0 has a typical size that is larger than both OCs and YMCs, and is similar to that of several MW GCs of comparable mass.  In conclusion, we can say that VdB0 seems a remarkable cluster in several of its properties when compared to the other known disk clusters of the Milky Way and M31. In this paper we have presented the data reduction, data analysis and diagnostics that will be applied to the whole survey sample and that will allow us to put VdB0 and the other clusters in the more general context of the star cluster populations in the disk of spiral galaxies."
    },
    "0812/0812.0559_arXiv.txt": {
        "abstract": "We argue that WIMP dark matter can annihilate via long-lived ``WIMPonium\" bound states in reasonable particle physics models of dark matter (DM).   WIMPonium bound states can occur at or near threshold leading to substantial enhancements in the DM annihilation rate, closely related to the Sommerfeld effect.  Large ``boost factor\" amplifications in the annihilation rate can thus occur without large density enhancements, possibly preferring colder less dense objects such as dwarf galaxies as locations for indirect DM searches.  The radiative capture to and transitions among the WIMPonium states generically lead to a rich energy spectrum of annihilation products, with many distinct lines possible in the case of 2-body decays to $\\gamma\\gamma$ or $\\gamma Z$ final states.   The existence of multiple radiative capture modes further enhances the total annihilation rate, and the detection of the lines would give direct over-determined information on the nature and self-interactions of the DM particles. ",
        "introduction": "Introduction}  One of the most promising ways in which to probe the nature of weakly-interacting-massive-particle (WIMP) dark matter is provided by indirect detection experiments in which the annihilation products of dark matter in astrophysical contexts are observed. Indirect detection experiments such as AMS \\cite{ams}, ATIC \\cite{atic}, EGRET \\cite{egret}, GLAST \\cite{glast},  HEAT \\cite{heat}, HESS \\cite{hess}, INTEGRAL \\cite{integral}, PAMELA\\cite{pamela}, and VERITAS \\cite{veritas} look for signals ranging across energetic gamma rays, positron excesses, and anti-proton fluxes, and hints of deviations from background expectations now abound.  A common feature in the interpretation of these experiments is that the WIMP annihilation proceeds via simple, almost free-particle annihilation leading to a rate that depends on the WIMP relative velocity in only a very simple, essentially structureless way, and moreover, is directly related to the cross-section that led to the DM density at freeze-out. This assumption then implies that the primary astrophysical quantity determining the fluxes from DM annihilation is the local DM squared-density, $\\rho({\\bf x})^2$, with the velocity distribution of the DM being essentially irrelevant.  In addition the assumption of almost free-particle annihilation gives a relatively simple resulting energy spectrum for the annihilation products (though of course features such as steep falls at kinematic thresholds are generic). The dependence of the annihilation fluxes on $\\rho({\\bf x})^2$ has been widely employed in the suggested interpretations of the various experimental anomalies as, very frequently, a so-called ``boost-factor\" in the annihilation rate is required to match the observed flux, and this is assumed to come from local over-densities of DM.  In this paper we point out that annihilation of WIMP dark matter via intermediate long-lived ``WIMPonium\" bound states, $\\Om^{(n,\\ell)}$, is possible in many particle physics models of DM \\cite{ourtalks} (see \\cite{posp} for another recent discussion of WIMPonium).  As we argue below, the WIMPonium bound states can occur at or near threshold in which case they lead to potentially very substantial (factors of $10^3$ to $10^5$)  enhancements in the DM annihilation rate, closely related to the well-known Sommerfeld non-perturbative enhancement \\cite{sommerfeld} that has been applied to freeze-out calculations and indirect dark matter signals  \\cite{freezeout,higgsportal,posp,others,jomass}.  The existence of this $\\rho^2$- independent dynamically-induced ``boost factor\" has important implications:  First large amplifications can occur in low-velocity-dispersion systems without a large (eg, cusp-like) density enhancement \\footnote{Furthermore cusp-like dark matter distributions seem not to be favoured at present, see for example \\cite{nocusp}.}, possibly preferring colder less dense objects such as dwarf galaxies as more promising places to search for DM signals compared to the higher density but higher velocity galactic center \\footnote{See \\cite{subir} for a travel guide to the directions in the sky where DM annihilation signals may be expected}.  Second, a knowledge of the full phase space density of the DM is necessary to reliably compute the DM annihilation rate, which strengthens further the case for detailed realistic simulations of the DM distribution in our galaxy.  In addition, the deeply-bound WIMPonium spectrum is often quite involved, and the radiative capture to and transitions among the various states generically lead to a rich energy spectrum of annihilation products, with many distinct lines (in the case of decays to $\\gamma\\gamma$ or $\\gamma Z$ final states, similar to that found in WIMP annihilations \\cite{bergetal}) possible. Although the observability, or otherwise of these distinct lines depends on the DM model being studied, and especially the resolution of the detector, the existence of all the various radiative capture modes further enhances the total annihilation rate, and the detection of even just a few of the lines would give direct information on the interactions of the DM particles. ",
        "conclusions": "We have shown that theories of TeV-scale physics can have dark matter candidates whose annihilation proceeds via the formation of near-threshold WIMPonium bound states.  Depending upon the closeness-to-threshold of the weakest-bound state, these can lead to a substantial velocity-dependent amplification of the dark matter annihilation cross-section, preferring the lowest velocity dispersion environments all other factors being equal, and providing a new dynamical source of ``boost factors\" for indirect detection signals.  In addition, the amplified radiative capture to more deeply bound states, and the transitions among such bound states can both lead to a rich spectrum of discrete $\\ga$ lines which, if observed, would give striking confirmation of the mechanism.  During the preparation of this work, \\cite{posp} appeared. This paper also considers aspects of the phenomenology of WIMP bound states in the context of their annihilations and the consequences for dark matter indirect detection."
    },
    "0812/0812.0590_arXiv.txt": {
        "abstract": "Disk self-gravity could play an important role in the dynamic evolution of interaction between disks and embedded protoplanets. We have developed a fast and accurate solver to calculate the disk potential and disk self-gravity forces for disk systems on a uniform polar grid. Our method follows closely the method given by Chan et al. (2006), in which an FFT in the azimuthal direction is performed and a direct integral approach in the frequency domain in the radial direction is implemented on a uniform polar grid. This method can be very effective for disks with vertical structures that depend only on the disk radius, achieving the same computational efficiency as for zero-thickness disks.  We describe how to parallelize the solver efficiently on distributed parallel computers. We propose a mode-cutoff procedure to reduce the parallel communication cost and achieve nearly linear scalability for a large number of processors. For comparison, we have also developed a particle-based fast tree-code to calculate the self-gravity of the disk system with vertical structure. The numerical results show that our direct integral method is at least two order of magnitudes faster than our optimized tree-code approach. ",
        "introduction": "Well before extrasolar planets were discovered, \\cite{GoTr79,GoTr80} and \\cite{LiPa86a,LiPa86b} have speculated that tidal interactions between disks and embedded protoplanets would lead to planet migration. \\cite{Ward97} suggested that two different types of migration could occur. \\cite{NeBe03a,NeBe03b} studied numerically the effects of disk self-gravity in two-dimensional simulations of planet-disk interactions. \\cite{PiHu05a} reported by an analytical derivation that the disk gravity accelerates the planetary migration. Recently, simulations by \\cite{BaMa08} confirmed that the self-gravity indeed accelerates the type I migration. They implemented a 2D self-gravity solver in their code using the fast Fourier transform (FFT) method of \\cite{BiTr87}, which requires a logarithmic radial spacing ($\\log(r)$).  In Newtonian gravity, we can define the gravitational potential $\\Psi$ associated with the mass density, $\\rho$, by the volume integral \\begin{equation} \\Psi({\\bf x}) = -G\\int\\int\\int\\frac{\\rho({\\bf x'})}{|{\\bf x}-{\\bf x'}|}d^3x' \\label{eq_pot1} \\end{equation} over all space, where $G$ is the gravitational constant. Rewriting Eq. (\\ref{eq_pot1}) in differential form, we obtain Poisson's equation \\begin{equation} \\nabla^2\\Psi = 4\\pi G\\rho, \\label{eq_pot2} \\end{equation} with $\\Psi$ satisfying the boundary condition $\\Psi(\\infty)= 0$ at all times. The numerical ``Poisson solvers'' to Eq. (\\ref{eq_pot2}) can be classified into two categories: difference methods and integral methods. The difference methods solve Eq. (\\ref{eq_pot2}) directly by either finite-difference or finite-element methods with necessary boundary conditions. The known boundary conditions at infinity are usually not very useful if a finite domain is considered. User specified boundary conditions or Dirichlet boundary conditions obtained by direct summation are often required.  A key advantage of difference methods is that, generally, they are relatively fast once initialized. However, they have low or limited accuracy, and they often rely on the integral method to provide the boundary conditions.  The integral methods are to integrate Eq. (\\ref{eq_pot1}) directly. They have the advantage that the summation stops naturally at the domain boundaries. However an integral method has two difficulties in a practical implementation. One is that the integral has a point mass singularities (i.e. when ${\\bf x'}\\to {\\bf x}$ in Eq. (\\ref{eq_pot1})), which is often circumvented by introducing softening. The other difficulty is that they are computationally prohibitive for a large system. The computational cost can be significantly lowered if the FFT method can be used.  The gravitational field, $g=-\\nabla \\Psi$, is more convenient in many applications. It can be calculated via numerical derivatives, which generally have poor precision. To obtain high accuracy, we can integrate the field directly by \\begin{equation} g({\\bf x}) = G\\int\\int\\int\\frac{\\rho({\\bf x'})({\\bf x'}-{\\bf x})}{|{\\bf x}- {\\bf x'}|^3}d^3x'~~, \\label{eq_f1} \\end{equation} which shares the same two difficulties as calculating Eq. (\\ref{eq_pot1}).  In this paper, we present a method for computing the disk self-gravity for quasi-2D disk models. We consider a disk with the cylindrical grid $(r,\\phi,z)$. We assume that the scale height of the disk (semi-thickness) is only radius-dependent, i.e., $H= H(r)$, and it is quite small, $H(r)/r \\ll 1$ (namely geometrically thin disks). We also assume that the vertical structure of the density can be described by some function $Z(r,z)$ that is independent of $\\phi$ and time $t$, i.e., $$ \\rho(t,r,\\phi,z) = \\Sigma(t,r,\\phi)Z(r,z). $$ In this paper, we are particularly interested in the gravitational potential and field force at the $z=0$ plane. Although our method is valid for any function of $Z(r,z)$, we assume that the density vertically has a Gaussian distribution. Under this assumption, \\begin{equation} Z(r,z) = \\frac{1}{\\sqrt{2\\pi H(r)^2}}\\exp\\left(-\\frac{z^2}{2H(r)^2}\\right). \\label{eq_z} \\end{equation} With the presence of the vertical structure (\\ref{eq_z}), the migration rate of the protoplanet is reduced up to 50\\% (see \\cite{Koller04}). We assume that the disk has a constant sound speed $c_s$. Then from the scale height $$ \\frac{H}{r} = \\frac{c_s}{v_K}~~, $$ where $v_K$ is the Keplerian velocity $v_K(r) = \\sqrt{GM_\\star/r}$, we obtain $$ H(r) = \\frac{c_s r^{3/2}}{\\sqrt{GM_\\star}}~~. $$ Other profiles for the scale height are also easy to handle. For example, if the aspect ratio, $h=H/r$, is constant, then $H(r) = rh$. The numerical verification in \\S \\ref{sec:test} actually uses a constant $H$ in the whole domain.  Since we are interested in the potential and field in the $z=0$ plane, the easiest method to obtain $\\Psi$ is to integrate directly the equation \\begin{equation} \\Psi(t,r,\\phi) = \\int_{r_{\\min}}^{r_{\\max}}\\int_0^{2\\pi}\\Sigma(t,r',\\phi')r'dr'd\\phi' \\int_{-\\infty}^{\\infty}-\\frac{GZ(r',z')}{\\sqrt{r^2+{r'}^2- 2rr'cos(\\phi-\\phi')+{z'}^2}} dz' \\label{2d_eq1} \\end{equation} where $r_{\\min}$ and $r_{\\max}$ are the inner and outer radial boundaries, and $\\Sigma$ is the vertically integrated density. As in \\cite{Hure05}, we have employed two different coordinate systems in the radial direction: $r$ as the field grid and $r'$ as the source grid. For simplicity, we set $G=1$ and $M_\\star = 1$.  The rest of the paper is organized as follows. In \\S \\ref{sec:non-axi}, we describe the method given by \\cite{ChPO06} on how to solve Eq. (\\ref{2d_eq1}) using a direct integral via a Green's function method. Our method, though essentially follows the one given by \\cite{ChPO06}, differs in that we use a direct summation method on a uniform grid in the radial direction, which is the same grid used in our hydro code. In addition, we present two approaches to circumvent the singularity in Eq. (\\ref{2d_eq1}) and two methods to calculate the force field. For comparison, we have also implemented a 2D tree-code with a simplified 3D treatment to calculate the self-gravity for disks with vertical structures. In \\S \\ref{sec:test} we present an efficient parallel implementation scheme on distributed memory computers. We describe a new algorithm to calculate the gravity force at an arbitrary point. We also present numerical test results to compare different approaches. A few concluding remarks are given in \\S \\ref{sec:conc}. ",
        "conclusions": "}  In this paper, we have presented a fast and accurate solver to calculate the potential and self-gravity forces for the disk systems. This method is implemented on a polar grid, and FFT can be used in the azimuthal direction. The pre-calculation of the Green's function and its FFT play a major role in the algorithm to reduce the computational cost. We think it could be the main reason why the Green's function method is much faster than the particle tree-code method. We also presented an efficient method in implementing the solver on parallel computers. We find the computational cost for the self-gravity solver to be comparable to that of the hydro solver for large, highly resolved grids run on a distributed memory parallel architecture.  We also developed a 2D tree-code solver, which uses a relatively inexpensive model force to accurately account for the vertical structure of the disk.  Finally, we notice that if the disk vertical structure varies with time, our pre-calculation must be done every time step making our solver inefficient.  We have applied our self-gravity solver to simulations of disk-planet interaction system. Compared with the simulations without self-gravity, the total computation time is increased by only 30\\% for a parallel computation with 100 processors. We have also confirmed that the 2D self-gravity indeed accelerates the planet migration. These results will be reported elsewhere.  \\bigskip \\noindent {\\bf Acknowledgment:} We would like to thank Dr. C.-K. Chan for helpful discussion during his stay at Los Alamos.  We also thanks the referee for many useful comments. This research was performed under the auspices of the Department of Energy. It was supported by the Laboratory Directed Research and Development (LDRD) Program at Los Alamos. It is also available as Los Alamos National Laboratory Report, \\laur{07-5882}."
    },
    "0812/0812.4466_arXiv.txt": {
        "abstract": "Using the sample of long Gamma-ray bursts (GRBs) detected by \\swift-BAT before June 2007, we measure the cumulative distribution of the peak photon fluxes ($\\log N$--$\\log P$) of the \\swift\\ bursts. Compared with the \\batse\\ sample, we find that the two distributions are consistent after correcting the band pass difference, suggesting that the two instruments are sampling the same population of bursts. We also compare the $\\log N$--$\\log P$ distributions for sub-samples of the \\swift\\ bursts, and find evidence for a deficit (99.75\\% confident) of dark bursts without optical counterparts at high peak flux levels, suggesting different redshift or $\\gamma$-ray luminosity distributions for these bursts. The consistency between the $\\log N$--$\\log P$ distributions for the optically detected bursts with and without redshift measurements indicates that the current sample of the \\swift\\ bursts with redshift measurements, although selected heterogeneously, represents a fair sample of the non-dark bursts. We calculate the luminosity functions of this sample in two redshift bins ($z<1$ and $z\\ge1$), and find a broken power-law is needed to fit the low redshift bin, where $dN/dL \\propto L^{-1.27\\pm0.06}$ for high luminosities ($L_{peak} > 5\\times10^{48}\\lumin$) and $dN/dL \\propto L^{-2.3\\pm0.3}$ at for low luminosities, confirming the results of several studies for a population of low luminosity GRBs. ",
        "introduction": "The cumulative distribution of source intensities ($\\log N$--$\\log S$) is a useful tool for studying the source populations, especially when the redshifts of the sources are not measured. In the Gamma-ray burst (GRB) field, since the peak photon flux of the GRB is directly related to the detection threshold, the cumulative distribution of peak photon flux ($\\log N$--$\\log P$) is used in many studies. A number of these studies have made use of the \\batse\\ sample of more than 2,000 GRBs. This sample was first used to demonstrate the cosmological nature of the bursts (e.g., Fenimore et al.\\ 1993; Pendleton et al.\\ 1996), and then used to constrain the GRB populations (e.g., Kommers et al.\\ 2000; Lin et al.\\ 2004; Guetta et al.\\ 2005; Dai \\& Zhang 2005), though many assumed that GRB rate follows the star formation rate in the analyses. Since the launch of \\swift\\ (Gehrels et al.\\ 2004), the sample of the \\swift-BAT bursts has reached to 237 in June 2007 (Sakamoto et al.\\ 2008). With the sample size, it is now possible to compare the \\swift\\ and \\batse\\ samples to test whether the two instruments sample the same population of bursts. It is possible that they are different since the two instruments have different band passes and sensitivities. For example, Band (2006) showed that \\swift\\ is more sensitive to soft bursts. Moreover, there is still a large number of bursts without redshift measurements, e.g., dark bursts, and the $\\log N$--$\\log P$ distribution provides a means of studying their source population besides using pseudo-redshifts derived from spectral or timing properties (e.g., Norris 2002).  The fraction of the \\swift\\ bursts with redshift measurements has increased significantly compared to the \\batse\\ bursts. Although selected heterogeneously, it is tempting to measure the luminosity function of the \\swift\\ bursts using this sample (e.g., Liang et al.\\ 2007).  Besides issues with redshift selection effects, the \\swift\\ trigger efficiency has not been well studied, which presents an additional difficulty. We show that by studying the $\\log N$--$\\log P$ distributions of the \\swift\\ and \\batse\\ bursts, we can justify the usage of the heterogeneous redshift sample and set detection thresholds for measuring luminosity functions.  We present the luminosity functions using the heterogeneous redshift sample, where we adopt a cosmology of $H_0 = 70~\\rm{km~s^{-1}~Mpc^{-1}}$, $\\Omega_{\\rm m} = 0.3$, and $\\Omega_{\\Lambda}= 0.7$. ",
        "conclusions": "We compare the $\\log N$--$\\log P$ distributions for the \\swift\\ bursts and \\batse\\ bursts, and find that they are consistent after correcting the band pass differences as suggested by the K-S test results. This shows that the two instruments sample the same population of bursts, although they have different band passes. We can also compare the normalizations of the distributions directly, since the field of view and the observing time of the two samples are measured, and we find that they are consistent. Indirect comparison between the \\swift\\ and \\batse\\ samples has been performed in previous studies (Salvaterra \\& Chincarini 2007; Virgili et al.\\ 2009), where the two samples are found to be consistent with same theoretical models. Given that \\swift\\ is more sensitive than \\batse\\ for the soft bursts (Band 2006), it is puzzling that the two distributions are consistent.  Currently, we are comparing the on-board triggered \\swift\\ sample with the \\batse\\ off-line search sample. It is possible that the difference will become significant when comparing both the off-line search samples.  We also compare the $\\log N$--$\\log P$ distributions for the sub-samples of the \\swift\\ bursts with or without optical afterglow detections or redshift measurements.  We find them to be broadly consistent. However, at the high peak flux regime ($f_{peak} > 5~\\pflux$), we find an indication that there is a deficit (96\\% confident) of ``dark bursts'' (bursts without optical counterparts). This deficit is more evident at $f_{peak} > 10~\\pflux$, where we find the effect is significant at the 99.75\\% confidence level.  A comparison between the $\\log N$--$\\log P$ distributions for dark bursts and optically bright bursts can place important constraints on the origin of the dark bursts.  There are several models proposed for the dark bursts. They can be bursts with intrinsic normal optical-to-$\\gamma$-ray ratios but with large intrinsic or foreground optical extinction (e.g., Taylor et al.\\ 1998; Djorgovski et al.\\ 2001; Fynbo et al.\\ 2001), have intrinsic faint optical afterglows, i.e., low optical-to-$\\gamma$-ray ratios, (e.g., Groot et al.\\ 1998a; Frail et al.\\ 1999), are normal bursts but exist in extremely high redshifts where the Lyman break lands in the optical (e.g., Groot et al.\\ 1998b; Fynbo et al.\\ 2001), or a combination of several origins mentioned above. A different $\\gamma$-ray luminosity distribution or redshift distribution is needed to interpret the difference in the $\\log N$--$\\log P$ distributions. Since the $\\gamma$-ray/hard X-ray flux is not sensitive to the ISM absorption, the hypotheses with optical extinction or intrinsic optical faintness predict that the $\\log N$--$\\log P$ distributions should have a similar shape between dark and optically detected bursts.  For the high redshift scenario, it is possible for the dark bursts to have a different shape in the $\\log N$--$\\log P$ distribution from the normal, non-dark bursts.  Therefore, the deficit of the dark bursts at high flux levels suggests that a significant fraction of dark burst is from extreme high redshifts. In addition, it is possible that the bulk of the dark bursts are selected due to their low optical luminosity, and the deficit of dark bursts at high peak photon fluxes merely reflects the detection threshold of optical observations under the assumption of a constant optical-to-$\\gamma$-ray ratio.  We measure the luminosity function of the \\swift\\ bursts using the sample with redshift measurements. We find the luminosity function can be fit by a broken power-law with $dN/dL \\propto L^{-2.3\\pm0.3}$ at the low luminosity end, and $dN/dL \\propto L^{-1.27\\pm0.06}$ at the high luminosity end. We compare this result with previous measurements (Schmidt 2001; Norris 2002; Stern et al.\\ 2002; Firmani et al.\\ 2004; Guetta et al.\\ 2005; Liang et al.\\ 2007; Salvaterra \\& Chincarini 2007) and find a best match with the result from Liang et al.\\ (2007). The requirement for an additional component at the low luminosity end confirms the existence of the population of low luminosity GRBs claimed by several studies (e.g., Cobb et al.\\ 2006; Piran et al.\\ 2006; Soderberg et al.\\ 2006; Guetta \\& Della Valle 2007; Liang et al.\\ 2007; Virgili et al.\\ 2009). Although we model it as a power-law, it could be the tail of a Gaussian component. At the high luminosity range, the measured slope $dN/dL \\propto L^{-1.27\\pm0.06}$ is close to the prediction of the ``quasi-universal Gaussian jet'' ($dN/dL \\propto L^{-1}$, Zhang et al.\\ 2004). At the very high luminosity end ($L_{peak} > 10^{51}\\lumin$), the ``quasi-universal Gaussian jet'' predicts $dN/dL \\propto L^{-2}$ (Lloyd-Ronning et al.\\ 2004; Dai \\& Zhang 2004), which cannot be tested with the current data.  Unlike many of the previous studies, we do not make any assumption on the GRB rate, since there are recent studies suggesting that the GRB rate does not follow the star formation rate (e.g., Stanek et al.\\ 2006). Instead, we measure the average GRB luminosity function in a large redshift bin. We compare the LFs in two luminosity bins.  The $z\\ge1$ LF shows a drop at two low luminosity data points. This result can be affected by the uncertainties in the \\swift\\ trigger efficiency close to the \\swift\\ detection limit. If we neglect these two data points, the LFs for the two redshift bins are consistent. However, the large measurement uncertainties in the $z<1$ LF make it difficult to test whether the GRB rate follows the star formation rate."
    },
    "0812/0812.1844_arXiv.txt": {
        "abstract": "We have measured the dayside spectrum of HD 189733b between 1.5 and 2.5 $\\mu$m using the NICMOS instrument on the Hubble Space Telescope. The emergent spectrum contains significant modulation, which we attribute to the presence of molecular bands seen in absorption.  We find that water (H$_{2}$O), carbon monoxide (CO), and carbon dioxide (CO$_{2}$) are needed to explain the observations, and we are able to estimate the mixing ratios for these molecules.  We also find temperature decreases with altitude in the $\\sim 0.01<P<$ $\\sim 1$ bar region of the dayside near-infrared photosphere and set an upper limit to the dayside abundance of methane (CH$_{4}$) at these pressures. ",
        "introduction": "HD~189733b is a transiting hot-Jupiter planet in a 2.2-day orbit around a K2V stellar primary \\citep{bouchy05}. Due to the relatively large depth of the eclipse ($\\sim2.5$\\% at K band) and the bright stellar primary (Kmag = 5.5), this system was immediately recognized as an important target for atmospheric characterization observations, and its emission has been studied extensively durning secondary eclipse \\citep{deming05,knutson07,grillmair07,knutson08,charbonneau08}. Multicolor photometric observations revealed the presence of H$_{2}$O \\citep{tinetti07a} and the likely presence of CO \\citep{charbonneau08}, while optical transmission spectroscopy suggests scattering by high altitude haze \\citep{pont08}.  In addition, extensive theoretical work has also been done on the atmosphere of this planet by \\cite{fortney06,barman08,burrows08,showman08} and others.  Recently, the NICMOS camera on the Hubble Space Telescope (HST) was used during the primary eclipse of HD~189733b to obtain a near-IR transmission spectrum of the planet's atmosphere; these results showed the presence of H$_{2}$O and CH$_{4}$ (Swain et al. 2008; hereafter SVT08).  The near-IR transmission spectrum of HD~189733b probes the upper (P $\\sim 10^{-4}$ bar estimated from our models) regions of the of the atmosphere at the terminator.  In this paper, we report the results of HST observations of the dayside spectrum (derived from secondary eclipse measurements).  At near infrared wavelengths, the dayside portion of the atmosphere of HD~189433b is probed in deeper (P $\\sim 0.1$ bar estimated from our models) regions. ",
        "conclusions": "In summary, we have presented the first near infrared spectrum of light emitted by a hot-Jupiter type planet. Using a radiative transfer model, we determine that the molecules H$_{2}$O, CO and CO$_{2}$ are likely present on the dayside of HD~189733b, and we are able to estimate abundances for these species.  Although we cannot tightly constrain the slope of the temperature profile, the observed absorption features indicate that temperature decreases with altitude at pressures between 1 and 0.01 bar.  In addition, there are residual features in the spectrum that could be explained by including opacities of other species that are presently not well constrained.  This work represents an initial step in exploiting the power of spectral analysis (even at low resolutions of $R \\simeq 40$) in determining the chemical composition of extrasolar planetary atmospheres; it also illustrates the extraordinary potential of the NICMOS instrument for characterizing bright, transiting exoplanets. In a previous paper, we presented a transmission spectrum of this planet at the terminator (SVT08), in which methane is seen to be more abundant ($c$ = $5\\times10^{-5}$) in the higher, cooler regions of the terminator region atmosphere.  However, it is difficult to compare directly the previous terminator results with our current dayside results because they probe atmospheric regions with significantly different temperatures and altitudes on this highly irradiated planet. The development of sophisticated global circulation and chemistry models could significantly advance the state of the art in this respect."
    },
    "0812/0812.1305_arXiv.txt": {
        "abstract": "The proposed Transiting Exoplanet Survey Satellite (TESS) will survey the entire sky to locate the nearest and brightest transiting extrasolar planets with orbital periods up to about 36 days.  Here we estimate the number and kind of astrophysical false positives that TESS will report, along with the number of extrasolar planets.  These estimates are then used to size the ground-based follow-up observing efforts needed to confirm and characterize the planets.  We estimate that the needed observing resources will be about 1400 telescope-nights of imaging with 0.5m to 1m-class telescopes, 300 telescope-nights with 1m to 2m-class telescopes for the classification of the host stars and for radial velocity measurements with roughly 1 \\kms\\ precision, and 380 telescope-nights with 2m to 4m-class telescopes for radial velocity studies with precision of a few \\ms. Follow-up spectroscopy of the smallest planets discovered by TESS at the best possible velocity precision will be limited by the number of telescope nights available on 4m to 10-m class telescopes with instruments such as HARPS and HIRES, but the pay-off of such efforts will be the determination of masses for Super Earths with sufficient accuracy to distinguish rocky desert planets from water worlds. ",
        "introduction": "Photometric surveys to detect transits by extrasolar planets have recently become a productive means of locating these objects \\citep[e.g., ][]{koc98,alo04b,bak04,bag07,wal03}.  All such surveys must deal with astrophysical false positives, i.e., periodic transit-like decreases in stellar brightness that arise from stars orbited by other stars, as opposed to stars orbited by planets.  In most surveys these false positives outnumber those from planetary transits, often by factors of ten or more.  The proposed Transiting Exoplanet Survey Satellite (TESS) seeks to use an orbiting array of 6 small-aperture, wide-field CCD cameras (190 mm focal length f/1.5, with 18-degree square field of view) to survey the entire sky in a 2-year observing campaign, yielding a comprehensive list of transiting planets orbiting 2.5 million of the nearest stars. From a near-equatorial orbit, TESS will be able to observe the whole sky, and to detect objects with transit depths as small as $3 \\times 10^{-4}$ of the brightness of the parent star, with periods up to 36 days.  These sensitivities will allow detection of Neptune-sized planets transiting Sun-like stars, and Super Earths around stars with radii somewhat smaller than the Sun's. TESS was selected by NASA for Phase A in May 2008.  In order to size the ground-based follow-up effort needed to distinguish true planets from false positives, one must estimate the rates of false positives that TESS will produce.  This paper describes how we made such estimates, and what we learned.  The point of departure for our computations was the simulation code described by \\citet{bro03}, which should be consulted for information about the computational strategy.  Modifications to the code were needed to make estimates that were appropriate to TESS; the most important of these was the need to average the rates for line-of-sight triple star systems (two stars of which compose an eclipsing binary) over all Galactic latitudes.  Lesser modifications brought the code up to date in terms of the period and radius distributions of presently-known extrasolar planets. ",
        "conclusions": "As discussed above, several kinds of ground-based follow-up observations are feasible and useful, each with its own observing requirements and each able to identify different kinds of false positives.  In this section we summarize the role of these false-positive tests in the context of TESS. To estimate the amount of follow-up observing time that will be needed, we must assign values for the fraction of impostors of different sorts that can be identified using each observing method. These estimates have various justifications, but sometimes the best we can offer are guesses, based in experience, but requiring considerable extrapolation. The reader should therefore treat the numerical estimates in the following paragraphs with caution.  The high-quality TESS photometry itself will allow us to identify some EBs and physical triples from the variability seen in their out-of-transit light curves. A crude guess based on our experience with ground-based surveys is that perhaps 20\\% of EBs and 10\\% of triples can be identified in this way.  Comparing the centroid of the transit signal with that of the apparent host star, as described above, should be much more effective at sorting out line-of-sight triples, with a success rate that might reach 90\\%.  This is because it should be possible to determine the relative position of the image centroid in and out of transit to a small fraction of the image size, roughly the FWHM divided by the S/N of the transit detection, assumed to be at least 7. Locating the positional offset during transit to a third of the image confusion radius should reduce the area allowed for an accidental alignment by an order of magnitude  Given a good ephemeris for a transiting candidate from the TESS photometry, searching for line-of-sight triples using seeing-limited ground-based imaging will also be very effective, and should require relatively little observing time per target.  The eclipsing object in these apparent triples will usually be quite faint relative to the supposed planet host, but for this reason its eclipses must be relatively deep.  A few observations inside and outside of the expected transit times thus should suffice to learn whether the varying star is the supposed planet host star or a faint neighbor. The large number of candidates for this type of follow-up observation should allow a multiplex strategy, where several objects can be followed up simultaneously. Assuming the images are seeing limited, we estimate that a further 90\\% of the line-of-sight triples remaining from the previous steps can be rejected by these observations, using about 0.5 hour total time on a 0.5-m class telescope for each target.  A single spectrum of gravity-sensitive features such as the Mg b lines in the green can immediately identify host stars that are giants. Luminous giants can be distinguished from dwarfs with surprisingly poor S/N, as low as 10 to 20 per spectral resolution element of 10 \\kms\\ \\citep{lat08}. Subgiants, for example near spectral type G, are more difficult to distinguish from dwarfs, because their separation in radius/gravity/luminosity is smaller. The identification of short-period double-lined eclipsing binaries is often revealed by one or two high-resolution spectra with modest S/N, thus demonstrating that the transit-like light curves must be due to grazing eclipses in a system where the secondary is not too much fainter than the primary. Stars that are rotating much too rapidly to allow very precise radial velocity measurements can also be revealed by a single high-resolution spectrum at modest S/N.  We estimate the typical integration time for such spectra to be 10 to 20 minutes with 1m to 2m-class telescopes and modern spectrometers, and that this will allow discrimination of the roughly 15\\% of main-sequence EBs that show composite spectra, and most of the EBs and triple systems (both line-of-sight and physical) where the brightest star is a giant. Based on the relative numbers of giants and dwarfs among stars in TESS's magnitude range, we expect that about one-third of triple systems will be distinguishable in this way.  A few moderate-precision (i.e., 0.5 to 1 km s$^{-1}$) radial velocity measurements provide almost total rejection of the remaining undiluted eclipsing binaries.  We assume that each target will require 2-4 radial velocity measurements, and that each of these spectra will also require 10 to 20 minutes of time on a 1m to 2m-class telescope. Experience suggests that these data will allow correct identification of 95\\% of EBs, 75\\% of Giant EBs, and two-thirds of triples (both line-of-sight and physical).  Accurate multicolor ground-based light curves provide a means of identifying a substantial fraction of the remaining diluted EB systems, namely the ones with deeper transits and big enough color differences between the EB and diluting star (such light curves are also necessary for characterizing the properties of genuine planets), These observations are relatively expensive in observing time because continuous coverage of full events is needed to achieve the necessary photometric precision, and one requires time resolution good enough to map out the shape of ingress and egress.  Transit durations are typically 3 hours; to obtain satisfactory out-of-transit baselines, a typical transit event requires about 5 hours of telescope time, For relatively deep transit events (corresponding to giant planets), the precision achievable for single transit events with 0.5m to 1m-class telescopes will be adequate. Because of the demand for substantial telescope time, precise ground-based photometry is usually reserved for targets that have already been vetted by less demanding methods.  Ground-based photometry of single transit events has rarely achieved a photometric precision better than 1 mmag in sample times on the order of a few minutes.  In a few cases larger ground-based telescopes have been used to push below 1 mmag \\citep[e.g., ][]{joh08}), and even better limiting precision has been demonstrated using a 1m-class telescope to observe multiple events \\citep{win07}. TESS will detect transits that are an order of magnitude shallower than can routinely be achieved from the ground.  These will need to be followed up with space-based resources such as HST and JWST.  In a final step, remaining impostors may be revealed (and true planets characterized) by precise, multi-epoch, high-resolution Doppler spectroscopy.  The high spectral resolution and S/N required for these measurements demands large telescopes, specialized highly stable spectrometers, and, for faint targets, long exposure times.  To get 4 suitable spectra at the level of a few \\ms\\ on these objects will require on the order 2 hours on 2m to 4m-class telescopes, and in some cases (such as stars with significant rotation or weak lines due to low metallicity or high temperature) significantly more.  These observations are expensive and therefore are typically the last to be done, after all other means of rejecting false positives have been exhausted.  Last, it is worth noting that a few false positives are likely to slip through even the exhaustive vetting process just described, and be counted as planets. These are most likely to be special kinds of physical triple systems, in which the EB components are very faint relative to the brightest star. These will yield low-amplitude transits with light curves that are difficult to characterize from the ground. With a large brightness ratio between the brightest star and the EB primary, the orbital displacement of the EB spectrum lines will likewise be difficult to discern. At the same time, however, the blended line profiles may yield a Doppler signal that is plausible for a planet orbiting the brightest star. In short, given the vast number of double and multiple stars in the sky, one must expect that some of them will succeed in simulating planet-bearing single stars. It is difficult to estimate how often this will occur, but given the special circumstances required, we guess the number will be perhaps 20, or about 1\\% of the total sample of planets.  Table 2 shows the expected observing effort implied for the TESS experiment by the above hierarchy of follow-up observations.  Rows in the table correspond to the various observational tests that can be performed.  We assume that these are done in succession, with viable planet candidates passing to the next test, and likely false positives being removed from further consideration.  The first 5 columns show the number of candidate transiting-planet systems that remain, going into each stage of the process, separated into true planets and the 4 different kinds of false positive.  Column 6 shows the number of systems that must be observed at each stage, and column 7 shows the corresponding number of observing hours (and the size of telescope facilities required).   \\begin{deluxetable}{lccccccr} \\tabletypesize{\\scriptsize} \\tablecaption{Follow-Up Observing Resources Needed \\label{Table 2}} \\tablehead{ \\colhead{Observation} & \\colhead{Planets} & \\colhead{MS EB} & \\colhead{Giant EB} & \\colhead{LOS Triple} & \\colhead{Phys. Triple} & \\colhead{\\# of Targets} & \\colhead{Observing Time} } \\startdata Examine TESS & 1687 & 2111 &   40 & 7237 & 2080 & 13155 &    --   \\\\ Light Curves \\\\ \\\\ Seeing-Limited & 1687 & 1689 & 32 & 724 & 1872 & 6004 & 3002h 0.5m-1m \\\\ Imaging \\\\ \\\\ Single Spectrum & 1687 & 1689 & 32 &  72 & 1872 & 5352 & 803h 1m-2m \\\\ (Classification) \\\\ \\\\ 6 Radial Velocity & 1687 & 1435 & 10 & 43 & 1123 & 4298 & 2149h 1m-2m \\\\ Spectra (km s$^{-1}$ \\\\ \\\\ High-Quality & 1687 & 72 & 2 & 16 & 404 & 2181 & 10906h 0.5m-1m \\\\ Light Curve \\\\ \\\\ 4 High-Quality & 1687 & 29 & 1 & 6 & 162 & 1885 & 3769h 2m-4m \\\\ RV Spectra \\\\ \\enddata \\end{deluxetable}  The observations that reveal the largest numbers of false positives are detailed examination of TESS light curves and moderate precision radial velocity measurements.  In both of these stages the total number of false positives is reduced by more than a factor of 2. Intermediate steps are also important, however.  For instance, if spectroscopy is attempted before most of the line-of-sight triples have been weeded out, then one is likely to spend much observing time acquiring spectra of the wrong stars in these systems.  The bulk of the observing time goes to the last 2 stages of the process, namely obtaining high-quality light curves and very precise radial velocity measurements.  By the time the last stage is reached, our experience with ground-based surveys is that somewhat less than half of the candidates prove to be planets. The fraction of true planets should be larger than this for the TESS candidates, because in the relevant range of transit depth, the frequency of both main-sequence EBs and physical triples is a strongly decreasing function of depth, whereas that of planets is almost constant (see Fig. 1).  It is possible to realize some savings in the needed follow-up time by using existing multicolor photometry or astrometry (e.g. 2MASS or Tycho, \\citet{skr06,hog00}) to identify and discard giants from the TESS target list.  Systems in which the light is dominated by a giant are almost always false positives, and there is no need to spend extra observing time determining to which category of false positive they belong.  The total observing time that can be saved by a pre-selection that eliminates giants is fairly modest, however.  The largest contribution of such false positives comes from eclipsing binaries diluted by giants.  Such systems should have a rate of occurrence similar to the fraction of giants in the field. For the magnitude limit of TESS, this fraction is about one-third.  In any case, most of the systems containing giants are eliminated early in the sifting process, at or before the step that calls for one classification-quality spectrum.  These stages are relatively cheap in observing time.  It is the true planets that must be followed all the way to the end of the decision tree, incurring high observing time costs along the way.  Finally, photometry or astrometry suitable for identifying giants is unavailable for most of the fainter stars in the TESS sample, and these are of course by far more numerous than the brighter ones.  For example, only about 50,000 dwarfs out of TESS's sample of 2.5 million stars have Hipparcos parallaxes.  Essentially all of the TESS targets will have 2MASS magnitudes, but with only 2MASS data, the giant-dwarf distinction is clear only for M stars, which are a small fraction of the total.  An optimistic estimate of the improvement yielded by pre-selecting target stars results from assuming that all giant stars can be removed from the sample before the search for transiting objects begins.  In this case, we estimate the total observing effort needed for false-positive rejection will be about 5\\% smaller for imaging, and about 10\\% smaller for low-precision radial velocities, than if no target selection is done.  There would be virtually no change in the amount of required time for highly precise radial velocities. Pre-selection of targets is still worthwhile, because the total amounts of telescope time involved are rather large, and because the brighter targets (for which pre-selection is most practical) are the most interesting ones for later follow-up studies, particularly by JWST.  But for purposes of scoping the ground-based follow-up effort, early elimination of giants is relatively unimportant."
    },
    "0812/0812.4228_arXiv.txt": {
        "abstract": "{Polarimetric maps have been used for the characterization of the magnetic field in molecular clouds. However, it is difficult to determine the 3-dimensional properties of these regions from the projected maps. For that reason, numerical simulations can be used as benchmarks for polarimetric measurements, and eventually reveal more about the interplay of turbulence and the magnetic field lines. In this work we make a number of MHD numerical simulations of turbulent molecular clouds and created their synthetic dust emission polarization maps, varying the direction of the observer. We determined the correlation of intensity emitted and polarization degree for the simulated models. We were able to reproduce the decay of the polarization degree at denser regions without any assumption regarding the properties of the dusty component. The anti-correlation arises from the simple cancellation of the polarization vectors along the line of sight. This effect is amplified within denser regions as the magnetic field configuration becomes more complex. We studied the probability distribution function, the power spectrum, and the structure function of the polarization angles. This statistical analysis revealed strong differences depending on the turbulent regime (i.e. sub/supersonic and sub/super-Alfvenic). Therefore, these methods can be used on polarimetric observations to characterize the dynamics of molecular clouds. We also presented a modified Chandrashekhar-Fermi method to obtain the intensity of the local magnetic field. The proposed formulation showed no limitations regarding orientation or turbulent regime.}  \\resumen{ Mapas polarim\\'etricos han sido usados para caracterizar el campo magn\\'etico en nubes moleculares. Sin embargo, es dif\\'{\\i}cil determinar las propiedades tridimensionales de mapas de proyecci\\'on de dichas regiones. Por esta raz\\'on, simulaciones num\\'ericas pueden ser usadas como pruebas de las mediciones polarim\\'etricas, y eventualmente revelar m\\'as acerca de la relaci\\'on entre la turbulencia y las l\\'{\\i}neas de campo magn\\'etico. En este trabajo realizamos una serie de simulaciones MHD de nubes moleculares turbulentas y creamos mapas sint\\'eticos de emisi\\'on polarizada de polvo, en las cuales variamos la direcci\\'on del observador. Determinamos la correlaci\\'on de la intensidad de la emisi\\'on y el grado de polarizaci\\'on para los modelos. Pudimos reproducir el decaimiento del grado de polarizaci\\'on en zonas m\\'as densas sin suposici\\'on alguna de las propiedades de la componente de polvo. La anticorrelaci\\'on viene la simple cancelaci\\'on de los vectores de polarizaci\\'on a lo largo de la visual. El efecto es amplificado dentro de zonas de alta densidad ya que la configuraci\\'on del campo magn\\'etico se torna m\\'as compleja. Estudiamos la funci\\'on de distribuci\\'on de probabilidad, el espectro de potencias, y funciones de estructura de los \\'angulos de polarizaci\\'on. Estos an\\'alisis estad\\'\\i sticos revelan gandres diferencias dependiendo del r\\'egimen turbulento (es decir sub \\'o s\\'uper-s\\'onico, sub \\'o super-Alfv\\'enico). Entonces, estos m\\'etodos pueden ser utilizados en observaciones polarim\\'etricas para caracterizar la din\\'amica de las nubes moleculares. Presentamos tambi\\'en un m\\'etodo Chandrassekhar-Fermi modificado para obtener la magnitud del campo magn\\'etico local. La formulaci\\'on sugerida no muestra limitaciones respecto al r\\'egimen turbulento u orientaci\\'on. }  \\addkeyword{ISM: magnetic fields} \\addkeyword{techniques: polarimetric} \\addkeyword{methods: numerical} \\addkeyword{methods: statistical}   \\begin{document} ",
        "introduction": "\\label{sec:intro}  Giant molecular clouds in the interstellar medium (ISM) are believed to be threaded by large scale magnetic fields \\citep{sch98, cru99}. However, it is still not completely clear what is the role of the magnetic field in the dynamics of the ISM and what is its effect on the star formation process. Also, the ratio of the magnetic and turbulent energy in these environments is a subject of controversy \\citep{padoan02, girart06}. Polarimetric maps have been extensively used for the determination of the magnetic fields in several astrophysical environments. This technique represents the best tool to characterize the vector components of the magnetic field parallel to plane of the sky. For galaxies and the intracluster medium this technique can be used for polarization of synchrotron emission \\citep{hil00}.  For a given polarization map of an observed region, the mean polarization angle indicates the orientation of the large scale magnetic field. On the other hand the polarization dispersion gives clues on the value of the turbulent energy. This, as a consequence, can be used to determine the magnetic field component in the plane of sky. Chandrasekhar \\& Fermi (1953) introduced a method (CF method hereafter) for estimating the ISM magnetic fields based on the dispersions of the polarization angle and gas velocity. Simply, it is assumed that the magnetic field perturbations are Alfvenic and that the rms velocity is isotropic.  A promising approach to test this method is to create two-dimensional (plane of sky) synthetic maps from numerically simulated cubes. Ostriker, Stone \\& Gammie (2001) performed 3D-MHD simulations, with $256^3$ resolution, in order to obtain polarization maps and study the validity of the CF method on the estimation of the magnetic field component along the plane of sky. They showed that the CF method gives reasonable results for highly magnetized media, in which the dispersion of the polarization angle is $< 25^{\\circ}$. However, they did not present any other statistical analysis or predictions that could be useful to determine the ISM magnetic field from observations. Polarization maps from numerical simulations can also be used in the study of the correlation between the polarization degree and the total emission intensity (or dust column density). Observationally, the polarization degree in dense molecular clouds decreases with the total intensity as $P \\propto I^{-\\alpha}$, with $\\alpha = 0.5 - 1.2$ \\citep{goncalves05}. Padoan et al. (2001) studied the role of turbulent cells in the $P \\ vs \\ I$ relation using supersonic and super-Alfvenic self-gravitating MHD simulations. They found a decrease of the polarization degree with total dust emission within gravitational cores, in agreement with observations, if grains are assumed to be unaligned for $A_V > 3$. When the alignment was assumed to be independent on $A_V$, the anti-correlation was not observed. Recently, Pelkonen, Juvela \\& Padoan (2007) extended this work and refined the calculation of polarization degree introducing the radiative transfer properly. In that work, the decrease in the alignment efficiency arises without any {\\it ad hoc} assumption. The alignment efficiency decreases as the radiative torques become less important in the denser regions. However, it is still not clear the role of the magnetic field topology and the presence of multiple cores intercepted by the line of sight on the decrease of polarization degree.  In this work we attempt to extend the previously cited studies improving and applying the CF method to different situations. For that, we used both sub and super-Alfvenic models, to study the role of the magnetic field topology in the observed polarization maps. ",
        "conclusions": "In this work we presented turbulent 3-D high resolution MHD numerical simulations in order to study the polarized emission of dust particles in molecular clouds. We obtained synthetic dust emission polarization maps calculating the Stokes parameters $Q$, $U$ and $I$ assuming  a perfect grain alignment and that the dust optical properties are the same at all cells. Under these conditions, we were able to study the polarization angle distributions and the polarization degree for the different models and for different inclinations of the magnetic field regarding the LOS.  As main results, we obtained an anti-correlation between the polarization degree and the column density, which is in agreement with observations, with exponent $\\gamma \\sim -0.5$. This value is related to random cancellation of polarization vectors integrated along the LOS, while larger indices require extra physics, e.g. dependency of the dust alignment efficiency with the local density.  We showed that the overall properties of the polarization maps are related to the Alfvenic Mach number and not to the magnetic to gas pressure ratio. Also, we studied the PDF's and structure functions of the polarization angles, which showed a degeneracy of the results with the Alfvenic Mach number and the angle between the magnetic field and the LOS. Zeeman measurements of dense clouds may be useful to help remove this degeneracy as it could provide the magnetic field component along the LOS and, therefore, the inclination.  Finally, we presented a generalization of the CF method, which was showed to be  useful for: (i) the determination of the total magnetic field projected in the plane of sky, and (ii) the separation of the two components $B_{\\rm sky}$ and $\\delta B$;"
    },
    "0812/0812.3913_arXiv.txt": {
        "abstract": "A photon of wavelength $\\lambda\\sim 1\\um$ interacting with a dust grain of radius $a_p\\sim 1{\\rm mm}$ (a ``pebble'') undergoes scattering in the forward direction, largely within a small characteristic diffraction angle $\\theta_s\\sim \\lambda/a_p\\sim 100''$.  Though mm-size dust grains contribute negligibly to the interstellar medium's visual extinction, the signal they produce in scattered light may be detectable, especially for variable sources. Observations of light scattered at small angles allows for the direct measurement of the large grain population; variable sources can also yield tomographic information of the interstellar medium's mass distribution. The ability to detect brilliant pebble halos require a careful understanding of the instrument PSF. ",
        "introduction": "The size distribution of dust in the diffuse interstellar medium has traditionally been studied through the wavelength dependence of interstellar extinction and polarization, and observations of scattered light from reflection nebulae \\citep[cf.][]{Draine_2003a}. These studies have resulted in models \\citep[][]{Weingartner+Draine_2001a, Zubko+Dwek+Arendt_2004, Draine+Fraisse_2009}. where the bulk of the dust mass resides in grains with radii $a\\lesssim 0.5\\um$, with approximately 50\\% of the grain mass above and below a characteristic radius $\\sim0.15\\um$. To account for the observed extinction, these models tend to consume the bulk of available elements such as Mg, Si, and Fe in the ``observed'' grain population with $a\\ltsim0.5\\um$. Little is known regarding the dust grain size distribution above a grain size of order $\\sim 0.5\\um$, aside from the general expectation that grains above this size should contain at most a small fraction of the total dust mass.  Therefore, it was surprising when impact detectors on {\\it Ulysses} and {\\it Galileo} measured a flux of dust particles entering the heliosphere \\citep{Landgraf+Baggaley+Grun+etal_2000, Kruger+Landgraf+Altobelli+Grun_2007} that appeared to indicate that the interstellar grain size distribution extended to much larger grains, with approximately equal mass per unit logarithmic interval out to the largest sizes ($a\\approx 1\\um$) that could be detected. This finding was completely at odds with the conclusions drawn from studies of interstellar extinction. \\citet{Draine_2009} argued that the size distribution inferred for the dust particles approaching the heliosphere could not characterize average interstellar dust, but the situation remains unclear. It is further confounded by radar detection of $a\\approx 30\\um$ particles entering the Earth's atmosphere on solar-hyperbolic trajectories \\citep{Taylor+Baggaley+Steel_1996, Baggaley_2000, Baggaley_2004} implying that they are arriving from interstellar space. The mass flux in these particles exceeds the mass flux in $0.2 < a<1\\um$ particles inferred from the {\\it Ulysses} and {\\it Galileo} observations. If these particles are truly entering the solar system from interstellar space, it implies a total mass density in $a\\gtsim 0.5\\micron$ grains that, at least locally, exceeds that in the $a\\ltsim 0.5\\um$ grains in conventional models for interstellar grains. Such an abundance of very large grains would be difficult to understand given the limitations of interstellar abundances of elements from which such grains would be constituted.   In this {\\it Letter}, we show that very large interstellar grains, or ``pebbles,'' are detectable in scattered light, a phenomenon we refer to as ``brilliant pebbles.''  The basic idea is presented in \\S\\ref{sec:basic}.  Then, in \\S\\ref{sec:applications} we assess the observability of these scattered light halos, and show that detection is possible even if only a few percent of the interstellar grain mass is in mm-sized particles. Variable sources (e.g., gamma-ray bursts ) are particularly useful. In \\S\\ref{s: size_dist} we discuss scattering by a distribution of pebble sizes. We summarize our results in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We propose a method for the detection of extremely large ($\\sim 0.01-1$ cm) dust grains. When the ratio of photon wavelength to grain size $\\lambda/a\\ll 1$, scattering of light is highly peaked in the forward direction, with a characteristic scattering angle $\\theta_0\\sim\\lambda/\\pi\\,a\\sim 30(1\\,{\\rm mm}/\\apeb)\\arcsec$. For a variable source, the halo will lag with a characteristic time-scale $\\sim 10$ min for scattering by 1~mm pebbles at a distance $\\sim 1$ kpc.  We refer to this phenomena as ``brilliant pebbles.''  Detection of the brilliant pebble halo will be technically challenging.  Even for favorable assumptions regarding the pebble size distribution, an accurate determination of the telescope PSF is required in order to detect such a low surface brightness signal. Nevertheless, we show that a pebble halo could be detected using existing facilities if the mass in $\\sim$1~mm pebbles is more than a few percent of the total dust mass, as has been suggested by observations of high-velocity dust grains in interplanetary space and hyperbolic micrometeors (see \\S~1).  Detection is best accomplished using a well-baffled telescope with a freshly aluminized mirror, to minimize scattering by small-scale imperfections such as dust on the mirror.  A brief but bright optical burst would be ideal, as the direct light can fade while the time-delayed pebble halo persists, allowing the telescope PSF and the pebble halo to be separated in the time domain. Also, the pebble halo from a bright short period eclipsing eclipsing binary that is significantly reddened can be subtracted from the PSF in the time domain as well.  If a ``brilliant pebble'' halo is detected, the angular variation of the halo intensity will indicate the pebble size range and the form of the particle size distribution.  The halo will be very blue in the core, and moderately blue in the intermediate halo; the color of the intermediate halo provides an independent measure of the size distribution. If the source is suitably variable, the distance to the pebbles can be determined from the time delay of the halo relative to the point source.  Upper limits on brilliant pebble halos can provide valuable constraints on the size distribution of solid particles in the interstellar medium."
    },
    "0812/0812.3122_arXiv.txt": {
        "abstract": "\\noindent In the standard model of cosmology, dark matter and dark energy are presently the two main contributors to the total energy in the Universe. However, these two dark components are still of unknown nature, and many alternative explanations are possible. We consider here the so-called unifying dark fluid models, which replace dark energy and dark matter by a unique dark fluid with specific properties. We will analyze in this context recent observational data from supernov\\ae~of type Ia, large scale structures and cosmic microwave background, as well as theoretical results of big-bang nucleosynthesis, in order to derive constraints on the dark fluid parameters. We will also consider constraints from local scales, and conclude with a brief study of a scalar field dark fluid model. ",
        "introduction": "\\noindent In the standard model of cosmology, the total energy density of the Universe is dominated today by the densities of two components: the first one, called ``dark matter'', is generally modeled as a system of collisionless particles (``cold dark matter'') and consequently has an attractive gravitational effect like usual matter. The second one, generally refereed as ``dark energy'' or ``cosmological constant'' can be considered as a vacuum energy with a negative pressure, which seems constant today. The real nature of these two components remains unknown, and the proposed solutions to these problems are presently faced to numbers of difficulties.\\\\ The general approach to try to identify these two distinct components is to find constraints to better understand their behaviors. The usual method to find the best values of the different parameters of the model is to predict observations and to adjust the parameters to improve the accuracy of the predictions. In this paper, I consider the so-called dark fluid model \\cite{arbey3}, in which the dark matter and the dark energy are in fact different aspects of a same dark fluid. For simplicity reasons, it is considered in the following that this fluid is perfect, {\\it i.e.} the entropy variations and the shear stress can be ignored. I will examine data from the latest observations of supernov\\ae~of type Ia, cosmic microwave background (CMB), large scale structures, and the theoretical predictions from big-bang nucleosynthesis (BBN), and I will derive constraints on the dark fluid parameters. I will then describe the necessary behavior of a dark fluid at local scales. In a last paragraph, I will discuss the scalar field dark fluid model first presented in \\cite{arbey3}. ",
        "conclusions": "\\noindent Astrophysical and cosmological observations are usually considered in terms of dark matter and dark energy. As we showed here, they can also be interpreted differently, in terms of unifying dark fluids, which could advantageously replace models containing two dark components. In this article, we have derived general constraints on such models in light of recent cosmological data, and we have considered the specific example of scalar field dark fluids and showed that they can be in agreement with the obtained constraints."
    },
    "0812/0812.1850_arXiv.txt": {
        "abstract": "We present a new prediction of GeV $\\gamma$-ray emission from radio lobes of young AGN jets. In the previous work of Kino et al. (2007), MeV $\\gamma$-ray bremsstrahlung emission was predicted from young cocoons/radio-lobes in the regime of no coolings. In this study, we include cooling effects of bremsstrahlung emission and adiabatic loss. With the initial conditions determined by observed young radio lobes, we solve a set of equations describing the expanding lobe evolution. Then we find that the lobes initially have electron temperature of $\\sim$GeV, and they cool down to $\\sim$MeV by the adiabatic loss. Correspondingly, the lobes initially yield  bright bremsstrahlung luminosity in $\\sim$GeV range and they fade out. We estimate these $\\gamma$-ray emissions and show that  nearby young radio lobes could be detected with Fermi Gamma-ray Space Telescope. ",
        "introduction": "Jets in active galactic nuclei (AGNs) are one of the most powerful objects in the universe. AGNs with jets are known as radio-luminous sources and they have been called as radio-loud AGNs. There is increasing evidence that radio-loud AGNs significantly interact with ambient medium and the cosmological importance of radio-loud AGN feedbacks have been advocated (e.g., Best et al. 2007). Among a variety of radio-loud AGNs, a population with its linear size $LS<1~{\\rm kpc}$ has been called as compact symmetric objects (CSOs). The previous studies support a scenario in which CSOs propagate from $\\sim 100~{\\rm pc}$  scales thrusting away an ambient medium and they grow up to Fanaroff-Riley type II radio galaxies (FR IIs) (e.g., Fanti et al. 1995; Readhead et al. 1996a, b; O'Dea \\& Baum 1997; Stanghellini et al. 1998; Snellen et al. 2000; Dallacasa et al. 2000; Orienti et al. 2007). CSOs are thus recognized as newly born AGN jets, and they are crucial sources to explore physics of radio-loud AGNs in their early days.           On the contrary to various radio observations, high-energy emission observations of CSOs had been sporadic (Siemiginowska 2008 for review). Recently, XMM-Newton observations begin to detect X-ray emissions from some of the CSOs although these are not very bright in X-ray (Guainazzi et al. 2006; Vink et al. 2006; O'Dea et al. 2006). Although a few authors recently indicate possible $\\gamma$-ray emissions from CSOs (e.g., Stawarz et al. 2008), there is little observations of CSOs viewed in $\\gamma$-rays. Hence $\\gamma$-ray observations of CSOs may provide us new independent knowledges on the youth era of AGN jets.    Recently, Kino et al. (2007) (hereafter KKI07) showed a possibility of thermal MeV $\\gamma$-ray bremsstrahlung emission from the young radio-loud AGNs. In KKI07, as a first step, they focused on the simplest case in which cooling effects are not significant. These sources roughly correspond to medium-size symmetric objects (MSOs) with $1~{\\rm kpc}<LS<10~{\\rm kpc}$ and FR II galaxies with $LS > 10~{\\rm kpc}$. Since various cooling timescales in smaller sources tend to be shorter than the ones in larger sources, cooling effects in CSOs are expected to be more significant than the ones in MSOs and FR IIs. However little is known about the cooling processes in compact radio sources such as CSOs. In this study we will consistently solve a set of equations which describe young lobe expansions including the effects of bremsstrahlung emission and adiabatic loss together with initial conditions determined by observations.     The goals of this paper are (i) to solve a hydrodynamical evolution of expanding lobes with the cooling effects and (ii) to show a  new  prediction of $\\gamma$-ray emission from the lobes based on the results  of (i). In \\S 2, we show a set of basic equations which describes an expanding radio lobe inflated by its internal energy. In \\S 3, we show the resultant temperature evolution of radio lobes. New predictions of bright GeV $\\gamma$-ray emissions from young radio lobes is presented in \\S 4. Summary and discussion are given in \\S 5. ",
        "conclusions": "We have investigated the temperature and luminosity evolutions of radio lobes of CSOs which expands by their own thermal pressure. In the previous work of KKI07, MeV $\\gamma$-ray bremsstrahlung emission was predicted from young lobes in the regime of no coolings. In this work, we include cooling effects of bremsstrahlung emission and adiabatic loss. Below we summarize  the main results of the present work.    \\begin{enumerate}   \\item   We examine the evolution of $kT_{\\pm}$ together with the initial conditions determined by observed CSOs. By solving a set of equations describing the expanding lobe evolution, it is found that the lobes initially have electron temperature of $\\sim$GeV for $\\Gamma_{\\rm j}\\sim 10$, the lobes then cool down to MeV by the adiabatic loss. During the early phase,  $kT_{\\pm}$ is governed by the adiabatic loss alone. Since the adiabatic loss is more effective than bremsstrahlung cooling in any  case. In the late phase, $kT_{\\pm}$ asymptotically approaches to  the constant temperature of $\\Gamma_{\\rm j}m_{e}c^{2}/3= 1~(\\Gamma_{\\rm j}/10)~{\\rm MeV}$ which has been predicted in KKI07.     \\item  Thermal bremsstrahlung emission peaked about GeV-$\\gamma$ band is newly predicted in CSOs because $n_{\\pm}$ and $kT_{\\pm}$ in younger radio lobes are larger than those in older ones. These spectra can be brighter than the sensitivity of Fermi/LAT for nearby CSOs in the case of High-$P_{0}$ lobe. This means that young radio-lobes can be a new population of GeV-$\\gamma$ emitter in the universe. From Eq. (\\ref{eq:l_brem}) showing $L_{\\rm brem,0}$, we see that the radio lobes with larger $v_{0}$ (i.e., faster expansions) and/or the one with larger $n_{-}$ (i.e., larger $L_{\\rm j}$) yield brighter emission for given  $R_{0}$. As for a specific source, we stress the importance of young recurrent lobe 3C 84 (e.g., Asada et al. 2006). Since 3C 84 is located at the center of very nearby Perceus cluster ($z=0.018$), a deep observation of the Fermi/LAT collabolated with radio obserbations on it will give us tight constraints on the physics of young radio lobes. Near future mission of the VLBI Space Observatory Programme 2 (VSOP-2) with unprecedented high angular resolution (Tsuboi et al. 2009) would play important roles for direct constraint on $v_{0}$ and $R_{0}$ of HFPs.       \\item  Stawarz et al. (2008) investigated a variety of non-thermal $\\gamma$-ray emissions from lobes of CSOs. They claim that the predicted non-thermal emissions can be also detected by Fermi/LAT with an assumed certain electron acceleration efficiency. In GeV $\\gamma$ range, the dominant components are inverse-Compton scattered emissions of ultraviolet  and infrared photons. It is worth to note that even though the luminosity of thermal bremsstrahlung and inverse-Compton ones are comparable, the spectrum shape thermal component is quite different from non-thermal spectra. Hence it is straightforward to distinguish whether the emission is thermal or non-thermal one.   \\item  The observations of Fermi/LAT will be tests for some unresolved questions of AGN jets. Suppose the case that we exactly know $R_{0}$, $v_{0}$, and $n_{\\rm ext}$, and a source distance. If we observe the predicted thermal GeV-$\\gamma$ emission from CSOs, then it suppose the scenario in which CSOs have relativistic jets and their lobes thermally expand. If we do not detect it, it is attribute to lower $kT_{\\pm}$ and/or smaller  $n_{\\pm}$. Possible reasons are as follows; (1) the jet is mainly made of $e/p$ plasma with the same $L_{\\rm j}$, (2) the jet consists of  $e^{\\pm}$ plasma on the whole but with smaller $L_{\\rm j}$, (3) the lobe is expanded by energetic non-thermal particle and the actual $n_{\\pm}$ is  smaller, (4) the jet has non-relativistic speed which leads to the lower $kT_{\\pm}$, (5) other cooling processes could make $kT_{\\pm}$  lower. We remain them as our future investigations.           \\end{enumerate}"
    },
    "0812/0812.1582_arXiv.txt": {
        "abstract": "We present the latest velocities for ten multiplanet systems, including a re-analysis of archival Keck and Lick data, resulting in improved velocities that supersede our previously published measurements. We derive updated orbital fits for ten Lick and Keck systems, including two systems (HD 11964, HD 183263) for which we provide confirmation of second planets only tentatively identified elsewhere, and two others (HD 187123 and HD 217107) for which we provide a major revision of the outer planet's orbit. We compile orbital elements from the literature to generate a catalog of the 28 published multiple-planet systems around stars within 200 pc. From this catalog we find several intriguing patterns emerging: \\begin{itemize} \\item Including those systems with long-term radial velocity trends, at least 28\\% of known planetary systems appear to contain multiple planets; \\item  Planets in multiple-planet systems have somewhat smaller eccentricities than single planets; and \\item The distribution of orbital distances of planets in multiplanet systems and single planets are inconsistent: single- planet systems show a pileup at $P \\sim$ 3 days and a jump near 1 AU, while multiplanet systems show a more uniform distribution in log-period. \\end{itemize} In addition, among all planetary systems we find the following: \\begin{itemize} \\item There may be an emerging, positive correlation between stellar mass and giant-planet semimajor axis. \\item Exoplanets more massive than Jupiter have eccentricities broadly distributed across $0 < e < 0.5$, while lower mass exoplanets exhibit a distribution peaked near e = 0. \\end{itemize} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.0856_arXiv.txt": {
        "abstract": "Only a fraction of the theoretically predicted nonradial pulsation modes have so far been observed in $\\delta$ Scuti stars. Nevertheless, the large number of frequencies detected in recent photometric studies of selected $\\delta$ Scuti stars allow us to look for regularities in the frequency spacing of modes. Mode identifications are used to interpret these results.  Statistical analyses of several $\\delta$ Scuti stars (FG Vir, 44 Tau, BL Cam and others) show that the photometrically observed frequencies are not distributed at random, but that the excited nonradial modes cluster around the frequencies of the radial modes over many radial orders.  The observed regularities can be partly explained by modes trapped in the stellar envelope. This mode selection mechanism was proposed by Dziembowski \\& Kr\\'{o}likowska (1990) and shown to be efficient for $\\ell = 1$ modes. New pulsation model calculations confirm the observed regularities.  We present the s-f diagram, which compares the average separation of the radial frequencies ($s$) with the frequency of the lowest-frequency unstable radial mode ($f$). This provides an estimate for the $\\log g$ value of the observed star, if we assume that the centers of the observed frequency clusters correspond to the radial mode frequencies. This assumption is confirmed by examples of well-studied $\\delta$~Scuti variables in which radial modes were definitely identified. ",
        "introduction": "Recent observational campaigns carried out with earth-based telescopes or space missions concentrating on selected stars have revealed a rich spectrum of radial and nonradial modes covering a wide range in frequencies. This range in frequencies varies from star to star and depends on a variety of factors, not all of which are understood. The comparison of the observations with pulsation models concentrates mainly on the range of frequencies in which pulsational instability occurs as well as mode selection and frequency values for specific modes. In particular, the question of which modes are selected by the star is not solved at the present time. The problem is also a question of amplitude size: we cannot distinguish between stability and oscillations with amplitudes below the level of detectability. Consequently, extensive observational campaigns are needed to lower the observational threshold.  \\begin{figure} \\centering \\includegraphics[bb=35 50 750 500,width=85mm,clip]{Fig1.eps} \\caption{Observed and theoretical frequencies of FG Vir. Observed frequencies with amplitudes higher than 0.5 mmag are compared to the theoretical frequency spectrum of unstable modes with $\\ell$ = 0 to 2. The rotational splitting of modes for an equatorial velocity of 62.5 km~s$^{-1}$ was taken into account. The predicted spectrum is much denser than observed.} \\end{figure}  The question arises of whether the mixture of the excited radial and nonradial modes shows frequency peaks which are essentially randomly distributed over the range of unstable frequencies or whether they tend to form groups. An example of the latter is the regular spacing found in high-order nonradial pulsation (the asymptotic case), as detected for the Sun and white dwarfs. The main-sequence and post-main-sequence $\\delta$~Scuti stars in the classical instability strip, on the other hand, are nonradial and radial pulsators oscillating with low-order p (and g) modes. The frequencies of pulsation are not known to be regularly spaced. In Fig.~1, the observed frequencies of the  star FG Vir (Breger et al. 2005) are compared with the theoretical frequencies (model used by Breger \\& Pamyatnykh 2006). In this figure we have used only the observed modes with photometric amplitudes $\\geq$ 0.5 mmag in order to concentrate on $\\ell$ = 0 to 2 modes. The comparison shows that a mode and/or amplitude selection mechanism is active: not all the theoretically predicted frequencies are seen. While it is possible to compute whether a mode is unstable or not by means of linear theory (e.g., Stellingwerf 1979 for radial modes), the question of the selection of which overtones are excited to observable amplitudes still remains to be solved by nonlinear computations (e.g., Stellingwerf 1980). For most of observed modes with small amplitudes, mode identification in terms of the modal quantum numbers ($n$, $\\ell$, $m$) is not yet possible. This makes the discovery of the mode selection mechanisms more difficult.  In this paper we examine the frequency distribution of mostly nonradial modes in a number of well-observed $\\delta$~Scuti stars in order to search for regularities. ",
        "conclusions": "In the observed pulsation spectra of well-studied $\\delta$~Scuti stars, such as 44~Tau, FG~Vir and BL Cam, a regular distribution of frequencies can be found. The detected frequencies tend to cluster in groups around radial modes. The comparison of the observations with theoretical pulsation models reveals that the cluster pattern may partly be explained by trapped modes in the stellar envelope.  Following the assumption that the radial modes correspond to the center of the cluster, we construct the s-f diagram. It relates the two parameters $f$, the frequency of the lowest-frequency radial mode, and $s$, the mean spacing between the radial modes. Only linearly unstable modes are considered. The s-f diagram allows to infer the stellar $\\log g$ value and to determine the order of the radial mode responsible for the lowest-frequency cluster for stars in the $\\delta$~Scuti mass range. The diagram can be applied if clear patterns are seen in the frequency spectrum. We examined the uncertainties of such a determination in detail and found them to be sufficiently small for a reliable $\\log g$ determination of slowly or moderately rotating stars without chemical peculiarities."
    },
    "0812/0812.0251_arXiv.txt": {
        "abstract": "{} {We present new VLT ISAAC spectra  for 30 quasars, which we combine with previous data to yield a sample of 53 intermediate redshift ($z \\approx$  0.9 -- 3.0) sources. The sample is used to explore properties of prominent lines in the \\hb\\ spectral region of these very luminous quasars.} {We compare this data with two large low redshift ($z<$0.8) samples in a search for trends over almost 6dex in source luminosity.} {We find two major trends: (1) a systematic increase of minimum FWHM \\hb\\ with luminosity (discussed in a previous paper). This lower FWHM envelope is best fit assuming that the narrowest sources radiate near the Eddington limit, show line emission from a virialized cloud distribution, and obey a well defined broad line region size vs. luminosity relation. (2) A systematic decrease of equivalent width \\oiiiopt\\ (from W$\\approx$15 to $\\sim$1\\AA) with increasing source bolometric luminosity (from $\\log L_\\mathrm{bol} \\approx $ 43 to  $\\log L_\\mathrm{bol} \\approx$ 49). Further identified trends require discrimination between so-called Population A and B sources. We generate median composite spectra in six luminosity bins to maximize S/N. Pop. A sources show reasonably symmetric Lorentzian \\hb\\ profiles at all luminosities while Pop. B sources require two component fits involving  an unshifted broad and a redshifted very broad  component. Very broad \\hb\\ increases in strength with increasing $\\log L_\\mathrm{bol}$\\  while the broad component remains constant resulting in an apparent ``Baldwin effect'' with equivalent width decreasing from $W\\sim$80 to $\\sim$20 \\AA\\ over our sample luminosity range. The roughly constant equivalent width  shown by the \\hb\\ very broad component implies production in optically-thick, photoionized  gas. The onset of the redshifted very broad component appears to be a critical change that occurs near the Pop. A-B boundary at FWHM \\hb\\ $\\approx$ 4000 \\kms which we relate to a critical Eddington ratio ($\\approx 0.2 \\pm$0.1). % }{} ",
        "introduction": "Much remains to be defined in quasar research even if one restricts attention to the broad emission lines often used to define them. Their broad line spectra show considerable diversity complicating attempts to generate composite spectra and making estimates of intrinsic properties such as black hole mass unreliable. In addition, the  relationship between broad line emitting active galactic nuclei (Type 1 AGN) and  various classes of sources that do not show broad lines (e.g. Type 2 AGN, LINERs, Blazars, NLRGs) is still uncertain. It is difficult to imagine how advances in physical understanding can come until source phenomenology is clarified. We have focused on clarifying the phenomenology of Type 1 sources because of their relatively unambiguous broad line signature. Leaving aside reverberation studies making use of source variability, observational advances can be expected to come from two areas of spectroscopic investigation: (1) inter-line comparisons (from \\lya\\ to \\ha) in terms of relative intensity and profile shape \\citep[e.g., ][]{shangetal07}; (2) moderate dispersion and high S/N studies of the \\hb\\ spectral region in quasars with $z \\geq 1.0$. There is a long history of infrared spectroscopy of the Balmer lines in high redshift AGN \\citep[e.g., ][]{kuhretal84,espeyetal89,carswelletal91,hilletal93,bakeretal94,elstonetal94,evansetal98,murayamaetal98,murayamaetal99} but only recently has it become possible to obtain spectra with resolution and S/N comparable to those for optical spectra of low $z$ sources. IR spectroscopy of \\hb\\ at $z \\ga $1.0 enables one to use the same line (\\hb), the same rest frame determination (\\oiiiopt\\ or narrow component of \\hb) and the same reduction procedures for estimating the central black hole mass (\\mbh). Even so observations of the most numerous and most luminous quasars at $z \\approx$ 2 can only be made when the optical lines are redshifted into one of the IR bands of high atmospheric transmission.  Spectrographs on HST  have provided UV coverage of higher ionization broad lines in low-$z$\\ Type 1 sources (Seyfert 1 and quasars) with $\\sim$150 sources  having useful measures of \\civ\\ \\citep[e.g.,][]{brothertonetal94,marzianietal96,baskinlaor05,sulenticetal07}. Almost all of these sources now have matching optical spectroscopic measures of the \\hb\\ region \\citep[e.g., ][]{marzianietal03a,baskinlaor05,shangetal07}. Continuing with a litany of difficulties we note that even some of these UV observations show marginal S/N while  many do not cover the full range from Ly$\\alpha$ to 3500 \\AA. Optical coverage is usually not synoptical. This makes detailed intercomparisons of the strongest high and low ionization broad lines possible for only a few tens of objects at best. The lack of spectra covering the rest frame from Ly$\\alpha$\\ to H$\\alpha$\\ is unfortunate since line intensity and profile ratios provide a wealth of physical constraints. The ability to compare lines from ions of widely different ionization potentials has helped to elucidate two main emitting regions within the broad line region (BLR): (a) one responsible for the production of low-ionization lines (LIL) like \\feii, \\hb\\ as well as \\mgii\\ and (b) one that emits mainly lines (\\civ, \\heii) from ions of high ionization potential \\citep[HIL; e.g., ][]{collinetal88}. Meaningful line profile studies begin with observations of  \\civ and \\hb\\ as the most typical HIL and LIL.  Our attempt at clarifying type 1 AGN phenomenology involves the Eigenvector 1 formalism of \\citet{borosongreen92} and was later expanded into 4 dimensions (4DE1). The latter focused on four parameters including  measures of  \\hb\\ and \\feii\\ \\citep{sulenticetal00a,sulenticetal00b,sulenticetal07}. Changes in the width and relative strength of  \\hb\\ and \\feii\\ lines appear to be primarily related to Eddington ratio convolved with source orientation \\citep{shangetal03,marzianietal03b,yipetal04,boroson05,collinetal06}. 4DE1 studies also introduced the concept of two quasar populations A and B that maximize phenomenological differences and possibly identify sources with higher and lower Eddington ratios respectively. The width of broad \\hb\\ (FWHM or 2$^\\mathrm{nd}$ profile moment $\\sigma$) have been widely used as measures of velocity dispersion in the line emitting gas \\citep[see e.g.][]{petersonetal04,sulenticetal06,vestergaardpeterson06}, allowing what are thought to be the most reliable estimates of black hole mass in low $z$\\ quasars. This approach has motivated us to seek similar measures in sources with the highest possible redshift rather than using other lines as \\hb\\ surrogates.  In this paper we leave aside inter-line comparisons and focus on the \\hb\\ spectral range in low and in intermediate-to-high $z$\\ quasars ($z \\ga$ 1). Near IR spectra of the latter were obtained with the infrared spectrometer ISAAC on ESO VLT Unit 1. The high S/N and resolution of the new observations allow  a meaningful comparison with measures obtained for low $z$\\ sources. Low $z$\\ measures come from two samples: (1) an  atlas of bright sources with z$\\la$ 0.8 \\citep[][hereafter, ATLAS sample]{marzianietal03a}  and (2) a magnitude-limited sample of 321 SDSS quasars \\citep[{\\it m}$_\\mathrm{g} <$17.0, $z <$0.7 from ][hereafter SDSS sample]{zamfiretal08}.    In two previous papers  we focused on a search for luminosity effects involving \\hb, \\feiiopt, and \\oiiiopt\\ \\citep[][hereafter Paper I]{sulenticetal04} as well as on the use of \\hb\\ as a virial black hole mass estimator \\citep[][hereafter Paper II]{sulenticetal06}. This paper presents observations and reductions for 30 additional sources ($z= 1.08- 3.09$; \\S \\ref{obs}).  We briefly discuss  the  \\hb, \\feii\\ and narrow line measures (\\S \\ref{imm}) and then show that the addition of 27 new sources to the VLT sample reenforces the luminosity trends described in Paper I (\\S \\ref{further}). We then consider the distribution of VLT sources in the optical plane of 4DE1. Median 4DE1 and luminosity binned composite spectra are discussed in the context of the 4DE1 Population A-B concept \\cite[][]{sulenticetal00a,sulenticetal08}. In \\S \\ref{disc} we use the binned composite spectra to decompose \\hb\\ and make estimates of black hole mass \\mbh\\ and Eddington ratio \\lledd\\ for Pop. A and B sources. ",
        "conclusions": "We presented VLT-ISAAC spectroscopic observations for 30 intermediate $z$\\ and high-luminosity HE quasars. Combined with previous data we have a sample of 53 objects ($z$ = 0.9 -- 3.0) which we compare with two large low-$z$ samples \\citep{marzianietal03a,zamfiretal08}. We find few correlations/trends between broad emission line properties and source luminosity. We previously proposed an empirical limit at FWHM(\\hb)~$\\approx 4000$~\\kms\\ to distinguish between two populations A-B of quasars with very different spectroscopic properties. Sources with FWHM(\\hb) $<$ 4000 \\kms\\ show Lorentzian \\hb\\ profiles while those above this limit are best fit with double Gaussian models. This low $z$/ low luminosity dichotomy is also found for the high luminosity ISAAC sources. Elementary \\lledd\\ computations suggest that  the dichotomy may be explained if a critical  \\lledd\\ is associated with a BLR structure change, but the issue of the Pop.A/B boundary deserves further scrutiny. Once the population A/B dichotomy is taken into account the phenomenology of the \\hbbc\\ emitting region, as well as inferences about BLR structure are quite similar over a 5 dex luminosity range that includes some of the most luminous known quasars.  This overall scenario is described in more  detail by the following results of the present paper: \\begin{enumerate}  \\item A minimum requirement for quantifying and interpreting BLR properties in Type 1 AGN involves the Pop. A--B dichotomy. Spectroscopic averaging without such a discrimination appears to be equivalent to discussing stellar properties without consideration of the OBAFGKM spectroscopic sequence  \\item Pop. A sources show a minimum FWHM \\hb\\ that increases with source luminosity  from $\\sim$500-1000 \\kms\\ to 3000-3500 \\kms. The best boundary to this lower envelope favors an exponent a $\\la$ 0.67 in the Kaspi relation.  Pop. A sources span a FWHM \\hb\\ range of $\\sim$ 4000 \\kms\\ driven by source orientation and virial motions in an accretion disk. The virial assumption implies a \\mbh\\ range of 4dex ($\\log$ \\mbh = 6--10) for this high accreting population. {  In the ISAAC sample studied in this paper Pop. A sources show $-0.7 \\la \\log $ \\lledd $\\la 0.18$.}  \\item \\hb\\ emission in Pop. B sources involves both unshifted broad (\\hbbc) and redshifted very broad line (\\hbvbc) components. Pop. B shows a ``super'' Baldwin effect with $W$(\\hbbc) decreasing by a factor 4 with increasing $L_\\mathrm{bol}$. At the same time $W$(\\hbvbc) remains almost constant implying that it arises in a photo-ionized optically thick medium. FWHM(\\hbbc) component shows no luminosity dependence after correction for \\hbvbc\\ broadening. This yields an \\mbh\\ range similar to that for Pop. A. {Pop. B is a lower accreting population with $-1.1  \\la  \\log$ \\lledd $\\la -0.2$\\ for the objects of this paper. One should  bear in mind the caveats of \\S \\ref{bhmass} in analyzing the \\lledd\\ limits, and that 80\\%\\ of Pop. B sources of the ISAAC sample show $\\log$ \\lledd$\\la -0.4$. Considering the blurring by errors in \\lledd\\ estimates, and results of previous papers, intrinsic \\lledd\\ ranges could be $\\approx$ 0.01 -- 0.2 (Pop. B) and $\\approx$ 0.2 -- 1 (Pop. A), with a separation at \\lledd $\\approx 0.2 \\pm 0.1$.  Pop. B sources in the SDSS sample seem to remain a population of lower Eddington ratio  even if the extreme correction to FWHM  implied by  removal of the \\hbvbc\\ is applied (\\S \\ref{disc:mbh}). }  \\item Most Type 1 AGN do not show broad line profiles consistent with simple accretion disk models. Perhaps Lorentzian Pop. A profiles can be accommodated with   thin/slim disk emission models (if disk emission is very extended) but Pop. B cannot. The redshifted \\hbvbc\\ in Pop. B sources can not be considered part of a double peak signature since the asymmetry and the centroid shift at $\\frac{1}{4}$\\ intensity (near the profile base is usually too large to be attributed to a gravitational redshift. Arp 102B -- the prototypical ``double-peaked'' Balmer line  emitter -- shows a redshift at line base that is modest, and consistent with the value expected for gravitational redshift at the inner emitting radius of the accretion disk according to the model of \\citet{chenetal89}. This is however not the case for many  Pop. B objects that have been included in so-called double-peaked source compilations   \\citep[i.e. as accretion disk candidates: e.g., ][]{stratevaetal03}.  \\item \\hbvbc\\ is the dominant broad line component at high  luminosity (from $\\sim$ 25\\% to 60\\% of the  \\hb\\ flux over our $L_\\mathrm{bol}$ range) for Pop. B sources. The geometry and kinematics of the VBLR are yet unclear. The \\hbvbc\\ redshift $\\Delta v \\approx$ 1--2000 \\kms\\ is unlikely to be gravitational in nature. The increase in dominance with redshift favors a connection to the hypothesized infall involved with the \\mbh\\ growth. Another possibility involves photon downshifting via some form of scattering \\citep[e.g.][and references therein]{laor06}.  \\end{enumerate}"
    },
    "0812/0812.0121_arXiv.txt": {
        "abstract": "The Greisen-Zatsepin-Kuzmin cutoff (GZK cutoff) predicted at the Ultra High Energy Cosmic Ray (UHECR) spectrum has been observed by the HiRes and Auger experiments. The results put severe constraints on the effect of Lorentz Invariance Violation (LIV) which has been introduced to explain the absence of the GZK cutoff indicated in the AGASA data. Assuming homogeneous source distribution with a single power-law spectrum, we calculate the spectrum observed on the Earth by taking photopion production, $e^+e^-$ pair production, and the adiabatic energy loss into account. The effect of LIV is also taken into account in the calculation. By fitting the HiRes monocular spectra and the Auger combined spectra, we show that the LIV parameter is constrained to $\\xi=-0.8^{+3.2}_{-0.5}\\times10^{-23}$ and $0.0^{+1.0}_{-0.4}\\times10^{-23}$ respectively, which is very consistent with strict Lorentz Invariance up to the highest energy. ",
        "introduction": "A number of extensions of the standard model suggest that Lorentz Invariance (LI) is only a low-energy approximation and it may be deformed at very high energies, e.g., approaching the Planck Scale $\\sim 10^{28}$ eV \\cite{Colladay:1996iz,Sarkar:2002mg,Stecker:2004pg,Li:2007vz}. In a simple form shown by Coleman and Glashow, Lorentz Invariance Violation (LIV) can be expressed as a modified energy-momentum relation $E^2=m^2+p^2+\\xi p^2$, under the assumption that LI is violated perturbatively in the context of conventional quantum field theory \\cite{Coleman:1998ti}. It can also be interpreted in terms of different maximal attainable velocities for different particles. Modified energy-momentum relation with LIV terms proportional to the cube of the momentum, or a higher power, are also considered \\cite{Morgan:2007dj}. Although LI has been confirmed at accelerators up to $2$ TeV for protons\\cite{Kahniashvili:2008va}, 104.5 GeV for electrons and 300 GeV for photons\\cite{Hohensee:2008xz}, it is still possible to see LIV in astrophysical processes with much higher energy, especially in the Ultra High Energy Cosmic Rays (UHECRs) with energies above $10^{18}$ eV \\cite{Stecker:2004xm}.  Since the gyration radius of UHECRs is larger than the height of our Galaxy in the Galactic magnetic field, UHECRs are generally thought to be of extragalactic origin. Propagating through intergalactic space, the UHECRs will interact with the Cosmic Microwave Background (CMB) photons, which results in energy and flux depletion. In particular, the photomeson process will induce a suppression in the spectrum above $(3-6)\\times 10^{19}$ eV and lead to the well-known Greisen-Zatsepin-Kuzmin (GZK) cutoff \\cite{Greisen:1966jv,Zatsepin:1966jv}. The spectrum of UHECRs can be calculated theoretically by assuming the source distribution and injection energy spectrum. The energy loss processes for UHECRs propagating in the intergalactic space include photopion production and $e^+e^-$ pair production when interacting with CMB, as well as the adiabatic energy loss due to the expansion of the Universe \\cite{Scully:2000dr,Berezinsky:1988wi, Berezinsky:2002nc}. Because of the short mean-free path of photopion production, the spectrum of UHECRs above $6\\times 10^{19}$ eV falls sharply and results in the GZK cutoff.  However, measurements of the UHECR spectrum have led to great confusion in the last decade. The result of the Akeno-AGASA experiment clearly shows an extension of the spectrum beyond the GZK cutoff \\cite{Shinozaki:2004nh}. To account for the AGASA data beyond the GZK cutoff, LIV has been introduced \\cite{Coleman:1998ti}. Even a LIV parameter as small as $\\xi \\approx 3 \\times 10^{-23}$ may lead to the removal of GZK cutoff\\cite{Coleman:1998ti,Stecker:2004xm,Bietenholz:2008ni}. However, the measurements by HiRes \\cite{Abbasi:2002ta} and Yakutsk \\cite{Egorova:2004cm} seem to show the existence of the GZK cutoff.  Recently, the situation tends to be clear. Having accumulated data for years, the HiRes Collaboration confirms their previous result and observes the GZK cutoff with a $5\\sigma$ standard deviation \\cite{Abbasi:2007sv}. The Pierre Auger Collaboration gives results consistent with HiRes and rejects a single power-law spectrum above $\\sim 10^{19}$eV at the $6\\sigma$ confidence level \\cite{Yamamoto:2007xj,Abraham:2008ru}. As confirmation of the GZK cutoff, severe constraints can be placed on the effect of LIV.  In this paper, we investigate how the LIV can be constrained according to the latest HiRes and Auger data. We give details of the method to calculate the UHECR spectrum with LI in Sec. II. Then LIV is introduced to our calculation of the spectrum of UHECRs in Sec. III. The modified spectrum with different LIV parameters and the constraints on the LIV parameters are also shown in this section. Section IV gives conclusions and discussions. ",
        "conclusions": "The recent observations of the UHECR spectrum by HiRes \\cite{Abbasi:2007sv} and Auger \\cite{Yamamoto:2007xj} show the existence of GZK suppression. In this work we use these observational data to test the LIV model and set constraint on the LIV parameter. The composition of UHECRs is assumed to be pure proton. We then solve the propagation equation of UHE protons in an expanding universe after incorporating the LIV effect in the energy loss rate of protons. A minimum $\\chi^2$ fit to the observational data above $10^{18}$ eV is adopted to derive the source spectrum of UHE protons and the LIV parameter. We find that the current data can limit the LIV parameter for pions to the level $\\sim 3\\times 10^{-23}$. The standard model with the GZK cutoff is very consistent with the observational data.  Only the LIV on the photopion production process is considered in this work. The pair production occurs at lower energy and the LIV effect might be less significant. To incorporate the LIV effect into the $e^+e^-$ production process, the analysis will be more complicated. In addition, we need to restrict the treatment to the case $\\xi_\\pi \\gg \\xi_N\\approx 0$. This has been shown to be due to the weakness of the presentation of the theory. We are unable to determine whether the perturbation term comes from the initial proton or the final state proton, which have different energy in the LS \\cite{Alfaro:2002ya}.  In this work the composition of UHECRs is assumed to be pure proton. However, the composition of UHECRs is poorly known from the experimental point of view. The observations of HiRes show that the UHECRs are proton dominant \\cite{Abbasi:2004nz}, while the results from the Auger experiment indicate a mediate mass composition \\cite{Unger:2007mc}. The determination of the UHECR composition depends on the interaction model and has a relatively large uncertainty at present. If the UHECRs are heavy nuclei dominant, the main process during the propagation in the CMB photon field is photo-disintegration, which will lead to the change of the composition and energy of primary cosmic rays and, accordingly, form a GZK-like spectrum \\cite{Aloisio:2008pp,Aloisio:2008uc,Hooper:2008pm}.  We can make a rough estimate of the LIV effect in such a case, by assuming the primary composition of UHECRs to be iron and considering the process $^{56}Fe+\\gamma(CMB)\\rightarrow^{55}Mn+p$. The analysis is greatly simplified by assuming $\\xi_{Mn} \\approx \\xi_p \\approx 0$, while the physics is not changed as only the difference of the $\\xi$'s in the initial and final states is relevant for the kinematics\\cite{Coleman:1998ti}. On the one hand, the interaction of photo-disintegration is possible only if $m_{Fe}^2+4\\epsilon E_{Fe}+\\xi_{Fe}E_{Fe}^2\\ge(m_{Mn}+m_p)^2$, resulting in $\\xi_{Fe}\\ge-4\\epsilon^2/[(m_{Mn}+m_p)^2-m_{Fe}^2] \\approx-2\\times10^{-25}[\\epsilon/\\epsilon_0]^2$ for CMB energy $\\epsilon_0=2.35\\times10^{-4}$eV. On the other hand, the spontaneous fragmentation $^{56}Fe\\rightarrow^{56}Mn+p$ should be forbidden for iron with energy lower than $\\sim 3\\times10^{20}$eV since cosmic rays with such energies have been detected. This condition requires $m_{Fe}^2+\\xi_{Fe}E_{Fe}^2<(m_{Mn}+m_p)^2$, which gives $\\xi_{Fe}<1.2\\times10^{-23}$ for $E_{Fe}=3\\times10^{20}$eV. Therefore, we get $-2\\times10^{-25}\\le \\xi_{Fe} \\le 1.2\\times10^{-23}$. It should be noted that the above estimates are quite rough. It will be more complicated to use the UHECR spectrum to study the LIV effects for heavy nuclei because there will be a chain of nuclei species taking effect in the interactions. Further detailed analysis is needed to probe the LIV if the UHECRs are proved to be heavy nuclei in future experiments."
    },
    "0812/0812.2915_arXiv.txt": {
        "abstract": "Given the existence of the $M_{\\rm BH}-\\sigma$ relation, models of self-regulated black hole (BH) growth require both a fuel supply and concomitant growth of the host bulge to deepen the central potential, or else the system will either starve or immediately self-regulate without any sustained activity. This leads to a generic prediction that the brightest quasars must be triggered in major mergers: a large fraction of the galaxy mass must be added/converted to new bulge mass and a galactic supply of gas must lose angular momentum in less than a dynamical time. Low-luminosity active galactic nuclei (AGN), in contrast, require little bulge growth and small gas supplies, and could be triggered in more common non-merger events.  This leads to the expectation of a characteristic transition to merger-induced fueling around the traditional quasar-Seyfert luminosity divide (growth of BH masses above/below $\\sim10^{7}\\,\\msun$). We compile and survey a number of observations in order to test several predictions of such a division, including: (1) A transition to bulge-dominated hosts (which any major merger remnant, regardless of difficult-to-observe tidal features, should be). (2) A transition between ``pseudobulges'' and ``classical'' bulges hosting the remnant BHs: pseudobulges are formed in secular processes and minor mergers, whereas classical bulges are relics of major mergers. (3) An increase in the amplitude of small-scale clustering (increased halo occupation of small group environments) where mergers are more efficient. (4) Different redshift evolution, with gas-rich merger rates rising to redshifts $z>2$ while secular processes are relatively constant in time. (5) An increasing prominence of post-starburst features in more luminous systems. Our compilation of observations in each of these areas provides tentative evidence for the predicted division around the Seyfert-quasar threshold, and we discuss how future observations can improve these constraints and, in combination with the tests here, break degeneracies between different fueling models. ",
        "introduction": "\\label{sec:intro}  The discovery of tight correlations between black hole (BH) mass and properties of the host galaxy spheroid, including spheroid luminosity \\citep{KormendyRichstone95}, stellar mass \\citep{magorrian}, velocity dispersion \\citep{FM00,Gebhardt00}, concentration and profile shape \\citep{graham:concentration,graham:sersic}, have fundamental implications for the growth of black holes and -- given the \\citet{soltan82} argument which implies that most black hole mass was assembled in luminous quasar phases \\citep[for more recent versions of this calculation, see e.g.][]{salucci:bhmf,yutremaine:bhmf,hopkins:bol.qlf,shankar:bol.qlf} -- corresponding quasar activity.  In order to explain not just the normalization but also the small (factor $\\lesssim2$) observed scatter in these correlations \\citep{tremaine:msigma}, models generically require some sort of self-regulated growth \\citep[e.g.][]{silkrees:msigma, burkertsilk:msigma,wyitheloeb:sam, dimatteo:msigma,hopkins:lifetimes.methods, hopkins:bhfp.obs,hopkins:bhfp.theory, murray:momentum.winds,thompson:rad.pressure}. Because the BH represents such a small fraction ($\\sim10^{-3}$) of the galaxy mass, it is difficult (if not impossible) to avoid at least the occasional rare situation in which $\\gtrsim100$ times the implied BH mass, from e.g.\\ the $M_{\\rm BH}-\\sigma$ relation, in cold gas is stripped of angular momentum and should accrete rapidly (given either e.g.\\ Bondi or thin-disk accretion), giving rise to BHs well off the observed correlation and orders of magnitude outside the scatter. This concern has led models to invoke some form of feedback from the BH as a self-regulation mechanism: regardless of the details, once the BH reaches some critical mass, if a small fraction ($\\lesssim5\\%$) of the radiant energy ($\\sim0.5\\%$ of the accretion energy) or momentum can couple to the nearby gas, then simple scalings show that gas will be unbound and not able to accrete. The equilibrium between the coupled energy/momentum and the depth of the central potential therefore sets the maximum black hole mass.  In addition to the nature and small scatter of the observed BH-host correlations, there is increasing evidence for a process of this nature. \\citet{hopkins:bhfp.obs,hopkins:bhfp.theory} pointed out that these scenarios generically imply that the most basic correlation is really between BH mass and central potential depth/spheroid binding energy, and demonstrated in the observations a $\\sim3\\,\\sigma$ preference for such an association (i.e.\\ the only relationship with no residual correlations remaining between BH mass and other host galaxy properties). \\citet{aller:mbh.esph} found a similar result in an independent analysis of the observations modeling the central potential in generality, including the contribution of disks and halos.  Moreover, although it remains unclear which feedback processes are the most important for this self-regulation, numerous such mechanisms are observed to be ubiquitous in quasar populations, including intermediate-scale winds, accretion disk outflows, jets, ionization fronts, Compton heating, radiation pressure, and shocks \\citep[see e.g.][]{weymann:BALs,starkcarlson:m82.nlr,baum:radio.outflow.old, colbert:seyfert.outflows.1,colbert:seyfert.outflows.2,colbert:seyfert.outflows.3, laor:warm.absorber,crenshaw:nlr, elvis:outflow.model,baum:radio.outflows,dekool:large.outflow.1, cecil:jet.superbubbles, levenson:seyfert.starburst.winds, ruiz:nlr.kinematics,kaspi:xr.seyfert.kinematics,walter:m82.outflow, pounds:qso.outflow.1,pounds:qso.outflow.2,biller:agn.hot.gas.outflow, whittlewilson:jet.ism.interaction,Steenbrugge:outflow.mdot, rupke:outflows,veilleux:winds,green:qso.outflow.in.lensing, rice:nlr.kinematics, gabel:large.outflow,krongold:warm.absorber.outflow.rate, McKernan:agn.fb.outflow.rates}.  In addition, although the exact normalization and scatter may evolve with redshift \\citep[potentially related to evolution in host galaxy properties setting the central potential, or evolution in feedback physics; see e.g.][]{croton:msigma.evolution, hopkins:bhfp.theory}, observations increasingly indicate that BH-host correlations are in place even at high redshifts \\citep{shields03:msigma.evolution,walter04:z6.msigma.evolution, merloni:magorrian.evolution,adelbergersteidel:magorrian.evolution,woo06:lowz.msigma.evolution, peng:magorrian.evolution,shankar:implied.msig.from.gal.ages} and in irregular or disturbed systems \\citep{borys:xray.ulirgs, alexander:smg.bh.masses,kim:msig.scatter.obs}, further demanding some form of self-regulation (as opposed to e.g.\\ coincidence or fine-tuning from non self-regulating BH growth mechanisms).  If BHs grow in a self-regulated manner, then their evolution must be thought of as fundamentally determined by the growth of their host spheroid. If the spheroid is not large enough -- if the central potential is no deeper -- then a system that has already regulated its growth will simply enter into a high-feedback mode and immediately self-regulate when new gas is channeled into the central regions, {\\em regardless of how large a gas supply is fed to the BH}. If this were not the case, again the problem of explaining the observed tightness of the BH-host correlations would present itself. {\\em ``Growing the monster'' requires not just a fuel supply, but also concomitant bulge growth}, in order to raise the limit to which the BH can accrete before self-regulating.  Given the frequency of e.g.\\ minor mergers and other processes that can remove angular momentum from at least some of the gas in the central regions of a galaxy \\citep{lin:merger.fraction,woods:tidal.triggering, maller:sph.merger.rates,barton:triggered.sf, woods:minor.mergers,stewart:mw.minor.accretion, fakhouri:halo.merger.rates,younger:minor.mergers,hopkins:disk.survival}, simple cosmological calculations (as well as a survey of the abundance of massive bulges in the Universe) suggest that the {\\em bigger} problem which must be overcome in order to allow BH growth may be the growth of bulges, {\\em not} the removal of angular momentum from gas. In any event, the two cannot be considered separately.  This consideration leads to a new framework in which Seyfert and quasar fueling mechanisms must be evaluated. It has long been pointed out that major, gas-rich galaxy-galaxy mergers represent a viable means to fuel at least some quasars \\citep{stockton:3c273.interaction, heckman84:qso.mergers, sanders88:quasars,sanders88:warm.ulirgs, stockton87:qso.fuzz,stocktonridgway:qso.merger.id,hutchingsneff:qso.host.imaging, sanders96:ulirgs.mergers,bahcall:qso.hosts, canalizostockton01:postsb.qso.mergers,hutchings:redqso.lowz, guyon:qso.hosts.ir,dasyra:pg.qso.dynamics, bennert:qso.hosts}, and indeed it is generally accepted that although all quasars may not be mergers, all gas-rich mergers will induce AGN activity at some level \\citep[see references above and e.g.][]{hopkins:qso.all}. These encounters also have the advantage that, when two spirals merge, the disks are converted into spheroid and the bulge further grows via centrally concentrated star formation in a merger-induced starburst \\citep{kormendysanders92,mihos:cusps, mihos:starbursts.94, mihos:starbursts.96, hibbard.yun:excess.light, robertson:fp,naab:gas,cox:kinematics, hopkins:cusps.mergers, hopkins:cusps.ell,hopkins:cores,hopkins:cusps.fp}. This deepens the central potential considerably -- so not only does such a merger strip gas of angular momentum at all levels and provide a fuel source to the BH, but it also increases the binding energy of the inflowing material, meaning that the BH {\\em must} grow larger by a significant factor in order to self-regulate and ultimately lie on the observed BH-host correlations.  There are, however, many alternative mechanisms by which a nuclear BH could be fueled, and these will generically be more common in gas-rich (especially low-mass) galaxies. These include minor mergers (mass ratios $\\sim1:10$ or so), angular momentum loss from stellar+gas bar systems in disks \\citep[for a review see][]{jogee:review}, the stochastic accretion of cold molecular clouds that happen to be near the BH or on scattered trajectories \\citep{hopkins:seyferts}, and Bondi-Hoyle spherical accretion of hot gas from the diffuse, pressure-supported atmosphere in the central bulge \\citep[see e.g.][and references therein]{allen:jet.bondi.power,best:radio.loudness}. However, these mechanisms are much less efficient (or even completely inefficient) at growing the bulge. Although it is possible to get some growth of the bulge and central potential from e.g.\\ bars in isolated galaxies, most numerical experiments and cosmological simulations \\citep{combes:pseudobulges, naab:minor.mergers,bournaud:minor.mergers, athanassoula:peanuts,debattista:pseudobulges.b,naab:etg.formation,younger:minor.mergers, hopkins:groups.ell,hopkins:disk.survival} as well as observations \\citep{oniell:bar.obs,sheth:bar.frac.evol, jogee:bar.frac.evol,driver:bulge.mfs, marinova:bar.frac.vs.freq,barazza:bar.colors} suggest that it is becomes difficult to get significant growth in these channels above a stellar mass $\\sim10^{10}\\,\\msun$, corresponding to a hosted BH mass $\\sim10^{7}\\,\\msun$, and the more massive bulge population will be increasingly dominated by systems whose growth primarily came in major mergers \\citep[for more detailed discussion, see][]{kormendy.kennicutt:pseudobulge.review, maller:sph.merger.rates,hopkins:groups.ell}. This BH mass of $\\sim 10^{7}\\,\\msun$ corresponds to a maximum bolometric luminosity of $\\sim 0.3\\times10^{12}\\,L_{\\sun}$ ($1.3\\times10^{45}\\,{\\rm erg\\,s^{-1}}$) when radiating at the Eddington limit\\footnote{The Eddington limit giving $L_{\\rm bol} = 3.3\\times10^{4}\\,L_{\\sun}\\,(M_{\\rm BH}/\\msun) = 1.29\\times10^{38}\\,{\\rm erg\\,s^{-1}}\\,(M_{\\rm BH}/\\msun)$.}  Moreover, these distinctions are reinforced by a simple scaling: the minimum bolometric luminosity associated with ``true'' quasars is $\\sim 10^{12}\\,\\lsun$ \\citep{soifer:iras} (this translates to a typical optical magnitude $M_{\\rm B}<-23$; the traditional Seyfert/quasar division introduced by \\citet{schmidt:pg.qsos} as an observational, not physical division). Fueling a Seyfert ($L_{\\rm bol}\\le10^{12}\\,L_{\\sun}$) for a \\citet{salpeter64} time $t_{S}=4.2\\times10^{7}\\,{\\rm yr}$ (the $e$-folding time for exponential growth at the Eddington limit, given a typical radiative efficiency expected of BH accretion $\\epsilon_{r}\\equiv L_{\\rm bol}/(\\dot{M}\\,c^{2}) \\sim 0.1$), similar to the lifetime of a high-accretion rate episode suggested by various observational constraints \\citep[clustering, sizes of radio jets, and the transverse proximity effect; for a review see][]{martini04}, requires $\\lesssim10^{7}\\,M_{\\sun}$ worth of accreted material (unsurprising, given that this luminosity corresponds to a BH of around this mass at its Eddington limit). Allowing for some reasonable efficiency of fueling (say $\\sim10\\%$ of the mass within the BH radius of influence or vicinity of the BH) this is still comparable to the mass in a single or a few giant molecular clouds ($\\sim10^{7}-10^{8}\\,\\msun$). The limits become less stringent as one considers lower luminosities (more typical Seyferts with $L_{\\rm bol}\\sim 10^{11}\\,L_{\\sun}$ requiring only $\\sim10^{6}-10^{7}\\,M_{\\sun}$ worth of gas). There are many processes that could sufficiently disturb the gas supply in the central regions of the galaxy so as to produce such an event -- for example, even in the Milky Way, a sufficient random perturbation (with mass of this magnitude) -- a near pass to the BH of another giant molecular cloud complex or massive star cluster -- could torque such a mass (already near the BH in molecular clouds) into accretion \\citep[see e.g.][and references therein]{genzel:gal.center.review}. Given the relative ease of triggering and ubiquity of such systems, there is no reason to invoke a rare and violent mechanism such as major mergers -- in fact, the natural expectation is that more mundane, far more frequent processes should dominate.  However, fueling a bright quasar (say $L_{\\rm bol}\\sim 10^{14}\\,L_{\\sun}$ for $M_{B}\\approx-28$) for a similar duration requires $\\gtrsim10^{9}-10^{10}\\,M_{\\sun}$ in gas -- i.e.\\ channeling an entire typical galaxy's supply of gas onto the BH itself. Given the efficiency of star formation when dense gas is accumulated in the central regions of a galaxy, almost any reasonable model demands more like $\\sim10^{11}\\,M_{\\sun}$ worth of gas within the central $\\sim50-100$\\,pc (even for objects that are fainter, but still quasars, at $L_{\\rm bol}\\gg 10^{12}\\,L_{\\sun}$). Given that the timescale needed is comparable to or shorter than the dynamical time of the galaxy itself, a massive, galaxy-wide perturbation is needed (essentially the entire supply of galactic gas must be stripped of angular momentum and free-fall to the center in less than a single rotational period of the disk), and the only mechanism expected in standard cosmologies to be capable (let alone expected) to produce such a perturbation is a major merger (any disk instability or minor merger, by the nature of the amplification of that instability or dynamical friction, respectively, could neither channel the required mass in an absolute sense nor do so in less than a few disk orbital periods).  This leads to the idea, motivated by detailed models for accretion in mergers and isolated systems, that there is some characteristic host bulge mass/BH mass and corresponding quasar luminosity (which happens to correspond to the observationally defined Seyfert-quasar division) below which these more ubiquitous mechanisms dominate AGN fueling (being more common and requiring less bulge growth to deepen the central potential in this mass regime).  Above this division, less violent mechanisms are simply inefficient (they may still happen, but they do not sufficiently raise the bulge mass, so BHs quickly self-regulate and do not experience any significant lifetime of high-Eddington ratio growth) and the population requires more extreme mechanisms such as major mergers to build the most massive bulges and (corresponding) BHs.  Here, we outline observational constraints that can be used to test this idea, and provide constraints on where non-major merger related fueling mechanisms do or do not dominate the BH and AGN populations. These various accretion processes have different observational signatures, particularly when we recognize that, in a self-regulated growth scenario, they must be treated not just as BH growth mechanisms but simultaneously grow bulges. In \\S~\\ref{sec:morph}-\\ref{sec:colors} below, we outline a few such tests and illustrate some preliminary constraints from previous works, and in \\S~\\ref{sec:discussion} we summarize the results.  For ease of comparison, we convert all observations to bolometric luminosities given the appropriate bolometric corrections from \\citet{hopkins:bol.qlf} \\citep[see also][]{elvis:atlas,richards:seds}. We adopt a $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, $H_{0}=70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ cosmology and a \\citet{salpeter:imf} stellar initial mass function (IMF), and normalize all observations and models appropriately (note that this generally affects only the exact normalization of quantities here, not the qualitative conclusions). All magnitudes are in the Vega system. ",
        "conclusions": "\\label{sec:discussion}  In self-regulated scenarios for the growth of black holes, growing the bulge (deepening the central potential around the BH) is as much a requirement of sustained bright, high Eddington ratio quasar activity as is actually getting a fuel supply of cold, low angular momentum gas into the vicinity of the BH. Moreover, fueling a Seyfert requires something like the mass of a few giant molecular clouds near the BH, whereas fueling a bright quasar requires channeling a galaxy's worth of cold gas into the central $\\sim10-100$\\,pc. This generically leads to the expectation that non-merger related models for quasar fueling become inefficient around a characteristic BH mass of $\\sim1-3\\times10^{7}\\,\\msun$ (host bulge mass $\\sim10^{10}\\,\\msun$), corresponding to the characteristic bolometric luminosity of these BHs near Eddington, $\\sim10^{12}\\,\\lsun$ -- remarkably near the traditional Seyfert-quasar divide \\citep[$M_{B}\\approx-23$;][]{schmidt:pg.qsos}. Essentially, this is a restatement of what has become the conventional wisdom and is increasingly established by observation in the field of elliptical galaxy/bulge formation: massive (predominantly ``classical'') bulges are formed in major mergers, whereas low-mass bulges (predominantly ``pseudobulges'' in low-mass disks with types Sb and later) have large populations formed by non-merger mechanisms (i.e.\\ bars and instabilities, or bar-like activity triggered in sufficiently minor mergers).  We outline five simple tests to distinguish these modes of fueling and determine whether (and if so where) such a transition occurs. In each case, there is at least tentative evidence for the scenario we outline, although more detailed observations as a function of AGN luminosity and redshift are needed to develop more robust constraints. True quasars tend to have spheroidal hosts, and their relic BHs live in classical bulges; they exhibit characteristically strong small-scale clustering; their redshift evolution from $z\\sim0-2$ is similar to that of the gas-rich merger population; and their colors indicate their nature as post-starburst systems. Seyferts, on the other hand, tend to live in a broader mix of hosts or in disk-dominated galaxies, and their relic BHs exist in pseudo-bulges; they exhibit weaker small-scale clustering, consistent with random galaxies of similar mass; their redshift evolution is weak -- their abundance remains constant or even decreases with redshift, similar to what is observed in the fraction of barred disks and predicted for stochastic (quiescent) fueling mechanisms; their star formation histories are suggestive of blue, star-forming galaxies.  It is important to emphasize that none of these constraints uniquely or individually implies e.g.\\ a merger origin for bright quasars. For example, the morphologies could reflect their simply being re-ignited by e.g.\\ minor mergers, gas cooling, or stellar mass loss in evolved ellipticals; but this is difficult to reconcile with their clustering (why they do not trace random, more evolved ellipticals in such a case), colors (why they appear post-starburst), or redshift evolution (why they would not grow in abundance at late times, as the cooling rates, stellar mass loss rates, and number density of such ellipticals are all higher). Their classical bulge relics could have been transformed from ``initial'' pseudo-bulge hosts at the time of the quasars by dissipationless (dry) mergers\\footnote{In practice, this is probably not possible without new dissipation to effectively re-build the bulge, because conservation of phase space densities in dissipationless mergers means that pseudobulges -- which characteristically have low central densities corresponding to fitted Sersic indices to their mass profiles with $n_{s}\\lesssim2$ -- cannot be dissipationlessly re-merged and produce something with the high central densities characteristic of classical bulges (reflected in Sersic indices $n_{s}\\gtrsim4$).}, but again this fails to explain their clustering and colors, and makes the prediction that there should be very little true classical bulge population at intermediate redshifts. There are many ways to populate galaxies as subhalos that reproduce the clustering signature on small scales \\citep[see the discussion in][]{padmanabhan:qso.smallscale.interp}, but those that correspond to a reasonable fueling hypothesis: e.g.\\ that AGN are random spirals in a certain mass range in groups, or rapidly star-forming objects/LIRGs, fail to reproduce the other observations considered. Redshift evolution and star formation histories are similarly degenerate under various fueling models, but again it is the particular combination of features that is difficult to reproduce.  Together, then, these argue for a scenario with a reasonably rapid (and not strongly redshift-dependent) transition between major merger and non-major merger fueled populations around the traditional Seyfert-quasar divide. If true, this has a number of corresponding implications for the demographics of BHs and AGN: BHs less massive than $\\sim1-3\\times10^{7}\\,M_{\\sun}$ (and their corresponding host bulges $\\lesssim10^{10}\\,M_{\\sun}$, typically Sb and later-type galaxies) would be built up primarily via non-merger mechanisms, whereas more massive black holes and bulges -- the ``classical'' systems generally associated with early-type galaxies (Sa and earlier) -- would be built up primarily via mergers. In an integral sense, then, these merger-built systems dominate the total mass density of black holes (with most of the mass density in $\\sim 10^{8}\\,M_{\\sun}$ black holes near the break in the black hole mass function), constituting $\\sim80\\%$ of the mass in BHs \\citep[see e.g.][]{marconi:bhmf}. Correspondingly (by the Soltan 1982 argument), these systems should dominate the total luminosity density of quasars and AGN (integrated over all redshifts). Given the observed evolution of the quasar luminosity function, a redshift-independent division at the quasar-Seyfert divide implies that the luminosity density is dominated by mergers at high redshifts, with a transition to dominance of non-merger induced fueling below redshifts $z\\sim1$. Despite their relative lack of importance in the integrated AGN luminosity density, because of the sensitivity to lower-redshift AGN, this suggests that a large (perhaps dominant) fraction of the X-ray background is contributed by these non-merger related systems. There are clearly a large number of important implications for such a distinction.  A major caveat to this comparison is that, depending on the selection depth, wavelength, and methodology, the Seyfert population recovered could have very different properties, biasing it to be more similar to or different from the quasar population. This is because, unlike at bright quasar luminosities which effectively require massive BHs near Eddington, there are two ways to achieve, in practice, Seyfert luminosities: a low-mass BH at high Eddington ratio and a high-mass BH at low Eddington ratio. High mass systems at low Eddington ratio could be the long-lived remnants of mergers, slowly decaying in AGN luminosity after some initial, bright quasar phase or re-activated by e.g.\\ accretion of new gas in cooling flows or minor mergers, exhibiting characteristically early-type, classical-bulge dominated hosts, with post-starburst populations that might be young for ellipticals but are older than those seen in spiral galaxies. Small-scale clustering signatures in such populations may remain if they were merger-triggered, but it depends significantly on their long-term evolution. The number density evolution of such systems, being cosmologically long-lived and not directly tied to an instantaneous merger rate (even if they are merger-triggered) will reflect, as we predict above, a relatively constant order unity duty cycle. Being at low Eddington ratio, and therefore dim relative to their hosts in optical bands (and possibly radiatively inefficient), such systems may be more commonly seen in X-ray or narrow-line surveys.  On the other hand, surveys of broad-line or optical/IR dominant Seyferts may be more likely to pick out systems at high Eddington ratio, i.e.\\ low-mass, gas-rich systems that might be triggered by minor mergers, bar or disk instabilities, or stochastic encounters with molecular clouds. These systems will be disky, gas-rich, and pseudobulge-dominated, with no small-scale clustering preference, a constant order unity duty cycle, and more ongoing star formation. These distinctions are discussed in detail in \\citet{hopkins:seyfert.bimodality}; here, we have attempted to develop a robust set of predictions that, together, can distinguish AGN populations connected to mergers from either of these ``quiescent'' modes of fueling Seyferts (or from the most likely Seyfert population, some mix of the two). But in several of the tests proposed here, these different fueling modes at low luminosities will manifest differently, and thus they should be considered means to constrain and discriminate among different low-luminosity populations, as well."
    },
    "0812/0812.4778_arXiv.txt": {
        "abstract": "Within the context of standard cosmology, an accelerating universe requires the presence of a third `dark' component of energy, beyond matter and radiation. The available data, however, are still deemed insufficient to distinguish between an evolving dark energy component and the simplest model of a time-independent cosmological constant. In this paper, we examine the cosmological expansion in terms of observer-dependent coordinates, in addition to the more conventional co-moving coordinates. This procedure explicitly reveals the role played by the radius $R_{\\rm h}$ of our cosmic horizon in the interrogation of the data. (In Rindler's notation, $R_{\\rm h}$ coincides with the `event horizon' in the case of de Sitter, but changes in time for other cosmologies that also contain matter and/or radiation.) With this approach, we show that the interpretation of dark energy as a cosmological constant is clearly disfavored by the observations. Within the framework of standard Friedman-Robertson-Walker cosmology, we derive an equation describing the evolution of $R_{\\rm h}$, and solve it using the WMAP and Type Ia supernova data. In particular, we consider the meaning of the \\emph{observed} equality (or near equality) $R_{\\rm h}(t_0)\\approx ct_0$, where $t_0$ is the age of the Universe. This empirical result is far from trivial, for a cosmological constant would drive $R_{\\rm h}(t)$ towards $ct$ (where $t$ is the cosmic time) only once---and that would have to occur right now. Though we are not here espousing any particular alternative model of dark energy, for comparison we also consider scenarios in which dark energy is given by scaling solutions, which simultaneously eliminate several conundrums in the standard model, including the `coincidence' and `flatness' problems, and account very well for the fact that $R_{\\rm h}(t_0)\\approx ct_0$. ",
        "introduction": "Over the past decade, Type Ia supernovae have been used successfully as standard candles to facilitate the acquisition of several important cosmological parameters. On the basis of this work, it is now widely believed that the Universe's expansion is accelerating.$^{1,2}$ In standard cosmology, built on the assumption of spatial homogeneity and isotropy, such an expansion requires the existence of a third form of energy, beyond the basic admixture of (visible and dark) matter and radiation.  One may see this directly from the (cosmological) Friedman-Robertson-Walker (FRW) differential equations of motion, usually written as \\begin{equation} H^2\\equiv\\left(\\frac{\\dot a}{a}\\right)^2=\\frac{8\\pi G}{3c^2}\\rho-\\frac{kc^2}{a^2}\\;, \\end{equation} \\begin{equation} \\frac{\\ddot a}{a}=-\\frac{4\\pi G}{3c^2}(\\rho+3p)\\;, \\end{equation} \\begin{equation} \\dot\\rho=-3H(\\rho+p)\\;, \\end{equation} in which an overdot denotes a derivative with respect to cosmic time $t$, and $\\rho$ and $p$ represent, respectively, the total energy density and total pressure. In these expressions, $a(t)$ is the expansion factor, and $(r,\\theta,\\phi)$ are the coordinates in the comoving frame. The constant $k$ is $+1$ for a closed universe, $0$ for a flat universe, and $-1$ for an open universe.  Following convention, we write the equation of state as $p=\\omega\\rho$. A quick inspection of Eq.~(2) shows that an accelerated expansion ($\\ddot a>0$) requires $\\omega<-1/3$. Thus, neither radiation ($\\rho_r$, with $\\omega_r=1/3$), nor (visible and dark) matter ($\\rho_m$, with $\\omega_m\\approx 0$) can satisfy this condition, leading to the supposition that a third `dark' component $\\rho_d$ (with $\\omega_d<-1/3$) of the energy density $\\rho$ must be present. In principle, each of these contributions to $\\rho$ evolves according to its own dependence on $a(t)$.  Over the past few years, complementary measurements$^3$ of the cosmic microwave background (CMB) radiation have indicated that the Universe is flat (i.e., $k=0$), so $\\rho$ is at (or very near) the ``critical\" density $\\rho_c\\equiv 3c^2H^2/8\\pi G$. But among the many peculiarities of the standard model is the inference, based on current observations, that $\\rho_d$ must itself be of order $\\rho_c$. Dark energy is often thought to be the manifestation of a cosmological constant, $\\Lambda$, though no reasonable explanation has yet been offered as to why such a fixed, universal density ought to exist at this scale. It is well known that if $\\Lambda$ is associated with the energy of the vacuum in quantum theory, it should have a scale representative of phase transitions in the early Universe---many, many orders of magnitude larger than $\\rho_c$.  Many authors have attempted to circumvent these difficulties by proposing alternative forms of dark energy, including Quintessence,$^{4,5}$ which represents an evolving canonical scalar field with an inflation-inducing potential, a Chameleon field$^{6,7,8}$ in which the scalar field couples to the baryon energy density and varies from solar system to cosmological scales, and modified gravity, arising out of both string motivated, or General Relativity modified actions,$^{9,10,11}$ which introduce large length scale corrections modifying the late time evolution of the Universe. The actual number of suggested remedies is far greater than this small, illustrative sample.  Nonetheless, though many in the cosmology community suspect that some sort of dynamics is responsible for the appearance of dark energy, until now the sensitivity of current observations has been deemed insufficient$^7$ to distinguish between an evolving dark energy component and the simplest model of a time-independent cosmological constant $\\Lambda$. This conclusion, however, appears to be premature, given that constraints on the universe's expansion arising from the observed behavior of our cosmic horizon has not yet been fully folded into the interrogation of the current data. The purpose of this paper is to demonstrate that a closer scrutiny of the available measurements, if proven to be reliable, can in fact already delineate between evolving and constant dark energy theories, and that a simple cosmological constant $\\Lambda$, characterized by a fixed $\\omega_d\\equiv \\omega_\\Lambda=-1$, is disfavored by the observations. ",
        "conclusions": "Our main goal in this study has been to examine what we can learn about the nature of dark energy from a consideration of the cosmic horizon $R_{\\rm h}$ and its evolution with cosmic time. A principal outcome of this work is the realization that the so-called ``coincidence\" problem in the standard model is actually more severe than previously thought. We have found that in a $\\Lambda$CDM universe, $R_{\\rm h}\\rightarrow ct_0$ only once, and according to the observations, this must be happening right now. The unlikelihood of this occurrence is an indication that dark energy almost certainly is not due to a cosmological constant. Other issues that have already been discussed extensively in the literature, such as the fact that the vacuum energy in quantum theory should greatly exceed the required value of $\\Lambda$, only make this argument even more compelling. Of course, this rejection of the cosmological-constant hypothesis then intensifies interest into the question of why we don't see any vacuum energy at all, but this is beyond the scope of the present paper.  \\begin{figure}[h] \\center{ \\rotatebox{-90}{\\resizebox{6cm}{!}{\\includegraphics{fig8.eps}}}\\\\ \\caption{Same as Figure~5, except for a ``scaling solution\" in which $\\rho_d\\propto\\rho_m$ and $\\omega_d=-2/3$. This type of universe provides the best fit to the Type Ia supernova data,$^{28}$ with a reduced $\\chi^2\\approx 1.001$ for $180-1=179$ d.o.f. The best fit corresponds to $\\Delta=25.26$.} } \\end{figure}  Many alternatives to a cosmological constant have been proposed over the past decade but, for the sake of simplicity, we have chosen in this paper to focus our attention on scaling solutions. The existence of such cosmologies has been discussed extensively in the literature, within the context of standard General Relativity, braneworlds (Randall-Sundrum and Gauss-Bonnett), and Cardassian scenarios, among others.$^{31,32,33}$  Our study has shown that scaling solutions fit the Type Ia supernova data at least as well as the basic $\\Lambda$CDM cosmology, but they go farther in simultaneously solving several conundrums with the standard model. As long as the time-averaged value of $\\omega$ is less than $-1/3$, they eliminate both the coincidence and flatness problems, possibly even obviating the need for a period of rapid inflation in the early universe.$^{34,35}$  But most importantly, as far as this study is concerned, scaling solutions account very well for the observed fact that $R_{\\rm h}\\approx ct_0$. If $\\langle\\omega\\rangle=-1/3$ exactly, then $R_{\\rm h}(t)=ct$ for all cosmic time, and therefore the fact that we see this condition in the present Universe is no coincidence at all. On the other hand, if $\\langle\\omega\\rangle<-1/3$, scaling solutions fit the Type Ia supernova data even better, but then we have to accept the conclusion that the Universe is older than the horizon time $t_h\\equiv R_{\\rm h}/c$. According to our calculations, which produce a best fit to the supernova data for $\\langle\\omega\\rangle=-0.47$ (corresponding to a dark energy equation of state with $\\omega_d=-2/3$), the age of the Universe should then be $t_0\\approx 16.9$ billion years. This may be surprising at first, but the fact of the matter is that such an age actually solves other major problems in cosmology, including the (too) early appearance of supermassive black holes, and the glaring deficit of dwarf halos in the local group of galaxies.  When thinking about a dynamical dark energy, it is worth recalling that scalar fields arise frequently in particle physics, including string theory, and any of these may be appropriate candidates for this mysterious new component of $\\rho$. Actually, though we have restricted our discussion to equations of state with $\\omega_d\\ge -1$, it may even turn out that a dark energy with $\\omega_d<-1$ is providing the Universe's acceleration. Such a field is usually referred to as a Phantom or a ghost. The simplest explanation for this form of dark energy would be a scalar field with a negative kinetic energy.$^{36}$ However, Phantom fields are plagued by severe quantum instabilities, since their energy density is unbounded from below and the vacuum may acquire normal, positive energy fields.$^{37}$ We have therefore not included theories with $\\omega_d<-1$ in our analysis here, though a further consideration of their viability may be warranted as the data continue to improve.  On the observational front, the prospects for confirming or rejecting some of the ideas presented in this paper look very promising indeed. An eagerly anticipated mission, SNAP,$^{38}$ will constrain the nature of dark energy in two ways. First, it will observe deeper Type Ia supernovae. Second, it will attempt to use weak gravitational lensing to probe foreground mass structures. If selected, SNAP should be launched by 2020. The Planck CMB satellite, an already funded mission, probably won't have the sensitivity to measure any evolution in $\\omega_d$, but it may be able to tell us whether or not $\\omega_d=-1$.  Finally, we may be on the verge of uncovering a class of sources other than Type Ia supernovae to use for dark-energy exploration. Type Ia supernovae have greatly enhanced our ability to study the Universe's expansion out to a redshift $\\sim 2$. But this new class of sources may possibly extend this range to values as high as 5--10. According to Ref.~39, Gamma Ray Bursts (GRBs) have the potential to detect dark energy with a reasonable significance, particularly if there was an appreciable amount of it at early times, as suggested by scaling solutions. It is still too early to tell if GRBs are good standard candles, but since differences between $\\Lambda$CDM and dynamical dark energy scenarios are more pronounced at early times (see Figs.~5 and 8), GRBs may in the long run turn out to be even more important than Type Ia supernovae in helping us learn about the true nature of this unexpected ``third\" form of energy."
    },
    "0812/0812.2018_arXiv.txt": {
        "abstract": "Concurrent observations of Lyman continuum (LyC) and Lyman $\\alpha$ (\\lya) emission escaping from star-forming systems at low redshift are essential to understanding the physics of reionization at high redshift ($z \\ga$ 6).  Some have suggested reionization is dominated by numerous small galaxies with LyC escape fractions $f_e \\sim$ 10\\%, while others suggest mini-quasars with higher $f_e$ might also play a role.  At  $z >$ 3, direct observation of LyC leakage becomes progressively more improbable due to the increase of intervening Ly limit systems, leaving \\lya\\ as the primary diagnostic available to the James Webb Space Telescope for exploring the epoch of reionization.  If a quantitative relationship between escaping LyC and \\lya\\ emission can be established at low $z$, then the diagnostic power of \\lya\\ as a LyC proxy at high $z$  can be fully realized.  Past efforts to detect $f_e $ near $z \\approx$ 3 have been fruitful but observations at low redshift have been less so.  We discuss the sensitivity requirements for detecting LyC leak in the far- and near-UV as a function of redshift 0.02 $< z \\lesssim$ 3 and $f_e \\ge$ 0.01 as estimated from UV luminosity functions.  UV observations are essential to understanding of the physics of LyC escape and the ultimate goal of identifying the source(s) responsible for reionization. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.0427_arXiv.txt": {
        "abstract": "We present MIPS observations of the cluster A3266. About 100 spectroscopic cluster members have been detected at 24 \\micron. The IR luminosity function (LF) in A3266 is very similar to that in the Coma cluster down to the detection limit $L_{IR}\\sim10^{43} \\ergss$, suggesting a universal form of the bright end IR LF for local rich clusters with $M\\sim10^{15}M_{\\sun}$. The shape of the bright end of the A3266-Coma composite IR LF is not significantly different from that of nearby field galaxies, but the fraction of IR-bright galaxies (SFR $> 0.2M_{\\sun}$ yr$^{-1}$) in both clusters increases with cluster-centric radius. The decrease of the blue galaxy fraction toward the high density cores only accounts for part of the trend; the fraction of red galaxies with moderate SFRs ($0.2~M_{\\sun}$ yr$^{-1} <$ SFR $< 1~M_{\\sun}$ yr$^{-1}$) also decreases with increasing galaxy density. These results suggest that for the IR bright galaxies, nearby rich clusters are distinguished from the field by a lower star-forming galaxy fraction, but not by a change in $L_{IR}^{*}$. The composite IR LF of Coma and A3266 shows strong evolution when compared with the composite IR LF of two $z\\sim0.8$ clusters, MS 1054 and RX J0152, with $L^{*}_{IR}\\propto (1+z)^{3.2^{+0.7}_{-0.7}},\\Phi^{*}_{IR}\\propto (1+z)^{1.7^{+1.0}_{-1.0}}$ . This $L^{*}_{IR}$ evolution is indistinguishable from that in the field, and the $\\Phi^{*}_{IR}$ evolution is stronger, but still consistent with that in the field. The similarity of the evolution of bright-end IR LF in very different cluster and field environments suggests either this evolution is driven by the mechanism that works in both environments, or clusters continually replenish their star-forming galaxies from the field, yielding an evolution in the IR LF that is similar to the field. The mass-normalized integrated star formation rates (SFRs) of clusters within 0.5$R_{200}$ also evolve strongly with redshift, as $(1+z)^{5.3}$. ",
        "introduction": "Galaxy clusters are unique laboratories to study the environmental effects on galaxy evolution. A high density environment could potentially alter the star formation (SF) in galaxies through many different physical mechanisms, e.g., ram pressure stripping, galaxy-galaxy interactions, and strangulation of the gas reservoir of galaxies \\citep[see][and references therein]{Boselli06}. As an extreme case of high density environment, clusters show clear evidence of less star formation compared to the field at both local and higher redshifts \\citep{Kennicutt83,Balogh97,Hashimoto98,Poggianti99}, suggesting environmental mechanisms play a role in transforming their star forming properties \\citep{Christlein05}. However, it is still not clear which mechanism, or mechanisms, operate and where the transformation happens in the first place. Is it an end product of galaxy `preprocessing' in the group environment \\citep{Zabludoff96,Zabludoff98,Lewis02,Gomez03} or is it caused by some mechanism specifically related to the cluster environment \\citep[c.f., ][]{Berrier08}? Quantitative and unbiased star formation measurements in clusters and their comparison to field galaxies are keys to answering these questions.  The luminosity function (LF) provides such a quantitative tool to study galaxy properties. It has been widely used to compare galaxies in clusters with those in the field at many different wavelengths. However, some of the results are controversial. For example, some studies claim that cluster LFs show little variation across a wide range of cluster properties \\citep{Colless89,Rauzy98,DePropris03,Popesso06}, while others find them to depend on cluster richness, Bautz-Morgan type \\citep{Bautz70}, or distance from the cluster center \\citep{Dressler78,Garilli99,Lopez97,Hansen05,Barkhouse07}. Another question is whether cluster LFs differ from those in the field. De Propris et al. (1998), \\citet{Christlein03}, \\citet{Cortese03} and \\citet{Bai06} found cluster LFs to be indistinguishable from field LFs, but others suggest that they have brighter characteristic magnitudes and different faint end slopes \\citep{Valotto97,Goto02,Yagi02,DePropris03}.  One of the major reasons for these apparent contradictions is that these studies are based on different wavelength ranges and therefore probe different properties of the cluster galaxy population. The optical and near-IR luminosities of galaxies measure the old stellar mass, while the blue and UV band luminosities are more sensitive to the young stellar population. Since the LFs of different types of galaxies are different \\citep{Madgwick02}, the LFs based on different wavelengths are biased toward different sub-populations of cluster galaxies. \\citet{DePropris03} found that although the B-band LF of early-type galaxies in clusters is different from that of the field, the B-band LF of late-type galaxies in clusters is very similar to that of the field. A similar result was also found in the R-band for star forming and quiescent galaxies by \\citet{Christlein03}. \\citet{Cortese05} found that the UV LF of star-forming galaxies in a nearby cluster is consistent with the field and that the difference between the cluster and field LFs is due to the quiescent galaxies.  Although UV and B-band luminosities are sensitive to star-forming galaxies, they can be heavily attenuated by dust and therefore bias the LF. In contrast, the total IR luminosity ($L_{IR}$, 8-1000 \\micron) shows a good correlation with the star formation rate (SFR) of galaxies unaffected by extinction \\citep{Kennicutt98}. \\citet{Bai06} studied the IR LF of the Coma cluster and concluded that its shape shows no significant difference from that of the field IR LF. However, this result is drawn from the study of one rich cluster and its validity in other clusters needs to be tested.  The study of the evolution of LFs with redshift can provide crucial constraints on models for the formation and evolution of galaxies. In particular, the evolution of the IR LFs will help us understand the star formation history. By comparing Coma with two high redshift clusters, MS~1054-03 and RX~J0152 \\citep{Marcillac07} , \\citet{Bai07} found a strong evolution of the IR LFs for luminous galaxies in both density and luminosity (\\lstar) from $z\\sim0$ to $\\sim0.8$. The evolution rate is consistent with that found in field IR LFs and suggests that this evolution pattern is probably not caused by cluster environmental mechanisms, but by processes similar to those that regulate the field IR LF evolution. However, because the comparison was based on only one low redshift cluster, Coma, there are possible systematic uncertainties in this evolution trend, especially if the IR LFs of local rich clusters have a large cluster-to-cluster variation. It has become clear that the Coma Cluster, once thought to be the archetype of a well-relaxed cluster, has many small substructures \\citep{Mellier88,White93, Neumann03}. Although we did not find any significant difference between the IR LF of the infalling NGC 4839 group and that elsewhere in the Coma cluster, it is still not clear how these substructures may affect the overall IR LF. To study how different dynamical states may affect the IR LF and to quantify cluster-to-cluster variations, we need to study more low redshift clusters.  To test the universality of the IR LF in local clusters and provide a reliable baseline for both the cluster-field comparison and for evolution in clusters, we present the IR LF of A3266 in this paper. A3266 is a nearby rich cluster ($z=0.06$) with a mass similar to Coma ($\\sim10^{15}M_{\\sun}$). It has a large sample of spectroscopically confirmed cluster members (more than 300), making it possible to obtain a IR LF complete down to $L_{IR}\\sim10^{43}$ \\ergss. It is very bright in the X-ray ($L_{X}\\approx10^{45}$ \\ergss). Its X-ray morphology and temperature map, as well as a comparison of the core and overall velocity dispersions, all suggest a recent major merger \\citep{Quintana96,Henriksen00,Sauvageot05,Finoguenov06}. In \\S 2, we present the data analysis for this cluster. In \\S 3, we calculate the IR luminosities and SFRs for the cluster members. In \\S 4, we determine the IR LF for A3266, compare it with Coma's, and look at the dependence of the IR galaxy properties with cluster-centric radius. In \\S 5, we discuss the implications of the results and study the evolution of the integrated SFRs. In \\S 6, we summarize the main conclusions. Throughout this paper, we assume a $\\Lambda$CDM cosmology with parameter set $(h,\\Omega_{0},\\Lambda_{0}) = (0.7,0.3,0.7)$. ",
        "conclusions": "Despite its complicated dynamical state, the IR LF in A3266 is very similar to that of the Coma Cluster, suggesting a universal form of the bright end IR LF ($L_{IR}>10^{43}$ \\ergss) in nearby rich clusters. The shape of this cluster IR LF, at least at the bright end, is not significantly different from the field IR LF at the same redshift. The fraction of the IR-bright galaxies with $L_{IR}>10^{42.7}$ \\ergss\\ (SFR$> 0.2~M_{\\sun}~yr^{-1}$) down to $M_{R}\\le -20.15$ is $40^{+3}_{-4}\\%$ in A3266 and $31^{+4}_{-5}\\%$ in Coma. There are few IR-bright galaxies within a projected radius of $r<0.3$ Mpc, and even they may not actually lie within the cores. The fraction of IR-bright galaxies in these two clusters increases linearly with radius, as $(0.24\\pm0.02) + (0.36\\pm0.06)R/R_{200}$. The decrease of the blue galaxy fraction toward higher density regions only accounts for part of the trend, specifically, for the IR-bright galaxies with SFR$>1M_{\\sun}~yr^{-1}$; the fraction of red galaxies with moderate SFR (0.2$~M_{\\sun}~yr^{-1}<SFR<1~M_{\\sun}~yr^{-1}$) also decreases with increasing galaxy density. These results show that the cluster environment is characterized by a decrease of IR-bright galaxy fraction toward the cluster center, but not by a change in $L_{IR}^{*}$. If some cluster environmental mechanism is responsible for the difference, this mechanism must be confined to the central regions (e.g., $< 0.5 R_{200}$) of the cluster and probably acts relatively quickly. These properties would be compatible with a process like ram pressure stripping, but not with a slower-acting one such as starvation.   When comparing the A3266-Coma composite IR LF with that of two rich clusters at $z\\sim0.8$, MS~1054 and RX~J0152, we find a strong evolution characterized by $L^{*}_{IR}\\propto (1+z)^{3.2^{+0.7}_{-0.7}},\\Phi^{*}_{IR}\\propto (1+z)^{1.7^{+1.0}_{-1.0}}$. This $L^{*}_{IR}$ evolution is indistinguishable from that in the field, and the $\\Phi^{*}_{IR}$ evolution is stronger, but still consistent with that in the field. The similarity of the evolution of bright-end IR LF in very different cluster and field environments suggests either this evolution is driven by the mechanism that works in both environments, or clusters continually replenish their star-forming galaxies from the field, yielding an evolution in the IR LF that is similar to the field.  The mass-normalized integrated SFR ($\\Sigma SFR/M$) within 0.5$R_{200}$ ($\\sim1.5$ Mpc) of A3266 and Coma are similar and show an evolution of $(1+z)^{5.3\\pm1.2}$ compared with those of the two high-$z$ clusters. The evolution derived from these four rich clusters lies under a more scattered evolution trend traced by a larger cluster sample. The large scatter indicates that in addition to evolution, cluster star formation may correlate with cluster masses, IGM properties, and perhaps other parameters. However, for the four rich clusters studied by MIPS, neither the IR LFs nor $\\Sigma SFR/M$ within $0.5R_{200}$ show a dependence on cluster mass, suggesting the mass dependence, if any, is weak for clusters with $M>\\sim 4\\times10^{14}~M_{\\sun}$."
    },
    "0812/0812.3081_arXiv.txt": {
        "abstract": "In this paper we describe the photometric calibration of data taken with the near-infrared Wide Field Camera (WFCAM) on the United Kingdom Infrared Telescope (UKIRT). The broadband $Z$$Y$$J$$H$$K$ data are directly calibrated from 2MASS point sources which are abundant in every WFCAM pointing. We perform an analysis of spatial systematics in the photometric calibration, both inter- and intra-detector and show that these are present at up to the $\\sim$5\\,per cent level in WFCAM. Although the causes of these systematics are not yet fully understood, a method for their removal is developed and tested. Following application of the correction procedure the photometric calibration of WFCAM is found to be accurate to $\\simeq$1.5\\,per cent for the $J$$H$$K$ bands and 2\\,per cent for the $Z$$Y$ bands, meeting the survey requirements. We investigate the transformations between the 2MASS and WFCAM systems and find that the $Z$ and $Y$ calibration is sensitive to the effects of interstellar reddening for large values of $E(B-V)'$, but that the $J$$H$$K$ filters remain largely unaffected. We measure a small correction to the WFCAM $Y$-band photometry required to place WFCAM on a Vega system, and investigate WFCAM measurements of published standard stars from the list of UKIRT faint standards. Finally we present empirically determined throughput measurements for WFCAM. ",
        "introduction": "This paper describes the photometric calibration of the United Kingdom Infrared Telescope (UKIRT) Wide Field Camera (WFCAM) $Z$$Y$$J$$H$$K$ passbands spanning $0.84-2.37\\,\\mu$m. The WFCAM photometric system, based on synthetic colours, is described in Hewett et al. (2006).  WFCAM is currently mounted on UKIRT about 75\\,per cent of the time, the majority of which is used to perform a set of five public surveys, under the umbrella of the UKIRT Infrared Deep Sky Survey: UKIDSS, (Lawrence et al. 2007, Warren et al. 2007), as well as a number of campaign projects and other smaller programmes.  All WFCAM data are processed and calibrated by the Cambridge Astronomical Survey Unit (CASU) using an automated pipeline (Irwin et al. 2008). Due to the large data rate of WFCAM (150--230 Gigabytes of raw image data are taken per night), a reliable and robust method for photometric calibration is required. The primary photometric calibrators are drawn from stars in the Two Micron all Sky Survey (2MASS, Skrutskie et al. 2006), present in large numbers in every WFCAM exposure (the median number is around 200 per pointing). The 2MASS calibration has been shown to be uniform across the whole sky to better than 2\\,per cent accuracy (Nikolaev et al. 2000).  In this paper we investigate the accuracy and homogeneity of the WFCAM calibration, by analysis of data taken in the first two years of science operations. The survey requirement is to provide a photometric calibration in broad band filters accurate to 2\\,per cent with a goal of 1\\,per cent. To test whether the calibrated photometry passes the requirements, we consider the following tests:  \\begin{itemize} \\item for repeat observations of non-variable stars (of sufficient signal-to-noise), the measured $rms$ (root mean square) should be no worse than 0.02 magnitudes. \\item the calibration should be robust against the effects of reddening, i.e. the calibrated frame zeropoints, measured on photometric nights, should have a standard deviation $\\leq 0.02$ magnitudes as a function of reddening. \\item the photometric calibration should be uniform across the sky, and not depend on the observed stellar population or the number of stars in the image. Specifically, the calibrated frame zeropoints (on photometric nights) should have a 1$\\sigma$ scatter of $\\leq 0.02$ magnitudes even in sparse regions of sky, e.g. at high Galactic latitude. \\end{itemize}  We demonstrate that for the vast majority of observations with WFCAM, these tests are passed, and we quantify the conditions under which they are not.  The paper is organised as follows: Section~\\ref{method} summarises how the WFCAM data are calibrated within the pipeline; in Section~\\ref{repeatability} we examine the repeatability of WFCAM photometry, quantify residual spatial systematics, and investigate the limits of the calibration in non-photometric conditions; in Section~\\ref{2mcomparison} we investigate the relationship between the 2MASS and WFCAM photometric systems and investigate the effects of interstellar reddening; in Section~\\ref{bias} we investigate the extent of bias in the WFCAM calibration arising from the overestimation of flux in faint sources in 2MASS; in Section~\\ref{secoffsets} we compare the WFCAM system to a {\\it Vega} system where A0 stars should have zero colour in all passbands; in Section~\\ref{ukirtfs} we compare the WFCAM photometry with published data for faint standards measured with UKIRT; Finally, in Section~\\ref{throughput} we derive the throughput of WFCAM. At the time of writing, UKIDSS has completed its fourth Data Release (DR4). ",
        "conclusions": "\\begin{enumerate}  \\item The CASU pipeline photometric calibration of WFCAM data, using 2MASS sources within each WFCAM pointing, has achieved a photometric accuracy better than 2\\,per cent for the UKIDSS Large Area Survey released in DR3 (Data Release 3) and later, based on repeat measurements.  \\item By stacking spatially binned WFCAM data on monthly timescales we have shown that photometric calibration residuals are present at the 1\\,per cent level. We have derived and tested a method for removing these residuals, and show that the photometric calibration for the WFCAM $J$$H$$K$ filters can reach an accuracy of 1.5\\,per cent. We attribute these residuals to scattered light in the WFCAM twilight flatfield frames. These corrections have been applied to the UKIDSS DR4.  \\item Even observations taken through thin cloud (where transmission is $>$ 80\\,per cent) can be calibrated to $\\sim$ 2\\,per cent using 2MASS calibrators.  \\item Monte-Carlo sampling the 2MASS stars observed by WFCAM, suggests that as long as $\\sim$ 50 calibrators fall within the field-of-view, then the WFCAM calibration is robust to 1\\,per cent in $Y$$J$$H$$K$ and 2\\,per cent in $Z$. 99\\,per cent of the sky has more than 50 2MASS calibrators in a single WFCAM pointing.  \\item The WFCAM calibration incorporates colour restrictions on secondary standards and modified colour equations to mitigate against the effects of Galactic extinction. Towards regions of significant reddening, the calibration meets the 2\\,per cent requirement for $E(B-V)' \\leq 2.0$ for the $J$$H$$K$ bands, but should be treated with caution above $E(B-V)' = 0.2$ for the $Z$ band and to a lesser extent for the $Y$ band.  \\item We find that the calibration of WFCAM from the shallower 2MASS survey is free from the effects of Eddington Bias, thanks to a careful choice of a signal-to-noise cut (SNR$=10$) for the 2MASS sources.  \\item We identified a small offset between the WFCAM $Y$ photometry and an idealized Vega photometric system in the second UKIDSS Data Release (DR2). A correction of $\\Delta_Y=-0.08$ magnitudes has now been included in the WFCAM calibration and applied from UKIDSS DR3, in the sense that the $Y$ band sources are now 0.08 magnitues fainter than in previous releases.  \\item We find that the WFCAM photometric system is within 0.02 magnitudes of the UKIRT MKO system (Leggett et al. 2006) for the $J$$H$$K$ filters, and is Vega-like to within 2\\,per cent (i.e. stars selected to have A0 colours in SDSS have WFCAM colours $<=0.02$ across all passbands).  \\end{enumerate}"
    },
    "0812/0812.4945_arXiv.txt": {
        "abstract": "We present the general properties of the Active Galactic Nuclei (AGNs) and discuss the origin and structure of jets that are associated to a fraction of these objects. We then we address the problems of particle acceleration at highly relativistic energies and set limits on the luminosity of AGN jets for being origin of UHECRs. \\vspace{1pc} ",
        "introduction": "Most of the galaxies of the Local Universe shine by effect of stellar and interstellar gas emissions, predominantly in the optical band. Typically, the emitted spectra are characterized by stellar absorption lines and emission lines by HII regions. These sources are called {\\it Normal Galaxies}. A small fraction of galaxies, about $1 \\%$, do not follow this behavior. In fact, they show strong and broad emission lines, consistent with velocity dispersion of the emitting gas attaining several thousand of kilometers per second. Most remarkably, the non-thermal emission coming from a central nucleus, with size $\\sim 10^{-2}$ pc, dominates over the thermal one coming from stars and interstellar gas and extends well beyond the optical band, from radio to gamma rays (Fig.~\\ref{fig:1}). These sources are called {\\it Active Galaxies} and host {\\it Active Galactic Nuclei} (AGNs) at their centers. \\begin{table*}[htb] \\caption{} \\label{tab:1} \\newcommand{\\m}{\\hphantom{$-$}} \\newcommand{\\cc}[1]{\\multicolumn{1}{c}{#1}} \\renewcommand{\\tabcolsep}{2pc} % \\renewcommand{\\arraystretch}{1.2} % \\begin{tabular}{@{}ll} \\hline Radio quiet AGNs & Radio loud AGNs  \\\\ \\hline Seyfert I galaxies (Sey 1) (BLR, $\\sigma \\sim 10^4$ km \\ s$^{-1}$) & Radio galaxies\\\\ Seyfert II galaxies (Sey 2) (NLR, $\\sigma \\leq 10^3$ km \\ s$^{-1}$) & Radio quasars\\\\ Radio quiet quasars (QSOs) & BL Lac Objects\\\\ ~  &  Optically Violent Variables (OVVs)\\\\ \\hline \\end{tabular}\\\\[2pt] \\end{table*} \\begin{figure}[htb] \\vspace{9pt} \\includegraphics[width=0.45\\textwidth]{prieto4.eps} \\caption{Spectral Energy Distribution of the Centaurus A core. The continuous line is a synchrotron plus Synchrotron Self-Compton (SSC) model \\cite{Prie07}.} \\label{fig:1} \\end{figure} AGNs are not all the same, on the contrary they can be extremely different in their properties. It resulted convenient to classify AGNs according to their radio power, in fact they can be separated into two distinct classes: Radio Quiet and Radio Loud AGNs. Typically, the luminosity in the GHz band of radio loud AGNs exceeds the radio quiet ones by about three orders of magnitudes. Observationally, Seyfert 1 show broad emission lines in the spectrum (Broad Line Region, BLR, with velocity dispersion $\\sigma \\sim 10^4$ km \\ s$^{-1}$), while for Seyfert 2 the spectral lines are narrower (Narrow Line Region, NLR, $\\sigma \\leq 10^3$ km \\ s$^{-1}$). Table~\\ref{tab:1} summarizes the objects belonging to the two classes.  \\begin{figure}[htb] \\includegraphics[width=0.45\\textwidth]{agnmodel.eps} \\caption{The unified AGN model} \\label{fig:2}       % \\end{figure}  This complex ``zoology\" has been interpreted by \\cite{Urry95} within the so called Unified Model for AGNs. According to this picture (Fig. \\ref{fig:2}), radio quiet AGNs have no jets (or very weak ones) while the radio loud ones have jets; the angle between the line of sight and the plane of the obscuring torus determines the observed properties of AGNs. We will discuss the main properties of radio loud AGNs, in particular of the radio galaxies, since the presence of a jet may be crucial for accelerating cosmic ray at the highest energies.  In Section 2 we review the main properties of radio galaxies and in Section 3 we will discuss the particle acceleration from AGN jets. The main outcomes are summarized in Section 4. ",
        "conclusions": "We have reviewed the main properties of AGNs, their separation into radio-loud and radio-quiet sources and the implications of this classification on the problem of cosmic ray acceleration from these objects. We have then examined, on quite general bases, the conditions for UHECRs production and set observational limits to the radio luminosity of radio sources in order to be able to be sources of UHECRs."
    },
    "0812/0812.1567_arXiv.txt": {
        "abstract": "{} {Within the central parsec of the Galaxy, several tens of young stars orbiting a central supermassive black hole are observed. A subset of these stars forms a coherently rotating disc. Other observations reveal a massive molecular torus which lies at a radius $\\sim1.5\\mathrm{pc}$ from the centre. In this paper we consider the gravitational influence of the molecular torus upon the stars of the stellar disc.} {We derive an analytical formula for the rate of precession of individual stellar orbits and we show that it is highly sensitive upon the orbital semi-major axis and inclination with respect to the plane of the torus as well as on the mass of the torus.} {Assuming that both the stellar disc and the molecular torus are stable on the time-scale $\\gtrsim6 \\mathrm{Myr}$, we constrain the mass of the torus and its inclination with respect to the young stellar disc. We further suggest that all young stars observed in the Galactic Centre may have a common origin in a single coherently rotating structure with an opening angle $\\lesssim 5\\degr$, which was partially destroyed (warped) during its lifetime by the gravitational influence of the molecular torus.}{} ",
        "introduction": "Near infrared observations of the central parsec of the Galaxy that were made over the past decade have brought new views of the environment in the vicinity of a supermassive black hole. They revealed a numerous population of young massive stars which may be distinguished into at least two different groups: within a distance $\\lesssim0.03\\mathrm{pc}$ from the centre there are found more than ten so called S-stars orbiting the supermassive black hole on apparently randomly oriented and highly eccentric ($e\\gtrsim0.8$) orbits. They appear to be standard OB main sequence stars (Ghez et al.~2003, Eisenhauer et al.~2005) which is in contradiction with strong tidal forces that prevent stellar formation at the place. Unless these stars mimic their age, it is likely that they have migrated to the centre from larger distances.  Further away, at $0.03\\mathrm{pc} \\lesssim r \\lesssim 0.5\\mathrm{pc}$, nearly up to one hundred young stars were detected (see Paumard et al.~2006 for one of the most recent reviews). These stars are classified mainly as post main sequence OB supergiants and Wolf-Rayet stars. According to the evolutionary phase, their age is estimated to be $6\\pm2\\mathrm{Myr}$. Levin \\&~Beloborodov~(2003) pointed out that a substantial fraction of these stars form a coherently rotating disc (usually referred to as a `clockwise' stellar disc or CWS). It is a flaring disc with an opening angle $\\approx15\\degr$ with a rather sharp inner edge at $0.03\\pc$ and it extends up to radius of $\\approx0.3\\pc$. The radial column density profile of the CWS decreases approximately as $r^{-2}$, i.e. most of the stars are concentrated at the inner edge. The mean plane of the disc can be determined by two angles: inclination $i^\\prime\\approx127\\degr$ with respect to the plane of the sky and longitude of the ascending node $\\Omega^\\prime\\approx99\\degr$ (measured from the North direction; see Paumard et al.~2006 for a detailed description of the convention). Levin \\& Beloborodov~(2003) suggested that this disc-like pattern is a consequence of a stellar formation in a self-gravitating accretion disc.  Further analyses (Genzel et al.~2003, Paumard et al.~2006) indicate the presence of another coherent stellar system which is usually referred to as the `counter-clockwise' stellar disc (CCWS). This structure is assumed to be formed by fewer ($\\lesssim15$) stars, it is narrower in the radial extent being concentrated around $r\\gtrsim0.15\\pc$ and has a larger opening angle $\\approx20\\degr$. The existence of the CCWS disc is a matter of an ongoing debate (e.g. Lu et~al.~2007), nevertheless, even if it is accepted as an explanation of the origin of another subset of young stars in the Galactic Centre, there would still remain more than twenty stars not belonging to any of the two stellar discs and, therefore, without a satisfactory explanation of their origin.  The gravitational potential in the considered region is dominated by the supermassive black hole of mass $\\mbh\\approx3.5\\times10^6\\msun$ (Ghez et al.~2003). It is surrounded by a roughly spherical cluster of late-type stars. The radial density profile is well fitted with a broken power-law $\\rho(r)\\propto r^{-\\beta}$ with index $\\beta=1.19$ below $r=0.22\\pc$ and $\\beta=1.75$ above the break radius (Sch\\\"odel et al.~2007). Its mass, $M_\\mathrm{c}$ within $1\\pc$ is comparable to the mass of the black hole.  The central region is surrounded by a molecular torus (circum-nuclear disc; CND) which lies at the outer edge of the black hole's sphere of influence ($\\rcnd\\approx 1.5\\pc$). Its mass estimated from the radio observations of ionised molecular gas is $\\mcnd\\approx10^6\\msun$ (Christopher et al.~2005). This massive structure defines a non-spherical component of the gravitational field in the central parsec.  In this paper we investigate the influence of the CND upon the dynamical evolution of the disc-like stellar structures. In the subsequent section we briefly review the dynamics in the perturbed Keplerian potential. In Sec.~\\ref{sec:results} we apply the results on the motion of the stellar discs --- we present constraints on some parameters of the CND and the CWS determined from their gravitational interaction and we also give suggestions on the dynamics of the whole system of young stars over its lifetime. Conclusions and discussion of our results are given in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} The massive molecular torus (CND) surrounding the central parsec of the Galactic Centre causes precession of the orbits of young stars which move at distances $0.03\\mathrm{pc}\\lesssim r \\lesssim 0.3\\mathrm{pc}$ around the supermassive black hole. The rate of the precession depends on the orbital parameters as well as on the orientation and mass of the CND. This rate is comparable to the lifetime of the young stars for a wide range of parameters and, therefore, this process should be taken into consideration in attempts to determine the relation between initial and current values of their orbital parameters. We have shown that $\\mcnd\\gtrsim0.3\\mbh$ would destroy any coherently rotating disc-like stellar structure within $6\\mathrm{Myr}$, provided the inclination of most of the orbits with respect to the CND deviates by more than $5\\degr$ from $90\\degr$. In other words, the stability of the stellar disc within its lifetime poses constraints on its inclination with respect to the CND and on the mass of the CND.  We further suggest that {\\em most if not all\\/} young stars observed in the Galactic Centre may have been formed in a single, initially coherently rotating structure, presumably via fragmentation of a thin self-gravitating gaseous disc. Within this hypothesis, the orientation of the stellar disc was nearly exactly perpendicular with respect to the CND. Its `core', represented by the CWS nowadays, remained nearly untouched by the precession. On the other hand, stars that were formed at the outer parts of the disc and/or slightly off its mean plane, or that were scattered out of it via two-body encounters, have undergone a more rapid precession of their orbits, i.e. they apparently don't belong to the stellar disc any longer. We have shown that within the $1\\sigma$ uncertainty of their current velocities there exist such parameters of the stellar orbits that would have had their angular momenta collinear about $6\\mathrm{Myr}$ ago. Due to the high sensitivity of the precession upon the orbit inclination with respect to the CND and the uncertainty in the observed parameters of the stellar orbits, the procedure described in the previous Section cannot provide robust constraints on the parameters of the model. Therefore, the concept of a single warped disc of young stars in the Galactic Centre may be considered as being viable, but not proven yet. Our hypothesis, however, gives an explicit prediction on a specific pattern of the normal vectors of the stellar orbits which may be determined from future, more accurate observations: all of them are assumed to be found close to the circumference perpendicular to the normal vector of the CND.  Let us emphasise that the gravitational influence of the CND leaves stronger imprints on the dynamics of stars more distant from the centre. Hence, we suggest that these stars deserve further attention from the observational point of view. Improved measurements of their kinematical state may bring a new light into the question of the formation and dynamical evolution of the population of the young stars in the Galactic Centre. Beside a generic demand on better constraints on orbital parameters from the observational side, there is also a room for improvements of the model itself. Its most important (and computationally rather expensive) modification will probably lie in an improved treatment of the evolution of the individual orbits, which would take into account gravity of the stellar disc itself.  As a final remark let us note that the strict constraints on the mutual (perpendicular) orientation of the stellar disc and the CND raises a question about the dynamics of gas from which the young stars were formed. It is likely that the parent gaseous disc had to be nearly perpendicular to the CND, so that it would not be destroyed via differential precession before it gave birth to the numerous stellar population. Such an initial orientation is statistically not very probable, opening the question whether it can be a generic result of dissipative (hydro)dynamics in the resonant external potential."
    },
    "0812/0812.2943_arXiv.txt": {
        "abstract": "We present the earliest ultraviolet spectrum of a gamma-ray burst (GRB) as observed with the $\\it Swift$-UVOT.  The \\GRB\\ spectrum was observed for 50 seconds with the \\uvg\\ starting 251 seconds after the $\\it Swift$-BAT trigger. During this time the GRB was $\\approx 13.4$ mag ($u$-filter) and was still rising to its peak optical brightness. In the \\uvg\\ spectrum we find a damped Ly-$\\alpha$ line, Ly-$\\beta$, and the Lyman continuum break at a redshift $z = 2.05 \\pm 0.01$.  A model fit to the Lyman absorption implies a gas column density of log $N_{\\hbox{HI}} = 22.0 \\pm 0.1$~cm$^{-2}$, which is typical of GRB host galaxies with damped Ly-$\\alpha$ absorbers. This observation of \\GRB\\ demonstrates that for brighter GRBs ($v \\approx 14$ mag) with moderate redshift ($0.5 < z < 3.5$) the UVOT is able to provide redshifts, and probe for damped Ly-$\\alpha$ absorbers within 4-6 minutes from the time of the $\\it Swift$-BAT trigger. ",
        "introduction": "\\label{sec1} Gamma-ray bursts (GRBs) are the most luminous cosmic explosions known, and their afterglows extend throughout the electromagnetic spectrum, from X-ray to radio wavelengths. Spectra of GRB afterglows in the ultraviolet (UV) to optical wavelength range are particularly important for the study of GRBs, because the spectral positions of absorption lines from the surrounding environment provide the redshifts, and hence the distances and luminosities of GRBs. Under exceptionally favourable conditions, prompt, very near-UV GRB spectra have been obtained from the ground (i.e. {\\it VLT}-UVES, which in principle is sensitive to wavelengths longward of 3000~\\AA, observed a GRB in $\\approx 10$ minutes). However, such observations are rare due to atmospheric limitations, time of occurrence, and the source position in the sky. Shorter wavelength UV GRB spectra have so far only been taken with the {\\it HST}/STIS \\citep{Smette}, at least 3 days after the GRB trigger. In this letter we present the first prompt UV spectrum of a GRB. The spectrum was obtained with the \\uvg\\ of the {\\it Swift} Ultra-Violet/Optical Telescope \\citep[UVOT;][]{Roming}.  Since the launch of the {\\it Swift} satellite \\citep{swift} in 2004, the UVOT has been making prompt photometric observations of GRBs in the optical and UV, simultaneously with the X-ray and gamma-ray observations taken by the X-Ray Telescope \\citep[XRT;][]{Burrows} and the Burst Alert Telescope \\citep[BAT;][]{Barthelmy}, respectively. The UVOT photometric system \\citep{Poole} provides photometry in three UV and three optical bands, and in one $white$-filter which is sensitive over the range of 1700--8000~\\AA.  Two grisms are also included in the UVOT filter wheel to enable spectra of GRBs to be obtained. Until recently they have not been used in the GRB observing sequence because early in the mission the data indicated that most GRB afterglows were too faint for grism observations, even at early times \\citep{Roming2006}. However, after four years of {\\it Swift} observations, our GRB sample is large enough that we can meaningfully reassess the frequency of GRB afterglows which are bright enough for grism spectroscopy, and ascertain the optimum timing for their inclusion in the automated observing sequence. The catalog of GRB afterglows observed with the UVOT in the first two-and-a-half years \\citep{Roming2008} shows that UVOT observes $\\approx 90$ each year, of which $\\approx 25$ are detected by UVOT in at least one exposure. In its first four years of operation, {\\it Swift} observed 5 GRBs (GRB~050525A, GRB~050922C, GRB~061007, GRB~080319B and GRB~080810) brighter than 14th magnitude in $v$. It is GRBs like these that should provide good spectra through the UVOT grisms.  The \\uvg\\ was introduced into the automated UVOT GRB response on the 7th Oct, 2008. In this latest sequence, the grism spectrum is 50 seconds long, proceeded by an initial 150 second $white$-filter finding chart and followed by a second finding chart in the $u$ filter. With a  median {\\it Swift} slew time of 86 seconds \\citep{Roming2008}, the grism observation typically begins 225-275 seconds after the burst trigger, and covers the wavelength range from $\\approx$~1700-5800~\\AA. The onboard software autonomously replaces the grism exposure with a $u$-filter exposure for bursts which are faint in the BAT, and therefore unlikely to be viable grism targets.  In the following sections we present the first grism image of a GRB taken with the UVOT \\uvg, the resulting spectrum of \\GRB, and its corresponding analysis.  \\begin{figure} \\includegraphics[width=88.0mm,angle=0]{figure_1.eps} \\caption{The lightcurve of \\GRB\\  shows that the spectrum was taken during the rise. The photometry point derived from the spectrum, indicated with an arrow, has a larger uncertainty than those of the lenticular filters.} \\label{fig2} \\end{figure} ",
        "conclusions": ""
    },
    "0812/0812.1806_arXiv.txt": {
        "abstract": "Geodesic motion determines important features of spacetimes. Null unstable geodesics are closely related to the appearance of compact objects to external observers and have been associated with the characteristic modes of black holes. By computing the Lyapunov exponent, which is the inverse of the instability timescale associated with this geodesic motion, we show that, in the eikonal limit, quasinormal modes of black holes in any dimensions are determined by the parameters of the circular null geodesics. This result is independent of the field equations and only assumes a stationary, spherically symmetric and asymptotically flat line element, but it does not seem to be easily extendable to anti-de Sitter spacetimes.  We further show that (i) in spacetime dimensions greater than four, equatorial circular timelike geodesics in a Myers-Perry black hole background are unstable, and (ii) the instability timescale of equatorial null geodesics in Myers-Perry spacetimes has a local minimum for spacetimes of dimension $d\\geq6$. ",
        "introduction": "Geodesics in black hole spacetimes have been extensively studied, both in four and higher dimensional spacetimes, with and without a cosmological constant. Geodesics may display a rich structure and they convey important information on the background geometry. Among the different kinds of geodesic motion, circular geodesics are especially interesting.  For instance, the binding energy of the last stable circular timelike geodesic in the Kerr geometry is related to the gravitational binding energy that can be radiated to infinity, and it can be used to estimate the spin of astrophysical black holes through observations of accretion disks \\cite{Narayan:2005ie,Zhang:1997dy,Shapiro:1983du}.  It was shown many years ago that null geodesics also play an important role. The optical appearance of a star undergoing gravitational collapse depends crucially on the circular unstable null geodesic, which also explains an exponential fade-out of the collapsing star's luminosity \\cite{podurets,amesthorne}. Null geodesics are also very useful to explain the characteristic modes of a black hole -- the so-called quasinormal modes (QNMs) \\cite{Kokkotas:1999bd,Nollert:1999ji}. These ``free'' modes of vibration can be interpreted in terms of null particles trapped at the unstable circular orbit and slowly leaking out \\cite{pressringdown,goebel,Ferrari:1984zz,mashhoon,Berti:2005eb}.  The real part of the complex QNM frequencies is determined by the angular velocity at the unstable null geodesic; the imaginary part is related to the instability timescale of the orbit, a quantity which is seldom considered in geodesic studies, with some noteworthy exceptions (see e.g.~\\cite{Bombelli:1991eg,Cornish:2003ig,Pretorius:2007jn,PerezGiz:2008yq,Levin:2008yp,Grossman:2008yk,Steklain:2008gz}). Furthermore, there is some evidence \\cite{Pretorius:2007jn} that unstable circular orbits could yield information on phenomena occurring at the threshold of black hole formation in the high-energy scattering of black holes, a process of interest in fundamental physics for a variety of reasons \\cite{Sperhake:2008ga,Shibata:2008rq}.  In this work we clarify some aspects of the relation between unstable null geodesics, Lyapunov exponents and quasinormal modes. In Section \\ref{LyapunovSec} we derive a simple formula for the principal Lyapunov exponent $\\lambda$ in terms of the second derivative of the effective potential for radial motion $V_r$: \\be\\label{LyapunovGeneralIntro} \\lambda=\\sqrt{\\frac{V_r''}{2 \\dot{t}^2}}\\,, \\ee where $t$ is coordinate time. Throughout this work, a dot denotes a derivative with respect to proper time and a prime stands for derivative with respect to areal radius $r$. The result above is valid for a wide class of spacetimes and geodesics, including stationary spherically symmetric spacetimes and equatorial orbits in the geometry of higher-dimensional, rotating (Myers-Perry) black hole solutions.  In Section \\ref{sphericallysym} we show that the relation between QNMs and unstable circular null geodesics is quite general, being valid in the eikonal limit for any static, spherically symmetric, asymptotically flat spacetime. More specifically, we show that the angular velocity $\\Omega_c$ at the unstable null geodesic and the Lyapunov exponent, determining the instability timescale of the orbit (see for instance \\cite{Bombelli:1991eg,Cornish:2003ig}) agree with analytic WKB approximations for QNMs \\cite{Mashhoon:1983,Schutz:1985km,Iyer:1986np}: \\be\\label{QNMLyapunovIntro} \\omega_{\\rm QNM}=\\Omega_c\\,l-i(n+1/2)\\,|\\lambda|\\,, \\ee where $n$ is the overtone number and $l$ is the angular momentum of the perturbation. The WKB results are formally valid only in the eikonal regime ($l\\gg 1$), but they seem to yield surprisingly accurate predictions even for low values of $l$ \\cite{Iyer:1986nq,Berti:2005eb}. A simple derivation of the Lyapunov exponent for spherically symmetric, asymptotically flat spacetimes, patterned after the original QNM calculation by Ferrari and Mashhoon \\cite{Ferrari:1984zz,mashhoon}, is provided in Appendix \\ref{sec:calcmashhoon}. For the important case of a $d$-dimensional Schwarzschild-Tangherlini \\cite{Tangherlini:1963bw} black hole solution we find that the critical exponent defined by Pretorius and Khurana \\cite{Pretorius:2007jn} can be determined analytically to be \\be \\gamma\\equiv \\frac{\\Omega_c}{2\\pi \\lambda} = % \\frac{1}{2\\pi\\sqrt{d-3}}\\,. \\label{tanghIntro} \\ee By exploring the connection between QNMs and null geodesics, we also find a simple analytical result for the quasinormal frequencies of near-extremal Schwarzschild-de Sitter black holes in $d=4$: \\be\\label{qnadsgeoIntro} \\omega_{\\rm QNM}=\\kappa_+\\left[l-i\\,(n+1/2)\\right]\\,, \\ee where $\\kappa_+$ denotes the surface gravity. In the eikonal limit, the above result agrees with that found in \\cite{Cardoso:2003sw}.  In Section \\ref{formulation} we analyze the higher-dimensional rotating black hole solutions found by Myers and Perry \\cite{Myers:1986un}. In $d=5$ we can compute $\\lambda$ analytically. The Lyapunov exponent goes to zero as one approaches extremality in $d=4,5$ spacetime dimensions. However, no such behavior is observed for $d>5$: the Lyapunov exponent (normalized by the orbital frequency) has a local minimum, which may be related to a possible instability of the system first suggested by Emparan and Myers \\cite{Emparan:2003sy}. In Appendix \\ref{sec:timelikemyers} we study in some detail timelike circular geodesics in the Myers-Perry spacetime and show that equatorial circular orbits are always unstable for $d>4$. Finally, in Appendix \\ref{SAdS} we discuss issues in generalizing our results to non-asymptotically flat spacetimes. ",
        "conclusions": "We have shown that for all spherically symmetric spacetimes, in a geometrical optics approximation, QNMs can be interpreted as particles trapped at unstable circular null geodesics and slowly leaking out.  The leaking timescale is given by the principal Lyapunov exponent, for which we obtained a fairly simple expression, Eq.~(\\ref{LyapunovGeneral}), in terms of the second derivative of the effective radial potential for geodesic motion. This simple, intuitive relation between QNMs and circular null geodesics is valid for all asymptotically flat, spherically symmetric black hole spacetimes.  Some aspects of our investigation deserve further analysis. The interpretation of QNMs in terms of unstable circular null geodesics is valid in all generality only for spherically symmetric, asymptotically flat spacetimes. Once we break the azimuthal degeneracy, things get a bit more complex. For instance, in $d=4$ it is known that equatorial geodesics can account for the $l=|m|$ modes of Kerr-Newman black holes \\cite{Ferrari:1984zz,Berti:2005eb,Berti:2005ys}, but it is unclear whether this analogy can be extended to modes with $l\\neq |m|$.  Perhaps modes with $l \\neq m$ can be explained in terms of more general (e.g., non-equatorial) geodesics. Besides improving our intuitive physical understanding of ringdown radiation, a deeper exploration of this analogy could have important implications for the interpretation of numerical simulations of black hole binary mergers and their use in gravitational-wave data analysis \\cite{Berti:2007fi,Hanna:2008um,Baker:2008mj}.  Another limitation of our results concerns their extension to non asymptotically flat (e.g., AdS) backgrounds.  In the large damping limit, a certain class of QNMs has been associated with {\\it radial} geodesics \\cite{Fidkowski:2003nf,Amado:2008hw,Festuccia:2008zx}. It would be very interesting to extend this analysis to the eikonal (large-$l$) limit. A possible starting point could be the $2+1$ dimensional Ba\\~nados-Teitelboim-Zanelli (BTZ) black hole \\cite{Banados:1992wn}, for which QNM frequencies are known analytically \\cite{Cardoso:2001hn}. Quite apart from the geodesic analogy, the large-$l$ limit is interesting {\\it per se}.  It turns out that the imaginary part of Schwarzschild-anti-de Sitter (SAdS) QNMs decreases with $l$ \\cite{Festuccia:2008zx}. Thus, if excited considerably, large-$l$ modes could dominate the black hole's response to perturbations. A more thorough investigation of the large-$l$ limit of QNMs is necessary.  Interesting physical phenomena could occur in (hypothetical) spacetimes for which timelike circular geodesics have a frequency equal to (or larger than) the frequency of unstable null geodesics. This would raise the interesting possibility of exciting QNMs by orbiting particles, possibly leading to instabilities of the spacetime.  It would be interesting to find general conditions under which spacetimes possess stable null geodesics; stable null geodesics may also be associated with instability (or marginal stability) of spacetimes."
    },
    "0812/0812.3032_arXiv.txt": {
        "abstract": "{} {We obtained phase-resolved spectroscopy of the accreting millisecond X-ray pulsar SAX J1808.4-3658 during its outburst in 2008 to find a signature of the donor star, constrain its radial velocity semi-amplitude ($K_2$), and derive estimates on the pulsar mass.  } {Using Doppler images of the Bowen region we find a significant ($\\ge$8$\\sigma$) compact spot at a position where the donor star is expected. If this is a signature of the donor star, we measure $K_{em}$=248$\\pm$20 km s$^{-1}$ (1$\\sigma$ confidence) which represents a strict lower limit to $K_2$. Also, the Doppler map of He\\,II $\\lambda4686$ shows the characteristic signature of the accretion disk, and there is a hint of enhanced emission that may be a result of tidal distortions in the accretion disk that are expected in very low mass ratio interacting binaries.} {The lower-limit on $K_2$ leads to a lower-limit on the mass function of $f$($M_1$)$\\ge$0.10$M_\\odot$. Applying the maximum $K$-correction gives 228$<$$K_2$$<$322 km s$^{-1}$ and a mass ratio of 0.051$<$$q$$<$0.072.} {Despite the limited S/N of the data we were able to detect a signature of the donor star in SAX\\,J1808.4$-$3658, although future observations during a new outburst are still warranted to confirm this. If the derived $K_{em}$ is correct, the largest uncertainty in the determination of the mass of the neutron star in SAX\\,J1808.4-3658 using dynamical studies lies with the poorly known inclination.} ",
        "introduction": "Low-Mass X-ray Binaries (LMXBs) are systems in which a compact object (a neutron star or a black hole) is accreting matter, via Roche lobe overflow, from a low mass ($<$1$M_{\\odot}$) star. Some LMXBs exhibit sporadic outburst activity but for most of their time remain in a state of low-level activity (White et~al. 1984); we will refer to these systems as transients.  In April 1998, a coherent 2.49 ms X-ray pulsation was discovered with the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) satellite in the transient SAX J1808.4$-$3658 (Wijnands \\& van der Klis 1998, Chakrabarty \\& Morgan 1998). This was the first detection of an Accreting Millisecond X-ray Pulsar (AMXP). Seven more of these systems have been discovered since then; all these systems are transients, have orbital periods in the range between 40 min and 4.3 hr and spin frequencies from 1.7 to 5.4 ms.  These findings directly confirmed evolutionary models that link the neutron stars of LMXBs to those of millisecond radio pulsars via the spinning up of the neutron star due to accretion during their LMXB phase (e.g. Wijnands 2006).  To date, six episodes of activity were detected from SAX J1808.4$-$3658 with a 2-3 year recurrence cycle. During the 1998 outburst a detailed analysis of the coherent timing behaviour showed that the neutron star was in a tight binary system with a 2.01 hr orbital period (Chakrabarty \\& Morgan 1998; Hartman et~al. 2008).  The mass function derived from X--ray data ($4\\times 10^{-5} M_{\\odot}$) and the requirement that the companion fills its Roche lobe led to the conclusion that it must be a rather low mass star, possibly a brown dwarf (Chakrabarty \\& Morgan 1998; Bildsten \\& Chakrabarty 2001).  \\begin{figure*}[t]\\begin{center} \\psfig{figure=SAX1808_average.ps,width=5.5cm, angle=-90} \\caption{Average flux calibrated spectrum of SAX\\,J1808.4$-$3658. We have labelled the most important features that are present. The strong absorption features around 3940 and 3980 \\AA~ are due to diffuse interstellar bands. \\label{average}} \\end{center}\\end{figure*}  We present here optical spectroscopy of SAX J1808.4$-$3658 obtained during the 2008 outburst.  One of the key issues in dynamical studies is to measure the radial velocity of the companion star ($K_2$) and use this value to constrain the optical mass function of the system. During quiescence, these goals can be achieved by tracing the absorption features originating in the photosphere of the companion star. However, the intrinsic faintness of the low mass companion stars in AMXPs makes such an analysis very difficult.  To overcome this problem, Steeghs \\& Casares (2002) have shown that during phases of high mass accretion rates, Bowen blend lines emitted by the irradiated face of the companion star can be used. A precise measurement of $K_2$ represents the only way to determine the optical mass function of the system and ultimately constrain the neutron star mass. ",
        "conclusions": "The observations presented here provide evidence for the detection of the signature of the irradiated donor star in SAX\\,J1808.4$-$3658. Clearly a similar experiment, but at higher S/N and spectral resolution, must be carried out again during a future outburst to obtain more than a single orbital period of data and resolve the narrow lines. This will not only allow us to unambiguously claim the presence of narrow components in the spectra of SAX\\,J1808.4$-$3658, but also measure the rotational broadening of the narrow components to further constrain $K_2$ via the relation in Wade \\& Horne (2003).  With these data we have shown a promising way forward to constrain $K_2$ and in combination with more quiescent data, we should be able to better constrain the inclination and thereby obtain the mass of the neutron star in SAX\\,J1808.4$-$3658.   \\begin{figure}\\begin{center} \\psfig{figure=SAX1808_donor.ps,width=5.5cm, angle=-90} \\caption{Blow-up of the Bowen region for the average spectrum (top) and the average in the rest-frame of the donor star (bottom). The bottom spectrum clearly shows the narrow components. Indicated are the most important N\\,III ($\\lambda$4634/4640) and C\\, III ($\\lambda$4647) lines. \\label{donor}} \\end{center}\\end{figure}"
    },
    "0812/0812.1037_arXiv.txt": {
        "abstract": "Linear theory provides a reasonable description of the velocity correlations of biased tracers both perpendicular and parallel to the line of separation, provided one accounts for the fact that the measurement is almost always made using pair-weighted statistics. This introduces an additional term which, for sufficiently biased tracers, may be large.  Previous work suggesting that linear theory was grossly in error for the components parallel to the line of separation ignored this term. ",
        "introduction": "Experiments will soon measure correlations in the peculiar velocity field traced by massive galaxy clusters.  The linear theory description of these correlations as traced by dark matter particles whose motions are driven by gravitational instability from an initially Gaussian random field, has been available for about 20 years (Gorski 1988). The extension of this formalism to biased tracers, within the framework of perturbation theory, is in Fisher (1995); Sheth et al. (2001b) derive similar results within the context of the Halo model (see Cooray \\& Sheth 2002 for a review); and Reg\\\"os \\& Szalay (1996) study this problem in the context of peaks (see Desjacques 2008 for corrections to their expressions). All these analyses yield consistent results.  However, measurements of the velocity correlations of massive haloes in simulations have given the impression that, although biased linear theory provides a reasonable description of the correlations between the velocity components that are perpendicular to the line of separation $\\Psi_\\perp$, it is wildly discrepant for the components which are parallel to the line of separation $\\Psi_{||}$ (Croft \\& Efstathiou 1994; Peel 2006), except on scales larger than about 100 Mpc.  This discrepancy has been attributed to nonlinear effects, such as those described by Colberg et al. (2000) and Sheth \\& Diaferio (2001).  The main purpose of this short note is to show that, in fact, linear theory does indeed provide a good description of $\\Psi_{||}$, provided one accounts for the fact that the measurement is actually pair weighted, and so the mean streaming motions may not be negligible. For dark matter, this mean is sufficiently small that it can be ignored, but ignoring the mean streaming of massive halos towards one another is a very bad approximation. We show that keeping this term in the theoretical calculation provides substantially better agreement with the measurements.  Section~\\ref{theory} provides the linear theory expressions for the velocity statistics of interest, Section~\\ref{measurement} presents a comparison of these expressions with simulations, and a final section summarizes our results.  An Appendix provides some technical details of the calculation. ",
        "conclusions": "The mean streaming motions of massive halos can be as large as 1000~km~s$^{-1}$ on scales of order $10h^{-1}$~Mpc.  In contrast, the one dimensional velocity dispersions of these halos is of order 300~km~s$^{-1}$. Nevertheless, these motions are rather well described by biased linear theory.  The fact that the mean streaming motions of halos can be substantially larger than the halo velocity dispersions has an important consequence:  they must not be ignored when making the linear theory estimate of velocity correlations $\\Psi_{||}$ and $\\Psi_\\perp$ for pair-weighted statistics. Ignoring this contribution leads one to predict that $\\Psi_{||}$ is positive; i.e., objects stream along with one another, whatever their mass. The measurements show that while this is true for the lower mass halos (they are still quite massive!), $\\Psi_{||}$ is negative for the most massive halos.  This means that the most massive halos move towards, rather than along with, one another; this is consistent with linear theory provided one accounts for the fact that the measurements are pair-weighted, so one should not ignore the contribution of the mean streaming motions.  We note that, although they do not say so explicitly, halo model analyses of halo motions (Sheth et al. 2001a) have correctly included the effects of this mean streaming.  It is the appearance of more recent studies which ignore this effect (e.g. we had communicated the importance of this term to Peel 2006, but his figures do not include it) which prompted our work. Because our analysis shows that linear theory can be used down to scales which are significantly smaller than previously thought, we hope that our analysis will aid in the interpretation of data from the next generation of surveys.  For deep surveys, the gain in accuracy is likely to be modest. This is because sky surveys typically estimate the correlation between radially projected velocities, separated by some angle $\\theta$ on the sky.  This quantity is related to ours by \\begin{equation} \\Psi_{12} = \\Psi_\\perp \\cos\\theta + (\\Psi_{||} - \\Psi_\\perp)\\,f(\\theta,r_1,r_2), \\end{equation} where $\\Psi_{||} \\equiv \\sigma_v^2\\,\\psi_{||}$, $\\Psi_{\\perp} \\equiv \\sigma_v^2\\,\\psi_{\\perp}$, both are evaluated at scale $r = \\sqrt{r_1^2 + r_2^2 - 2 r_1r_2\\cos\\theta}$, and \\begin{equation} f(\\theta,r_1,r_2) = \\frac{(r_1^2 + r_2^2)\\cos\\theta - r_1r_2(1 + \\cos^2\\theta)} {r_1^2 + r_2^2 - 2 r_1r_2\\cos\\theta}. \\end{equation} This can also be written as \\begin{equation} \\Psi_{12}(r) = \\Psi_\\perp(r)\\, \\frac{r_1r_2}{r^2} \\,\\sin^2\\theta + \\Psi_{||}(r)\\,\\left(\\cos\\theta - \\frac{r_1r_2}{r^2}\\,\\sin^2\\theta\\right), \\end{equation} showing that our extra term is most important for pairs which are close in angle but lie at different distances along the line of sight. For very deep surveys ($r_1$ and $r_2$ greater than 100 Mpc, say) the observable is dominated by $\\Psi_\\perp$ for all scales where the addition of our extra term is important.  Nevertheless, we emphasize that the term is present, and its inclusion greatly improves comparison with simulations, thus strengthening ones confidence in the validity of linear theory.  Although accounting for mean streaming motions results in substantially better agreement with the simulations, there remains room for improvement.  Our treatment of how $\\langle v_{||}(r)\\rangle$ should depend on halo mass on small scales is rather ad hoc. Methods motivated by perturbation theory are beginning to provide more detailed prescriptions for this term (e.g. Smith et al. 2008). Also, in perturbation theory, the power spectrum of peculiar velocities is expected to be slightly suppressed relative to the linear theory prediction, even on relatively large scales (e.g. Bernardeau et al. 2002; Cooray \\& Sheth 2002; Pueblas \\& Scoccimarro 2008). Although including this effect is beyond the scope of this work, we note that accounting for it will lower the theory curves slightly, reducing the small discrepancy in the panels on the right of Figure~\\ref{xivkms}. Incorporating these refinements into our analysis is the subject of work in progress."
    },
    "0812/0812.4576_arXiv.txt": {
        "abstract": "Earlier studies of grain alignment dealt mostly with interstellar grains that have strong internal relaxation of energy which aligns grain axis of maximum moment of inertia with respect to rain's angular momentum. In this paper, we study the alignment by radiative torques for large irregular grains, e.g., grains in accretion disks, for which internal relaxation is subdominant. We use both numerical calculations and the analytical model of a helical grain introduced by us earlier. We demonstrate that grains in such a regime exhibit more complex dynamics. In particular, if initially the grain axis of maximum moment of inertia makes a small angle with angular momentum, then radiative torques can align the grain axis of maximum moment of inertia with angular momentum, and both axis of maximum moment of inertia and angular momentum are aligned with the magnetic field when attractors with high angular momentum (high-J attractors) are available. For the alignment without high-J attractors, beside the earlier studied attractors with low angular momentum (low-J attractors), there appears new low-J attractors. The former and later cases correspond to the alignment with long axes perpendicular and parallel to the angular momentum, respectively. In addition, we also study the alignment of grains in the presence of strong internal relaxation, but induced not by a radiation beam as in earlier studies, instead, induced by a complex radiation field, that can be decomposed into dipole and quadrupole components. We found that in this situation, the parameter space $q^{max}$, for the existence of high-$J$ attractors in trajectory maps is more extended, which entails higher degrees of polarization expected. Our obtained results are useful for modeling polarization arising from aligned grains in molecular clouds and accretion disks. ",
        "introduction": "Magnetic fields play a crucial role in many astrophysical processes (e.g., star formation, accretion disks, cosmic ray transport). An important and easily available way to map magnetic fields is to observe the polarization of emission or absorption by aligned grains (Goodman et al. 1995; Hildebrand et al. 2000, 2002; Crutcher et al. 2004). The reliable interpretation of polarimetry in terms of magnetic fields requires, however, a solid understanding of the alignment of grains. This is the major motivation of work in the field of grain alignment. An additional motivation to understand when grains get aligned comes from the necessity of separating the polarization by dust grains from the polarized Cosmic Microwave Background (CMB) signal (see Lazarian 2008).  Grain alignment is a complex area of studies related to many physical effects of different time scales (see review by Lazarian 2007). It is also the research area associated with the names of Lyman Spitzer and Ed Purcell, who decisively contributed to understanding of a number of fundamental physical processes of grain alignment.  A number of mechanism has been known to induce grain alignment (see Lazarian 2007). However, recently grain helicity has been discussed as the necessary element of successful alignment. Note that the idea of grains having twist and therefore interacting differently with left and right polarized photons can be traced back to the pioneering work by Dolginov \\& Mytrophanov (1976). A more recent work in the area showed that a similar effect is present when irregular grains interact with a gaseous flow (Lazarian \\& Hoang 2007b).\\footnote{Such flows arise naturally in the reference of a grain as the grain interacts with ubiquitous MHD turbulence (Lazarian \\& Yan 2002; Yan \\& Lazarian 2003; Yan, Lazarian, \\& Draine 2004).}  In this paper we concentrate on studying grain alignment by radiative torques (hereafter RATs). However, the intrinsic similarity of the radiative and mechanical torques acting on helical grains (see discussion in Lazarian \\& Hoang 2007b) makes our present results applicable to the alignment by mechanical torques. Note that for the alignment of helical grains, the magnetic field acts only through inducing the Larmor precession. Therefore the alignment may potentially occur with both longer and shorter grain axes perpendicular to the magnetic field. Observations of interstellar dust indicate that the alignment with the long axes perpendicular to the magnetic field is dominant in the diffuse gas. Therefore we shall call this alignment {\\it right alignment}, and the opposite type of alignment we shall call {\\it wrong alignment}. Different alignment processes produce either right or wrong alignment and the goal of the grain alignment theory is to predict correctly both the direction and the degree of the expected alignment.  As dust grain scatters or absorbs photons it experiences radiative torques. These torques can be stochastic or regular. Stochastic RATs arise from, for instance, a spheroidal grain randomly emitting or absorbing photons. The latter process, for instance, was invoked by Harwit (1970) in his model of grain alignment based on grains being preferentially spun up in the direction perpendicular to the photon beam. However, Purcell \\& Spitzer (1971) showed that the randomization arising from the same grain emitting thermal photons makes the achievable degree of grain alignment negligible. In fact, they showed that stochastic RATs arising from grain thermal emission is an important process of grain randomization. More recently, grain thermal emission was analyzed as a source of the excitation of grain rotation as well as its damping in relation to the rotation of tiny spinning grains that are likely to be responsible for the so-called anomalous foreground emission (Draine \\& Lazarian 1998).  RATs were first discussed by Dolginov (1972) in terms of chiral, e.g. hypothetical quartz grains. The idea is that starlight passing through such grains would spin them up. Later, Dolginov \\& Mytrophanov (1976) considered an irregular grain model of two spheroids twisted to each other, made of more accepted materials, e.g. silicate grains, and claimed that these grains will be both spun up and aligned by {\\it regular RATs}. This work was, unfortunately, mostly ignored for 20 years.  Grain alignment by RATs drew much attention when Bruce Draine modified the Discrete Dipole Approximation Code(Draine \\& Flatau 2004, hereafter DDSCAT) and enabled later researchers to calculate RATs for irregular grains. \\footnote{The only work on RATs before that was Lazarian (1995) which corrected some of the points in the Dolginov \\& Mytrophanov (1976) study, but underestimated the importance of RATs.}  Draine \\& Weingartner (1996, hereafter DW96) treated RATs as a kind of pinwheel torque, and they found that RATs can accelerate grains to suprathermal rotations. Later, Draine \\& Weingartner (1997, hereafter DW97) introduced many essential elements of the modern treatment of RATs, for instance, phase trajectory maps. They confirmed Dolginov \\& Mytraphanov (1976) hypothesis that RATs can produce the alignment by their own. However, the important moments of grain dynamics, so-called crossovers, were not treated correctly in the paper. Crossovers, are special events in the dynamics of grains rotating subject to regular torques. During crossovers grain angular velocity approaches its minimal value and models of crossovers caused by pinwheel torques in Spitzer \\& McGlynn (1979) and in Lazarian \\& Draine (1997) envisaged grain flipping with the direction of the angular momentum staying the same. DW97, however, assumed that the RAT crossovers are different and during the RAT-induced crossovers angular momentum changes its direction to the opposite. This treatment of crossovers resulted in grains having cyclic trajectories, which were, in fact, an artifact of the adopted treatment, as was shown later (Lazarian \\& Hoang 2007a, henceforth LH07a). In fact, LH07a showed that instead of cyclic trajectories, one gets situations when RATs tend to {\\it slow} grains down.  A more general analysis for dynamics of RAT alignment was done in Weingartner \\& Draine (2003, henceforth WD03), where thermal fluctuations within grain were accounted for. These fluctuations were shown by Lazarian (1994, see also Lazarian \\& Roberge 1997) to induce grain thermal wobbling, which amplitude gets larger as the grain angular momentum approaches its thermal values.  For their choice of parameters, WD03 found that an attractor with low angular momentum, so-called, low-J attractor, appears on the trajectory phase maps.\\footnote{Conventionally, an attractor with angular momentum $J$ larger than the thermal value is called a high-J attractor, and the attractor with $J$ of the order of thermal value is called low-J attractor.} The study, however, was limited in terms of parameter space explored. In fact, it was a numerical study of alignment of a single grain subjected to the electromagnetic wave of a single frequency at a single incident direction. Therefore implications of the study for the problem of the RAT alignment were difficult to evaluate.  We feel that the numerical work above done with DDSCAT should be treated as an experimental attack on the problem of the RAT alignment. The limitations of such an approach are self-evident. For instance, to obtain predictions one should analyze many hundreds of different grain shapes and many wavelengths and many incident directions. The study in DW97 included 3 shapes and resulting torques looked very different, which made one wonder what causes the alignment\\footnote{In relation to this one of us (AL) recalls that Lyman Spitzer after studying the work on RATs suggested that one should seek for some simple trigonometric fits to torques that would produce the alignment.}. At the same time, unfortunately, the analytical study in Dolginov \\& Mytrophanov (1976) was in error.  The first analytical model (henceforth AMO) of RATs consistent with the numerical calculations was proposed in LH07a. This study provided simple analytical expressions for the RATs (see also \\S 2). In particular, the functional forms of the torques acting on grains of different shapes are similar, with the difference between the grains of different shapes amounting to a single parameter, termed $q^{max}$-ratio, which is the ratio of the magnitudes of two first components of torques in the grain symmetry system of reference where the radiation direction is the symmetry axis. AMO allowed theoretical predictions and the analytical treatment of the RAT alignment. As a result, studies of a parameter space seized to be an insurmountable  problem. In LH07a, similar to DW97, the thermal fluctuations were intentionally disregarded and in such a setup the low-$J$ attractors, rather than cyclic trajectories, were reported. In fact, it became clear that RATs, over a large part of the parameter space, {\\it spin down} rather than {\\it spin up} grains. Using AMO, LH07a obtained the values $q^{max}$ for which grains have only low-$J$ attractors and when they have both low-$J$ and high-$J$ attractors simultaneously.  AMO was extensively used in all our papers that followed. For instance, in Hoang \\& Lazarian (2008a, henceforth HL08a), we studied the RAT alignment in the presence of thermal fluctuations using both AMO and DDSCAT for an extended sample of grain shapes, radiation direction and wavelength. We found that {\\it irregular} grains do not stop completely, but rotate at a rate, which is comparable with the rate of thermal rotation at the dust temperature, which agrees with an earlier example in WD03. Importantly, HL08a considered the RAT alignment in the presence of gaseous bombardment and reported a new effect, namely, the transfer between the low-$J$ and high-$J$ attractors, in situations when the both low-$J$ and high-$J$ attractors were present simultaneously\\footnote{Note that the parameter space for the existence of low-$J$ and high-$J$ attractors are similar and known through the LH07a study.}. Thus, counterintuitively, the collisions were found to {\\it increase} the alignment.  The most important practical implication of the LH07a and HL08a studies was that, over a large range of the parameter space of $q^{max}$-ratio and angles between the magnetic field and radiation direction, the grains rotate thermally and therefore, as a result the aforementioned thermal fluctuations of grains, and exhibit the reduced alignment of grain axes with respect to ${\\bf J}$. This opened possibilities of estimating the expected alignment and comparing it to the alignment inferred from observations. Importantly enough, Lazarian \\& Hoang (2008, henceforth LH08) reported, that in the presence of superparamagnetic inclusions (Jones \\& Spitzer 1967, Mathis 1986, Bradley 1994, Martin 1994, Goodman \\& Whittet 1994), the high-$J$ attractors {\\it always} exist. This finding entails a very non-trivial effect, namely, that, in the presence of both superparamagnetic inclusions, grains, that otherwise would rotate subthermally, get into the state of fast suprathermal rotation. According to HL08a, this also means that {\\it all} grains eventually get into the state of fast rotation and {\\it perfect} alignment. From the point of view of observations, this opens prospects of testing the existence of superparamagnetic inclusions by measuring the degree of polarization.  This paper continues our studies in LH07a and HL08a. The major issues we address below are (1) the effect of internal relaxation, and (2) the effect of complex radiation field on the grain alignment.  The first issue is important in terms of the alignment of large grains, e.g. grains within proto-planetary accretion disks. The alignment of such grains was assumed in a recent modeling of polarization from T-Tauri accretion disks in Cho \\& Lazarian (2007). Grains which are important for such a modeling may, however, should be sufficiently large that the internal relaxation time for them is longer than the alignment time. Therefore, how grains get aligned in the absence of efficient internal relaxation is an important question.  The second issue is vital for modeling the polarization from both disks and dark clouds. The latter modeling which used the RAT alignment was undertaken in Cho \\& Lazarian (2005), Pelkonen et al. (2007) and Bethell et al. (2007). All these studies assumed rather naive models of alignment. In particular, the actual modeling would require taking into account the actual complex radiation field that is  being experienced by the grain, rather than to assume that the radiation is coming from a point source.  In what follows, in \\S 2  we briefly describe the AMO, grain torque-free motion, and calculations of RATs. In \\S 3 we compare the Barnett and nuclear relaxation times with the radiative alignment time and identify grain size when the internal relaxation can be disregarded. In \\S 4 we study the properties of RATs and the resulting RAT alignment for large grains in the absence of the internal relaxation. Grain alignment by RATs arising from dipole and quadrupole components of radiation fields is studied in \\S 5. We discuss our findings in \\S 6, and our summary is presented in \\S 7. ",
        "conclusions": "In the rest of this section we provide a brief account of our accomplishments in the paper and present our outlook on the further work in the field of grain alignment.  \\subsection{Alignment in environments different from ISM: large grains}  Traditionally, grain alignment was the topic of interstellar medium (ISM) research. The gap between the studies of polarized radiation from aligned  dust in environments other then the ISM and the theory of grain alignment got so wide that the researchers outside the ISM domain sometimes write papers about aligned dust and do not refer to {\\it any} theoretical work on grain alignment. However, there is ample evidence of grain alignment in non-ISM environments (see Tamura et al. 1999; Rosenbush et al. 2007; Hough et al. 2007).  In many of these environments, e.g., for dust in comets and circumstellar dust, the radiation field is stronger than in the ISM, thus RATs are stronger. One difference that one faces there is that grains may be substantially larger. As we demonstrated in the paper, for larger grains the internal relaxation gets slower, which makes one wonder whether the RAT alignment happens very differently from the ISM case. Our study shows that the direction of alignment of the angular momentum $\\bJ$ with respect to the magnetic field $\\bB$ is similar to that in the presence of strong internal relaxation. This alignment can be perfect with high-J and/or low-J attractors, depending on the factor $q^{max}$. However, the alignment of grain axes with respect to $\\bJ$ depends on the initial angle between them, more precisely, on the initial value of dimensionless parameter $p$. If this angle is smaller than a particular value, $p_{4}$, defined in the main text, which is a function of $q^{max}$, the grain axis of maximum moment of inertia  $\\ba_{1}$ gets aligned along $\\bJ$, and both $\\ba_{1}$ and $\\bJ$ become aligned with $\\bB$. HL08b showed that pinwheel torques (e.g., H$_{2}$ formation, isotropic radiative torque, etc.) increase fast with the grain size. As a result, large grains can be spun up to suprathermal rotation with $J\\gg J_{th}$. The direct effect of suprathermal rotation is the alignment of $\\ba_{1}$ and $\\bJ$.  Table 1 compares our results for grains without and with internal relaxation using the AMO. We can see that for high-J attractors, the alignment is the same in both cases. We note that the results in the later case were shown to be consistent with the alignment of irregular grains obtained by DDSCAT. However the general case of alignment of large irregular grains in the absence of internal relaxation requires further studies to confirm the predictions obtained with the AMO. Further studies should consider also the effect of pinwheel torques on the alignment in the absence of internal relaxation.  \\subsection{Utility of AMO}  The analytical model (AMO) was introduced in LH07a to describe the case of alignment in the presence of {\\it perfect} alignment of ${\\bf J}$ with the axis of maximal moment of grain inertia. It provided excellent representation of the alignment of irregular grains within this approximation. Later in HL08a and HL08b we applied AMO to grains in the presence of thermal fluctuations, which at small values of $J$ partially randomize the alignment of ${\\bf J}$ with the aforementioned axis. Nevertheless, AMO happened to do a nice job in these cases, adequately describing the behavior of irregular grains.  The case of no internal alignment is the extreme case of testing AMO. For instance, AMO has only one helicity axis, while multiple helicity axes are possible for an irregular grain. It is interesting, that even for this case, our limited testing is indicative of the AMO's utility.  \\subsection{Towards modeling of polarized radiation}  Similar to the polarimetric work of non-ISM observers, the modeling of polarization arising from aligned grains has been mostly developing without much connection to the theory of grain alignment. Exceptions from this rule include Cho \\& Lazarian (2005, 2007), Pelkonen et al. (2007), Bethell et al. (2007), and Falceta-Golcaves et al. (2008). However, this modeling was done on the basis of somewhat ad hoc alignment prescriptions, which should be improved as the theory gets predictive.  A step towards a more reliable polarization modeling is done in this paper where we considered grain alignment induced by the dipole and quadroupole components\\footnote{The alternative to modeling the RAT alignment by the multipoles of the radiation field is to study the alignment numerically at every point of the data cube using the radiation field at the particular point. Naturally, this would entail much more intensive computations.}  of the radiation field. The earlier studies assumed that the radiation was coming from a particular direction, which is, for instance, not the case for most of the grains in starless cores or accretion disks. In the latter cases it is proper to decompose the radiation field in multipoles and consider the effects of the individual components. Our study shows that the parameter space for having high-$J$ attractors differs for the alignment by the different multipole components.  Moreover we identified the range of torque ratio, $q^{max}$, for which the presence of dipole and quadroupole components of radiation field results in the alignment with high-$J$ attractors. The required range $q^{max}$ is fulfilled for irregular grains studied in HL08a.  \\subsection{Credit to Dolginov \\& Mytrophanov 1976}  In our paper we showed that the analytical results on RATs in Dolginov \\& Mytrophanov (1976) were incorrect (see Fig. \\ref{f15}). Therefore, naturally, our present results on the RAT alignment obtained for grains in the absence of internal relaxation, which was also the assumption in Dolginov \\& Mytrophanov (1976) study, differ from those in the study. This should not, however, undermine the pivotal significance of the the Dolginov \\& Mytrophanov paper\\footnote{Incidentally, the same paper discusses for the first time the Barnett effect in the application to insterstellar grains. This induced the all-important notion of fast Larmor precession of grains and helped E. Purcell to discover the effect of Barnett relaxation.}. This paper is important as it {\\it discovered} RATs and discussed the possibility of irregular grains being aligned by RATs, even if it failed to describe it quantitatively. In addition, the notion of grain helicity, which is the central concept of AMO, can be traced back to Dolginov \\& Mytrophanov (1976) work. Moreover, as we mentioned in the introduction, the problem was so hard that the papers that followed Dolginov \\& Mytraphanov (1976), in all its complexity, were not able to capture correctly the physics of the RAT alignment either.  \\subsection{Other processes}  Our paper, for the sake of simplicity, does not consider the effect of the pinwheel torques. Such torques, e.g. torques arising due to H$_2$ formation over catalytic sites over grain surface (Purcell 1979), were discussed in the framework of the RAT alignment in Hoang \\& Lazarian (2008b) for grains with strong internal relaxation. A study of effect of these torques on large grains, for which the effects of internal relaxation are reduced, will be done elsewhere. In addition, in our paper we considered the RAT alignment in respect to magnetic field. As we discussed in LH07a, for a sufficiently slow rate of Larmor rotation, the alignment can happen also in respect to the radiation field. If, however, grains have superparamagnetic inclusions the rate of rotation increases substantially. This, potentially, provides another way of testing whether grains have or do not have superparamagnetic inclusions."
    },
    "0812/0812.1486_arXiv.txt": {
        "abstract": "{Multi-frequency VLBI observations allow studies of the continuum spectrum in the different parts of the parsec scale jets of AGN, providing information on the physical properties of the plasma and the magnetic fields in them. Since VLBI networks cannot be scaled, the range of spatial frequencies observed differs significantly between the different observing frequencies, which makes it difficult to obtain a broadband spectrum of the individual emission features in the jet. In this paper we discuss a model-fitting based spectral extraction method, which can significantly relieve this problem. The method uses {\\it a priori} knowledge of the source structure, measured at high frequencies, to allow at lower frequencies the derivation of the sizes and flux densities of even those emission features that have mutual separations significantly less than the Rayleigh limit at the given frequency. We have successfully used this method in the analysis of 5-86\\,GHz VLBA data of 3C\\,273. The spectra and sizes of several individual jet features were measured, thus allowing derivation of the magnetic flux density and the energy density of the relativistic electrons in the different parts of the jet. We discuss the results, which include e.g. a detection of a strong gradient in the magnetic field across the jet of 3C\\,273.}  \\FullConference{The 9th European VLBI Network Symposium on The role of VLBI in the Golden Age for Radio Astronomy and EVN Users Meeting\\\\ September 23-26, 2008\\\\ Bologna, Italy}  \\begin{document} ",
        "introduction": "Despite the large steps that have been taken in our understanding of the nature of relativistic outflows in active galactic nuclei during the last three decades, we still remain ignorant of many of the key issues of the subject: details of the launching mechanism, matter content and the origin of the high energy emission are among the hitherto poorly understood aspects of relativistic jets \\cite{mar06}. One of the reasons for slow progress in these questions is our poor ability to measure some of the key physical parameters of jets.  There are basically two ways to determine the energy of the relativistic particles in jets: one is to assume equipartition between the magnetic and particle energy densities, the other is to measure the size and spectral turnover of the synchrotron self-absorbed emission features. The latter gives magnetic flux density and the normalization factor of the electron energy distribution separately, but it requires a measurement of the 2-D spectrum of parsec scale jets using multi-frequency VLBI data \\cite{mar87}. Such measurements are possible, yet challenging. In this paper, we describe a new method for the extraction of the spectral information from a multi-frequency VLBI data set and apply this method to a set of six-frequency VLBA observations of the quasar 3C\\,273 ($z$=0.158). Throughout the paper, we use a cosmology with $H_0$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$ = 0.27, and $\\Omega _\\Lambda$ = 0.73. This corresponds to a linear scale of 2.7 pc mas$^{-1}$ for 3C\\,273. ",
        "conclusions": ""
    },
    "0812/0812.1353_arXiv.txt": {
        "abstract": "{ In order to reveal the long-term evolution of relativistic jets in active galactic nuclei (AGNs), we examine the dynamical evolution of variously-sized radio galaxies [i.e., compact symmetric objects (CSOs), medium-size symmetric objects (MSOs), Fanaroff-Riley type I\\hspace{-.1em}I radio galaxies (FRI\\hspace{-.1em}Is)]. By comparing the observed relation between the hot spot size and the linear size of radio source with a coevolution model of hot spot and cocoon, we find that the advance speed of hot spots and lobes inevitably show the deceleration phase (CSO-MSO phase) and the acceleration phase (MSO-FRI\\hspace{-.1em}I phase). The deceleration is caused by the growth of the cross-sectional area of the cocoon head. Moreover, by comparing the hot spot speed with the sound speed of the ambient medium, we predict that only CSOs whose initial advance speed is higher than 0.3-0.5c can evolve into FRI\\hspace{-.1em}Is. } ",
        "introduction": "According to the standard model for radio-loud AGNs, a pair of relativistic jets transport away the bulk kinetic energy from the central compact region close to a central supermassive black hole (SMBH) to $l_{\\rm h}\\geq 10\\, {\\rm kpc}$ scale radio lobes, where $l_{\\rm h}$ is the distance from the centre of the galaxies (Blandford \\& Rees 1974). However, how young radio loud AGNs evolve into powerful extended radio sources (e.g., FRI\\hspace{-.1em}I radio galaxies) is one of the primary problems in astrophysics. In order to clarify this issue, it is important to understand the nature of the small, young progenitors of FR I\\hspace{-.1em}Is. Young and compact radio sources such as CSOs ($l_{\\rm h} <$ 1kpc) and MSOs ($l_{\\rm h}=1-10$ kpc) have been discovered by many authors (e.g., Wikinson et al. 1994; Fanti et al. 1995; Readhead et al. 1996). The bright CSOs and MSOs show the FRI\\hspace{-.1em}I like morphology and possess hot spots (reverse shock) which are the signatures of supersonic expansions. The recent observations of several CSOs and MSOs have indicated that their age ($\\sim 10^{2-5}\\,{\\rm yr}$) is shorter than the typical age of FRI\\hspace{-.1em}Is (e.g., Parma et al. 1999), which implies that CSOs and MSOs are possible candidates as the progenitors of FRI\\hspace{-.1em}Is (e.g., Owsianik, Conway \\& Polatidis 1998; Taylor et al. 2000; Giroletti et al. 2003; Polatidis \\& Conway 2003; Nagai et al. 2006).  A number of authors have investigated the long-term evolution of extragalactic radio sources in different ways (e.g., Begelman \\& Cioffi 1989: hereafter BC89; Fanti et al. 1995; Begelman 1996; Readhead et al. 1996; Kaiser \\& Alexander 1997; Snellen, Schilizzi \\& van Langevelde 2000; Perucho \\& Mart{\\'{\\i}} 2002; Kawakatu \\& Kino 2006, hereafter KK06). A constant advance speed of hot spots, or a constant aspect ratio of the cocoon (i.e., a self-similar evolution) have often been assumed for a dynamical evolution of radio-loud AGNs for simplicity. {\\it Do these assumptions reflect the actual evolution of radio sources ?} In order to answer this question, it is essential to build up an appropriate model of radio sources without assuming the constant aspect ratio and the constant advance speed of hot spots such as that presented by KK06.  Recent observations suggested that the power law index for the evolution of hot spot size changes at the transition between the interstellar medium and intergalactic medium, i.e., $l_{\\rm h}\\sim 1-10\\,{\\rm kpc}$ (Jeyakumar \\& Saikia 2000: hereafter J00; Perucho \\& Mart{\\'{\\i}} 2003: hereafter PM03). Since hotspots are identified as the reverse-shocked region of the decelerating jet, the evolution of hot spot size could reflect the dynamical growth of radio sources of various scales. However, it was hard to derive the dynamical evolution of radio sources from previous work (J00 and PM03) because of lack of spatial resolution, the observational bias and small sample of radio sources. In this study, we first compile sample of CSOs, MSOs and FRI\\hspace{-.1em}Is sources larger than in previous works, by considering the observational bias and being careful about the data quality. Then, from the direct comparison with KK06, we make clear a long-term evolution of advance speed of hot spots (Kawakatu, Nagai \\& Kino 2008 for details: hereafter KNK08).  \\begin{figure} \\includegraphics[width=80mm,height=80mm]{Kawakatu_f1.eps} \\caption{ The relation of hot-spot size ($r_{\\rm HS}$) and hot-spot distance from the core ($l_{\\rm h}$). Crosses with arrows indicate upper limit. Solid line corresponds to the best-fit for the sources where $10^{-3}$~kpc$\\le l_{\\rm h}\\le1$~kpc whereas broken line corresponds to that for the sources where $1$~kpc$\\le l_{\\rm h}\\le10^{3}$~kpc. Note that upper-limit data were not included in the fittings. } \\label{label1} \\end{figure} ",
        "conclusions": "We have investigated the relation between CSOs, MSOs and FRI\\hspace{-.1em}Is, by comparing the coevolution model of hot spots and a cocoon (KK06) with the observed $r_{\\rm HS}-l_{\\rm h}$ relation reflecting the dynamical evolution of radio sources. We find that the advance speed of hot spots and lobes strongly decelerate when the jets pass through the ambient medium in host galaxies (i.e., the CSO-MSO phase), while the jets accelerate outside host galaxies (i.e, the MSO-FRI\\hspace{-.1em}I phase). The reason of deceleration is the growth of the cross-sectional area of cocoon head. The predicted deceleration of hot spots seems to be consistent with the recent observation show that the outflow velocity of the ionized gas around radio lobes may decelerate as $l_{\\rm h}$ increaes (Labiano 2008). Furthermore, by comparing the hot spot speed with the sound speed of the ambient medium, we predict that only CSOs whose initial advance speed is higher than 0.3-0.5c can evolve into FRI\\hspace{-.1em}Is. This indicates that the origin of the FRI/FRI\\hspace{-.1em}I dichotomy could be related to the difference in the initial advance speed (Kawakatu, Kino \\& Nagai 2008 in preparation)."
    },
    "0812/0812.0163_arXiv.txt": {
        "abstract": "The UKIRT Infrared Deep Sky Survey (UKIDSS) is the first of a new generation of infrared surveys. Here we combine the data from two UKIDSS components, the Large Area Survey (LAS) and the Galactic Cluster Survey (GCS), with 2MASS data to produce an infrared proper motion survey for low mass stars and brown dwarfs. In total we detect 267 low mass stars and brown dwarfs with significant proper motions. We recover all ten known single L dwarfs and the one known T dwarf above the 2MASS detection limit in our LAS survey area and identify eight additional new candidate L dwarfs. We also find one new candidate L dwarf in our GCS sample. Our sample also contains objects from eleven potential common proper motion binaries. Finally we test our proper motions and find that while the LAS objects have proper motions consistent with absolute proper motions, the GCS stars may have proper motions which are significantly under-estimated. This is due possibly to the bulk motion of some of the local astrometric reference stars used in the proper motion determination. ",
        "introduction": "The study of the lowest mass stars and brown dwarfs is one of the most active areas of current Galactic research. The identification of samples of such objects allows both the low mass/substellar luminosity and mass functions to be constrained. These can in turn be used to constain models of star and brown dwarf formation. Additionally such samples may produce interesting single or multiple objects, whose spectroscopic properties could inform atmospheric models for brown dwarfs and giant planets.  Until relatively recently the majority of proper motion surveys have concentrated on optical data. Among the first to use proper motion as a tool to select faint, nearby populations was Luyten. His work using photographic plates culminated in two large scale catalogues of high proper motion stars, the Luyten Half Arcsecond catalogue (LHS - Luyten 1979a) and the New Luyten Two-Tenths catalogue (NLTT - Luyten 1979b). Since that time much work has gone into identifying objects Luyten missed and filling in areas of poor coverage. The majority of recent wide-field proper motion surveys have used digitised photographic plate data, mostly in optical filters. Such surveys include Lepine \\& Shara (2008), Hambly et al. (2004), Pokorny et al. (2003) and Finch et al. (2007). Infrared proper motion surveys are an ideal instrument for producing samples of low mass stars and brown dwarfs. These objects are cool $T_{eff}<3500K$ and hence are intrinsically faint and red. As a result they are most easily detected in the infrared. The first infrared proper motion survey was conducted by Deacon, Hambly \\& Cooke (2005), this used SuperCOSMOS (Hambly et al. 2001) scans of UKST $I$ plates along with data from the 2 Micron All Sky Survey (2MASS - Skrutskie et al., 2006) to produce a sample of 144 low mass stars and brown dwarfs with proper motions above half an arcsecond per year. In the later paper (Deacon \\& Hambly 2007) the sample was expanded to more than 7000 objects with the minimum proper motion limit reduced to 0.1''/yr. Recently Looper et al. (2008) used available overlap areas in the 2MASS data to detect two new L dwarfs using an infrared proper motion survey. Metchev et al. (2008) used 2MASS data along with those from the partially infrared Sloan Digital Sky Survey (Adelman-McCarthy et al. 2006) to identify two new T dwarfs and 22 new L dwarfs.  Recent large-scale infrared surveys such as the DEep Near Infrared Survey (DENIS - Epchtein et al., 1997)) and 2MASS took place at the end of the last decade and the start of this one. They produced CCD quality infrared data over tens of thousands of square degrees to depths of 16.5 and 15.8 respectively in the $J$ band. Both surveys have led to large samples of late type objects being identified and classified (Kirkpatrick et al. 1999, Burgasser et al. 2002, Delfosse et al. 1997). The generation of infrared surveys following DENIS and 2MASS is led by the UKIRT Infrared Deep Sky Survey (UKIDSS - Lawrence et al. 2007). This suite of surveys using WFCAM (Casali et al. 2007) on the UK Infrared Telescope (UKIRT) provides both large scale surveys for studies of galactic and extragalactic populations and deep pencil-beam surveys to examine the population of high redshift galaxies. The two surveys of the most interest to the study of low mass stars and brown dwarfs are the Large Area Survey (LAS) and the Galactic Clusters Survey (GCS). The LAS will have a final area of 4000 square degrees, will be in four passbands ($Y$, $J$, $H$ and $K$) to a depth of $J$=19.6 and will also provide a second $J$ band epoch for proper motion measurements. It covers an area away from the Galactic plane and is intended for the study of high redshift quasars and cool dwarf/subdwarfs in the field. This survey has already produced interesting results, both in terms of individual cool objects (Warren et al. 2007) and the population of low mass objects (Pinfield et al. 2008). The GCS concentrates on ten star forming regions and open clusters. In addition to the four colours used in the LAS, $Z$ band data is available and a second $K$ epoch will provide proper motion data for the clusters. So far studies such as Lodieu et al. (2006) and Lodieu et al. (2007) have used the data released to produce mass functions for Upper Sco and the Pleiades respectively. The GCS can be used also for the study of objects lying in the foreground of the target clusters (Lodieu et al., in preparation). The UKIDSS data is released through the WFCAM Science Archive (WSA - Hambly et al. 2008) and is available to all European Southern Observatories (ESO) members. The current release is Data Release 4 (DR4). Additionally astronomers from non-ESO countries can access the data after an 18 month delay. The soon to be operational Visual and Infrared Survey Telescope for Astronomy (VISTA - Emerson et al. 2004) will provide an even faster infrared surveying capability than WFCAM. It will cover the whole southern hemisphere in two bands ($J$ and $K$) including several thousand square degrees with additional $Y$ and/or $H$ band photometry. ",
        "conclusions": "Our attempt to utilise the currently available UKIDSS data for a proper motion survey has successfully detected all the expected, previously known single L and T dwarfs in our LAS survey area. Additionally we have found seven objects in the LAS and one in the GCS that are good candidate L dwarfs. Finally, initial spectral follow--up of our sample has classified two previously known objects. Rough distance estimates put these as being within 25pc. In further iterations of this work with future UKIDSS data releases we will attempt to tie our proper motions to absolute proper motions, avoiding the offsets in the Galactic Clusters Survey sample. This work could easily be extended to cover the upcoming surveys by the VISTA telescope."
    },
    "0812/0812.2811_arXiv.txt": {
        "abstract": "We prove that in Einstein-Maxwell theory the inequality $(8\\pi J)^2+(4\\pi Q^2)^2 < A^2$ holds for any sub-extremal axisymmetric and stationary black hole with arbitrary surrounding matter. Here $J, Q$, and $A$ are angular momentum, electric charge, and horizon area of the black hole, respectively. ",
        "introduction": "For a single rotating, electrically charged, axisymmetric and stationary black hole in vacuum (described by the Kerr-Newman family of solutions), the angular momentum $J$, the electric charge $Q$, and the horizon area $A$ are restricted by the inequality \\begin{equation}\\label{ineq} p_J^2+p_Q^2 \\le 1 \\quad\\textrm{with}\\quad p_J := \\frac{8\\pi J}{A},\\quad p_Q := \\frac{4\\pi Q^2}{A}. \\end{equation} Equality in \\eqref{ineq} holds if and only if the Kerr-Newman black hole is extremal. That is to say, \\begin{equation}\\label{ineq_strict} p_J^2+p_Q^2 < 1 \\end{equation} holds for any non-extremal Kerr-Newman black hole.  As was shown in \\cite{Ansorg}, the equality $p_J^2+p_Q^2=1$ holds more generally in Einstein-Maxwell theory for axisymmetric and stationary \\emph{degenerate}\\footnote{Degeneracy of an  axisymmetric and stationary black hole is defined by vanishing surface gravity $\\kappa$.} black holes \\emph{with surrounding matter}. Moreover, it was conjectured in \\cite{Ansorg} that inequality (\\ref{ineq}) is still valid if the black hole is surrounded by matter (i.e. if it is not a member of the Kerr-Newman family).  Inequality (\\ref{ineq_strict}) was proved in \\cite{Hennig} for axisymmetric and stationary black holes with surrounding matter in pure Einsteinian gravity (without Maxwell field). In that article, emphasis was put on ``physically relevant'' configurations by assuming the black hole to be \\emph{sub-extremal}. This condition requires the existence of trapped surfaces (i.e. surfaces with a negative expansion of outgoing null geodesics) in every sufficiently small interior vicinity of the event horizon, see \\cite{Booth}. Here, we consider again sub-extremal axisymmetric and stationary black holes with arbitrary surrounding matter, but provide a proof of (\\ref{ineq_strict}) which is valid in the full Einstein-Maxwell theory.  The idea of the proof relies on showing that a black hole \\emph{cannot} be sub-extremal for $p_J^2+p_Q^2 \\ge 1$. In order to prove this, we study the Einstein-Maxwell equations in a vicinity of the black hole horizon. It turns out that a reformulation can be found which states that an appropriate functional $I$ (to be defined below) must always be greater than or equal to $1$. In this way, we encounter a \\emph{variational problem}, and the corresponding solution provides a proof of inequality (\\ref{ineq_strict}). As will be shown below, this variational problem can be treated with methods from the calculus of variations.  This paper is organized as follows. In Sec.~\\ref{coords}, we introduce appropriate coordinates which are adapted to the subsequent analysis. Moreover, we list the Einstein-Maxwell equations and the corresponding boundary and regularity conditions in these coordinates. In Sec.~\\ref{calc}, we express the ingredients $p_J$ and $p_Q$, which appear in the inequality (\\ref{ineq_strict}), in terms of metric and electromagnetic potentials. We formulate the variational problem mentioned above in Sec.~\\ref{reform} and solve it in Sec.~\\ref{solution}. Finally, we conclude this paper with a discussion on physical implications of inequality (\\ref{ineq_strict}), see Sec.~\\ref{discussion}. In an appendix, we establish a connection to degenerate black holes. ",
        "conclusions": "\\label{discussion}  With techniques from the calculus of variations, we have shown that the inequality $p_J^2+p_Q^2<1$ holds for axisymmetric and stationary sub-extremal black holes with surrounding matter in full Einstein-Maxwell theory.  In particular, we have proved the inequality $I[S,T,U]\\ge 1$ for the functional $I$ defined in \\eqref{functional}. As $I$ could not directly be seen to have a local minimizer, we introduced a family of approximating functionals $\\Ie$ which could be shown to have one.   Together with a theorem for \\emph{degenerate} black holes in \\cite{Ansorg}, we can deduce the following. \\begin{theorem}\\label{theorem} Consider Einstein-Maxwell spacetimes with vanishing cosmological constant. Then, for every axisymmetric and stationary sub-extremal black hole with arbitrary surrounding matter we have the inequality $$(8\\pi J)^2 +(4\\pi Q^2)^2 < A^2.$$ If the axisymmetric and stationary black hole is degenerate, the equation $$(8\\pi J)^2 +(4\\pi Q^2)^2 = A^2$$ holds. \\end{theorem}  Observe that the assumptions for the result in \\cite{Ansorg} which has been used here have been weakened, see appendix~\\ref{appendix}.  Theorem \\ref{theorem} provides a remarkable relation between the \\emph{geometrical} concept of the existence of trapped surfaces and the \\emph{physical} black hole properties described by rotation rate $p_J$ and charge rate $p_Q$. We see that ``physically reasonable'' (sub-extremal) black holes cannot rotate ``too fast'' and cannot be charged ``too strongly''.  Finally, our results shed new light on the notions of \\emph{sub-extremality} and \\emph{extremality} of axisymmetric and stationary black holes. Any sub-extremal black hole in the sense of \\cite{Booth} (the notion of which we have used throughout this paper) is also sub-extremal in the sense that $p_J^2+p_Q^2<1$. In fact, $p_J^2+p_Q^2=1$ holds in the degenerate limit, for which reason we may call these black holes ``extremal''.   \\begin{acknowledgement} We would like to thank Herbert Pfister for many valuable discussions and John Head for commenting on the manuscript. This work was supported by the Deutsche Forschungsgemeinschaft (DFG) through the Collaborative Research Centre SFB/TR7 ``Gravitational wave astronomy'' and by the International Max Planck Research School for ``Geometric Analysis, Gravitation and String Theory''. \\end{acknowledgement} \\appendix"
    },
    "0812/0812.0814_arXiv.txt": {
        "abstract": "We report an over-density of bright sub-millimetre galaxies (SMGs) in the 0.15~deg$^2$ AzTEC/COSMOS survey and a spatial correlation between the SMGs and the optical-IR galaxy density at $z\\lsim$~1.1. This portion of the COSMOS field shows a $\\sim3\\sigma$ over-density of robust SMG detections when compared to a background, or ``blank-field'', population model that is consistent with SMG surveys of fields with no extragalactic bias. The SMG over-density is most significant in the number of very bright detections (14 sources with measured fluxes $S_{1.1mm} > 6$~mJy), which is entirely incompatible with sample variance within our adopted blank-field number densities and infers an over-density significance of $\\gg 4\\sigma$. We find that the over-density and spatial correlation to optical-IR galaxy density are most consistent with lensing of a background SMG population by foreground mass structures along the line of sight, rather than physical association of the SMGs with the $z\\lsim$~1.1 galaxies/clusters. The SMG positions are only weakly correlated with weak-lensing maps , suggesting that the dominant sources of correlation are individual galaxies and the more tenuous structures in the region and not the massive and compact clusters. These results highlight the important roles cosmic variance and large-scale structure can play in the study of SMGs. ",
        "introduction": "Foreground structure, cosmic variance, and source environment can affect the observer's perception and interpretation of the source population being probed in a particular survey. For example, gravitational lensing by massive foreground clusters affects both the observed flux of sources and the areal coverage of the survey in the source plane. These aspects of gravitational lensing have been utilised to probe the very faint sub-millimetre galaxy (SMG) population below the confusion limit imposed by the high density of faint SMGs relative to the survey beam size \\citep[e.g.][]{smail97,chapman02a,cowie02,smail02,knudsen06,wilson08b}. The measured (sub)millimetre fluxes of sources found in the direction of very massive clusters can also be affected by, and confused with, the signal imposed through the Sunyaev-Zel'dovich effect on the cosmic microwave background \\citep[e.g.][]{wilson08b}. Furthermore, surveys can be affected by foreground structures with high galaxy-densities, which increase the likelihood of galaxy-galaxy lensing by intervening galaxies and complicate counterpart identification at other wavelengths \\citep{chapman02b,dunlop04}.  Spectroscopic observations have shown that the vast majority of SMGs with detectable radio counterparts lie at an average redshift of $z\\sim$~2.2 \\citep{chapman05}, while spectroscopic \\citep{valiante07} and photometric \\citep[e.g.][]{younger07} analysis put many radio-faint SMGs at even higher redshifts. The average SMG is unlikely to be found at $z\\lsim1$, however it remains to be seen if the $z\\sim1$ SMG population can be locally enhanced due to large-scale structure and cosmic variance. Some evidence exists for increased number densities of SMGs in mass-biased regions of the $z \\gsim 1$ Universe. Surveys towards several $z \\sim 1$ clusters \\citep{best02,webb05} find a number density of SMGs in excess of the blank-field counts that can not be explained by gravitational lensing alone. This implies that some of the SMGs are physically associated with the clusters, although the number statistics are small and the lensing could be underestimated \\citep{webb05}. Similar over-densities have been found towards high-redshift radio galaxies \\citep{stevens03,debreuck04,greve07} and $z > 5$ quasars \\citep{priddey08}, where lensing of background sources is less likely to be an issue. Spectroscopic observations have also found common redshifts amongst SMGs in the SSA22 and HDF fields, suggesting physical over-densities of SMGs at redshifts of 3.1 and 2.0, respectively \\citep{chapman05}. Together, these surveys suggest that these massive dusty starbursts are prominent in moderate and high-redshift cluster/proto-cluster environments.  In this paper, we analyse the density and distribution of SMGs in the AzTEC/COSMOS survey \\citep{scott08}. The AzTEC/COSMOS survey covers a region within the COSMOS field \\citep{scoville07a} known to contain a high density of optical-IR galaxies and prominent large-scale structure at $z \\lsim 1.1$ \\citep{scoville07b}, including a massive $M \\sim 10^{15} M_\\odot$ cluster at $z \\approx 0.73$ \\citep{guzzo07}. In \\S~\\ref{sec:counts} we present the 1.1~mm source counts for the AzTEC/COSMOS field, revealing a strong over-density of bright SMGs compared to the blank-field. We explore the nature of this over-density through examination of the spatial correlation between SMGs and the known large-scale structures (\\S~\\ref{sec:LSScorr}). Positive correlation between SMG positions and low-redshift large-scale structure has been previously detected statistically in 3 disjoint SCUBA surveys \\citep{almaini03,almaini05}. We now present a wide-field investigation of such correlations using the AzTEC/COSMOS survey, which has advantages in its contiguous size (0.15~deg$^2$), broad range of low-redshift environments, and the availability of deep multiband imaging and reliable photometric redshifts \\citep{ilbert08}.  This is the second paper describing the 1.1~mm results of the AzTEC/COSMOS survey. Paper I \\citep{scott08} presented the data-reduction algorithms, AzTEC/COSMOS map and source catalogue, and confirmation of robustness of the AzTEC/JCMT data and pointing. Additionally, seven of the brightest AzTEC/COSMOS sources have had high-resolution follow-up imaging at 890~$\\micron$ using the Sub-Millimetre Array (SMA) and are discussed in detail in \\citet{younger07}. Spitzer IRAC colours of these SMGs, and others, are discussed in \\citet{yun08}. ",
        "conclusions": "The central 0.15~deg$^2$ of the AzTEC/COSMOS survey shows a significant over-density of bright 1.1~mm-detected SMGs when compared to the background population inferred by other surveys. We find that this over-density cannot be explained as sample variance of the blank-field SMG population. The SMG positions are significantly correlated with the $z\\lsim$~1.1 optical/IR galaxy density on the sky, which is believed to be in the foreground of nearly all AzTEC/COSMOS SMGs. Both the spatial correlation and the AzTEC/COSMOS SMG number counts are consistent with gravitational amplification of the blank-field SMG population. The lack of strong correlation to the weak-lensing maps of \\citet{massey07} indicates that this amplification is primarily due to weak-lensing by the large-scale structure as opposed to lensing by the compact and massive clusters in the field. SMGs detected in a different part of the COSMOS field by the 1.2~mm COSBO survey are also spatially correlated to the $z\\lsim$~1.1 structure, however, this correlation is dominated by two compact structures (likely clusters) in the field. The lack of significant large-scale structure (i.e. lensing opportunities) in the rest of the COSBO survey region results in COSBO number counts that are consistent with the blank-field \\citep{bertoldi07} -- a strong contrast to the significant SMG over-density and rich foreground structure found in the nearby AzTEC/COSMOS field."
    },
    "0812/0812.2020_arXiv.txt": {
        "abstract": "The observed intergalactic \\ovi\\ absorbers at $z>0$ have been regarded as a significant reservoir of the ``missing baryons''. However, to fully understand how these absorbers contribute to the baryon inventory, it is crucial to determine whether the systems are collisionally ionized or photoionized (or both). Using the identified intergalactic \\ovi\\ absorbers as tracers, we search for the corresponding X-ray absorption lines, which are useful for finding the missing baryons and for revealing the nature of the \\ovi\\ absorbers. Stacking the {\\sl Chandra} grating spectra along six AGN sight lines, we obtain three spectra with signal-to-noise ratios of 32, 28, and 10 per 12.5 m\\AA\\ spectral bin around the expected \\ovii~K$\\alpha$ wavelength.  These spectra correspond to \\ovi\\ absorbers with various dynamic properties.  We find no detectable \\neix, \\ovii, \\oviii, \\nvii, or \\cvi\\ absorption lines in the spectra, but the high counting statistics allows us to obtain firm upper limits on the corresponding ionic column densities (in particular $N_{\\rm OVII}\\lsim10~N_{\\rm OVI}$ on average at the 95\\% confidence level). Jointly analyzing these non-detected X-ray lines with the averaged \\ovi\\ column density, we further limit the average temperature of the \\ovi-bearing gas to be $T\\lsim10^{5.7}$ K in collisional ionization equilibrium.  We discuss the implications of these results for physical properties of the putative warm-hot intergalactic medium and its detection in future X-ray observations. ",
        "introduction": "\\label{sec:intro}  Hydrodynamic simulations of cosmic large-scale structure indicate that, in the local universe ($z\\lsim1$), the shock-heated gas in the tenuous warm-hot intergalactic medium (WHIM) at temperatures $T\\sim10^{5-7}$ K contains a substantial amount of baryonic matter \\citep{cen99, dav01}, providing a possible solution for the so-called ``missing baryons'' problem (e.g., \\citealt{per92, fuk98}). In a collisionally ionized gas at these temperatures, the most abundant heavy elements (e.g., C, N, O, and Ne) are in their high ionization (e.g., Li-, He-, and H-like) states \\citep{sut93}, whose K- and L-shell transitions are in the X-ray and far-ultraviolet (far-UV) wavelength bands, respectively.  The existence of the WHIM has been suggested through detections of X-ray emission from high-density regions near galaxies, groups, and clusters (e.g., \\citealt{wang97, wang04, fin03, sol07, man07}; see also \\citealt{dur08} for a review), but these regions are expected to contain only a small portion of the WHIM gas. The majority is located in low-density cosmic-web filaments \\citep{dav01}, whose emission is hard to detect with current X-ray and far-UV telescopes.  These filaments are most easily probed through absorption lines imprinted on spectra of background sources.  With the high sensitivity spectrographs aboard the {\\sl Hubble Space Telescope} (HST) and the {\\sl Far Ultraviolet Spectroscopic Explorer} (FUSE), the intergalactic \\ovi\\ absorption lines (at rest-frame wavelengths $1031.93$ \\AA\\ and 1037.62 \\AA) have routinely been detected in many background AGN spectra (e.g., \\citealt{sav98, tri00a, tri00, shu03, pro04, dan05, coo08}). However, the nature of these \\ovi\\ absorbers is still under debate. In a survey of the intergalactic medium (IGM) absorption lines, \\citet{dan08} found a good correlation in column density and a similar power-law slope of the column density distribution $d{\\cal{N}}/dz$ of \\ovi\\ and \\nv, which are distinct from those of \\HI, \\ciii, and \\siiii.  These features, along with the theoretical expectations of high post-shock temperatures and long cooling times of the WHIM, motivated them to argue for a multiphase nature of the IGM. In this picture, \\ovi\\ and \\nv\\ trace the canonical WHIM at temperatures of $10^{5-6}$ K, while the low ionization ions like \\cii-\\civ\\ and \\siii-\\siiv\\ are predominantly photoionized. But in two other independent surveys that include a number of common sight lines to those used by \\citet{dan08}, \\citet{tri08} and \\citet{tho08} found that 30-40\\% of the \\ovi\\ absorbers have velocity centroids that are well aligned with those of the associated \\HI\\ absorbers and that some absorber temperatures inferred from the line widths, $T < 10^5$ K, are well below that expected in the canonical WHIM. \\citet{tri08} also showed that these velocity-aligned absorbers can be naturally explained by photoionization models. These latter findings are consistent with previous investigations along individual sight lines based on the kinematic and chemical properties of the high- and low-ionization absorbers (e.g., \\citealt{tri00, sav02, sem04, pro04, leh06}). However, the UV data usually do not have high enough quality to uniquely constrain the ionization mechanism; \\citet{tri08} also showed that the \\ovi\\ lines could alternatively arise in hot interface layers on the surface of low-ionization clouds if the \\ovi\\ systems are always multiphase media. Current data cannot rule out this possibility, and many \\ovi\\ absorbers are fully consistent with this hypothesis (see \\S~4.2 in \\citealt{tri08}). So far, the strongest indication of the collisional ionization origin of the \\ovi\\ absorbers is the detection of \\neviii\\ absorption lines (at rest-frame wavelengths 770.4 \\AA\\ and 780.3 \\AA) toward HE~0226-4110 at $z_{\\rm a}=0.2070$ and toward 3C~263 at $z_{\\rm a}=0.3257$ \\citep{sav05, nar08}. However, a systematic search toward other sight lines has failed to find similar systems \\citep{leh06}.  The claimed detections of the WHIM in the X-ray band are still controversial.  In the spectrum of PKS~2155-304 observed with the {\\sl Chandra} Advanced CCD Imaging Spectrometer (ACIS), \\citet{fang02} reported an \\oviii\\ K$\\alpha$ absorption line at $z=0.0553$, claiming the first detection of the WHIM in the X-ray. This detection has not been confirmed, although it cannot be ruled out either, by the {\\sl Chandra} High Resolution Camera (HRC) observations and the subsequent {\\sl XMM-Newton} Reflection Grating Spectrometer (RGS) observations \\citep{cag04, wil07, fang07}.  \\citet{nic05} reported detections of two $z>0$ \\ovii\\ WHIM systems in the {\\sl Chandra} spectra of Mrk~421, which were not confirmed from the {\\sl XMM-Newton} RGS observations of the same source \\citep{kaa06, ras07}.  \\citet{kaa06} also questioned the statistical significance of the Mrk~421 detections in the {\\sl Chandra} spectra. The other claimed $z>0$ X-ray WHIM absorptions are mostly detected at marginal significance (e.g., \\citealt{mat03}). Very recently, \\citet{buote09} have reported a $3\\sigma$ detection of \\ovii\\ K$\\alpha$ absorption affiliated with a large-scale structure (the Sculptor Wall). The \\ovii\\ and \\oviii\\ absorption lines at $z\\simeq0$, which may be partially due to the WHIM in the Local Group (e.g., \\citealt{nic02, fang03}), have been relatively well detected, but these absorptions are severely confused by the contributions from the Galactic hot gas close to (or within) the Milky Way disk \\citep{sem03, yao05, yao07, wang05, fang06, bre07, yao08}.  With the controversial X-ray detections of the WHIM and the debatable nature of the observed far-UV absorbers, two key questions still remain open: (1) Does the WHIM exist? (2) If it does, what are its physical properties?  The highly ionized X-ray absorption lines are believed to be useful for finding the WHIM gas and for probing the nature of the \\ovi\\ absorbers.  The \\ovii\\ ion can trace gas over a broad temperature range, and its column density is expected to be $\\gsim10$ times higher than that of \\ovi\\ in a shock-heated gas with shock temperatures $T\\gsim10^{5.7}$ K (for overdensity $\\delta\\sim10-100$; e.g., \\citealt{fur05}).  Moreover, cosmological simulations predict that hot \\ovi\\ absorbers should have comparably strong (or even stronger) affiliated \\ovii\\ absorption lines (see, e.g., Figures 8-10 in \\citealt{cen06} and Figure 14 in \\citealt{chen03}). However, because of the limited sensitivity and spectral resolution of the current X-ray observatories like {\\sl Chandra} and {\\sl XMM-Newton}, searching for X-ray lines in the WHIM is currently a difficult task.  For instance, the column densities of the X-ray absorbing ions, even for the most promising one (\\ovii ), are expected to be small with $N_{\\rm OVII}\\approx 10^{15}~\\rm{cm^{-2}}$ (e.g., \\citealt{chen03}). For the most part, the background AGN are also relatively faint in the X-ray band.  Consequently, very long exposures (see \\S~\\ref{sec:dis}) are required to collect enough photons to conduct a blind search for X-ray absorption lines from the WHIM along a random line of sight. These difficulties contribute to the controversies and frustrations surrounding the reported X-ray detections.  In this work, we search for the X-ray absorption lines of \\ovii, as well as \\oviii, \\neix, \\nvii, and \\cvi, by extensively exploring the archived {\\sl Chandra} grating observations.  We use the identified \\ovi\\ absorbers as tracers to avoid uncertainties of a blind search. We also develop a technique to stack all the applicable X-ray observations and obtain spectra with unprecedented counting statistics.  The paper is organized as follows. In \\S~\\ref{sec:data}, we list the identified \\ovi\\ absorbers and the corresponding {\\sl Chandra} observations utilized in this work, and describe the data reduction process. In \\S~\\ref{sec:method}, we present our search method and stack the X-ray observations in the rest frame of the observed \\ovi\\ absorbers. We search for the X-ray absorption lines in the stacked spectra, present the results of the data analysis in \\S~\\ref{sec:results}, and discuss the implications of our results in \\S~\\ref{sec:dis}. ",
        "conclusions": "\\label{sec:dis}  We have presented a search for X-ray absorption lines produced in the WHIM by using the identified far-UV \\ovi\\ absorbers as tracers for stacking the archived {\\sl Chandra} observations. The three final stacked spectra, corresponding to all, complex, and strong \\ovi\\ absorbers, have signal-to-noise ratios of $\\sim32$, 28, and 10, respectively, per 12.5 m\\AA\\ spectral bin around the \\ovii\\ K$\\alpha$ wavelength.  There are no detectable X-ray absorption lines at the expected wavelengths in these spectra.  We have obtained upper limits to EWs of the K$\\alpha$ lines of \\ovii, \\oviii, \\neix, \\nvii, and \\cvi\\ and their column densities. Combining these non-detected lines with the average \\ovi\\ column density, we have also derived an upper limit to the temperature of the \\ovi-bearing gas.  The X-ray measurements of the \\ovi\\ systems with different physical properties in principle can reveal the ionization mechanism of the intervening gas.  While most observed \\ovi\\ absorbers could arise in interfaces between the low- and high-ionization phases (e.g., \\citealt{dan08, tri08}), compared to the velocity well-aligned simple \\ovi\\ absorbers, the dynamically complex \\ovi\\ systems provide more direct evidence for the multiphase nature of the absorbers (e.g., \\citealt{tri01, tri08, shu03, sem04, sav05}).  Cosmological simulations of large-structure formation also predict that the \\ovi\\ systems with larger equivalent widths are more likely to be collisionally ionized \\citep{cen01, fang01}. If the \\ovi\\ absorbers utilized in this work are indeed multiphase and trace the shock-heated IGM, the hot \\ovii-bearing gas is expected to be either surrounding the cool absorbers (traced by the absorption lines of Ly$\\alpha$, \\cii, \\siii, etc.)  or itself containing the observed \\ovi\\ or a combination of both. The column densities of \\ovii\\ (and/or \\oviii) are expected to be at least an order of magnitude more than that of \\ovi.  But the current X-ray observations only constrain the $N_{\\rm OVII}/N_{\\rm OVI}\\lsim10$ on average (Tables~\\ref{tab:ovi} and \\ref{tab:upper}), which does not strongly favor the collisional ionization scenario of the \\ovi\\ absorbers for shock temperatures $T>10^{5.7}$ K \\citep{fur05}.  However, these data still cannot rule out the photoionization scenario, in which $N_{\\rm OVII}/N_{\\rm OVI}$ is expected to be $\\lsim3$ (see Fig.~6 in \\citealt{tri01} and \\citealt{fur05}).  There are several possibilities or their combinations that may cause the non-detections of the X-ray lines:  1. Some of the \\ovi\\ systems may be photoionized or partly photoionized.  The purely photoionized systems may not contribute to the baryon inventory (\\S~\\ref{sec:intro}) since they could have already been counted through Ly$\\alpha$ absorbers.  Since the density of the IGM is very low, photoionization could play an important role in producing the highly ionized oxygen species (especially \\ovi \\footnote{Higher ionization stages are harder to produce by photoionization, but we note that \\citet{hel98} have argued that in the IGM, even \\ovii\\ and \\oviii\\ are predominantly \\emph{photoionized} by the X-ray background.}) in the WHIM besides the gravitational shocks (e.g., \\citealt{chen03, tri08}). In this case, the $N_{\\rm OVII}/N_{\\rm OVI}$ is expected to be lower than that in a solely collisionally ionized gas.  Simulations show that, while the EWs of \\ovi\\ and \\ovii\\ are fairly well correlated in the putative WHIM in a broad overdensity range ($\\delta\\sim10-100$), for the \\ovi-absorbers with $EW$(\\ovi)$>34$ m\\AA, there are $\\lsim20\\%$ of the associated \\ovii-absorbers with $EW$(\\ovii)$>2$ m\\AA\\ (see Figs.~13-15 in \\citealt{chen03} and Table~\\ref{tab:upper}).  2. Thermal properties of the WHIM may vary among different systems and sight lines. The \\ovii\\ and/or \\oviii\\ absorptions are expected to be strong in some systems and weak in others, but the ``strong'' X-ray absorption signals that originate from the shock-heated gas could have been diluted in the stacked spectra.  3. The characteristic temperature of the hot gas may be too high (e.g., $T>3\\times10^6$ K, as observed in the intragroup and intracluster medium) to produce observable \\ovii\\ or \\oviii.  In this case, the interface between hot and cool IGM could still be present to be responsible for the observed \\ovi\\ absorbers, but the average temperature upper limit ($\\log[T{\\rm (K)}]<5.7$) derived from the ratio of $N_{\\rm OVII}/N_{\\rm OVI}$ (\\S~\\ref{sec:results}) only applies to the vicinity of the \\ovi-bearing gas. This may be the case for the $z=0.0553$ absorber along the PKS 2155-304 sight line where the \\oviii\\ line was detected but the \\ovii\\ was not (\\S~\\ref{sec:intro} and references therein).  4. Lastly, the \\ovi\\ absorbers might only trace the warm part of the WHIM at temperatures $T\\lsim3\\times10^5$ K, and the hot ($T\\sim10^6$ K) IGM is not necessarily always co-located with the warm gas. If this is the case, because of the small ionization fraction of \\ovi\\ at high temperatures, the \\ovi\\ absorption lines produced in the hot gas might not have been detected along most of the sight lines. The possible large dispersion velocity of the hot gas could also make the \\ovi\\ absorption feature shallower and harder to detect.  The high sensitivity of the Cosmic Origins Spectrograph, which is scheduled to be installed on {\\sl HST} in spring 2009, will provide a great opportunity to search for such weak \\ovi\\ absorbers (at $z>0.12$) and to examine this possibility.  The results obtained in this work also have important implications regarding future observations of the WHIM with the current and the future X-ray observatories.  With the still limited spectral resolution [even in the proposed {\\sl International X-ray observatory} (IXO), $E/\\Delta E\\sim3,000$, which is still $\\sim10-15$ times less than those of current far-UV instruments], it will be challenging to conduct a blind search for the X-ray absorption line ``forest'' produced in the WHIM. A more effective strategy is still to use the identified \\ovi\\ and/or Ly$\\alpha$ absorbers as references, as implemented in this work (also see \\citealt{chen03, dan05, dan08}). However, if the six sight lines (Table~\\ref{tab:ovi}) used in this work fairly sample the WHIM filaments over the whole sky, the gas only contains $\\lsim10^{15}~{\\rm cm^{-2}}$ of \\ovii\\ in column density and the \\ovii\\ K$\\alpha$ EW $\\lsim2.5$ m\\AA\\ in absorption on average (Table~\\ref{tab:upper}). To detect a weak line with EW$\\sim2$ m\\AA\\ in an AGN spectrum with a flux of $4\\times10^{-12}~{\\rm ergs~s^{-1}~cm^{-2}~keV^{-1}}$ around the line centroid, our 1,000-run simulations for an exposure of 50 Ms with the {\\sl Chandra} ACIS-LETG only yield 375 detections at $EW/\\Delta EW\\gsim3\\sigma$ significance level. In comparison, for the {\\sl IXO} with $E/\\Delta E\\sim3,000$ and effective area of $\\sim3,000~{\\rm cm^2}$, our simulations indicate that a 100 ks observation will have a 90\\% probability of detecting such a weak line at $\\gsim3.5\\sigma$ significance level."
    },
    "0812/0812.4817_arXiv.txt": {
        "abstract": "We report the results  of 3D spectroscopic observations  of Mrk 493 (NLS1 galaxy) with the integral-field spectrograph MPFS of the SAO RAS 6-m telescope. The difference in the slope of the optical continuum emission intensity across the nucleus part { and an extensive continuum emission region}  is detected. { The emission in  lines (H$\\alpha$, H$\\beta$, [OIII], etc.) coincides with a composite nuclear region: an AGN plus a circum-nuclear star-forming ring observed in the HST UV/optical images. The [SII] emission region tends to be up to 1kpc around the center. The H$\\alpha$ and H$\\beta$ could be decomposed into three components (broad $\\sim$ 2000 km/s. intermediate $\\sim$ 700 km/s and narrow $\\sim$ 250 km/s). We found that  width  ($\\sim$ 750 km/s) of the Fe II  lines correspond to the intermediate component, that may indicate a non-BLR origin of the Fe II lines, or that a large fraction of the Fe II emission arise in the outher parts of the BLR.  }  The weak broad component detected in the H$\\alpha$, H$\\beta$ and He$\\lambda$4686  may come from the unresolved central BLR, { but also  partly produced by violent starburst in the circum-nuclear ring. Moreover, diagnostic diagrams clearly show presence of the HII regions (not a Sy 1 nucleus) in the NLR of Mrk 493.} ",
        "introduction": "Mrk 493, a narrow-line Seyfert 1 galaxy (NLS1), is known as an Active Galactic Nucleus with a strong Fe II emission (Osterbrock and Pogge 1985, Crenshaw et al. 1991). NLS1s have been recognized as a distinct type of Seyfert nuclei. The optical emission-line properties of NLS1s can be summarized as  (e.g., Osterbrock \\& Pogge 1985, Komossa 2008): (i) The Balmer lines are only slightly broader than the forbidden lines such as the [O III] $\\lambda$5007 (typically less than 2000 km s$^{-1}$). (ii) The [O III] $\\lambda$5007/H$\\beta$ intensity ratio  is smaller than 3. (iii) They present the  strong Fe II emission  which are often seen in Sy 1s but generally not in Sy 2's. Moreover, the X-ray spectra of NLS1's are very steep (e.g. Boller et al. 1996, Leighly 1999ab) and highly variable (Boller et al. 1996; Leighly 1999a; Gliozzi et al. 2007). Also, it was found that NLS1s have less massive black hole masses than Sy 1s (Mathur 2000, Wang \\& Zhang 2007) and could be young Sy 1s (Botte et al. 2004). Even though NLS1s have been known for almost more than 20 years, it is still not clearly understood what NLS1s are in the context of the current AGN unified model (Nagao et al. 2001).   One of the interesting observational facts connected with NLS1 is that in their spectra a strong Fe II emission is present. It was shown by  Boroson \\& Green (1992) that  a strong anti-correlation between the strengths of the ${\\rm [OIII]}$ and Fe II line intensities exists in the optical  spectra,  where NLS1s are showing the strongest Fe II and weakest ${\\rm [O III]}$ emission. The observed line widths and absence of the forbidden emission suggests that the Fe II lines are formed in the dense broad-line region (BLR), but photo-ionization models cannot account for all of the Fe II emission. The 'Fe II discrepancy' remains unsolved, though models which consider non-radiative heating with an overabundance of iron are promising (Joly 1993, Collin \\& Joly 2000). Also Mathur (2000) proposed that the strong Fe II emission observed in NLS1s may be connected with their large accretion rates and that a collisional ionization origin of Fe II is possible.  Here we report 3D spectroscopic observations of Mrk 493. The aims of the observations were  (i) to investigate the characteristics of the Fe II emitting region and possible their connection with other emitting regions (H$\\alpha$, H$\\beta$, [OIII], ...) and (ii) to map the circum-nuclear emitting gas in this type of AGN. In this paper we accepted a distance to Mrk 493 of 125.2 Mpc (for H$_0=75$ km\\,s$^{-1}$\\,Mpc$^{-1}$), where 1$''$ on the sky corresponds to 610 pc at the galaxy distance. ",
        "conclusions": "  (i) {The continuum and [SII] line emitting regions in Mrk 493 seem to be very extensive ([SII] up to 1 kpc, and continuum $>$1kpc), and they are larger than the  [OIII], Fe II and Balmer emission line regions.}   (ii)  It seems that the strong Fe II emission is not coming from the BLR ($W_{BLR}\\sim$2000 km/s), but from a region with velocity dispersion around 800 km/s. { This is in an agreement with some earlier investigation given by Popovi\\'c et al. (2004) and recent ones given by Kuehn et al. (2008) and Hu et al. (2008) that a large fraction of the Fe II emission may come from  an ILR.}  (iii)  {  There is no   changes of the continuum and emission  lines (instead He II$\\lambda$4686)  between two observations (in the three-year period), that may also indicate non-BLR origin of  lines. Otherwise, the BLR in Mrk   493 should be a): very extensive $\\sim$ 1pc and/or b) very stable.}  (iv) The source of ionization of the narrow Balmer and [OIII] line emitting regions in Mrk 493 seems to not  be the nucleus, but the star-forming processes;  (v) { The position  of the line and continuum regions coincides with a composite nuclear region: an AGN plus a circum-nuclear star-forming ring observed in the HST UV/optical images, but it seems that the emitting source of the narrow line and continuum is probably the circum-nuclear star-forming ring, but here remains a question about the mechanism of the Fe II production in a such  ring.}"
    },
    "0812/0812.1329_arXiv.txt": {
        "abstract": "We investigate the optical emission-line flux ratios of narrow-line regions, in order to determine whether the formation of AGN jets requires specific accretion conditions. We find that bright compact radio galaxies, which are powerful radio galaxies in the early stage of the jet activity, exhibit systematically larger flux ratios of [O{\\sc i}]$\\lambda$6300/[O{\\sc iii}]$\\lambda$5007 and smaller flux ratios of [O{\\sc iii}]$\\lambda$5007/[O{\\sc iii}]$\\lambda$4363 than radio-quiet (RQ) Seyfert 2 galaxies. Comparing the observed line ratios with photoionization models, it is found that the difference in the flux ratio of low- to high-ionization lines (e.g., [O{\\sc i}]$\\lambda$6300/[O{\\sc iii}]$\\lambda$5007) can be well understood by the difference in the spectral energy distribution (SED) of ionizing sources. Powerful young radio-loud (YRL) AGNs favor SED without a strong big blue bump, i.e., a radiatively inefficient accretion flow (RIAF), while RQ AGNs are consistent with the models adopting SED with a strong big blue bump, i.e., a geometrically thin, optically thick disk. These findings imply that the formation of powerful AGN jets requires the accretion disk with harder ionizing SED (i.e., a RIAF). We discuss the obscuring structure of YRL AGNs as a plausible origin of the difference in flux ratios of [O{\\sc iii}]$ \\lambda$5007/[O{\\sc iii}]$\\lambda$4363. ",
        "introduction": "The formation of relativistic jets is one of the fundamental issues in active galactic nuclei (AGNs) physics. The question how the jets form is relevant to supermassive black hole (BH) cases as well as X-ray binaries (XRBs; see reviews by Fender 2003). Thanks to their short dynamical timescale, XRBs in various states have been observed in detail, and it is found that individual sources occupy particular accretion state with various X-ray luminosity (e.g., Fender et al. 2004; Remilland \\& McClintock 2006). These observations suggested that at the low luminosity (i.e., the low mass accretion rate) the jet power is a monotonic function of the accretion luminosity. On the other hand, the jet power may change abruptly by a large factor, following transitions between the soft state (a spectra characterized by the black body components) and hard state (a spectra characterized by the power-law components) at the high luminosity (i.e., the high mass accretion rate). By analogy to XRBs, it is expected that the power of AGN jets also depends on the mass accretion rates in a similar way. For example, Rawlings \\& Saunders (1991) found that for the extended radio galaxies whose projected linear size ($LS$) is larger than $10$ kpc the jet power is proportional to the optical narrow line luminosity of NLRs which are photoionized by the nuclear continuum radiation (see also Willott et al. 1999). Ghisellini \\& Celotti (2001) claimed that the less powerful radio galaxies such as Fanaroff-Riley type I radio galaxies (FR Is) exhibit the low accretion rate, while the powerful radio galaxies such as Fanaroff-Riley type I\\hspace{-.1em}I radio galaxies (FR I\\hspace{-.1em}Is) show the high accretion rate. This work is consistent with a scenario (e.g., Marchesini, Celotti, \\& Ferrarese 2004; Wold et al. 2007; Hardcastle et al. 2007) that FR Is have a radiatively inefficient accretion disk (RIAF) (e.g., Narayan \\& Yi 1995), while FR I\\hspace{-.1em}Is have a standard accretion disk (e.g., Shakura \\& Sunyaev 1973).   In investigating the disk-jet connection in AGNs, however, we should keep in mind the relatively longer life time of the extended radio sources, i.e., FR Is and FR I\\hspace{-.1em}Is. The reasons are as follows. The time-scale of the transition to the different physical states of the accretion disk is $\\Delta t({\\rm XRBs})\\approx$ $10^{-2}-10^{-1}\\, {\\rm yr}$ (a few days $-$ a few weeks change) for XRBs (e.g., Fender et al. 2004). Assuming $M_{\\rm BH}({\\rm XRBs})=10M_{\\odot}$ and $M_{\\rm BH}({\\rm AGNs})= 10^{8}M_{\\odot}$, the transition time-scale for AGNs is $\\Delta t({\\rm AGNs})\\approx 10^{5-6}\\, {\\rm yr}$ because of $\\Delta t \\propto M_{\\rm BH}$. Thus, the physical states of the accretion disk may change in less than the typical age of extended radio galaxies, i.e., $t_{\\rm age}\\approx 10^{7-8}$ yr (e.g., Parma et al. 1999; Bird, Martini, \\& Kaiser 2008). In other words, the observed relation between the jet power and the accretion power for the extended radio galaxies may not be proving the direct disk-jet connection in AGNs. Thus, in order to investigate the disk-jet connection in AGNs more directly, it is essential to examine the physical states of the accretion disk in younger radio galaxies.  The young and compact radio galaxies such as compact symmetric objects (CSOs; $LS <$ 1kpc) and medium-size symmetric objects (MSOs; $LS=1-10$ kpc) have been discovered by several authors (e.g., Wikinson et al. 1994; Fanti et al. 1995; Readhead et al. 1996). Among these compact radio sources, the bright objects show the FR I\\hspace{-.1em}I-like morphology along with hot spots (reverse shocks), which are considered to be a signature of supersonic expansions (Owsianik, Conway \\& Polatidis 1998; Taylor et al. 2000; Tschager et al. 2000; Giroletti et al. 2003; Polatidis \\& Conway 2003; Gugliucci et al. 2005; Nagai et al. 2006; Gugliucci et al. 2007; Luo et al. 2007). In contrast, the faint compact radio sources do not show hot spots (e.g., Kunert-Bajraszewska et al. 2005; Giroletti 2007) and some of them show the FR I-like morphology (Giroletti 2007). Since the age of CSOs and MSOs is less than $\\sim 10^{5}\\,{\\rm yr}$, they are recognized as newly born young AGN jets (e.g., Fanti et al. 1995; Readhead et al. 1996, O'Dea \\& Baum 1997; Stanghellini et al. 1998; Snellen et al. 2000; Dallacasa et al. 2000; Orienti et al. 2007). Thus, CSOs and MSOs are adequate targets to examine the disk-jet connection in AGNs since the time lag between the optical and radio activity is relatively small ($\\Delta t({\\rm AGNs}) \\lesssim 10^{5-6}\\, {\\rm yr}$). In this paper, we will use bright CSOs and MSOs (hereafter powerful young radio-loud [YRL] AGNs) to investigate the disk-jet connection in AGNs.  In order to probe whether the formation of AGN jets requires any specific accretion conditions, it is essential to determine the physical states of the accretion disk of powerful YRL AGNs. To this end, it can be useful to examine a systematic difference in the SEDs produced by accretion disks of powerful YRL AGNs and radio-quiet (RQ) AGNs. However, it is challenging to derive the intrinsic SEDs of powerful YRL AGNs from the hard X-ray observations, because of the contamination of AGN jets and the obscuration of the accretion disk by the dusty torus. Alternatively, the optical emission-line diagnostics of narrow-line region (NLR) can be used to distinguish the physical states of the accretion disk, since the NLR is thought to be photoionized by the nuclear radiation for RQ Seyfert galaxies (e.g., Yee 1980; Shuder \\& Osterbrock 1981; Evans et al. 1999) and even for extended radio galaxies (e.g., Villar-Martin et al. 1997; Nagao et al. 2006). However, very little is currently known about the optical emission-line properties of powerful YRL AGNs. By examining the difference of optical line ratios between powerful YRL AGNs and the extended radio galaxies, Morganti et al. (1997) found that powerful YRL AGNs show the characteristics very similar to that of extended radio galaxies. However, it is not clear how the physical state of accretion disks (especially SEDs of accretion disks) is different between the powerful YRL AGNs and the RQ AGNs. Answering this question will provide crucial information on the disk-jet connection in AGNs.  The main goal of this paper is to investigate whether there is the difference in SEDs between the powerful YRL AGNs and Seyfert 2 galaxies (i.e., RQ AGNs), by utilizing the optical emission-line diagnostics of NLRs. We collected samples of powerful YRL AGNs and RQ AGNs (\\S 2), and generated various photoionization models with different input SEDs (\\S 3). In \\S 4, we compared the observational data with models. On the basis of our findings, Discussion on accretion conditions for the formation of powerful AGN jets is presented in \\S5, and Conclusions follow in \\S 6. ",
        "conclusions": "We examine the optical narrow emission-line flux ratios of NLRs, in order to determine whether the formation of powerful AGN jets requires specific accretion conditions. A sample of powerful YRL AGNs, which are at the early stage of the jet activity, and radio-quiet Seyfert 2 galaxies are selected. We summarize our main results.  \\begin{enumerate} \\item Powerful YRL AGNs exhibit larger flux ratios of [O{\\sc i}]$\\lambda$6300/[O{\\sc iii}]$\\lambda$5007 than RQ Seyfert 2 galaxies. By comparing the observed line ratios with the photoionization model calculations, we find that the difference between powerful YRL AGNs and Seyfert 2 galaxies can be well understood by the difference in SEDs of ionizing radiation from the accretion disk. The powerful YRL AGNs favor SED without a strong BBB such as the ADAF, while the line ratios of Seyfert 2 galaxies are consistent with the models adopting SED with a strong BBB (a geometrically thin, optically thick disk). In contrast, the observed difference in the line ratios is difficult to explain by the difference in the contribution of shocks caused by the AGN jets.  \\item The line ratio [O{\\sc iii}]$\\lambda 5007$/[O{\\sc iii}]$ \\lambda$4363 of powerful YRL AGNs is systematically smaller than that of RQ AGNs. If we adopt a scenario that [O{\\sc iii}]$\\lambda$4363 originates in the dense gas clouds, which are located at the inner regions compared with [O{\\sc iii}]$\\lambda$5007, it is inferred that the obscuring torus around powerful YRL AGNs would be thinner than that around Seyfert galaxies.  \\end{enumerate}"
    },
    "0812/0812.4951_arXiv.txt": {
        "abstract": "The findings of a nine orbit calibration plan carried out during {\\it HST} Cycle 15, to fully determine the NICMOS camera 2 (2.0\\micron) polarization calibration to high accuracy, are reported. Recently Ueta et al. and Batcheldor et al. have suggested that NICMOS possesses a residual instrumental polarization at a level of $1.2-1.5\\%$. This would completely inhibit the data reduction in a number of GO programs, and hamper the ability of the instrument to perform high accuracy polarimetry. We obtained polarimetric calibration observations of three polarimetric standards at three spacecraft roll angles separated by $\\sim60\\degr$. Combined with archival data, these observations were used to characterize the residual instrumental polarization in order for NICMOS to reach its full potential of accurate imaging polarimetry at $p\\approx1\\%$. Using these data, we place an 0.6\\% upper limit on the instrumental polarization and calculate values of the parallel transmission coefficients that reproduce the ground-based results for the polarimetric standards. The uncertainties associated with the parallel transmission coefficients, a result of the photometric repeatability of the observations, are seen to dominate the accuracy of $p$ and $\\theta$. However, the updated coefficients do allow imaging polarimetry of targets with $p\\approx1.0\\%$ at an accuracy of $\\pm0.6\\%$ and $\\pm15\\degr$. This work enables a new caliber of science with {\\it HST}. ",
        "introduction": "Polarimetry is a powerful observational tool that augments and complements the capabilities of imaging, photometry and spectroscopy. While the latter allow the determination of spatial distributions, chemical composition and dynamics, polarimetry lets us observe the nature of magnetic fields, object orientation, scattering, and the properties of interstellar particles in general. It also allows us to probe the nature of emission mechanisms (e.g., synchrotron/thermal), and to investigate the geometry of unresolved sources.  The value of polarimetry to elucidate the physical nature of light-scattering particles and their environments in astrophysical systems has been repeatedly demonstrated in, for example, active galactic nuclei \\citep{1985ApJ...297..621A,1995ApJ...452L..87C}, galactic and extra-galactic magnetic fields \\citep{1995RvMA....8..185W}, radio galaxies \\citep{1998A&A...331..475F,2001A&A...366....7V}, and circumstellar disks \\citep{2001ApJ...553L.189K}. However, because of ``beam depolarization'', results derived from polarimetric imaging observations may be biased when obtained at low spatial resolution. For example, compact (but spatially resolved) objects that exhibit circular symmetry or bi-conical structures, like active galactic nuclei (AGN) and stars surrounded by orbiting circumstellar dust, contain polarized elements that partially reduce, if not fully cancel, each other when spatially averaged. If the object is unresolved, then the polarization vectors sum to {\\it zero}; in any case, beam depolarization will always underestimate the intrinsic polarization.  NICMOS is the {\\it only} near infrared (NIR) instrument capable of the high resolution, high fidelity polarimetry needed to examine the scattering geometry and materials, in detail, for many types of astronomical objects. To date, there have been a multitude of studies carried out using NICMOS direct imaging polarimetry \\citep[e.g.][]{2000ApJ...544..269C,2000MNRAS.313L..52T,2000ApJ...536L..89S,2002ApJ...576..429S,2002ApJ...574...95S,2003AJ....126..848S, 2005ApJ...634.1146M,2006ApJ...642..339S} and there are now more ambitious programs to measure the characteristics of AGN and regions within circumstellar protoplanetary debris disks, that typically exhibit polarizations of less than 5\\% \\citep{2004MNRAS.350..140S,2005AAS...207.1013S}. However, there have been recent reports of a systematic residual instrumental polarization of $\\sim1.5\\%$ in some objects \\citep{2005AJ....129.1625U}, which although acceptable for highly polarized targets, seriously compromises studies of low polarization objects. A comprehensive study of this residual polarization has been carried out by \\citet[][hereafter B06]{2006PASP..118..642B}. Using all of the available standard polarization data in the archives, B06 indeed find a residual excess polarization of $\\approx1.2\\%$ in the NIC2 camera. The extent of this excess remains unknown for NIC1.  Unfortunately, prior to the new calibration reported in this paper, since the advent of the NICMOS cooling system, NCS \\citep{2004AdSpR..34..543S}, installed during SM3B in 2002, observations of only one polarized and one unpolarized standard have been obtained, and each at only two celestial orientation angles. Hence, it has been impossible, with those limiting data, to entirely characterize the NICMOS residual polarization. Polarization measurements at three well separated roll angles are required to remove the dependence of the measured Stokes parameters on the relative transmission of each of the filters.  \\begin{deluxetable}{lcccc} \\tablecaption{Details of Observations \\label{tab:obs}} \\tablewidth{0pt} \\tablehead{ \\colhead{Target}&\\colhead{DATE-OBS}&\\colhead{ORIENTAT}&\\colhead{Camera}&\\colhead{T (s)} } \\startdata VR 84c          & 2007-05-03 & 182.2\\degr & NIC1 &  3.6 \\\\ & 2007-08-03 & 299.3\\degr &      &      \\\\ & 2006-12-15 &  59.2\\degr &      &      \\\\ & 2007-05-03 & 181.4\\degr & NIC2 & 14.0 \\\\ & 2007-08-03 & 298.6\\degr &      &      \\\\ & 2006-12-15 &  58.4\\degr &      &      \\\\ VSS VIII-13     & 2006-08-04 & 180.3\\degr & NIC1 &  1.6 \\\\ & 2006-06-25 & 114.3\\degr &      &      \\\\ & 2007-04-11 &  54.3\\degr &      &      \\\\ & 2006-08-04 & 179.6\\degr & NIC2 &  7.0\t\\\\ & 2006-06-25 & 113.6\\degr &      &      \\\\ & 2007-04-11 &  53.6\\degr &      &      \\\\ HD331891        & 2006-08-27 & 273.8\\degr & NIC1 &  1.6 \\\\ & 2006-07-01 & 338.4\\degr &      &      \\\\ & 2007-04-15 &  38.3\\degr &      &      \\\\ & 2006-08-27 & 273.1\\degr & NIC2 &  7.0 \\\\ & 2006-07-01 & 337.6\\degr &      &      \\\\ & 2007-04-15 &  37.5\\degr &      &      \\\\ \\enddata \\tablecomments{Each observation was performed at three different roll angles and for the POL0, POL120 and POL240 polarizers. In total, NIC2 was dithered around 36 pointings per orbit, with 12 pointings per polarizer. The DATE-OBS parameter gives the observation date in YYYY-MM-DD format. ``ORIENTAT'' refers to the celestial position angle of the image +y axis (degrees east of north). ``T'' gives the exposure times in seconds.} \\end{deluxetable}  This paper reports the findings of a nine orbit {\\it HST} calibration plan (Cycle 15), to fully investigate the residual polarization in NICMOS. NICMOS carries two cameras with polarimetry optics designated NIC1 and NIC2 (a third camera for wide-field and grism imaging has no polarimetry capability). NIC1 provides an 11''$\\times$11'' field-of-view (FOV) with an image scale of 43.1 milli-arcsecs (mas) per pixel and 1.045\\micron~ broadband (0.475\\micron~ FWHM; $R=\\Delta\\lambda/\\lambda = 45\\%$) polarizing filters. NIC2 (the primary focus of this paper) provides a 19\\farcs3$\\times$19\\farcs2 FOV with 75.8 mas pixels and a medium bandwidth polarimetric passband of 1.994\\micron~ (0.202\\micron~ FWHM; $R=10\\%$). Both optical channels critically sample the respective point spread functions (PSFs) in their polarimetric passband. For additional details of the NICMOS instrument see \\cite{1998ApJ...492L..95T} and the NICMOS Instrument Handbook \\citep{2007hsti.book.....B}. The two cameras have non common path optics from the instrument's field divider mirror assembly to their respective re-imaged focal plane detectors. As polarimetric imaging in the two cameras is (uniquely) carried out at different wavelengths, there is no reason to expect the same level of residual instrumental polarization in each camera. With this in mind we carried out observing plans to push the NIC2 calibration to its instrumental limit, given the intrinsically superior polarimetric performance of the NIC2 camera over NIC1 \\citep{2000PASP..112..983H}. Due to the nature of NIC1, we confine the results from this camera to Appendix~\\ref{app:a}. In \\S~\\ref{obs} we describe the observing strategy employed and in \\S~\\ref{red} we explain the data reduction procedures. A detailed examination of the photometry is completed in \\S~\\ref{phot}. The results and methods are presented in \\S~\\ref{ptheta} before being discussed in \\S~\\ref{dis} and concluded in \\S~\\ref{cons}. A recommended observing plan for NICMOS imaging polarimetry with NIC2 is included in \\S~\\ref{dis2}. ",
        "conclusions": "\\label{cons}  In a non-ideal imaging polarimeter, such as NICMOS, it is essential to observe several polarimetric standards, at three well separated position angles through each polarizer, in order to full characterize the instrumental polarization. The additional roll angles remove the uncertainties in the (unknown) parallel transmission coefficients and allow rotation of any instrumental polarization with respects to the equatorial frame. Using this technique we have placed an upper limit to the NIC2 instrumental polarization of 0.6\\%. With a known value for the Stokes $I$ parameter, the instrumental polarization can be transformed into the $Q,U$ plane and be subtracted. New parallel transmission coefficients can then be determined numerically by comparing $p$ and $\\theta$ with that of the apriori well-determined calibration standards. Following this approach we have determined the $t_0$ and $t_{120}$ coefficients to be $0.883\\pm0.004$ and $0.837\\pm0.004$ respectively. As with the previous calibrations of the NIC2 polarimeter, we held $t_{240}$ constant at 0.9667. The $t_{120}$ coefficient is consistent with the previous calibration, but our knowledge of $t_0$ has been improved, resulting in a change in its previously determined value by $\\sim0.5\\%$. Such a small change in $t_0$ does not warrant the re-analysis of previous NIC2 imaging polarimetry data of {\\it highly} polarized ($p>5\\%$) targets, but is significant for targets with intrinsically low polarization fractions.  As we use 13 determinations of $t_k$, we are able to assign $1\\sigma$ uncertainties to the calibration coefficients. Propagating these uncertainties through the polarimetric analysis, we find that they dominate all other sources of error. Applying these adjusted values of $t_k$ to this (and the archived) calibration data, we find that NIC2 is now capable of confidently detecting polarizations at a level of $\\approx1.0\\%$. The uncertainties associated with such measurements are $\\pm0.6\\%$ and $\\pm15\\degr$ in $p$ and $\\theta$, for sources at the $\\approx1\\%$ level of intrinsic polarization. This is the first time that such a level of accuracy has been achieved with the NICMOS polarization calibration. This improved calibration opens a new domain for observational investigations with {\\it HST} by enabling very high precision polarimetry of intrinsically very low polarization sources."
    },
    "0812/0812.3576_arXiv.txt": {
        "abstract": "While the PAMELA collaboration has recently confirmed the cosmic ray positron excess, it is interesting to review the effects of dark matter (DM) subhalos on the predicted antimatter signals. We recall that, according to general subhalo properties as inferred from theoretical cosmology, and for DM with constant annihilation cross section, the enhancement cannot be $\\gtrsim$ 20 for the antimatter yield. This bound is obviously different from that found for $\\gamma$-rays. We also recall some predictions for supersymmetric benchmark models observable at the LHC and derived in the cosmological N-body framework, showing in the meantime the existing discrepancy between profiles derived from N-body experiments and the current observations of the Milky Way. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.0269_arXiv.txt": {
        "abstract": "We investigate the gas-phase and grain-surface chemistry in the inner 30\\,AU of a typical protoplanetary disk using a new model which calculates the gas temperature by solving the gas heating and cooling balance and which has an improved treatment of the UV radiation field. We discuss inner-disk chemistry in general, obtaining excellent agreement with recent observations which have probed the material in the inner regions of protoplanetary disks. We also apply our model to study the isotopic fractionation of carbon. Results show that the fractionation ratio, $^{12}$C/$^{13}$C, of the system varies with radius and height in the disk. Different behaviour is seen in the fractionation of different species. We compare our results with $^{12}$C/$^{13}$C ratios in the Solar System comets, and find a stark contrast, indicative of reprocessing. ",
        "introduction": "Studying the chemistry of the inner regions of protoplanetary disks (PPDs) is a step into the unknown. Until recently observations of molecules in PPDs could not probe the warmer, inner regions of the disk where planets form. Detections of CO, HCO$^+$, H$_2$CO, C$_2$H, CS, SiO, HNC, CN and HCN \\citep{dut97,kas97,qi03,thi04,sem05} at millimetre wavelengths can only tell us about the physics and chemistry in the cold conditions at radii greater than $\\sim$50\\,AU because of the limits of millimetre-wave sensitivity and spatial resolution. We have to go to the infrared to investigate the hotter material, and there comparatively few molecules have been seen. Initially H$_2$, CO \\citep{naj03,bri03,bla04} and H$_2$O \\citep{car04} were detected. Recently, the \\textit{Spitzer Space Telescope} has added OH, C$_2$H$_2$, CO$_2$ and HCN to the tally \\citep{car08,lah06} and a few detections of very large polycyclic aromatic hydrocarbon (PAH) molecules have been made \\citep{gee06}. The Keck Interferometer has also brought us detections of hot HCN and C$_2$H$_2$ \\citep{gib07}. The regions of the disk close to the star are of immense interest because they are where terrestrial (and larger) planets begin to form, as well as comets. Molecules in the inner 30 or 40\\,AU could be the building blocks of the kinds of pre-biotic and biotic molecules we see in the Solar System today.  In the absence of observational evidence, we are left with theoretical work. This is valuable in its own right for understanding the processes in these regions, and for predicting observations for future proposals and even future technologies, such as ALMA\\footnote{The Atacama Large Millimeter Array, due for completion in 2012 (\\url{www.alma.info})}, which will be able to probe these hidden inner disks. Previous modelling attempts of these very inner regions of PPDs (R$<$10\\,AU) have been few \\citep{mar02,mil03,ilg04,ilg06,agu08}, but have shown that these inner regions of PPDs are rich in molecules, including some complex molecules, such as benzene \\citep{woo07}.  In this paper we present chemical models of a protoplanetary T Tauri-stage disk and a Minimum Mass Solar Nebula (MMSN), paying particular attention to the fractionation of carbon-bearing species. The fractionation in a particular molecule is defined as the abundance of the $^{12}$C-bearing variant over the $^{13}$C-bearing variant\\footnote{When there is more than one $^{13}$C per molecule, the fractionation ratio is taken to be [$^{12}$C]/[$^{13}$C]}. Fractionation in carbon gives us a method of labelling: identifying, for instance, different regions of the disk and also identifying the different evolutionary processes involved in the formation of the disk -- each process leaves its isotopic signature. Thus initially we study a typical protoplanetary disk, the physical aspects of which are explained in Sect.~\\ref{sec:modelphys}. We include a treatment of gas heating and cooling, which is expanded upon fully in Appendix~\\ref{sec:heatcool}. In Sect.~\\ref{sec:modelchem} we give a description of the chemical network, and the various chemical and photo-fractionation mechanisms including isotope-exchange reactions. Finally, we present our findings on carbon fractionation in protoplanetary disks, and set our results in both a Solar System and interstellar context. ",
        "conclusions": "We have shown by means of a chemical model of a protoplanetary disk that the fractionation ratio of carbon, $^{12}$C/$^{13}$C, varies according to position in the disk. The fractionation ratio is governed by temperature, which affects the rate and direction of chemical exchange reactions, and incident UV radiation, which affects self-shielding molecules such as CO. This produces a picture of the fractionation in a disk which is stratified. Certainly in the upper region of the disk, fractionation timescales are faster than mixing timescales, meaning that this stratified picture should persist if mixing were to be considered in the model. We also considered a Minimum Mass Solar Nebula model in which the chemistry in our inflowing gas packets has longer to evolve. Extremely good agreement was seen between observations of Solar System comets and the fractionation ratio in ices in the midplane of the disk model. In general, our chemical model has excellent agreement with recent observations of T Tauri disks, both in terms of chemical abundance and location in the disk."
    },
    "0812/0812.4411_arXiv.txt": {
        "abstract": "{ We present the results of a combined investigation of the spectral and kHz QPO evolution around the Z-track in GX\\th 5-1 based on high-quality {\\it Rossi-XTE} data. In spectral analysis, we find that the Extended ADC emission model provides very good fits to all of the spectra, and the results point clearly to a model for the nature of the Z-track in this source, in agreement with previous results for the similar source GX\\th 340+0. In this model, at the soft apex of the Z-track, the mass accretion rate $\\dot M$ is at its minimum and the neutron star has its lowest temperature; but as the source moves along the normal branch, the luminosity of the Comptonized emission increases, indicating that $\\dot M$ increases and the neutron star gets hotter. The measured flux $f$ of the neutron star emission increases by a factor of ten becoming super-Eddington, and we propose that this causes disruption of the inner disk and the formation of jets. In flaring, the luminosity of the dominant Comptonized emission from the accretion disk corona is constant, while the neutron star emission increases, and we propose for the first time that flaring consists of unstable nuclear burning on the neutron star, supported by the agreement between the measured mass accretion rate per unit area $\\dot m$ at the onset of flaring and the theoretical critical value at which burning becomes unstable. There is a striking correlation between the frequencies of the kHz QPO and the ratio of the flux to the Eddington value: $f/f_{\\rm Edd}$, suggesting an explanation of the higher frequency QPO and of its variation along the Z-track. It is well known that a Keplerian orbit in the disk at this frequency corresponds to a position some distance from the neutron star; we propose that the oscillation always occurs at the inner disk edge, which moves radially outwards on the upper normal and horizontal branches as the measured increasing radiation pressure increasingly disrupts the inner disk. ",
        "introduction": "The Z-track sources form the brightest group of low mass X-ray binaries (LMXB) containing a neutron star, with luminosities at or above the Eddington limit. They are characterised by having three distinct branches in a hardness-intensity diagram: the horizontal branch (HB), the normal branch (NB) and the flaring branch (FB) (Hasinger \\& van der Klis 1989) showing that major physical changes take place within the sources but the nature of these has not been understood. It has been widely thought that the physical and spectral changes are driven by change of a single parameter along the Z-track, presumably the mass accretion rate (Priedhorsky et al. 1986), assumed to increase monotonically in the direction HB - NB - FB. However, as discussed below, the evidence for this is rather limited. Moreover the variation of the X-ray intensity does not obviously support this as the intensity does not increase monotonically in the direction HB - NB - FB, but decreases on the normal branch. Arguably the most important feature of the Z-track sources is the detection of radio emission showing jets to be present but only in the upper normal and horizontal branches (Penninx 1989). The presence of jets was dramatically demonstrated by extended radio observations of the Z-track source \\hbox{Sco\\th X-1} (Fomalont et al. 2001) which revealed radio condensations moving away from the source with velocity $v/c$ of 0.45. Thus the Z-track sources uniquely provide the possibility of determining the physical conditions within the sources on the part of the Z-track where radio is detected so telling us the conditions needed for the launching of jets. Apart from this an understanding of the Z-track sources is essential to the basic understanding of LMXB in general. Extensive work has been carried out on the timing properties of Z-track sources (van der Klis et al. 1987; Hasinger \\& van der Klis, 1989) and revealed the existence of quasi-periodic oscillations (QPO) which change along the Z-track. However, analysis of the timing properties has not provided an explanation of the Z-track phenomenon. Spectral analysis is more likely to reveal the nature of the physical changes taking place as directly showing changes in the emission components during the spectral evolution along the Z-track, but spectral studies of the sources have been hindered by lack of agreement over the emission model to be used.  \\subsection{Spectral studies of the Z-track sources}  The nature of X-ray emission in LMXB has been controversial essentially because two radically different types of model for the continuum are capable of describing the spectra in which a dominant Comptonized component and also a thermal component are clearly present. Analysis of the dipping class of LMXB led to the ``Extended ADC'' model in which the dominant component is Comptonized emission from the accretion disk corona (ADC) plus thermal emission from the neutron star (Church \\& Ba\\l uci\\'nska-Church 1995, 2004). Both spectral and timing analysis showed the Comptonizing ADC to be a region extending over a substantial fraction of the inner accretion disk (typically 15\\%) and thin \\hbox{($H/r$ $<<$ 1)} as discussed in more detail below. The other model is the Eastern model comprising multi-colour blackbody emission from the inner accretion disk plus Comptonized emission from a central region close to the neutron star (Mitsuda et al. 1989). Previous spectral fitting of the Z-track sources has mostly been based on use of the Eastern model as summarized below. However, there is now a body of evidence that the Comptonizing region is extended (Sect. 1.2) and is not a small central region. Moreover, when the Eastern model was applied to the Z-track sources, it was difficult to interpret the variation of spectral parameters and a physical explanation was often not given.  An extensive investigation of spectral evolution along the Z-track in several sources using {\\it Exosat} data was carried out by Schulz et al. (1989) and by Schulz and Wijers (1993) using a model with neutron star blackbody emission plus Comptonization. Hasinger et al. (1990) fitted Ginga data on Cyg X-2 with both Western and Eastern models. In the Western model the main change was an increase of blackbody temperature on climbing the normal branch. In the Eastern model, there was a systematic change in the Comptonization y-parameter on the normal branch. Because of the ability of both models to fit, it was not thought possible to distinguish between physical models. Consideration of multi-wavelength data and of the timing properties of the source led the authors to conclude that the mass accretion rate $\\dot M$ increased monotonically in the direction HB - NB - FB although  X-ray intensity decreased on the NB. A possible explanation was that a thickening of the inner disk obscured the X-ray emission. The conclusion about $\\dot M$ led Hasinger et al., following Hasinger (1988) and Lamb (1989), to propose a scenario in which $\\dot M$ was small on the HB so that a thin disk existed together with a magnetosphere in which radio emission was generated. Increasing $\\dot M$ resulted in the magnetosphere being enveloped by a vertically expanding disk at the hard apex, suppressing radio emission. At the soft apex $\\dot M$ exceeds the Eddington value so that strong radiation pressure led to mass ejection above and below the disk. Asai et al. (1994) fitted {\\it Ginga} data on the NB and FB spectra of GX\\th 5-1 with both Western and Eastern models, but using the Eastern model in its original form consisting of disk blackbody plus blackbody from the neutron star. This work concentrated on the emission feature found at 10 keV.  Spectral fitting of Cyg\\th X-2 was carried out by Done et al. (2002) using the Eastern model assuming that the seed photons for Comptonization come from the neutron star or inner disk and so should be described by a blackbody spectrum. It was also assumed that the central Comptonizing cloud would illuminate the disk so that reflection should be present (although in the Z-track sources the height of the inner disk is much greater than that of the neutron star). The spectral fitting indicated an increasing flux of the disk blackbody in the lower normal and flaring branches where it was assumed that $\\dot M$ increased; however, the disk temperature appeared constant so it was argued that the inner disk radius must increase. The Comptonized flux also remained constant when it might be expected to increase and it was not clear what caused the spectral parameters to vary in the way found. Agrawal \\& Sreekumar (2003) used the Eastern model to fit {\\it Rossi-XTE} spectra of GX\\th 349+2 in the NB and FB using the form {\\sc diskbb + comptt}, i.e. disk blackbody plus Comptonization with the seed photons having a Wien distribution. The inner disk temperature first increased on the NB then decreased in flaring; the inner disk radius first decreased then increased, suggesting that the inner disk edge was moving to larger radial positions, and the authors proposed that there was thus an outflow of material . The electron temperature of the Comptonizing plasma first increased on the NB then decreased in flaring while the optical depth of the region first decreased then increased. It was suggested that the initial heating caused a density decrease and reduction in optical depth; on the FB the outflow proposed added density to the hot central corona causing the increase of optical depth on this branch.  Spectral fitting of broadband {\\it BeppoSAX} data was carried out for GX\\th 17+2, GX\\th 349+2 and Cyg\\th X-2 (di Salvo et al. 2000, 20001, 2002). In the case of Cyg\\th X-2, the Eastern model was applied together with a power law to represent a high energy tail, and it was argued that all three sources could be fitted by this model. In Cyg\\th X-2 spectral evolution  consisted of an increase of disk blackbody temperature and decrease of blackbody radius moving from the HB to the NB, suggesting that the inner disk radius was shrinking although an explanation of this was not proposed. There were changes in temperature of the Comptonizing plasma and on the normal branch, the inferred radius of this region fell to 2 - 3 km, incompatible with the neutron star being the source of seed photons unless there was a marked change in geometry from the normally assumed spherical Comptonizing cloud of the Eastern model. D'Ai et al. (2007) investigated the Z-track in Sco\\th X-1 using extensive numbers of observations with {\\it RXTE}. The Eastern model in the form used by di Salvo et al. was used. The disk  blackbody temperature was highest on the flaring branch as expected on the standard assumption that $\\dot M$ is largest here. The electron temperature of the Comptonizing plasma increased around the Z having its highest values in flaring, while at the same time the optical depth decreased; however, the reasons for these changes were not clear.  It is difficult to find consistency in the spectral fitting results and with the difficulty of interpreting the evolution of spectral parameters along the Z-track it is fair to say that there is no general concensus on the nature of the Z-track. Much effort was devoted to producing a theoretical model of of accretion in the inner disk in LMXB which proposed detailed physical changes taking place along the Z-track (Lamb 1989; Miller 1990;  Psaltis et al. 1995). This involves a magnetosphere at the inner accretion disk, and changes in the geometry and extent of the magnetosphere. The model assumes the main element of the Eastern model that is, a small central Comptonizing region. In the present work we take the approach that the evidence favours an extended Comptonizing region and we test the hypothesis that the Extended ADC model may provide a physically reasonable and consistent explanation of the Z-track source GX\\th 5-1.  \\subsection{The extended ADC}  Our work on the dipping class of LMXB previously led to our proposal of the extended ADC model comprising blackbody emission from the neutron star plus Comptonized emission from an extended, thin ADC existing as a hot layer above the accretion disk. In the dipping sources with inclination angles between 65$\\degmark$ and 85$\\degmark$, reductions in X-ray intensity take place at the orbital period generally thought to involve absorption in the bulge in the outer disk where the accretion flow impacts. Spectral evolution cannot be described in terms of absorption of a single emission component, but is more complex strongly constraining allowable models. It provides direct evidence for the extended nature of the Comptonized emission as this spectral component is removed only slowly during dipping showing the emitter to be extended. Secondly, the ADC size can be measured directly by the technique of dip ingress timing since the ingress time depends on the size of the major emission component: the Comptonized emission. This has shown that the radial extent of the ADC is typically 50\\th 000 km (Church \\& Ba\\l uci\\'nska-Church 2004) also showing that the ADC, like the disk is in general thin, with height $H$ at radial position $r$ having $H/r$ $<<$ 1. The extended ADC model was shown to give very good fits to the complex spectral evolution in dipping in many observations of the dipping LMXB (Church et al. 1997, 1998a,b, 2005; Ba\\l uci\\'nska-Church et al. 1999, 2000; Smale et al. 2001; Barnard et al. 2001). In addition, it was shown able to fit the spectra of all types of LMXB in a survey of LMXB made with {\\it ASCA} (Church \\& Ba\\l uci\\'nska-Church 2001) and so was able to fit sources of all inclination angles.  The evidence favouring the extended ADC is thus strong and it is difficult to avoid the conclusion that the main assumption of the Eastern model is invalid.  The main difference between the spectral forms of the extended ADC model and the Eastern model relates to the seed photons. In the Eastern model it is usually assumed that these originate on the neutron star and can be described by a single temperature blackbody or by the Wien approximation. In the extended ADC model the major source of seed photons must be the thermal emission of the disk below the ADC. Because of the temperature gradient $T(r)$ in the disk, the outer portions of the ADC will be above relatively cool parts of the disk providing a large population of soft seed photons ($kT$ $\\sim$ 0.1 keV or less) and the seed photon spectrum is very different from a blackbody at $\\sim$1 keV. Comptonization is well represented by a power law extending to energies lower than 0.1 keV, with a cut-off at high energies related to the maximum electron energy (see Church \\& Ba\\l uci\\'nska-Church 2004 for a full discussion). The use of a simple blackbody for the neutron star emission in applying the extended ADC model reflects the general lack of evidence for modification of the blackbody in the atmosphere of the neutron star (Church et al. 2002). Taking into account many spectral studies of LMXB, including the spectra of X-ray bursts, the only evidence for modification is in a minority of bursts, with no evidence in non-burst emission. Church et al. showed that whether modification occurs depends critically on the electron density in the neutron star atmosphere and that this is poorly known.  \\subsection{A model for the Z-track}  We previously applied the extended ADC model to high quality {\\it Rossi-X-ray Timing Explorer} data from the Z-track source GX\\th 340+0 (Church et al. 2006), and found that it fitted the spectra well at all positions on the Z-track, and also clearly suggested a physical model for the phenomenon. In this model, the soft apex between NB and FB is a state of the source in which the intensity and  the mass accretion rate were low. Moving onto the NB, there was a large increase in X-ray intensity (count s$^{-1}$) and an increase in the total broadband luminosity by about a factor of two caused by an increase in the luminosity $L_{\\rm ADC}$ of the Comptonized emission of the ADC. It was argued that this increase indicated that $\\dot M$ was increasing on the NB and was supported by a marked increase in the blackbody temperature of the neutron star from $\\sim$1 keV to $\\sim$2 keV and even higher values on the HB. This heating of the neutron star means a large increase in $T^4$ and so the radiation pressure became very high. There was a simultaneous decrease in emitting area on the neutron star. In this situation, it is appropriate to compare the flux emitted per unit area of the star with the Eddington value (see Sect. 4.1), and the flux increased from a low value till on the HB, it was three times super-Eddington. It was argued that this would disrupt the inner disk deflecting the accretion flow into the vertical direction leading to jet formation, and explaining why jets are detected on the upper NB and HB where the radiation pressure was high. In flaring it was found that $L_{\\rm ADC}$ was constant implying that $\\dot M$ was constant although the neutron star luminosity increased. The mass accretion rate per unit area of the neutron star closely agreed with the theoretical value for the onset of unstable nuclear burning, and it was proposed thus flaring was due to this.  In the present work, we apply the extended ADC model to high-quality {\\it Rossi-XTE} observations of the Z-track source GX\\th 5-1 to test the hypothesis that the same physical explanation can explain a second Z-track source. Spectral and timing studies were made using the same selections of data allowing the evolution of kHz QPO frequency to be directly correlated with the spectral evolution. This revealed a striking dependence of kHz QPO frequency on the neutron star radiation pressure. In the next section kHz QPO are briefly reviewed.   \\subsection{The kHz QPO}  Quasi periodic oscillations (QPO) have been the subject of intense study for many years, and a major discovery of the {\\it Rossi X-ray Timing Explorer} was that of the twin kilohertz QPO  in LMXB. Predictions of millisecond variability pre-date this considerably and Sunyaev (1973) argued that hot clumps of material orbiting in the inner disk around a black hole or neutron star would lead to quasiperiodic variability on millisecond timescales. Kilohertz QPO have now been detected in many LMXB (see reviews of van der Klis 2000, 2006), most of these displaying two peaks at frequencies such as at 600 and 900 Hz, and in the Z-track sources both frequencies increase during movement along the HB towards the hard apex. Rather similar QPO have also been seen in a number of black hole binaries.  In almost all QPO models the higher frequency peak is associated with the orbital frequency at some distance into the accretion disk away from the neutron star; for example, a frequency of 900 Hz corresponds to a radial position $r$ of about 18 km. In the sonic point beat frequency model (Miller et al. 1998) radiation drag reduces the angular momentum in the disk within a few stellar radii of the neutron star leading to a sonic point where the radial velocity becomes high (Lamb 1989; Miller 1991; Miller \\& Lamb 1993, 1996). The upper frequency $\\nu_2$ is associated with inhomogeneities orbiting at the sonic point in the inner disk, which are irradiated by X-rays from hot regions rotating on the neutron star leading to a beating of the upper frequency with the spin frequency $\\nu_{\\rm spin}$ so producing the lower frequency QPO. Although earlier observations found the frequency difference $\\Delta \\nu$ constant within errors consistent with the model, improved data showed that $\\Delta \\nu$ varied around the Z-track. The detection of twin kHz QPO and X-ray burst oscillations in the accreting millisecond pulsar SAX\\th J1808.4-3658 (Wijnands et al. 2003) led to the conclusion that $\\Delta \\nu$ $\\sim$ $\\nu_{\\rm spin}$ in some sources, but $\\nu_{\\rm spin}$/2 in other sources, inconsistent with the model, leading Lamb \\& Miller (2003) to propose a relativistic disk-spin resonance model.  The relativistic precession model (Stella \\& Vietri 1998) identifies $\\nu_2$ with motion of inhomogeneities at the inner disk edge and the lower frequency $\\nu _1$ with precession of the inhomogeneities. The disk resonance model (Abramowicz \\& Klu\\'zniak 2001) recognizes the preference of the kHz QPO to have a ratio of 3:2 in their frequencies suggesting a non-linear resonance between normal modes of the disk. This model predicted that twin peaks should also exist in black hole binaries, and the discovery of a second peak in GRO\\thinspace J1655-40 (Strohmayer 2001) was a success of the model. Comparison with data (Abramowicz et al. 2003; Klu\\'zniak et al. 2007) for both neutron star and black hole systems supports the model: black hole systems accurately follow the 3:2 ratio; LMXB show some scatter but a histogram of the ratio peaks strongly at 3:2. However, Belloni et al. (2005) have argued that in the case of \\hbox{Sco\\th X-1,} the peak in the distribution at a ratio of 1.5 was only marginally significant.   \\subsection{The assumed monotonic variation of mass accretion rate}  The assumption of $\\dot M$ increasing in the direction HB - NB - FB generally, but not universally made, is not supported by our previous results for GX\\th 340+0, and so we briefly review evidence for and against this.  The evidence for this was based originally on a multi-wavelength campaign on Cyg\\th X-2 (Hasinger et al. 1990) involving X-rays, ultraviolet, optical and radio data. Results from {\\it IUE} (Vrtilek et al. 1990) had increasing intensity for Z-track movement HB - NB - FB implying an increase of $\\dot M$; however, strong variability interpreted as flaring may have been due to X-ray dipping in which case there is no strong correlation of UV with track position.  The properties of most power spectral features correlate well with position on the Z-track; for example, the frequencies of the HBO (horizontal branch oscillations), i.e. low frequency QPO, increase around the Z-track in the direction HB - NB from $\\sim$5 to $\\sim$60 Hz (e.g. Homan et al. 2002). Similarly, NBO (normal branch oscillations) are seen on the NB and in two sources, Sco\\th X-1 and GX\\th 17+2, oscillations are seen in the FB, but only close to the soft apex, and there is a continuous increase of QPO frequency from lower NB to FB (Casella et al. 2006), i.e. in crossing the soft apex between these branches. Similarly, the frequencies of kilohertz QPO increase along the HB where they are observed (e.g. van der Klis 2000). It is general practice to display QPO properties as a function of arc distance $S$ measured along the Z-track in the direction HB - NB - FB, and this shows that QPO properties are generally well-correlated with $S$ (e.g. Dieters \\& van der Klis 2000) implying a dependence on a single parameter, i.e. $\\dot M$. It is also generally though that the soft apex represents the position at which $\\dot M$ increases to the Eddington value (Lamb 1989). In addition, the frequencies of HBO and kHz QPO in many models are directly related to $\\dot M$ (see Homan et al. 2002).  However, in recent years, some authors are beginning to doubt the general assumption of a monotonic increase in $\\dot M$. Some results conflict with the assumption: e.g. in GX\\th 17+2, the frequency of the HBO first increases, but then {\\it decreases} on the normal branch (Wijnands et al. 1996; Homan et al. 2002). In Sco\\th X-1 in the neighbourhood of the soft apex there is a rapid increase of QPO frequency (Casella et al. 2000; Dieters \\& van der Klis 2000) with $S$ which is inconsistent with the assumption that $S$ is a measure of $\\dot M$ as the change is over a very short part of the Z-track. Moreover the approximate constancy of NBO frequency in general over most of the NB does not indicate a dependency on $\\dot M$. Also, in GX\\th 17+2, the properties of the X-ray bursts sometimes observed do not depend on Z-track position (Kuulkers et al. 2002).  In our previous work on GX\\th 340+0 (Church et al. 2006) we proposed that $\\dot M$ increased between the soft apex and the hard apex, while on the FB, $\\dot M$ was constant. The X-ray intensity, the measured broadband luminosity of the source and the luminosity of the Comptonized emission all increased substantially on the NB strongly suggesting that $\\dot M$ is not decreasing. The argument is often made that the X-ray intensity cannot be a good indicator of $S$ and thus of $\\dot M$ (e.g. Dieters \\& van der Klis 2000) because the intensity increases on both the NB and FB, moving away from the soft apex. However, our work on GX\\th 340+0 shows this not to be a secure argument, since the intensity increases on the NB because $\\dot M$ increases, while on the FB, $\\dot M$ is constant, but strong nuclear burning on the neutron star causes the intensity increase. One aim of the present work is to test whether similar results are found in GX\\th 5-1.  \\subsection{GX\\th 5-1}  GX\\th 5-1 is the second brightest persistent LMXB (after Sco\\th X-1), discovered in 1968 (Fisher et al. 1968; Bradt et al. 1968) as a bright X-ray source in the Sagittarius region and classified as a Z-track source by Hasinger \\& van der Klis (1989).  A three-year study of the source was made using {\\it Ginga} (van der Klis et al. 1991) which revealed an extended horizontal branch and an upcurving at the end suggesting a change of behaviour, effectively a fourth branch (Lewin et al. 1992). Analysis of {\\it Exosat} and {\\it Ginga} data revealed flaring branch data \\begin{figure*}[!t] \\begin{center}                                                         % \\includegraphics[width=64mm,height=140mm,angle=270]{lc64}        % \\caption{Background-subtracted deadtime-corrected PCA lightcurve of the November 1998 observation of GX\\th 5-1 with 64 s binning.} \\end{center} \\end{figure*} for the first time (Kuulkers et al. 1994) and a detailed investigation of correlated spectral and timing behaviour was made.  It was the first Z-track source found to have QPO (van der Klis et al. 1985) and  kilohertz QPO were detected by Wijnands et al. (1998). Radio emission has been detected from GX\\th 5-1, e.g. by Penninx (1989) implying the existence of jets, and it is known to be strongest on the horizontal branch (Berendsen et al. 2000), as for other Z-track sources. Analysis of 1987 {\\it Ginga} data (Asai et al. 1994) and {\\it ASCA} data (Asai et al. 2000) showed there was no evidence for an iron K line, but a Gaussian feature existed at 10 keV. To date no satisfactory explanation exists for this line. ",
        "conclusions": "We have shown that spectra selected along the Z-track in GX\\th 5-1 are well fitted using the extended ADC model, and that the spectral evolution is very similar to that we found in GX\\th 340+0 (Church et al. 2006). Thus the physical model proposed to explain the spectral evolution in that source is also able to explain GX\\th 5-1. In this model, the soft apex has the lowest value of $\\dot M$ and the strong physical changes on the normal branch are primarily due to an increase of $\\dot M$ and of neutron star temperature. The increase of temperature means that the radiation pressure {\\it must} be high in the neighbourhood of the neutron star. We propose that this super-Eddington radiation pressure combined with inner disk radiation pressure disrupts the inner disk diverting accretion flow into the vertical direction leading to jet formation. Thus high radiation pressure is a necessary but not necessarily sufficient condition for jet formation. On the horizontal branch, the decrease of intensity and $L_{\\rm ADC}$ indicates that $\\dot M$ decreases; however the blackbody temperature continues to increase and there are various ways of explaining this, such as $\\dot M$ decreasing in the disk while an enhanced accretion flow is still affecting the neutron star. Evidence is presented that the flaring branch consists of unstable nuclear burning on the neutron star given the constant $L_{\\rm ADC}$ and the increase of $L_{\\rm BB}$, and the good agreement between $\\dot m$ at the soft apex and the theoretical $\\dot m_{\\rm ST}$ below which burning is unstable. However, we should add the caution that while we now have evidence that the flaring branch can be explained in this way in the three Cyg-like Z-track sources, i.e. in GX\\th 340+0 (Church et al. 2006), in GX\\th 5-1 (present paper) and in Cyg\\th X-2 (Church et al. 2007), we do not claim that this is so in the other group of Z-track sources: the Sco-like sources. Preliminary evidence suggests that these are similar on the NB and HB, but that the FB may be different.  Radio emission is detected at the hard apex and horizontal branch showing that jets are present (Penninx 1989; Hjellming \\& Han 1995; Berendsen et al. 2000) where we find that radiation pressure is strongest and is non-existent on the FB and soft apex. It was, of course, previously suggested that radiation pressure may be important in jet formation (Bisnovatyi-Kogan \\& Blinnikov 1977; Begelman \\& Rees 1984) and that the conical opening in the inner accretion disk may perform some degree of collimation (Lynden-Bell 1978).  The present results do not support the standard view that $\\dot M$ increases monotonically in the direction HB - NB - FB, but have $\\dot M$ constant in flaring and increasing from the soft apex to hard apex. Moreover, in the standard view the source becomes super-Eddington at the soft apex whereas the present work indicates the opposite, that this happens at the hard apex.  By combining timing and spectral analysis we have shown that the variation of the kHz QPO along the Z-track is highly correlated with $f/f_{\\rm Edd}$, and the major change in the QPO frequencies takes place when this ratio equals unity, and the radiation pressure becomes large. This suggests that the higher frequency kHz QPO is an oscillation at the inner edge of the disk, i.e. the edge formed by the action of radiation pressure removing a large section of the unperturbed disk, and that the edge moves to larger radial positions when the radiation pressure increases. A major feature of kHz QPO is that they are essentially a feature of the horizontal branch, fading away at the upper end of the normal branch and this is consistent with our model for the Z-track as the high radiation pressure of the neutron star revealed by spectral fitting is also a feature of the horizontal branch. Fig. 11 shows that when $f/f_{\\rm Edd}$ has fallen to $\\sim$0.5, kHz QPO are no longer seen. Thus there is an indication that the kHz QPO only exist when part of the inner disk is blown away. The smallest Keplerian radius we measure has a value of 18 km on the upper normal branch and if QPO continued to exist on the normal branch, the Keplerian radius would decrease further. However, as the radiation pressure falls there would be no truncation of the disk at an edge acting as a site for the oscillation, and this is consistent with the lack of kHz QPO on the lower normal branch and flaring branch.  Thus there is a simple explanation for the known occurrence of the higher frequency kHz QPO at a preferred radial position in the disk: that this is the disk edge, and the results suggest that the variation of frequency is caused by change in position of the edge.  The physics of the accretion flow in the inner disk is very uncertain (e.g. van der Klis 2000), and probably the most important question is whether the disk in general in LMXB meets the neutron star or is truncated at an innermost stable orbit, an Alfv\\'en radius or a sonic point, and this is relevant to the QPO models which invoke truncation. However, it is difficult to say whether truncation of the disk due to relativistic, magnetospheric or other effects actually takes place. The present work presents a simple observational result on the higher frequency QPO and does not directly say anything about the lower frequency oscillation, and it is not appropriate to comment on the detailed physics of the QPO models here. However, the present work indicates that the inner disk is not in general truncated, but can become truncated on the upper NB and HB. Disk truncation is due to high radiation pressure and {\\it we do not need} to invoke truncation because of an innermost stable orbit, or an Alfv\\'en radius external to the neutron star or a sonic point.  \\thanks{ This work was supported in part by the Polish KBN grant KBN-1528/P03/2003/25, by the Polish Ministry of Higher Education and Science grant no. 3946/B/H03/2008/34 and by PPARC grant PPA/G/S/2001/00052.}"
    },
    "0812/0812.3377_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\section*{The constants of nature}  Since the time of Newton, the constancy of the fundamental laws of nature has been undoubted. Comparing and reproducing experiments have been at the root of the scientific approach: A physical experiment which we perform today will have the same outcome as the same experiment performed tomorrow\\footnote{Neglecting experiments which incorporate probabilities, for instance quantum mechanical effects.}. Neglecting local gravitational effects, it should also not matter where we perform the experiment. Hence, it has been unquestionable for a long time that the laws of nature are constant over space and time. Moreover, Einstein formulated this space- and time independence of physics in his strong equivalence principle, making it an essential part of his theory of general relativity.  Today's view of this question is somewhat different, at least from theoretical aspects. Even though compelling evidence for changes in the laws of physics has up to now not been found, we have to admit that we are still lacking a profound test of this constancy. In the past, the laws of physics have only been thoroughly tested on time and length scales accessible by mankind, \\ie on timescales of years and on length scales that do not go beyond the size of our solar system\\footnote{Note that general relativity has furthermore not been tested on length scales smaller than about 1mm.}. Only recently astrophysics and cosmology have opened a door to test physics on immensely broader scales, reaching out to unimaginable length scales of several gigaparsecs and going back in time to the very beginning of our Universe.  This thesis will deal with probes of possible variations of constants throughout the whole accessible history of the Universe. In a first part, we will study one of the most distant (in time and space) events where physics can be applied and tested, primordial nucleosynthesis. It is the process during which the light elements of our Universe were formed and which happened when our Universe was only one minute old, extremely hot and dense. If physics was really subject to variations, primordial nucleosynthesis is a prime candidate for any studies of this kind. In a second step the obtained results will be combined together with further tests of varying constants at later times to derive a ``history of variations''. Finally, we will show how these results can be used to test models beyond standard physics which currently cannot be accessed directly by experiments.   \\section*{Outline}  In this thesis I will work out the influence of varying ``constants'' on the process of primordial nucleosynthesis and implied constraints to theories beyond standard physics. In the next chapter, I will give a short introduction to variations of physical constants, some historical remarks and theoretical motivations. Chapter \\ref{Chap:Cosmology} will lay the theoretical framework for our understanding of the Universe as a whole, explaining general relativity and the main laws of cosmology. In Chapter \\ref{chap:SUSYandGUT} I will introduce the concepts of supersymmetry and grand unified theories (GUTs) which are widely accepted as extensions of the Standard Model of particle physics. Chapter \\ref{chap:Quintessence} will introduce quintessence models which can yield variations of constants.  Part II will focus on the details of one particular process in the history of our Universe, \\BigBangNucleosynthesis\\ (BBN). Chapter \\ref{chap:BBN} will explain the standard process of BBN and the physics behind. Chapter \\ref{chap:BBNvaryingNucl} will introduce the possibility of varying constants in the process of BBN, and Chapter \\ref{chap:BBNvaryingFund} will demonstrate how one can relate the results to variations of the Standard Model parameters. Finally, in Chapter \\ref{chap:BBNobsVStheo}, the observed element abundances will be used to derive constraints on variations of fundamental parameters.  In Part III I will study relevant tests of varying constants from the Big Bang until today. Chapter \\ref{chap:ExpTestofVariations} gives an overview over tests of varying constants, and in Chapter \\ref{chap:VariationsFromBBNtoTodayinGUT} I will combine these tests within six different GUT models, showing how variations of constants can in principle be used to probe models of grand unification. Using the six GUT models, Chapter \\ref{chap:ProbingQuintessence} shows how models of quintessence can be probed under the assumption of grand unification.  Finally, in Chapter \\ref{chap:Conclusion}, I will sum up the findings of this thesis and give some final conclusions and outlook. \\\\ \\\\ \\\\ The work on this thesis has led to four main publications:  \\begin{itemize} \\item Michael Doran, Steffen Stern, Eduard Thommes,\\\\ Baryon Acoustic Oscillations and Dynamical Dark Energy,\\\\ JCAP 0704:015 (2007) \\cite{DST06} (not in focus of this thesis) \\item Thomas Dent, Steffen Stern, Christof Wetterich,\\\\ Primordial nucleosynthesis as a probe of fundamental physics parameters,\\\\ Phys.~Rev.~D 76, 063513 (2007) \\cite{DSW07} \\item Thomas Dent, Steffen Stern, Christof Wetterich,\\\\ Unifying cosmological and recent time variations of fundamental couplings,\\\\ Phys.~Rev.~D 78, 103518 (2008) \\cite{DSW08.1} \\item Thomas Dent, Steffen Stern, Christof Wetterich,\\\\ Time variation of fundamental couplings and dynamical dark energy,\\\\ Preprint arXiv:0809.4628, accepted by JCAP \\cite{DSW08.2} \\end{itemize}       \\chapter{Variation of ``constants''} ",
        "conclusions": ""
    },
    "0812/0812.0835_arXiv.txt": {
        "abstract": "{With the plethora of detailed results from heliospheric missions such as Ulysses and SOHO and at the advent of the first mission dedicated to in situ studies of neutral heliospheric atoms IBEX, we have entered the era of precision heliospheric studies. Interpretation of these data require precision modeling, with second-order effects quantitatively taken into account.} {We study the influence of the non-flat shape of the solar Lyman-$\\alpha $ line on the distribution of neutral interstellar hydrogen in the inner heliosphere and assess the  importance of this effect for interpretation of heliospheric in situ measurements.} {Based on available data, we (i) constructed a model of evolution for the solar Lyman-$\\alpha $ line profile with solar activity, (ii) modified an existing test-particle code used to calculate the distribution of neutral interstellar hydrogen in the inner heliosphere so that it takes the dependence of radiation pressure on radial velocity into account, and (iii) compared results of the old and new version.} {Discrepancies between the classical and Doppler models appear between $\\sim 5$ and 3~AU and increase towards the Sun from a few percent to a factor of 1.5 at 1~AU. The classical model overestimates the density everywhere except for a $\\sim 60^{\\circ}$ cone around the downwind direction, where a density deficit appears. The magnitude of the discrepancies appreciably depends on the phase of the solar cycle, but only weakly on the parameters of the gas at the termination shock. For in situ measurements of neutral atoms performed at $\\sim 1$~AU, like those planned for IBEX, the Doppler correction will need to be taken into account, because the modifications include both the magnitude and direction of the local flux by a few km/s and degrees, respectively, which, when unaccounted for, would introduce an error of a few km/s and degrees in determination of the magnitude and direction of the bulk velocity vector at the termination shock.} {The Doppler correction is appreciable for in situ observations of neutral H populations and their derivatives performed a few AU from the Sun.} ",
        "introduction": "\\label{sec:intro}  Development of models of the distribution of neutral interstellar hydrogen in the heliosphere began after the discovery of a diffuse interplanetary Lyman-$\\alpha$ glow \\citep{thomas_krassa:71, bertaux_blamont:71}, interpreted by \\citet{fahr:68} and \\citet{blum_fahr:70a} as caused by the scattering of solar Lyman-$\\alpha$ radiation on this gas flowing through the Solar System.  At the very early phase, the influence of processes going on at the heliospheric interface on the heliospheric gas was neglected. It was assumed that both the ionization due to solar output and the radiation pressure acting on the H atoms are stationary and spherically symmetric around the Sun and that their rates fall off proportionally to inverse square of heliocentric distance. It was further assumed that the inflowing neutral gas is monoenergetic, i.e., that before the encounter with the Sun the atoms move with identical, parallel-oriented velocities, equal to the macroscopic bulk velocity of the gas.  These assumptions formed the basis of the first model of density distribution of neutral interstellar gas around the Sun: the purely analytical ``cold'' model \\citep{fahr:68, axford:72}. It features axial symmetry about the axis of inflow of the gas and shows either a singularity at the downwind axis, when radiation pressure is too weak to compensate for solar gravity, or an empty  ``avoidance zone'' in the downwind region with a paraboloidal boundary surface, when radiation pressure overcompensates for solar gravity.  Lifting of the monoenergetic assumption (i.e. allowing for a finite temperature of interstellar gas, high enough to yield thermal velocity comparable to the bulk velocity of the gas), adoption of the distribution function of the gas far away from the Sun (``in infinity'') in the form of a Maxwellian shifted in the velocity space by the bulk velocity vector, and assumption that the gas is collisionless on the distance scale comparable to the size of the heliosphere, all allowed us to use the Boltzmann equation to describe the problem of the interaction of neutral interstellar gas with the solar environment. Its solution brought the ``hot model'' of the gas distribution in the inner heliosphere \\citep{thomas:78, fahr:78, fahr:79, wu_judge:79a, lallement_etal:85b}. This model required numerical integration of the distribution function, but the function itself was still analytical. The ``hot model'' in its various implementations became the canonical model of neutral interstellar gas distribution in the inner heliosphere, so further on we will refer to it as ``the classical hot model''.  Further modeling of the interaction and distribution of neutral interstellar hydrogen near the Sun focused on two main topics. On one hand, a lot of effort was made toward simulating and understanding processes going on at the boundary between the expanding solar wind and the incoming, partially ionized interstellar gas, in the region referred to as the heliospheric interface -- and this issue will not be addressed in the present paper: the reader is referred to recent reviews by \\citet{baranov:06a, baranov:06b, izmodenov:06a, izmodenov_baranov:06a}. On the other hand, development of the hot model continued, aimed at a more realistic description of the distribution of neutral interstellar hydrogen in the inner heliosphere, suitable for quantitative, and not only qualitative interpretation of heliospheric measurements.  In the first shot, the assumption that the ionization rate is spherically symmteric was eliminated, when \\citet{lallement_etal:85a} described the latitudinal modulation of the charge exchange rate with a one-parameter formula $1 - A\\,\\sin^2 \\phi$, implementing it in the CNRS model of the heliospheric gas distribution. This allowed the authors to vary the equator-to-pole contrast of the ionization rate, but required keeping the width and range of the enhanced ionization band fixed.  A different extension of the hot model was proposed by \\citet{rucinski_fahr:89, rucinski_fahr:91}, who realized that the rate of ionization by electron impact is not proportional to the inverse square of solar distance. Electron ionization is particularly important for interstellar helium \\citep{mcmullin_etal:04a, lallement_etal:04b, witte:04}. The effects in hydrogen are only noticeable within a few AU from the Sun, where the density of hydrogen gas is already very reduced by earlier ionization and solar radiation pressure, overcompensating for solar gravity; therefore this aspect of the physics of neutral interstellar hydrogen has been neglected for quite a while since then, but is reintroduced in the most recent versions of the model \\citep{bzowski_etal:08a}.  In the next round in the development of the density model, gone was assumption of the invariability of radiation pressure and ionization rate. This aspect of the physics of neutral interstellar gas was implemented in a simplified numerical model by \\citet{kyrola_etal:94} and shortly afterwards in the Warsaw test-particle code, to which we will refer in this paper as the Warsaw model \\citep{rucinski_bzowski:95b, bzowski_rucinski:95a, bzowski_rucinski:95b, bzowski_etal:97}. Indeed, both the solar EUV flux \\citep{floyd_etal:02a} and the solar wind flux \\citep{king_papitashvili:05} vary considerably during the solar cycle. The result is a solar cycle variation of the solar radiation pressure, of the EUV ionization rate, and of the rate of charge exchange between neutral H atoms and solar wind protons. This in turn results in appreciable variations in the density and bulk velocity of neutral interstellar hydrogen within a dozen AU from the Sun.  The Warsaw time-dependent model was fully numerical: not only did the integration of the distribution function have to be performed numerically, as in the classical ``hot model'', but also the calculation of the integrand, i.e. of the distribution function itself, was no longer analytical. At this phase in the heliospheric research, substitute idealized models of the evolution of these parameters had to be used due to the lack of  sufficiently long time series of measurements of the solar wind speed and density, and of the solar EUV output. These aspects of the development of heliospheric gas models were discussed in greater detail in a review by \\citet{rucinski_bzowski:96}.  In the next move, the inner-heliospheric effects of the interaction of the solar wind and interstellar gas within the heliospheric interface were taken into account. Owing to the charge exchange between the atoms of interstellar gas and the heated and compressed plasma in front of the heliopause, the original population of neutral atoms is somewhat cooled and accelerated, and a new population of atoms appears. These new neutral atoms initially maintain the properties of the parent plasma population, but later on they decouple from this plasma and flow through the heliopause and inside the termination shock of the solar wind. The two neutral populations interact further with the ionized components, exchanging charge with the protons from the local plasma and producing new populations of neutral atoms. Hence, the processes in the heliospheric interface create a few distinct, collisionless populations of neutral atoms \\citep{osterbart_fahr:92, baranov_etal:91, baranov_malama:93}, as discussed in detail by \\citet{izmodenov:00} and \\citet{malama_etal:06}.  From the viewpoint of modeling the distribution of neutral interstellar hydrogen in the inner heliosphere, the most important aspect of the processes going on in the heliospheric interface is the modification of the distribution function at the termination shock \\citep{izmodenov_etal:01a}. Instead of the shifted Maxwellian with parameters homogeneous in space, \\citet{scherer_etal:99} adopted a sum of two Maxwellians, with non-isotropic temperatures, shifted by appropriate bulk velocity vectors. The two components of the new functions corresponded to the two thermal populations (primary and secondary) predicted by the Moscow Monte Carlo simulation of heliospheric interface and had parameters that depended on the offset angle $\\theta$ from the upwind direction. This new version of the Warsaw model was time dependent, but still axially symmetric. At that time, enough measurements had been published to attempt to introduce observation-based models of the radiation pressure and ionization rate to the simulations.  Axial symmetry was removed in the next round of the Warsaw model development, when an anisotropy of the ionization rate was introduced. \\citet{bzowski_etal:01a, bzowski_etal:02} allowed the ionization rate to change as a continuous function of the  heliographic latitude, with the latitudinal profile of the ionization rate continuously changing with the phase of solar cycle. As a result, the ionization field in the model became 2D (maintainining the axial symmetry about the solar rotation axis) and time-dependent; and since the gas inflow axis and solar rotation axis are inclined at an angle to each other, the model of neutral hydrogen distribution in the inner heliosphere became 3D and time-dependent.  The most recent extension of the Warsaw model is presented in this paper. We improve on the modeling of radiation pressure acting on individual H atoms, which now is not only a function of time, but also of the radial velocity of the atom. In the previous versions of the model it was assumed that the profile of the solar Lyman-$\\alpha$ line, responsible for the radiation pressure, is flat. In reality it not only is not flat, but shows considerable variations depending on the phase of solar activity \\citep{lemaire_etal:02, lemaire_etal:05}.  We take this into account, and in Sect. 2 we develop an observation-based model of the evolution of the line profile as a function of the wavelength- and disk-integrated flux. We use this model in a newly-developed code that simulates the density and higher moments of the distribution function of neutral interstellar hydrogen in the inner heliosphere. We discuss these calculations in Sect. 3. With the use of the newly-developed code, we assess modifications of the local gas density and flux with respect to the results of models neglecting the dependence of radiation pressure on the radial velocity. We also show possible implications for interpretation of heliospheric in situ measurements of neutral atoms and pickup ions. This discussion is provided in Sect. 4. Section 5 offers a summary of the results. ",
        "conclusions": "\\label{sec:conclu}  We compared the density, bulk velocity, and flux of neutral interstellar hydrogen in the inner heliosphere, calculated using either the classical hot model or a newly-developed model with the radiation pressure being a function of specific radial velocities of individual atoms. The conclusions are the following. \\begin{enumerate} \\item Differences between the Doppler and classical hot models are restricted to $\\sim 3$ to $\\sim 5$~AU from the Sun, depending on the angle off the upwind direction. The classical hot model overestimates the density of neutral hydrogen gas everywhere in the inner heliosphere apart from a $\\sim 60 \\degr$ cone around the downwind direction, where a density deficit is predicted. Generally, the density excess/deficit is a strong function of the offset angle from the upwind direction and of the heliocentric distance (Figs. \\ref{xg}, \\ref{xc}, and \\ref{xd}).  \\item The magnitude of density excess $q$ varies with the solar activity level and is higher at solar maximum (Figs. \\ref{xg} and \\ref{xc}), i.e. $v_r$-independent models stronger overestimate the local density during solar maximum.  \\item The density excess $q$ is a weakly increasing, almost linear function of the ionization rate in the inner heliosphere and a weak, almost bilinear function of the bulk velocity and temperature of the gas at the termination shock (Figs. \\ref{xf} and \\ref{xe}).  \\item The classical model shows a higher upwind-downwind amplitude of radial velocities at 1~AU than the Doppler model; the differences are about 10\\%. The excess of the absolute flux of neutral atoms, returned by the classical model, is close to 15\\% at 1~AU during solar minimum, i.e. at the time of observations by the forthcoming IBEX mission. The excess is much higher during solar maximum, but the absolute values of the fluxes are lower by a few orders of magnitude than during solar minimum (Figs. \\ref{xi},\\ref{xj},\\ref{xk},\\ref{xl}).  \\item The classical model returns a different deflection of the local upstream directions of the primary population of neutral interstellar H than the Doppler model. The deflections of the secondary population returned by the classical and Doppler models are practically identical, meaning that any interpretation of in situ measurements of these fluxes must be done very carefully, with an appropriate account of the solar line profile shape to avoid confusing deflections caused by the Doppler effect in the radiation pressure with, e.g., results of deformating the heliospheric interface by external magnetic field (Fig. \\ref{xm}).  \\item Discrepancies between the PUI flux at Ulysses in aphelion, returned by the Doppler and classical models, are about 10\\% and are almost insensitive to the phase of the solar cycle. \\end{enumerate}"
    },
    "0812/0812.1308_arXiv.txt": {
        "abstract": "{}{We report on the redshift of the lensing galaxy and of the quasar \\qj\\, and  on the lens model of the system.}{A deep VLT/FORS2 spectrum  and  HST/NICMOS-F160W   images  are   deconvolved.  From the images we derive the light profile of the lensing galaxy and an accurate relative astrometry for the system.  In addition we measure the flux ratio between the quasar images in the \\ion{Mg}{II} emission line to constrain the mass model.}{ From the spectrum we measure the redshift of the lensing galaxy ($z=0.317\\pm 0.001$) and of  the quasar ($z=1.294\\pm0.008$). Using the flux ratio in the lens model allows to discard the SIE as a suitable  approximation of the lens potential. On the contrary the truncated-PIEMD  gives a good fit to the lens and leads to a time delay of $\\Delta$t$_{A-B}$=-14.5$\\pm$0.1~days, with \\ho=73\\kmsmpc.}{Using the flux ratio to constrain the mass model favors the truncated-PIEMD over the SIE, while ignoring this constraint leaves the choice open. }{} ",
        "introduction": "The study of  gravitationally lensed quasars is relevant both to cosmology and  to the determination of the total mass in (lensing) galaxies. Time delays between multiple quasar images depend both on the mass distribution of the lensing galaxy and on the Hubble parameter, H$_0$ (Refsdal, 1964). The translation of an observed time delay into a value of \\ho\\, requires the modelling of the surface mass density of the lens.  For that purpose, geometric and photometric constraints are needed such as: (a) spectroscopic redshifts for both the source and the lens components and (b) accurate relative positions of the lensed images and of the lensing galaxy. Additional constraints to be used are the flux ratios (FRs) between the quasar images. Indeed, the FRs are equal to the magnification ratios at the image positions, and hence depend on the mass distribution of the lens. However, the contribution of the lens (lensing galaxy and environment) must be disentangled from other sources of flux variation such as:  differential galactic extinction, intrinsic quasar variability, microlensing and/or millilensing (de)magnification, all phenomena which have different time scales and wavelength dependencies (see e.g. Anguita et al. 2008 for more details). In the case of radio loud quasars this is rather straightforward, yet it does not apply to the whole quasar population. For optically selected quasar, the separation of the different sources of flux variation is more complex. In spite of this complexity, one must attempt it to derive values of the image FRs genuinely related to the strong lens. In the case of doubly imaged quasars (which constitute $\\sim$60\\% of the strongly lensed quasars known so far), it is particularly important to use the FR to constraint the lens mass model, as these systems are  naturally the least constrained. \\\\ In this paper we report on new observations of  the lensed quasar \\qj, and on a model of the lensing system. The quasar and the  datasets relevant to our study are presented in $\\S$~\\ref{theqso}. In $\\S$~\\ref{deconvsec}, the spectra and image deconvolutions are discussed. We model the lens in $\\S$~\\ref{simple}. Conclusions are given in $\\S$~\\ref{dandc}.  \\\\ We assume a WMAP type cosmology (Spergel et al. 2003, 2007): $H_0$=73\\kmsmpc, $\\Omega_m$=0.3 and $\\Omega_\\Lambda$=0.7. All magnitudes quoted in the paper are in the AB system. ",
        "conclusions": "\\label{dandc} We have measured the redshift of the lensing galaxy and improved the measurement of the quasar redshift in \\qj. While modelling the lens, we noticed that using only the quasar image positions to constrain the mass model does not allow to discriminate a mass density profile over another: in the case of \\qj, both a SIE and a TPIEMD appear as suitable approximations of the lensing potential. If we add the constraint of the quasar image FR (measured in the \\ion{Mg}{II} line), it becomes possible to discard the SIE as a suitable approximation of the lensing potential. The lens modelled as a TPIEMD is good. The time delay inferred with the TPIEMD is $\\Delta$t$_{A-B}$=-14.5$\\pm$0.1 days for $H_0$=73\\kmsmpc.\\\\ In conclusion, using the  FR to constrain the lens allows to select the most suitable galaxy profile and to measure its mass, as well as to predict the time delay with more confidence (or the Hubble constant if the time delay is known)."
    },
    "0812/0812.4231_arXiv.txt": {
        "abstract": "With a column density $\\log N($\\ovi$) = 14.95 \\pm 0.05$, the \\ovi\\ absorber at $z_{\\rm abs} \\simeq 0.2028$ observed toward the QSO PKS\\,0312--77 ($z_{\\rm em} = 0.223$) is the strongest yet detected at $z < 0.5$. At nearly identical redshift  ($z_{\\rm abs} \\simeq 0.2026$),  we also identify a Lyman limit system (LLS, $\\log N($\\hi$) = 18.22\\,^{+0.19}_{-0.25}$). Combining FUV and NUV spectra of PKS\\,0312--77 with optical observations of galaxies in the surrounding field ($15\\arcmin \\times 32\\arcmin$), we present an analysis of these absorbers and their connection to galaxies. The observed \\oi/\\hi\\ ratio and photoionization modelling of other low ions indicate the metallicity of the LLS is $[{\\rm Z/H}]_{\\rm LLS} \\simeq -0.6$ and that the LLS is nearly 100\\% photoionized. In contrast, the \\ovi-bearing gas is collisionally ionized at $T\\sim (3$--$10) \\times 10^5$ K as derived from the high-ion ratios and profile broadenings. Our galaxy survey reveals 13 ($0.3 \\la L/L_* \\la 1.6$) galaxies at $\\rho < 2 h^{-1}_{70}$ Mpc and $|\\delta v| \\la 1100$ \\km\\ from the LLS. A probable origin for the LLS is debris from a galaxy merger, which led to a  $0.7 L_*$ galaxy ($[{\\rm Z/H}]_{\\rm gal} \\simeq +0.15$) at $\\rho \\simeq 38 h^{-1}_{70}$ kpc. Outflow from this galaxy may  also be responsible for the supersolar ($[{\\rm Z/H}]_{\\rm abs} \\simeq +0.15$), fully ionized absorber at $z_{\\rm abs} \\simeq 0.2018$ ($-190$ \\km\\ from the LLS). The hot \\ovi\\ absorber likely probes coronal gas about the $0.7 L_*$ galaxy and/or ($\\sim$0.1 keV) intragroup gas of a spiral-rich system. The association of other strong \\ovi\\ absorbers with LLS suggests they trace galactic and not intergalactic structures. ",
        "introduction": "Connecting the QSO absorbers with their environments is crucial for using the absorbers to study the properties and evolution of galaxies, the intergalactic medium (IGM), and the galaxy-IGM interface.  The taxonomy of QSO absorbers is usually made according to their  \\hi\\ column densities ($N($\\hi$)$): the Ly$\\alpha$ forest \\citep{rauch98} with $\\log N($\\hi$) < 16$, the Lyman limit systems that are optically thick at the Lyman limit \\citep[LLS;][]{tytler82} with $16 \\le \\log N($\\hi$) < 20.3$, and the damped Ly$\\alpha$ absorbers \\citep[DLAs;][]{wolfe05} with $\\log N($\\hi$)  \\ge 20.3$. Given their  placement in the \\hi\\ column density hierarchy, the LLS likely represent the interface  between the tenuous, highly ionized Ly$\\alpha$ forest  and the dense, neutral DLAs.  Many observational studies have been undertaken to connect these absorbers to physical objects, such as galaxies, galaxy halos, intergalactic voids. At $z\\la 1$, analyses of the  absorber-galaxy relationship show that QSO absorbers are not distributed randomly with respect to galaxies, even for absorbers with the lowest \\hi\\ column densities detected so far \\citep[e.g.,][]{lanzetta95,tripp98,impey99,chen01,bowen02,penton02,stocke06,wakker08}. Strong  \\mgii\\ absorbers are  found within the extended halos (impact parameter $\\rho < 100$ kpc) of individual galaxies \\citep[e.g.,][]{bergeron91,steidel93}, while the DLA are generally associated with the main bodies of galaxies \\citep[e.g.,][]{chen03,rao03,wolfe05}. Although strong \\mgii\\ absorbers are believed to be tracers of LLS, there is a large scatter between $W_\\lambda($\\mgii$)$ and $N($\\hi$)$, not allowing a direct connection between these two quantities \\citep{churchill05,menard08}, but see also \\citet{bouche08}. Furthermore the origin, metallicity, and physical properties of the \\mgii-bearing gas is largely unknown because of the limited information available (e.g., metallicity, ionization and physical conditions).  This requires high spectral resolution UV space-based observations at $z\\la 1$ where several \\hi\\ Lyman series lines and metal lines in various ionization stages combined with galaxy redshift surveys and imaging can be acquired.  Currently, there are only three reported LLS where detailed information about their metallicity, ionization, and galaxy environment is available. Those are at $z = 0.08092 $ toward PHL\\,1811 \\citep{jenkins03,jenkins05}, $z =  0.16710$ toward PKS\\,0405--123 \\citep{chen00,prochaska04,prochaska06,williger06}, and $z = 0.09847 $ toward PKS\\,1302--102  \\citep{cooksey08}. These studies suggest that LLS can be either  metal-rich ($Z$$\\sim$$0.5$--$1 Z_\\odot$) or metal poor ($Z\\la 0.02 Z_\\odot$) and are generally associated with the extended  reaches ($\\sim30$--100 kpc) of individual galaxies. Combined with previous studies, these conclusions demonstrate that LLS are crucial for understanding the interaction between galaxies and their environments, an important ingredient in any cosmological simulations \\citep[e.g.,][]{bertone07,oppenheimer08a}.  In this paper, we present another detailed study of the relationship between a LLS observed at $z= 0.20258$ toward the QSO PKS\\,0312--77 ($z_{\\rm em} = 0.2230$) and galaxies.\\footnote{We also refer the reader to the master's thesis of Sarah S. Giandoni (New Mexico State University, 2005) for a complimentary analysis of the PKS0312--77 sightline.} This QSO was originally observed in the UV using the high-resolution mode of Space Telescope Imaging Spectrograph (E140M, E230M) onboard of the {\\em Hubble Space Telescope}\\ ({\\em HST}) to study the Magellanic Bridge  \\citep[see][]{lehner08}, a region of gas linking the Small and Large Magellanic Clouds (SMC, LMC). The STIS spectrum revealed very strong \\lya\\ and \\lyb\\ absorption lines, suggesting the presence of a LLS. This  was subsequently confirmed by {\\em Far Ultraviolet Spectroscopic Explorer}\\ ({\\em FUSE}) observations where the expected flux decrement at the Lyman limit  at $z\\approx 0.203$ is detected. We show that the \\hi\\ column density (that often lacks in low resolution surveys of Lyman thick absorbers) can be determined owing to the high spectral resolution of these spectrographs. Determining $N($\\hi$)$ is crucial as it allows us to determine the metallicity and ionization conditions of the LLS.   At a similar redshift, this sightline also reveals the strongest \\ovi\\ absorber discovered to date, with $\\log N($\\ovi$) = 14.95 \\pm 0.05$ (see \\S\\ref{ssec-metal}). Strong \\ovi\\ absorbers are rare and are unlikely to be related to the warm-hot intergalactic medium \\citep[WHIM,][]{cen99,dave99}, where a large fraction of the baryons could reside at $z<0.5$. \\citet{oppenheimer08} argued that strong \\ovi\\ collisionally ionized absorbers are related to the recycling of gas between the IGM and galaxies. Their conclusions are based on {\\sc Gadget-2} cosmological simulations that included a variety of wind models and input physics variations. The strong \\ovi\\ absorbers may therefore be higher redshift counterparts of \\ovi\\ absorption that probes the extended halos of galaxies, such as in our Galaxy \\citep{savage03} and the LMC \\citep{howk02,lehner07}. The strong \\ovi\\  absorbers may also probe cooling gas from intragroup medium revealed by soft X-rays observations or even possibly relatively cool (0.1--0.3 keV) intragroup gas \\citep{mulchaey96}. Recent surveys at low $z$ show indeed that, as for the \\hi\\ and low-ion absorbers, \\ovi\\ absorbers are not usually found in intergalactic voids but within $\\la 600$--800 kpc of galaxies \\citep{stocke06,wakker08}. These works, however did not address explicitly the origin(s) of very strong ($\\log N >14.5$)  \\ovi\\ absorbers. Hence this sightline provides the unique opportunity to study simultaneously the interaction of galaxies with their surroundings using very different tracers of gas-phases and energies and to test recent cosmological simulations.   The organization of this paper is as follows. After describing the observations and data reduction of the  PKS\\,0312--77 UV spectrum in \\S\\ref{sec-obs}, we present our analysis to estimate the redshifts, equivalent-widths ($W_\\lambda$), and column densities ($N$), of the absorbers in \\S\\ref{sec-anal}. In \\S\\ref{sec-abs} we determine the physical properties and abundances of the metal-line absorbers observed at $z \\approx 0.203$, while in \\S\\ref{sec-gal} we present our galaxies survey and discuss the relationship between the absorbers and galaxies. In \\S\\ref{sec-bigpic} we briefly discuss that current observations of the galaxy-IGM absorbers show compelling evidence that the LLS and strong \\ovi\\ absorbers trace the galaxy-intergalactic interface, i.e. the galactic environments on physical scale of tens to hundreds of kpc around galaxies. A summary of the main results is presented in \\S\\ref{sec-sum}. For the reader's information, all distances in this paper are physical separations derived from the angular diameter spaces, assuming a $\\Lambda$CDM cosmology with $\\Omega_m = 0.3$, $\\Lambda = 0.7$, and $H_0 = 70 $ \\km\\,Mpc$^{-1}$ (we use the notation $h_{70} = H_0/70\\, {\\rm km\\,s^{-1}\\,Mpc^{-1}}$).   \\begin{figure}[tbp] \\epsscale{1} \\plotone{f1.eps} \\caption{Lyman limit at  $z = 0.20258$ observed in the \\fuse\\ spectrum. The \\hi\\ lines used in our profile fitting are indicated (the Ly$\\gamma$ line is shown but not used because it is contaminated by an airglow emission line). Other absorption lines are either interstellar or other intervening IGM features.  From the flux decrement, we can place a limit on the \\hi\\ optical depth at the Lyman limit and hence on the \\hi\\ column density (the flux levels to estimate the the optical depth at the Lyman limit are shown with horizontal thick solid lines). For display purposes, the data are binned to the {\\fuse}\\ resolution, i.e. 1 pixel is about 20 \\km. \\label{fig-lls}} \\end{figure} ",
        "conclusions": "\\label{sec-sum} We have presented multi-wavelength observations of the absorbing material at $z\\approx 0.203$ along the QSO PKS\\,0312--77 and its field of view with the goal of exploring the properties of the Lyman limit system and the strongest \\ovi\\ absorber yet discovered in the low redshift Universe, and the connection between the absorbers and their environments (i.e. galaxies). The main results of our analysis are as follows:  1. Using $N($\\oi$)/N($\\hi$)$ combined with a photoionization model, we show that the LLS at $z = 0.20258$ has a metallicity of about $-0.6$ dex solar. At slightly lower redshift ($z \\approx 0.2018$, velocity separation of about $-190$ \\km), another absorber is detected with a much higher metallicity ($[{\\rm Z/H}] = +0.15$) according to our ionization models, implying that these two absorbers have different origins. The metallicity variation implies poor mixing of metals on galactic scale as observed in lower and higher redshift galactic halos.  2. The gas in both absorbers at $z \\approx 0.2018$ and $z = 0.20258$ is nearly 100\\% photoionized. But only  from $-70$ to $+150$ \\km\\ ($0.2023 \\la z \\la 0.2030$), extremely strong \\ovi\\ absorption ($W_{\\rm O\\,VI \\lambda 1032} = 493 \\pm 40$ m\\AA, $\\log N($\\ovi$) = 14.95 \\pm 0.05 $) is observed. Associated with the \\ovi, there are a detection of \\nv\\ and a tentative detection of \\svi.  At $z \\approx 0.2018$, narrow \\nv\\ absorption is detected, more consistent with the \\nv\\ originating from photoionized gas or in collisionally ionized gas far from equilibrium.  3. Using Cloudy photoionization models, we show that the high ions at $-70 \\la v \\la +150$ \\km\\ cannot be photoionized. CIE or non-equilibrium models can reproduce the observed high-ion ratios if the gas temperature is $T\\sim (3$--$10) \\times 10^5$ K. The broadenings of \\ovi\\ and \\nv\\ are consistent with such high temperatures. The high-ion profiles are broad, while the low-ion profiles reveal several narrow components. The full velocity extents of the low and high ions are, however, quite similar, as usually observed in galactic environments. If the gas of the LLS and \\ovi\\ absorber is cospatial, it is multiphase, with the photoionized gas embedded within the hot, collisionally highly ionized gas.  4. Our galaxy survey in the field of view of PKS\\,0312--77 shows that there are thirteen $0.3 \\la L/L_* \\la 1.6$ galaxies at $\\rho \\la 2 h^{-1}_{70}$ Mpc and velocity offset from the absorbers $|\\delta v| \\la 1100$ \\km, implying a group of  galaxies near the \\ovi\\ absorber at $z \\approx 0.203$. The closest galaxy (\\#1339, a $0.7 L_*$ galaxy) has only an impact parameter of $38 h^{-1}_{70}$ kpc and is offset by 16 \\km\\ from the LLS at $z = 0.20258$. There is no evidence of other $\\ga 0.1 L_*$ galaxies  within $100 h^{-1}_{70}$ kpc. From both a visual inspection of its morphology and its spectral classification (Sc or Irregular), galaxy 1339 appears to have resulted from a galaxy merger. Using diagnostics from the emission lines, we show that the metallicity of galaxy 1339 is supersolar ($[{\\rm Z/H}]_{\\rm gal} = +0.15 \\pm 0.05$) and that star formation occurs at a rate  $2\\pm 1$ M$_\\odot$ yr$^{-1}$.  5. Merger debris of galaxy 1339 is a very likely possibility for the origin of the LLS. Outflowing material from galaxy 1339 is also very probable the origin for the supersolar absorber at $z \\approx 0.2018$. The strong \\ovi\\ absorber may be a tracer  of a galaxy halo fountain. However, the presence of a group dominated by late-type galaxies and with no very bright early-type galaxy may as well suggest  that the \\ovi\\ absorber probes diffuse intragroup medium at $T \\la 10^6$ K.  6. Compiling our results with other studies, it is apparent that while the origin of the LLS is not unique (and likely includes galactic feedback, galactic halo, accreting material,...), they must play an important role in the formation and evolution of galaxies, and are among the best probes to study the galaxy-IGM interface over cosmic time. Strong \\ovi\\ absorbers associated with LLS appear good tracers of enrichment in galactic halos and intragroup medium rather than the WHIM itself."
    },
    "0812/0812.3411_arXiv.txt": {
        "abstract": "The radiation background above the ionization edge of \\ion{He}{2} varies strongly during and after helium reionization, because the attenuation length of such photons is relatively short ($\\la 40 \\Mpc$) and because the ionizing sources (quasars) are rare. Here we construct analytic and Monte Carlo models to examine these fluctuations, including, for the first time, those during the reionization era itself.  In agreement with detailed numerical simulations, our analytic model for the post-reionization Universe predicts order-of-magnitude fluctuations in the \\ion{He}{2} ionization rate $\\Gamma$.  Observations of the hardness ratio \\ion{He}{2}/\\ion{H}{1} show somewhat larger fluctuations, which may be due to more complicated radiative transfer effects.  During reionization, the fluctuations are even stronger.  In contrast to hydrogen reionization, our model predicts that regions with strong \\ion{He}{2} \\lya forest transmission should be reasonably common even during the beginning stages of reionization, because of strong illumination from nearby bright quasars.  Partly because of this, the mean ionizing background does not evolve strongly during and after helium reionization; it is roughly proportional to the filling fraction of \\ion{He}{3} regions.  On the other hand, regions full of \\ion{He}{2} and also ``fossil\" ionized regions that contain no (or few) active sources appear as strong IGM absorbers.  Their presence exaggerates the evolution of the hardness ratio, making it evolve more strongly than naively expected during the reionization era. ",
        "introduction": "\\label{intro}  To the vast majority of baryonic matter in the Universe, the most important radiation field is the metagalactic ionizing background.  As a result, a great deal of effort has gone into measuring its properties.  During most of cosmic history, the background evolved relatively slowly, but there were two major exceptions:  the reionization of hydrogen (at $z \\ga 6$) and of helium (with double ionization occurring at $z \\sim 3$).  During each of these episodes, the intergalactic medium (IGM) underwent a phase transition and the Universe became (mostly) transparent to the relevant ionizing photons, allowing the high-energy radiation field to grow rapidly.  Recently, these events have received a great deal of attention in both the observational and theoretical communities (see, e.g., reviews by \\citealt{barkana01, ciardi05-rev, fan06-review, furl06-review} about hydrogen reionization), and our understanding of the evolution of the ionizing radiation field has become increasingly sophisticated.  However, most such studies still treat the ionizing background as spatially \\emph{uniform}.  During reionization itself, this assumption is obviously wrong, because some regions are exposed to strong ionizing radiation while others remain neutral with no local illumination.  Moreover, even within the ionized regions, differences in the effective ``horizon\" to which ionizing sources can be seen induce large spatial fluctuations in the background.  Recent models have begun to address the variations expected during hydrogen reionization \\citep{mesinger08-ib, furl09-mfp}, demonstrating that they are crucial for interpreting measurements of that era and for understanding radiative feedback processes.  These fluctuations will persist even in the post-reionization era because of source clustering, stochastic fluctuations in the galaxy density, and radiative transfer effects.  Such variations are relatively small for the hydrogen-ionizing background (typically only a few percent of the mean value at $z \\sim 3$; \\citealt{zuo92a, zuo92b, fardal93, meiksin03, croft04}), principally because the mean free path of hydrogen-ionizing photons is extremely large ($\\ga 100 \\Mpc$; \\citealt{madau99-qso, faucher08-ionbkgd}) and because sources (galaxies) are relatively common.  Fluctuations are larger at $z \\ga 4$, when the mean free path is relatively small and galaxies are more highly clustered (e.g., \\citealt{meiksin03}).  These studies concluded that such fluctuations are only significant if sources are rare and only measurable with future high-precision observations.  Thus the usual assumption of a spatially constant background is reasonable.  On the other hand, fluctuations in the helium-ionizing background have received less theoretical attention (though see \\citealt{fardal98, maselli05, bolton06, meiksin07}).  Nevertheless, they are both larger and more observationally relevant than those for hydrogen.  One reason is that the IGM absorbs helium-ionizing photons more strongly than hydrogen-ionizing photons, leading to shorter attenuation lengths and larger fluctuations.  Second, \\ion{He}{2} has an ionization potential of 54.4 eV, a sufficiently high energy that quasars (with hard spectra) are required to ionize it.  As such, the sources are quite rare, and themselves quite variable, implying large random fluctuations in the background.  Moreover, helium remains singly-ionized until $z \\sim 3$, a regime that is relatively easy to observe.  A wealth of data now suggests dramatic evolution in the properties of intergalactic helium at about this time.  The strongest evidence comes from far-ultraviolet spectra of the \\ion{He}{2} \\lya forest along the lines of sight to several bright quasars at $z \\sim 3$ \\citep{jakobsen94, davidsen96, anderson99, heap00, smette02, zheng04, shull04, reimers04, reimers05, fechner06, fechner07}, which show a rapid increase in \\lya absorption by \\ion{He}{2} around that time -- although with substantial opacity fluctuations observed along several lines of sight \\citep{anderson99, heap00, smette02, reimers05}.  These are most likely a direct effect of inhomogeneities in the helium-ionizing background during (or just after) reionization.  Indirect measurements of the metagalactic ionizing background, which hardens as helium is reionized and the IGM becomes transparent to high-energy photons, suggests similar fluctuations after reionization.  In particular, measurements of the \\ion{He}{2}/\\ion{H}{1} ratio suggest that the helium-ionizing background fluctuates strongly at $z \\sim 2.6$, with nearly order-of-magnitude spatial variations on scales spanning a few to a few tens of Mpc \\citep{shull04, zheng04, fechner06, fechner07}.  Such strong fluctuations must be due to some combination of source count variations and radiative transfer effects.   \\citet{meiksin07} showed with an analytic model that the broad distribution of quasar luminosities, together with their sparseness, accounts for much of the observed variation.  \\citet{bolton06} also examined this regime with numerical simulations of quasar ionizing radiation, finding that the former effect could account for much, but not all, of the observed variance.  However, these simulations were limited to a relatively small box (less than a full attenuation volume).  \\citet{maselli05} found (again using numerical simulations) that radiative transfer effects can also cause substantial variations, although they were hampered by the assumption of spatially uniform ionizing sources.  \\citet{shull04} identified another contributing factor:  the broad distribution of quasar spectral indices, which directly affects the ionization rate.  Other observations of the \\ion{He}{2}/\\ion{H}{1} ratio suggest rapid time evolution as well \\citep{heap00}, which is qualitatively consistent with helium reionization occurring at $z \\sim 3$:  we would expect that the patchiness of the reionization process would lead to especially strong fluctuations between the bubbles of fully-ionized helium and the remaining \\ion{He}{2} lying between them \\citep{furl08-helium, mcquinn09}.  However, because of the difficulty of properly simulating helium reionization, these observations have not yet been quantitatively evaluated in light of modern reionization models.  The only previous study, by \\citet{tittley07}, showed that the complex radiative transfer during \\ion{He}{2} reionization could create wide variations in the hardness ratio.  These evolving variations should also be visible in metal lines.  For example, the ionization potentials of \\ion{Si}{4} and \\ion{C}{4} straddle that of \\ion{He}{2}, so their ratio should evolve during helium reionization.  \\citet{songaila98, songaila05} found a sudden break in that ratio at $z \\sim 3$ (see also \\citealt{boksenberg03}); modeling of the ionizing background from optically thin and optically thick metal line systems also shows a significant hardening at $z \\sim 3$ \\citep{vladilo03, agafonova05, agafonova07}.  However, other data of comparable quality show no evidence for rapid evolution \\citep{kim02, aguirre04}.  These contrasting conclusions again suggest that the ionizing background may itself be fluctuating strongly between different lines of sight.  Here we study spatial and temporal variations in the helium-ionizing background using a combination of analytic and simple Monte Carlo models.  Moreover, we examine both the relatively simple post-reionization limit (as in \\citealt{meiksin04}) and the behavior during reionization, when the fluctuations may be much stronger.  Extending the models to the reionization epoch allows us to predict the distribution of the hardness ratio, and hence the observability of the \\ion{He}{2} \\lya forest, during that era.  We will show that, in contrast to conventional wisdom, substantial transmission will remain throughout the bulk of helium reionization.  We describe our methods to compute the distribution of the amplitude of the ionizing background in \\S \\ref{method} and \\ref{attenuation}.  We describe our post-reionization results in \\S \\ref{postreion} and those during reionization in \\S \\ref{during-reion}.  Finally, we conclude in \\S \\ref{disc}.  In our numerical calculations, we assume a cosmology with $\\Omega_m=0.26$, $\\Omega_\\Lambda=0.74$, $\\Omega_b=0.044$, $H_0=100 h \\hunits$ (with $h=0.74$), $n=0.95$, and $\\sigma_8=0.8$, consistent with the most recent measurements \\citep{dunkley08,komatsu08}.   Unless otherwise specified, we use comoving units for all distances. ",
        "conclusions": "\\label{disc}  We have examined the evolution of the high-energy ionizing background during and after helium reionization.  Through a combination of analytic calculations and a Monte Carlo model, we have computed the probability distribution of $J$, the amplitude of the ionizing background at the ionization edge (which is a reasonable proxy for the ionization rate, $\\Gamma$).  After reionization is over, our model requires only the attenuation length of helium-ionizing photons and the quasar luminosity function (in the relevant energy range) as inputs.  We find that bright quasars remain rare enough at $z \\sim 2$--3 that $f(J)$ is quite broad, with a full-width-at-half-maximum comparable to $\\VEV{J}$, in agreement with \\citet{meiksin07}.  Currently, the distribution of the hardness ratio $\\eta = N_{\\rm HeII}/N_{\\rm HI}$ is the best way to measure such fluctuations \\citep{zheng04, shull04, fechner06, fechner07}.  The best-observed distribution \\citet{fechner06} is slightly broader than our models predict, especially toward high $\\eta$, but the overall agreement is good (see also \\citealt{meiksin07}).  Some of this additional broadening is undoubtedly due to more complex radiative transfer than we include here as well as variations in the spectral indices of the ionizing sources.  Note that there also appear to be observational biases affecting many of the estimates in the literature.  For example, the apparent optical depth method becomes unreliable in saturated lines \\citep{fechner06, fechner07}.  A more careful comparison of detailed theoretical models and the data is still needed in order to quantify the tension.  During reionization, quasars sit inside of discrete ionized bubbles; the ``walls\" between them are full of \\ion{He}{2} that blocks low-energy photons from more distant sources.  Thus $f(J)$ depends on the size distribution of these bubbles.  With a simple analytic model of this distribution, we have used a Monte Carlo method (that also includes attenuation) to generate $f(J)$ in this regime.  Initially, $f(J)$ is quite broad because of the range of bubble sizes (and because most bubbles are sufficiently small that a faint quasar can provide all of the local ionizing radiation field).  As the bubbles grow larger, they encompass more and more sources and the variance slowly declines.  Eventually, the distribution matches smoothly onto its post-reionization form:  after $\\bar{x}_{\\rm HeIII} \\sim 0.9$, the only significant change is the disappearance of the low-$J$ tail from isolated regions.  Thus we do \\emph{not} expect a sharp feature in the overall amplitude of the ionizing background at the completion of reionization.  We found that the mean ionizing background within illuminated regions remains roughly constant.  This contrasts with the situation during hydrogen reionization, where the same quantity increases in proportion to the characteristic bubble size.  The difference occurs because so few sources contribute to helium reionization:  only a few are needed to reach the mean post-reionization value inside a bubble.  During hydrogen reionization, on the other hand, the number of visible sources increases proportionally to the volume of the bubble.  Despite this smooth match onto the post-reionization Universe, there are fairly sharp observational probes of the reionization era.  Most importantly, we have shown that the fraction of pixels with large values of $\\eta$ ($\\ga 1000$) increases faster than linearly as $\\bar{x}_{\\rm HeII}$ increases.  Not only is the hardness ratio large inside of regions that have not yet been fully-ionized, but ionized bubbles without any active sources also rapidly recombine and become nearly opaque in \\ion{He}{2}.  This should make \\ion{He}{2} reionization \\emph{easier} to identify.  This contrasts with the post-reionization Universe, where $\\eta \\la 200$ everywhere in our model; this maximum value corresponds to the minimum ionizing background generated by distant quasars.  Our model shows that this should be a rather sharp cutoff (at least without shadowing and radiative transfer; \\citealt{tittley07}), because the probability of having zero bright sources within two attenuation lengths is very small.  Direct comparisons to data are difficult, however, because real observations inevitably contain a number of lower limits to $\\eta$ (corresponding to points where the \\ion{H}{1} optical depth is too small to measure).  A careful calibration to more detailed simulations is required to draw firm conclusions about the timing of helium reionization from the $\\eta$ distributions.  On the other hand, the high-$J$ tail of the distribution (corresponding to small $\\eta$) -- which easily reaches values several times the mean -- remains more or less intact throughout at least the latter half of reionization.  Thus, we expect that measurable pockets of transmission in the \\ion{He}{2} \\lya forest \\emph{will} persist into the reionization epoch at $z \\ga 3$:  a true \\citet{gunn65} trough will only appear when nearly all of the helium atoms are singly ionized.  This implies that ultraviolet spectra of higher-redshift quasars will continue to provide detailed information about this important transformational epoch, even at redshifts well beyond the nominal ``end\" of reionization.  Fortunately, dozens of promising targets have now been detected \\citep{zheng04-sdss, zheng08, syphers08}, and the recent installation of the Cosmic Origins Spectrograph and repair of the Space Telescope Imaging Spectrograph on the \\emph{Hubble Space Telescope} has provided two powerful instruments for such exploration.  Note that this also contrasts with hydrogen reionization, where variations in the ionizing background are small during the final phases \\citep{furl09-mfp}.  There, the huge number of galaxies responsible for reionization markedly decreases the importance of the high-$J$ tail.  Without such strongly-ionized regions, and with the high mean densities at these redshifts, the \\ion{H}{1} \\lya forest becomes completely saturated even inside of large ionized bubbles before reionization completes; the appearance of a \\citet{gunn65} trough at $z \\sim 6$ may or may not hint to us that reionization is ending at that time, but in either case the prospects for detecting substantial transmission at higher redshifts are dim.  That is not the case for helium, for which our model strongly predicts that regions of transmission will be relatively common throughout the bulk of that process.  Careful study of the \\ion{He}{2} forest during the reionization era may help to illuminate some of the physics of that time.  For example, comparing the hydrogen and helium absorption lines may help us to identify when thermal broadening is important to the forest (and hence when, and where, helium reionization photoheats the IGM; \\citealt{gleser05, furl08-igmtemp, mcquinn09, bolton08}).  The spatial scales over which the hardness ratio fluctuates will also constrain the growth of \\ion{He}{3} bubbles, and of the attenuation length.  These will no doubt be difficult measurements, but they will provide more direct information about the equation of state of the IGM than any other method.  Our model does not, of course, provide a complete picture of the ionizing background during helium reionization.  Most importantly, we only consider the effects of radiative transfer in the crudest possible manner (via a fixed attenuation length $r_0$).  Shadowing, fluctuations in the mean free path, and the large range of relevant frequencies and spectral indices will all modify the simple picture here \\citep{shull04, maselli05, tittley07}.  Another aspect that we ignore is quasar clustering, which is fairly strong \\citep{shen07, francke08}, although still probably sub-dominant to random fluctuations \\citep{mcquinn09}.  Another important detail that our model does not directly describe is the evolution of $r_0$.  We have argued that the attenuation length as seen by a single quasar does not evolve significantly (Fig.~\\ref{fig:rmax}); however, once a universal, more spatially constant background has built up, $r_0$ will converge to a more uniform value.  Reassuringly, these two approaches give similar mean free paths at $z \\approx 3$ (see also \\citealt{furl08-helium}), roughly where they must converge to each other at the end of reionization.  On the other hand, in our figures we have kept redshift constant for a straightforward comparison; if $r_0$ evolves rapidly there will be more evolution in $f(J)$ than we have shown here (although this is only important once the characteristic bubble size exceeds $r_0$).  In this paper, we have focused on how fluctuations in the helium-ionizing background affect the hardness ratio of forest lines.  Of course, the wide variation in $J$ even after helium reionization ends may be observable in other ways.  For example, it will affect the evolution of the mean optical depth in the \\ion{He}{2} forest:  strongly illuminated regions will allow substantial transmission in some regions even when most of the Universe is opaque (see, e.g., \\citealt{furl09-mfp} for a similar treatment for hydrogen reionization).  This will cause a more gradual evolution in the overall optical depth $\\tau_{\\rm eff}$ than predicted by naive reionization models.  Moreover, the high-energy radiation field can also be measured with metal lines.  The data so far are controversial:  some measurements imply a hardening in the radiation field at $z \\sim 3$ \\citep{songaila98, songaila05, agafonova07}, while others are consistent with no evolution \\citep{kim02, aguirre04}.  The complexities of radiative transfer during helium reionization may account for these \\citep{madau08}.  Alternatively, the strongly fluctuating ionizing background throughout reionization provides a natural explanation for these differences, although of course a quantitative estimate is still required."
    },
    "0812/0812.3282_arXiv.txt": {
        "abstract": "We derive the covariant formalism associated with the relativistic Sunyaev-Zel'dovich effect due to a non-thermal population of high energy electrons in clusters of galaxies. More precisely, we show that the formalism proposed by Wright in 1979, based on an empirical approach to compute the inverse Compton scattering of a population of relativistic electrons on CMB photons, can actually be re-interpreted as a Boltzmann-like equation, in the single scattering approximation. ",
        "introduction": "\\label{sec:intro}  The purpose of this paper is (i) to revisit the formulation of the relativistic Sunyaev-Zel'dovich (rSZ) distortion of the cosmic microwave background (CMB) due to a population of non-thermal high energy electrons in clusters of galaxies, and (ii) to provide a unified and more comprehensive framework to compute the rSZ effect.   Inverse Compton scattering of non-relativistic thermal electrons on cosmic microwave background (CMB) photons has been widely investigated in the literature since the 1960-1970's, notably in the context of the seminal works by Sunyaev and Zel'dovich on the SZ effect~\\cite{1972CoASP...4..173S,1980ARA&A..18..537S}, and in early studies on relativistic electron cosmic rays in high energy astrophysics (e.g.~\\cite{1965PhRv..137.1306J,1968PhRv..167.1159J,1970RvMP...42..237B}). The SZ effect generated by a population of  hot (still non-relativistic) electrons has also been addressed in detail since the 1980's (see e.g.~\\cite{1979ApJ...232..348W,1989ApJ...339..619T,1995ApJ...445...33R, 1995ARA&A..33..541R,1998ApJ...499....1C,1998ApJ...502....7I,1998ApJ...508....1S, 1997astro.ph..9065S,2001ApJ...554...74D}), while the SZ effect originating from relativistic electron cosmic rays has been discussed in e.g. Refs.~\\cite{1999ApJ...523...78M,1999PhR...310...97B,2000A&A...360..417E, 2003A&A...397...27C}. So far, these works on hot thermal or relativistic non-thermal electrons have used different approaches.    For example, in  Refs.~\\cite{1998ApJ...499....1C,1998ApJ...502....7I, 1998ApJ...508....1S,1997astro.ph..9065S,2001ApJ...554...74D}, the authors have computed the relativistic corrections to the standard thermal SZ predictions due to the hot trail of intracluster thermal electrons (including the effect from multiple scattering). They used a Fokker-Planck expansion of the collisional Boltzmann equation in terms of the electron temperature-to-mass ratio $T/m\\lesssim 0.1$, in which the full squared matrix amplitude of Compton scattering (that is given in any textbook, e.g.~\\cite{Peskin:1995ev}) was implemented for the relevant kinematics. Such an approach therefore not only encompasses but also extends the usual calculation of the standard thermal SZ effect, thus making this formalism very attractive. Nonetheless, it has not been applied yet to the case where there are two populations of electrons (one dominant semi-relativistic population and a subdominant relativistic population).   The authors of e.g. Refs.~\\cite{1979ApJ...232..348W,1989ApJ...339..619T,1995ApJ...445...33R, 1999ApJ...523...78M,2000A&A...360..417E,2003A&A...397...27C} have considered the SZ contributions coming from a population of thermal electrons plus an extra source of non-thermal electron cosmic rays, as we want to do. However, they used a more empirical approach based on a classical radiative transfer method proposed by Chandrasekhar~\\cite{1950ratr.book.....C}. Indeed, they computed the spectral shift in the CMB photon intensity due to Compton scattering processes with electrons, scattering \\emph{off} some photon frequencies (from energy $E$ to $E'$) and scattering \\emph{in} some others (from energy $E'$ to $E$). The squared matrix amplitude (that is the probability of frequency change) was expressed in the electron rest frame, leading to an expression close to its associated non-relativistic limit, such as the one originally used by Chandrasekhar. Multiple scatterings were eventually implemented by further using Chandrasekhar's radiative approach.    In Sect.~\\ref{sec:boltz}, we demonstrate that Chandrasekhar's radiative transfer method and the Boltzmann-like approach mentioned above are in fact formally equivalent in the single scattering approximation. They thus give the same results, which we show by using the full relativistic squared matrix amplitude instead of the photon frequency redistribution probability function derived by Wright~\\cite{1979ApJ...232..348W}. This therefore sketches a unified formalism for the computation of both hot thermal and relativistic non-thermal SZ effects, as subsequently found in~\\cite{2009arXiv0902.2595N}. For illustration, we derive the relativistic photon energy redististribution function in the CMB reference frame (therefore different from Wright's $P(s,\\beta)$ function) that can be used to compute the relativistic SZ effect generated by an additional population of high energy electron cosmic rays in a cluster. For the sake of completenes, we revisit in Sect.~\\ref{sec:compton} the full relativistic phase space and squared amplitude of the inverse Compton process in the relativistic limit. Our conclusions are drawn in Sect.~\\ref{sec:concl}. Numerical applications of this framework will be given in a companion paper~\\cite{barthes_etal_in_prep}. ",
        "conclusions": "\\label{sec:concl}  We have shown that the approach that Wright developed in Ref.~\\cite{1979ApJ...232..348W} was equivalent to using a Boltzmann-like equation in the single scattering approximation, already very well described in the literature (e.g.~\\cite{1998ApJ...499....1C, 1998ApJ...502....7I,1997astro.ph..9065S,2001ApJ...554...74D}). The same result has been recovered in a recent subsequent work by Nozawa and Kohyama~\\cite{2009arXiv0902.2595N}. Here, we have established the relativistic expressions valid in the CMB frame to employ within Wright's formalism, summarized in Eqs.~(\\ref{eq:boltz_to_rad}). They differ from Wright's expressions which are instead derived in the electron rest frame~\\cite{2009arXiv0902.2595N}, and which are probably more suited for numerical computations. Finally, we have recovered that the SZ shift at energy $E_k$ due to a relativistic population of electrons is given by an integral over a function which is still proportional to $I_{\\gamma}^0 (E_k)-I_{\\gamma}^0(t\\; E_k)/t^3$, where ${I_{\\gamma}}^0$ is the intensity of the pure black-body CMB spectrum and $t = E_{k'}/E_k$ is the ratio of scattered-to-focused frequencies."
    },
    "0812/0812.3557_arXiv.txt": {
        "abstract": "Measurements of Baryonic Acoustic Oscillations in galaxy surveys have been recognized as a powerful tool for constraining dark energy. However, this method relies on the knowledge of the size of the acoustic horizon at recombination derived from Cosmic Microwave Background Anisotropy measurements. This estimate is typically derived assuming a standard recombination scheme; additional radiation sources can delay recombination altering the cosmic ionization history and the cosmological inferences drawn from CMB and BAO data. In this paper we quantify the effect of delayed recombination on the determination of dark energy parameters from future BAO surveys  such as BOSS and WFMOS. We find  the impact to be small but still not negligible. In particular, if  recombination is non-standard (to a level still allowed by CMB data), but this is ignored, future surveys may incorrectly  suggest the presence of a redshift dependent dark energy component. On the other hand, in the case of delayed recombination, adding to the analysis one extra parameter describing deviations from standard recombination, does not significantly degrade the error-bars on dark energy parameters and yields unbiased estimates. This is due to the CMB-BAO complementarity. ",
        "introduction": "A key goal of modern cosmology is to investigate the nature of the dark energy component, responsible for the current accelerated expansion of the Universe. A variety of observables, from Cosmic Microwave Background (CMB) anisotropies  \\cite{wmap5komatsu,wmap5cosm}  to galaxy surveys \\cite{Tegmark:2003ud,Tegmark:2003uf,Cole:2005sx,Sanchez:2005pi,Tegmark:2006az}, will continue to be measured with increasingly refined accuracy in the next years thanks to a coordinated effort  of satellite, ground based and balloon-borne missions. Despite the fact that it accounts for about $70 \\%$ of the total energy density of the universe, dark energy is largely unclustered and is typically measured just by its effect on the evolution of the expansion history (i.e. the Hubble parameter) at redshift $z <2$. Since the cosmic expansion depends on other key parameters as curvature or matter density, the nature of dark energy can therefore be revealed only by combination of different observables and/or observations over a wide redshift range.  A recent development in the field of  large-scale galaxy surveys,  is to use the  measurement of Baryonic Acoustic Oscillations (BAO) \\cite{Eisenstein:2005su,Percival:2006gs,Percival:2007yw} signal  at $z<2$ as a standard ruler. Oscillations in the primordial photon-baryon plasma leave an imprint in the matter distribution. Since the frequency of these oscillations is related to the size of the sound horizon at recombination, which is well constrained by CMB measurements, it is possible to use the measurements of those oscillations at different redshifts as a standard ruler and therefore determine the rate of the cosmic expansion.  It has been found that a number of theoretical systematics including non-linear growth, non-linear bias, and non-linear redshift distortions can be efficiently modeled to minimize their contribution to uncertainties in the analysis of BAO observations \\cite{Eisenstein:2006nk,Seo:2007ns,Seo:2008yx}. As such, they are one of the key observables of the next decade in cosmology and it is therefore important to investigate how solid are the theoretical assumptions behind this method. A crucial assumption, indeed, is the possibility of having a very accurate meausurement of the size of the acoustic horizon at recombination. \\cite{whitenstein} show that the interpretation of low redshift BAOs is robust  if CMB correctly determines  the  baryon-to-photon ratio and the matter-radiation equality but not the matter baryon or photon densities alone. CMB measurements of these quantities is generally robust, but , in view of high-precision data, it is important to quantify any possible systematic effects introduced by deviations from the standard evolution of the early universe . Previous papers have studied, for example, the stability of the result under the hypothesis of an extra background of relativistic particles or early dark energy (\\cite{linder}). In this case, the epoch of matter-photons equality is shifted and this may bias the determination of size the acoustic horizon if unaccounted for.  In this paper we consider another possible mechanism that could, in principle, affect the BAO  interpretation as standard rulers if not modeled, namely a possible delay, or change in the recombination epoch that would result in a variation of the acoustic horizon. Several papers, recently, have indeed considered this hypothesis (see e.g. \\cite{Seager00,Naselsky02,Dorosh02,bms,bms2,gbms,Zhang:2006fr}). Dark matter decay or annihilation, black hole evaporation, or cosmic string decay can, for example, produce an extra background of resonance photons that could delay recombination (see \\cite{gbms} and references therein). More exotic mechanisms as, for example, variations in the fine structure constant or in the Newton constant, can also lead to a modified recombination epoch (\\cite{alpha,zalzahn}). We conclude that current constraints  on delayed recombination are compatible with a  change in the sound horizon which is  non-negligible when compared with the precision for future BAOs data.  This means that, if unaccounted for, delayed recombination could bias dark energy constraints from BAO. On the other hand, if delayed recombination is accounted for in the joint CMB-BAO analysis form forthcoming surveys,  error-bars on dark energy parameters are only degraded by  less than $10$ \\% compared to the standard analysis.  The paper is organized as follows: in the next section we discuss delayed recombination and its effect on the sound horizon. In Sec. III we review the Fisher matrix formalim, while in Sec. IV we forecast the impact of delayed recombination on dark energy parameters inference for several future surveys. Finally, in Sec. IV, we draw together our findings and discuss the conclusions and their implications for future work. ",
        "conclusions": "In this paper we have investigated how a biased determination of the sound horizon due to a incorrect recombination model affects cosmological parameters measurements  from future BAO data. We have shown that assuming a standard recombination model when the true model is different can bias the measure of  the CMB sound horizon and the best fit value of cosmological parameters such as the Hubble parameter. This shift propagates in  a shift on the values of $D_A(z)$ and $H(z)$ and hence on the best fit values of $w_0$ and $w_1$ determined from BAO  surveys. We have shown that a deviation from standard recombination which is still allowed by current CMB data, propagates in a bias in $w_0$ and $w_1$  comparable  to  their $1\\sigma$ statistical error from forthcoming surveys. An incorrect calibration of sound horizon due to a wrong assumption for the recombination model could bias estimation of dark energy parameters by 10\\% or more. In this case  future surveys may thus  incorrectly  suggest the presence of a redshift dependent dark energy component.  To recover an unbiased  best fit value  for dark energy parameters the correct recombination model must be used or must be described by the set of cosmological parameters used in the analysis. We have concentrated on the case of delayed recombination where extra sources of resonance (Lyman-${\\alpha}$) photons modify the evolution of the ionization fraction. This is well modeled by the addition of a single extra parameter. We have employed a Fisher matrix formalism to forecast errors and parameters biases for forthcoming CMB (Planck) and BAO (BOSS, WFMOS) surveys. For this data set combination, introducing the additional  degree of freedom describing deviations from standard recombination does not significantly degrade the error-bars on dark energy parameters and yields unbiased estimates.  This is due to the CMB-BAO complementarity: forthcoming data sets will thus disentangle effects of non-standard recombination from the effects of dark energy."
    },
    "0812/0812.2598_arXiv.txt": {
        "abstract": "We show here that the absorptive H$_\\alpha$ polarized line profile previously seen in many Herbig Ae/Be (HAeBe) stars is a nearly ubiquitous feature of other types of embedded or obscured stars. This characteristic 1\\% linear polarization variation across the absorptive part of the H$_\\alpha$ line is seen in Post-AGB stars  as well as RV-Tau, $\\delta$-Scuti, and other types. Each of these stars shows evidence of obscuration by intervening circumstellar hydrogen gas and the polarization effect is in the absorptive component, consistent with an optical pumping model. We present ESPaDOnS spectropolarimetric observations of 9 post-AGB and RV-Tau types in addition to many multi-epoch HiVIS observations of these targets. We find significant polarization changes across the H$_\\alpha$ line in 8/9 stars with polarization amplitudes of 0.5\\% to over 3\\% (5/6 Post-AGB and 3/3 RV-Tau). In all but one of these, the polarization change is dominated by the absorptive component of the line profile. There is no evidence that subclasses of obscured stars showing stellar pulsations (RV-Tau for Post-AGB stars and $\\delta$-Scuti for Herbig Ae/Be stars) show significant spectropolarimetric differences from the main class. Significant temporal variability is evident from both HiVIS and ESPaDOnS data for several stars presented here: 89 Her, AC Her, SS Lep, MWC 120, AB Aurigae and HD144668. The morphologies and temporal variability are comparable to existing large samples of Herbig Ae/Be and Be type stars. Since Post-AGB stars have circumstellar gas that is very different from Be stars, we discuss these observations in the context of their differing environments. ",
        "introduction": "Interpreting observed linear spectropolarimetric variations across a line profile, like H$_\\alpha$, has been difficult. Many phenomena such as disks, winds, magnetic fields, an asymmetric radiation field, and scattering asymmetries can produce linear polarization changes across a spectral line (cf. McLean 1979, Wood et al. 1993, Wood \\& Brown 1994, Harries 2000, Ignace et al. 2004, Vink et al. 2005a, Kuhn et al. 2007). We have observed a common morphological feature in the linearly polarized H$_\\alpha$ spectrum of most HAeBe stars which should be present in many other stellar systems. We confirm this expectation here and demonstrate the potential for learning about the near-star environment of embedded stars from high spectral resolution and good spectropolarimetric sensitivity ($<$0.1\\% accuracy).  There are now many detections of linear polarization signals associated with the absorptive component of the H$_\\alpha$ line profile. With high spectral resolution (R$>$10000) there are comparatively large amplitude (0.2-2\\%) signals in Herbig Ae/Be stars, Be's and other emission-line stars (Kuhn et al. 2007, Harrington 2008, Harrington \\& Kuhn 2009). The absorptive polarization effects are ubiquitous in the Herbig Ae/Be stars with roughly 2/3 of the 'windy' and 'disky' stars showing this linear polarization signature. The classical Be stars more typically showed a broad `depolarization' morphology (10/30) but about half of these show additional absorptive effects (4/10) and another sub-set (5/30) show complex linear polarization spectral morphologies. The presence of absorptive polarization effects in most of the obscured stars in this survey suggests that the phenomenon may be present in any obscured star. In the optical pumping model described by Kuhn et al. 2007, the polarized absorption effect is present anywhere there is a difference between the line of sight to the obscuring material and the direction from the material to the star. In order to explore this hypothesis, observations of other obscured stars have been performed here.  Post-AGB stars and the associated RV-Tau type stars are often heavily obscured and have a very dynamic circumstellar environment (see van Winckel 2003 for a review). Mass-loss, shells and circumstellar envelopes are common and characteristic of the very different environments of Post-AGB, Herbig Ae/Be and Be stars. There are lower resolution spectropolarimetric studies of many different objects (cf. Beiging et al. 2006, Oppenheimer et al. 2005, Oudmaijer \\& Drew 1999, Pereyra et al. 2009, Trammell 1994, Vink et al. 2002, 2003, 2005b). A comparison between spectropolarimetric morphologies and magnitudes can illuminate the important polarization mechanisms in different stellar environments.  \\begin{figure*} [!htb] \\includegraphics[width=0.23\\linewidth, angle=90]{f1a.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{f1b.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{f1c.eps} \\\\ \\includegraphics[width=0.23\\linewidth, angle=90]{f1d.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{f1e.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{f1f.eps} \\caption{\\label{fig:poagb} The ESPaDOnS archive H$_\\alpha$ spectropolarimetry for the Post-AGB stars. From left to right - V382 Aur (HD 46703), PS Gem (HD 52961), AG Ant (HR 4049), HD 108015, LN Hya (HR 4912) and 89 Her. See table \\ref{stars} for the log of observations. V832 Aur and PS Gem show a simple near-line-center excursion in a mostly-symmetric line profile. LN Hya with a similar line profile does not show a significant detection. HD 108015 has an asymmetric profile but also shows this single excursion at line center. AG Ant and 89 Her show P-Cygni type profiles but with vary different widths. 89 Her spectropolarimetry is dominated by an absorptive effects whereas AG Ant has a very narrow signature. } \\end{figure*} ",
        "conclusions": "Spectropolarimetric signatures have been detected in Post-AGB stars at a very high frequency - 8/9. The magnitude of the detected spectropolarimetric effects go from 0.5\\% up to over 3\\% and is quite similar to the polarization-in-absorption seen in previous studies. There very significant absorptive spectropolarimetric variability in 89 Her, AC Her and U Mon with possible variability in PS Gem. This is similar to the variability seen in the Herbig Ae/Be stars (AB Aurigae, MWC 120 and MWC 480) as well as SS Lep, an A-type star showing a strong P-Cygni type H$_\\alpha$ profile. There is no evidence that stellar pulsations (RV-Tau types for Post-AGB stars or $\\delta$-Scuti types for Herbig Ae stars) have any influence on spectropolarimetric morphology. The short (hour to week) time-scale variability of the polarized spectra of many of our target stars is consistent with a near-stellar location for the absorbing/polarizing gas -- as the optical pumping model requires.  The Herbig Ae/Be stars shown in Harrington \\& Kuhn 2009 did not fit any of the typical scattering morphologies and were dominated by absorptive polarimetric effects. The disk-scattering models produce double-peaked spectropolarimetric morphologies that must be broader than the entire H$_\\alpha$ emission line, contrary to what we observe. The classical depolarization signature produces a smooth broad line profile change due to unpolarized emission which dilutes the continuum polarization. The continuum is polarized by scattering in a flattened circumstellar envelope. This depolarization signature can be modified by absorption - circumstellar material outside the envelope and emission region can selectively absorb unpolarized stellar light or polarized envelope light modifying the depolarization effect. In this case, the circumstellar material must be outside the H$_\\alpha$ formation region and it must act on an effect that is formed over the entire line.    \\begin{figure*} [!h] \\begin{centering} \\includegraphics[width=0.35\\linewidth, angle=90]{f11a.eps} \\includegraphics[width=0.35\\linewidth, angle=90]{f11b.eps}  \\\\ \\includegraphics[width=0.35\\linewidth, angle=90]{f11c.eps} \\includegraphics[width=0.35\\linewidth, angle=90]{f11d.eps} \\\\ \\caption{\\label{scomp} Comparison spectropolarimetry with Herbig Ae/Be and other B-type stars. Spectropolarimetry taken with the HiVIS spectropolarimeter for the Be stars MWC 143 and $\\psi$ Per are shown in the top left and right panels. In both stars there are strong central absorptive components but a strong absorptive effect is only seen in $\\psi$ Per with a small deviation from the broad effect present in MWC 143. HiVIS spectropolarimetry for the Be star $\\alpha$ Col is seen in the bottom left. Though there is an obvious central absorptive component to the H$_\\alpha$ line, the spectropolarimetric signature spans most of the line width. ESPaDOnS H$_\\alpha$ spectropolarimetry for the Herbig Ae star MWC 480 binned to a regular wavelength grid and binned by an additional factor of four is shown in the bottoom right panel. The panel shows three independent measurements. A very clear polarimetric signature is seen in and around the absorptive components of the H$_\\alpha$ line. This is entirely typical of the Herbig Ae/Be polarization-in-absorption detections. } \\end{centering} \\end{figure*}   Figure \\ref{scomp} shows examples of spectropolarimetric morphologies for other Herbig Ae/Be stars and Be/Emission-Line stars. In Harrington \\& Kuhn 2009, all but one of the Herbig Ae/Be stars showed a detectable signature that correlates with the absorptive components of the line. Most effects were entirely contained inside absorptive components of the line. Roughly 2/3 of the HAe/Be stars with 'windy' or 'disky' H$_\\alpha$ profiles (strongly blue-shifted or central absorptive components respectively) showed signatures of similar magnitude and morphology as these Post-AGB stars. The Be stars typically (10/30) showed a broad depolarization type signature, as shown in figure \\ref{scomp} for MWC 143. However, 4 of these 10 detections showed more complex additional effects correlating with absorptive components of the H$_\\alpha$ line (e.g.   $\\psi$ Per in figure \\ref{scomp}). A smaller sub-set of B-type stars (5/30) showed smaller amplitude and more complex morphologies like $\\alpha$ Col. This star had a lower amplitude signature that was entirely stable over two consecutive nights (at two identical telescope pointings).  In the optical pumping theory of Kuhn et al. 2007, the absorptive polarization is greatest where the obscuring material is at a large star-cloud-telescope angle. In the classical P-Cygni profile, this occurs in the red-most region of the blue-shifted absorptive component. In the case of a thin circumstellar disk, this occurs at the extreme red and blue shifts of the central absorptive component of the line. In the very diverse group of stars presented here with P-Cygni profiles (AG Ant, 89 Her, SS Lep and AB Aurigae) the spectropolarimetric effects are clearly blueshifted and vary strongly with time. The effects are correlated with the presence of absorption as expected in this model.  There are no detailed spectropolarimetric models which yet account for the many phenomena thought to be present in circumstellar environments. However, simple morphological arguments can illuminate the general physical circumstellar properties. The observations presented here clearly show a remarkable similarity between the spectropolarimetric morphologies of Post-AGB and Herbig Ae/Be stars. The optical pumping model naturally reproduces the, apparently quite common, absorptive spectropolarimetric signatures (8/9 P-AGB). In that model, this signal originates from material along the line of sight to the star within 1 or 2 stellar radii of its photosphere (c.f. Kuhn et al. 2007). Though the model did not include non-photospheric emission sources, the geometry can be easily scaled up to a cloud obscuring an extended emission region with a similar scaling of the region producing the spectropolarimetric effect. This is consistent with the short time-scale variability seen in these targets. Post-AGB type stars are apparently similar to Herbig Ae/Be in that they show similar wavelength dependent absorptive polarization. The absorptive quality of the dominant linear polarization makes it unlikely that scattering models can account for this spectropolarimetric morphology. These observations extend our detection of the absorptive polarization spectropolarimetric effect to another class of stars and provide further evidence that line-of-sight absorption occurring near a star causes a linear spectropolarimetric signal at absorptive line profile wavelengths."
    },
    "0812/0812.2883_arXiv.txt": {
        "abstract": "IceCube is a cubic neutrino telescope under construction at the South Pole since the austral summer 2004/2005 with a total instrumented volume of the order of 1 ${\\rm km}^{3}$. At the moment it is taking data with 40 deployed strings. The full detector is expected to be completed in 2011 with up to 80 strings holding 60 digital optical modules (DOMs) each. The progenitor detector AMANDA has been operating at the same site since 1997 and is still functioning as a means to enhance neutrino effective area at energies below 100 GeV. A summary of science results and status of the project is presented. \\vspace{1pc} ",
        "introduction": "\\subsection{High-Energy Neutrino Astrophysics}  Neglecting gravity, neutrinos interact only by way of the weak interaction. This precludes absorption by matter and radiation fields which affect photons and therefore potentially allows observation of objects and processes that are not accessible to conventional gamma-ray astronomy. Targets of interest include the neighborhood of black holes and active galaxies, where photonic emissions are subject to absorption by ambient matter and extragalactic background light. The obvious disadvantage resulting from the restriction to weak interactions is the need for very large detector volumes to achieve at least a moderate event yield.  One of the longest-standing problems in astrophysics is the search for the origin of cosmic radiation \\cite{Halzen:2002pg}. Since the direction of charged cosmic rays, except at the very highest energies, effectively gets scrambled by magnetic fields, any information about their origin must come from electrically neutral particles, i.e. photons and neutrinos. Investigation of this issue is one of the main goals of Very High Energy (VHE) neutrino detectors, with instrumented volumes up to one cubic kilometer. These take advantage of Cherenkov radiation of charged particles produced in interactions of neutrinos with ambient matter. The large fiducial volume required for the detection of the comparatively low neutrino fluxes at TeV-PeV energies makes it necessary to take advantage of large natural occurrences of light-transparent media, such as the deep sea or polar ice sheets.  So far, no VHE neutrino signal from outside the solar system has been detected. This situation is expected to change after completion of a new generation of detectors with an instrumented volume of the order of 1 ${\\rm km}^3$ \\cite{Halzen:2007ip}.  \\begin{figure}[htb] \\includegraphics{icecube_scheme.pdf} \\caption{Schematic view of the IceCube detector, showing the three currently operating components: IceTop surface air shower array, AMANDA detector and InIce strings. As of 2008, IceTop and InIce are both completed to 50\\% (40/80).} \\label{fig:icecube_scheme} \\end{figure} ",
        "conclusions": "The transition from AMANDA to IceCube has so far been remarkably successful. First physics results have been published, validating the excellent performance of the new detector. As IceCube continues to grow with every deployment season, increases in both volume and angular resolution should soon allow results to significantly surpass those from AMANDA. An important symbolic milestone for IceCube was reached in 2008, as integrated exposure exceeded $1{\\rm km}^3{\\rm yr}$. Data becoming available over the course of the next few years will allow important physics investigations, such as a probe of the diffuse Waxman-Bahcall flux and exclusion or confirmation of various emission models for individual cosmic objects."
    },
    "0812/0812.2284_arXiv.txt": {
        "abstract": "We investigate the relationship between black hole mass and bulge luminosity for AGNs with reverberation-based black hole mass measurements and bulge luminosities from two-dimensional decompositions of {\\em Hubble Space Telescope} host galaxy images.  We find that the slope of the relationship for AGNs is $0.76 - 0.85$ with an uncertainty of $\\sim 0.1$, somewhat shallower than the $M_{\\rm BH} \\propto L^{1.0 \\pm 0.1}$ relationship that has been fit to nearby quiescent galaxies with dynamical black hole mass measurements.  This is somewhat perplexing, as the AGN black hole masses include an overall scaling factor that brings the AGN $M_{\\rm BH} - \\sigma_{\\star}$ relationship into agreement with that of quiescent galaxies.  We discuss biases that may be inherent to the AGN and quiescent galaxy samples and could cause the apparent inconsistency in the forms of their $M_{\\rm BH} - L_{\\rm bulge}$ relationships. ",
        "introduction": "Most galactic bulges are now believed to harbor a massive black hole. For AGNs, the evidence of the black hole is obvious from the activity. However, even in quiescent galaxies, the effect of the black hole can be detected from stellar and gas kinematics near the nucleus.  What remains to be determined is the formation mechanism for these massive black holes and their role in shaping and responding to the evolution of their host galaxies.  In an early review on the subject of quiescent galaxy black hole masses, \\citet{kormendy95} pointed out that the estimated central black hole masses of eight galaxies seemed to show a correlation with the host galaxy bulge luminosities (or, equivalently, bulge stellar masses).  \\citet{magorrian98} later investigated a sample of 32 nearby galaxies and their black hole masses, and confirmed that the estimated black hole mass in each galaxy was indeed proportional to the luminosity (or mass) of the host galaxy bulge, albeit with a scatter about the fit of approximately $\\pm 0.5$\\,dex.  Later studies claimed, in some cases, that there may be a difference in the black hole--bulge relationship for quiescent and active galaxies, or even for Seyfert galaxies and quasars (e.g., \\citealt{wandel99a}).  Subsequent investigations seem to have decreased these discrepancies through more sophisticated techniques of measuring the bulge luminosity such as 2-dimensional image decompositions (e.g., \\citealt{mclure01,wandel02} [hereafter W02]) or dynamical modeling of the host galaxy (e.g., \\citealt{haring04}).  We have recently completed a {\\em Hubble Space Telescope} campaign to image the host galaxies of AGNs with black hole masses from reverberation-mapping.  The images were acquired for the purpose of decomposing the surface brightness profiles of the host galaxies and creating ``nucleus-free'' images from to measure the host-galaxy starlight contributions to ground-based spectroscopic luminosity measurements of the AGNs (results described by \\citealt{bentz06a,bentz08b}).  Combining the bulge luminosities estimated from the surface brightness decompositions of these high-resolution images with the recently updated and homogeneously analyzed database of reverberation masses for these objects (\\citealt{peterson04,grier08}), we re-examine the relationship between black hole mass and host galaxy bulge luminosity for nearby AGNs. Throughout this work, we will assume a standard flat $\\Lambda$CDM cosmology with $\\Omega_{\\rm B} = 0.04$, $\\Omega_{\\rm DM} = 0.26$, $\\Omega_{\\Lambda} = 0.70$, and $H_0 = 70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$.    \\begin{deluxetable*}{lccccccc} \\tablecolumns{8} \\tablewidth{500pt} \\tablecaption{Black Hole Masses and Bulge Luminosities} \\tablehead{ \\colhead{Object} & \\colhead{$D_L$} & \\colhead{$E(B-V)$} & \\colhead{$m_{HST}$} & \\colhead{$m_V$} & \\colhead{$M_V$} & \\colhead{$\\log L$} & \\colhead{$\\log M_{\\rm BH}$}\\\\ \\colhead{} & \\colhead{(Mpc)} & \\colhead{(mag)} & \\colhead{(vegamag)} & \\colhead{(vegamag)} & \\colhead{(mag)} & \\colhead{($L_{\\odot}$)} & \\colhead{($M_{\\odot}$)}}  \\startdata Mrk\\,335      &\t113  &  0.035  & 16.16  & 16.26  &  $-19.00$  &  9.53\t      & $7.15 \\pm 0.11$  \\\\ PG\\,0026+129  &\t672  &  0.071  & 16.21  & 16.23  &  $-22.91$  &  11.09       & $8.59 \\pm 0.11$  \\\\ PG\\,0052+251  &\t740  &  0.047  & 17.81  & 17.92  &  $-21.42$  &  10.50       & $8.57 \\pm 0.09$  \\\\ Fairall\\,9    &\t209  &  0.027  & 15.13  & 15.20  &  $-21.40$  &  10.49       & $8.41 \\pm 0.10$  \\\\ Mrk\\,590      &\t115  &  0.037  & 15.59  & 15.68  &  $-19.63$  &  9.78, 9.99  & $7.68 \\pm 0.07$  \\\\ 3C\\,120\t      &\t145  &  0.297  & 16.64  & 16.72  &  $-19.08$  &  9.57\t      & $7.74 \\pm 0.21$  \\\\ Ark\\,120      &\t142  &  0.128  & 14.50  & 14.59  &  $-21.17$  &  10.40       & $8.18 \\pm 0.06$  \\\\ Mrk\\,79\t      &\t96.7 &  0.071  & 15.84  & 15.95  &  $-18.98$  &  9.52, 10.19 & $7.72 \\pm 0.12$  \\\\ PG\\,0804+761  &\t461  &  0.035  & 16.70  & 16.74  &  $-21.58$  &  10.56       & $8.84 \\pm 0.05$  \\\\ PG\\,0844+349  &\t287  &  0.037  & 16.75  & 16.81  &  $-20.48$  &  10.13       & $7.97 \\pm 0.18$  \\\\ Mrk\\,110      &\t155  &  0.013  & 18.08  & 18.16  &  $-17.79$  &  9.05\t      & $7.40 \\pm 0.11$  \\\\ PG\\,0953+414  &\t1170 &  0.013  & 17.64  & 17.82  &  $-22.53$  &  10.94       & $8.44 \\pm 0.09$  \\\\ PG\\,1211+143  &\t368  &  0.035  & 17.16  & 17.21  &  $-20.62$  &  10.18       & $8.16 \\pm 0.13$  \\\\ PG\\,1226+023  &\t758  &  0.021  & 15.49  & 15.61  &  $-23.79$  &  11.45       & $8.95 \\pm 0.09$  \\\\ PG\\,1229+204  &\t283  &  0.027  & 17.17  & 17.23  &  $-20.03$  &  9.94, 9.98  & $7.86 \\pm 0.21$  \\\\ PG\\,1307+085  &\t740  &  0.034  & 16.66  & 16.70  &  $-22.65$  &  10.99       & $8.64 \\pm 0.12$  \\\\ Mrk\\,279      &\t133  &  0.016  & 16.13  & 16.22  &  $-19.40$  &  9.69, 10.00 & $7.54 \\pm 0.11$  \\\\ PG\\,1411+442  &\t410  &  0.008  & 16.70  & 16.74  &  $-21.32$  &  10.46       & $8.65 \\pm 0.14$  \\\\ PG\\,1426+015  &\t395  &  0.032  & 15.36  & 15.34  &  $-22.64$  &  10.99       & $9.11 \\pm 0.13$  \\\\ Mrk\\,817      &\t138  &  0.007  & 17.62  & 17.71  &  $-17.99$  &  9.13\t      & $7.69 \\pm 0.07$  \\\\ PG\\,1613+658  &\t606  &  0.027  & 16.14  & 16.20  &  $-22.71$  &  11.02       & $8.45 \\pm 0.20$  \\\\ PG\\,1617+175  &\t522  &  0.042  & 16.36  & 16.34  &  $-22.25$  &  10.83       & $8.77 \\pm 0.10$  \\\\ PG\\,1700+518  &\t1510 &  0.035  & 17.54  & 17.67  &  $-23.22$  &  11.22       & $8.89 \\pm 0.10$  \\\\ 3C\\,390.3     &\t251  &  0.071  & 16.74  & 16.80  &  $-20.19$  &  10.01       & $8.46 \\pm 0.10$  \\\\ Mrk\\,509      &\t151  &  0.057  & 13.96  & 13.96  &  $-21.94$  &  10.71       & $8.16 \\pm 0.04$  \\\\ PG\\,2130+099  &\t283  &  0.044  & 18.73  & 18.79  &  $-18.47$  &  9.32\t      & $7.58 \\pm 0.17$  \\\\  \\enddata  \\tablecomments{For those objects with two bulge luminosities listed, the first is the luminosity of the largest-scale non-disk component, and the second is the luminosity of all non-disk components, including bars or inner bulges.  Black hole masses are from \\citet{peterson04}, except for PG\\,2130+099 \\citep{grier08}.}  \\end{deluxetable*} ",
        "conclusions": "The compendium of AGN black hole masses and bulge luminosities presented here offers several advantages over previous measurements compiled from the literature (i.e., \\citealt{wandel99a,wandel02,mclure01}).  The masses result from a homogeneous analysis of reverberation-mapping data \\citep{peterson04}, in contrast to many recent studies of AGN black hole -- bulge relationships where black hole masses are inferred from single-epoch spectral measurements (e.g., \\citealt{greene08,kim08}).  The bulge luminosities in this work are estimated from two-dimensional surface brightness decompositions of unsaturated high-resolution space-based images taken with the same instrument and the same filter.  The exceptions are the five objects that were imaged through the F547M filter using WFPC2, but the bandpass is very similar to the ACS F550M filter employed for the other objects ($\\lambda_{\\rm c} {\\rm (F550M)}=5580$\\,\\AA\\ versus $\\lambda_{\\rm c}{\\rm (F547M)}=5483$\\,\\AA, and $\\Delta \\lambda{\\rm (F550M)} = 547$\\,\\AA\\ versus $\\Delta \\lambda{\\rm (F547M)}=483$\\,\\AA).   \\begin{deluxetable}{lccc} \\tablecolumns{4} \\tablewidth{220pt} \\tablecaption{Fits to the $M_{\\rm BH} - L_{\\rm bulge}$ Relationship} \\tablehead{ \\colhead{Sample} & \\colhead{$K$} & \\colhead{$\\alpha$} & \\colhead{Scatter\\tablenotemark{a}}}  \\startdata  \\multicolumn{4}{c}{BCES} \\\\ \\hline \\\\ AGNs                            & $-0.02 \\pm 0.06$  & $0.80 \\pm 0.09$  & \\\\ AGNs (+ extra components )      & $-0.07 \\pm 0.08$  & $0.85 \\pm 0.11$  & \\\\ FF05 Ellipticals                & $0.42 \\pm 0.12$   & $1.43 \\pm 0.21$  & \\\\ FF05 Ellipticals ($-$ outliers) & $0.30 \\pm 0.10$   & $1.42 \\pm 0.24$  & \\\\  \\\\ \\hline \\\\ \\multicolumn{4}{c}{FITEXY} \\\\ \\hline \\\\  AGNs                            & $-0.05 \\pm 0.06$  & $0.76 \\pm 0.08$  & 0.38 \\\\ AGNs (+ extra components)       & $-0.13 \\pm 0.06$  & $0.80 \\pm 0.09$  & 0.44 \\\\ FF05 Ellipticals                & $0.15 \\pm 0.11$   & $1.11 \\pm 0.21$  & 0.73 \\\\ FF05 Ellipticals ($-$ outliers) & $0.14 \\pm 0.11$   & $1.18 \\pm 0.19$  & 0.64 \\\\  \\enddata  \\tablenotetext{a}{The fractional scatter, quantified as the fraction of the measurement value of $M_{\\rm BH}$ that must be added in quadrature to the error value in order to obtain a reduced $\\chi^2$ of $1.0$.}  \\end{deluxetable}    Unfortunately, there is not a similarly consistent sample of high-quality images from which the bulge luminosities can be estimated for the quiescent galaxies with dynamical black hole masses.  While high-quality observations do exist for these nearby and well-studied galaxies, they have not been obtained in a uniform fashion.  The recent work by \\citet{graham07} attempts to compensate for the different methods and analysis techniques employed in several publications (\\citealt{mclure02b} with an updated cosmology presented in \\citealt{mclure04,marconi03,erwin04}), all of which arrive at different values for the slope of the quiescent galaxy $M_{\\rm BH} - L_{\\rm bulge}$ relationship and/or different black hole masses predicted for a specific bulge luminosity. \\citeauthor{graham07} carefully updated and revised the samples of objects included in these studies and found that the differences of the best-fit parameters found by each study are mitigated, the scatter in the measurements is significantly decreased, and that $M_{\\rm BH} \\propto L^{1.0}$. Interestingly, he finds a somewhat shallower slope, $\\alpha \\approx 0.75 - 0.88$, when bulge luminosities are estimated by two-dimensional decompositions of $B$-band images and corrected for inclination-dependent dust extinction in the host galaxy disks.  Only 13 objects are included in that particular analysis (an updated form of the study presented by \\citealt{erwin04}) but it is intriguing nonetheless in its close agreement with the fit that we find for the AGNs.  While we expect that the AGN bulge luminosities in this may be slightly underestimated based on the conservative galaxy fits we employed, there is also a small $k$-correction introduced, as the average redshift of the 26 AGNs is $z \\approx 0.1$.  The portion of the SED observed through the medium-band $V$ filters employed in the {\\em HST} observations is fainter ($\\sim 0.1$\\,dex in luminosity) than if the galaxies were at $z=0$, assuming their stellar populations resemble that of the bulge template of \\citet{kinney96}.  Accounting for any such biases in the galaxy fitting or colors would intensify the apparent differences between the slopes measured for the AGNs and quiescent galaxies.  The black hole masses for the AGNs have already been scaled so that their $M_{\\rm BH} - \\sigma_{\\star}$ relationship is brought into agreement with the quiescent $M_{\\rm BH} - \\sigma_{\\star}$ relationship fit to the same sample of quiescent galaxies included in the FF05 study.  The differences may be mitigated if \\citet{marconi08} are correct in their suggestion that neglecting radiation pressure leads to systematic under-estimation of black hole masses from reverberation mapping data (see, however, \\citealt{netzer08}).   \\begin{figure} \\epsscale{1.2} \\plotone{f3.eps} \\caption{The $M_{\\rm BH} - L_{\\rm bulge}$ relationship for quiescent galaxies as determined for the elliptical galaxies from \\citeauthor{ferrarese05} (2005; filled triangles) compared to the $M_{\\rm BH} - L_{\\rm bulge}$ relationship for AGNs from this work.  The slope of the relationship appears shallower for AGNs than for quiescent galaxies.  The solid line is the same fit displayed in Figure~2 for the sample of AGNs in this work, and has a slope of $\\alpha = 0.80 \\pm 0.09$.  The short-dashed line is the BCES fit to the elliptical galaxies and has a slope of $1.43 \\pm 0.21$, while the long-dashed line is the FITEXY fit to the elliptical galaxies and has a slope of $1.11 \\pm 0.21$. } \\end{figure}   A separate concern has been raised by \\citet{yu02}, \\citet{bernardi07}, \\citet{lauer07a}, and \\citet{tundo07}, who present an apparent disagreement between the quiescent galaxy $M_{\\rm BH} - \\sigma_{\\star}$ and $M_{\\rm BH} - L_{\\rm bulge}$ relationships and suggest that the quiescent galaxy sample is biased towards galaxies with overly large velocity dispersions for their luminosities (see, however, \\citealt{graham08} who examines the role of bars in this issue).  There is no reason to suspect the AGN sample of having the same bias, as the mass measurements are made using flux variability techniques and not dynamical techniques.  Such a bias may help explain why some of the quiescent galaxies at the high luminosity end have black hole masses that are more than an order-of-magnitude larger than the active galaxies, although it may not completely resolve this disparity.  Finally, there may be no reason to expect that the $M_{\\rm BH} - L_{\\rm bulge}$ relationship is the same for the AGNs and quiescent galaxies in these samples, as there is only a modest number of objects in each sample and selection effects likely play an important role on both sides \\citep{lauer07b}.  Clearly, there remain several areas that are in need of investigation, any of which may shed light on the apparently inconsistent fits to the $M_{\\rm BH} - L_{\\rm bulge}$ relationship for AGNs and for quiescent galaxies.  As the $M_{\\rm BH} - L_{\\rm bulge}$ relationship is an important and widely used means of estimating black hole masses throughout cosmic history (e.g., \\citealt{marconi04,shankar04}), an accurate characterization of this relationship is necessary for understanding black hole growth and evolution as well as the interplay between black holes and their host galaxies."
    },
    "0812/0812.0281_arXiv.txt": {
        "abstract": "We investigate the infrared / radio correlation using the technique of source stacking, in order to probe the average properties of radio sources that are too faint to be detected individually.  We compare the two methods used in the literature to stack sources, and demonstrate that the creation of stacked images leads to a loss of information.  We stack infrared sources in the {\\it Spitzer} extragalactic First Look Survey (xFLS) field, and the three northern {\\it Spitzer} Wide-area Infrared Extragalactic survey (SWIRE) fields, using radio surveys created at 610~MHz and 1.4~GHz, and find a variation in the absolute strength of the correlation between the xFLS and SWIRE regions, but no evidence for significant evolution in the correlation over the 24-$\\umu$m flux density range 150~$\\umu$Jy -- 2~mJy.  We carry out the first radio source stacking experiment using 70-$\\umu$m-selected galaxies, and find no evidence for significant evolution over the 70-$\\umu$m flux density range 10~mJy -- 100~mJy. ",
        "introduction": "There is a well-known relationship between the infrared and radio flux densities of star-forming galaxies \\citep*[the `infrared / radio correlation', e.g.][]{Helou85,Appleton04} that is thought to be the result of both the infrared and radio emission from star-forming galaxies being related to the rate of massive star formation.  Recent studies of the infrared / radio correlation within nearby galaxies \\citep[e.g.][]{Murphy06} give weight to this argument, as the radio emission has been shown to represent a smeared out version of the infrared emission.  This has been attributed to the diffusion of cosmic ray electrons away from star-forming regions.  The correlation is seen to apply to a wide range of galaxies associated with star-formation, but not to those sources which are associated with Active Galactic Nuclei (AGN) activity \\citep[e.g.][]{deJong85,Sanders88,Sopp91,Roy98}.  Studies at high redshift are tentatively finding that the infrared / radio correlation does not deviate significantly from that seen in the local Universe, at least out to $z \\sim 1$ (\\citealp[e.g.][]{Garrett02,Appleton04,Frayer06}), and potentially to $z \\sim 3.5$ \\citep{Ibar08}.  The correlation has been shown to remain invariant over more than 5 orders of magnitude \\citep*[e.g.][]{Yun01}, however studies that rely on samples of sources which are detected at both infrared and radio wavelengths will naturally be mainly sampling luminous galaxies, and there has been relatively little work carried out on testing the infrared / radio correlation within fainter galaxies.  The technique of source `stacking' is well established as a method for combining data from many individual objects in order to study the statistical properties of sources which would otherwise be below the detection limit for a particular survey.  By constructing a list of source positions based on prior information obtained at another frequency, a series of small `cut-out' images can be created, centred on these source positions, of sources that may be below the noise level (and therefore undetected).  If $N$ of these images are combined, each with a local noise of $\\sigma$, then the noise level of the stacked image will be expected to decrease approximately as $\\sigma / \\sqrt{N}$, which quickly allows sources below the original detection threshold of the survey to be studied.  Stacking has previously been carried out using optical \\citep[e.g.][]{Zibetti05} and infrared \\citep[e.g.][]{Zheng06,Zheng07} images, as well as in the radio.  The large coverage area of the Faint Images of the Radio Sky at Twenty-cm \\citep*[FIRST;][]{Becker95} survey has led to a number of stacking experiments being carried out \\citep[e.g.][]{Wals05,deVries07,White07,Hodge08}, looking at the radio properties of quasars and LL-AGN, and some studies of star-forming galaxies have taken place using smaller, but deeper radio surveys \\citep[e.g.][]{Boyle07,Ivison07,Beswick08,Carilli08}.  \\renewcommand{\\thefootnote}{\\it{\\alph{footnote}}} \\begin{table*} \\begin{center} \\begin{minipage}{\\textwidth} \\begin{center} \\caption{A summary of the properties of the radio surveys used in this work.  The area and resolution of each survey, and the noise level at the centre of each image are given, along with the primary reference paper, which should be consulted for further details.} \\label{tab:surveydetails} \\begin{tabular}{lccccccl} \\hline Field & Instrument & Frequency & Area & Resolution & Position Angle & Noise level & Reference\\\\ & & & (deg$^{2}$) & (arcsec$^{2}$) & (deg) & ($\\umu$Jy~beam$^{-1}$) & \\\\ \\hline xFLS         & VLA  & 1.4~GHz & 4 & $5.0\\times5.0$ & 0  & 23 & \\citet{Condon03}\\\\ xFLS         & GMRT & 610~MHz & 4 & $5.8\\times4.7$ & 60 & 30 & \\citet{Garn07}\\\\ ELAIS-N1     & GMRT & 610~MHz & 9 & $6.0\\times5.0$ & 45 & 40 / 70\\footnote[1]{Deep / shallow regions respectively.} & \\citet{Garn08EN1}\\\\ Lockman Hole & GMRT & 610~MHz & 5 & $6.0\\times5.0$ & 45 & 60 & \\citet{Garn08LH}\\\\ ELAIS-N2     & GMRT & 610~MHz & 6 & $6.5\\times5.0$ & 70 & 90 & \\citet{Garn09EN2}\\\\ \\hline \\end{tabular} \\end{center} \\end{minipage} \\end{center} \\end{table*} \\renewcommand{\\thefootnote}{\\arabic{footnote}}  There have been two stacking studies of the infrared / radio correlation for galaxies which are detected in the infrared, but are below the detection limits of their radio surveys.  The strength of the infrared / radio correlation can be quantified by the logarithmic flux density ratio $q_{24}$ \\citep[][and see Section~\\ref{sec:VLAstackq}]{Appleton04}, and the results from these studies appear to be inconsistent, with \\citet{Boyle07} finding a significantly higher value of $q_{24}$ than is observed in sources which are detected in both wavebands, and \\citet{Beswick08} finding a significantly lower value of $q_{24}$, and a tentative evolution in the value of $q_{24}$ with 24-$\\umu$m flux density.  In this work we describe a stacking study of the infrared / radio correlation, using the sensitive {\\it Spitzer Space Telescope} \\citep{Werner04} observations of the {\\it Spitzer} extragalactic First Look Survey field (xFLS) and the {\\it Spitzer} Wide-area Infrared Extragalactic survey \\citep[SWIRE;][]{Lonsdale03} fields.  The three northern SWIRE fields are the European Large-Area {\\it ISO} Survey-North~1 (ELAIS-N1), -North~2 (ELAIS-N2) and Lockman Hole regions, and we take radio data from a 1.4-GHz Very Large Array (VLA) survey of the xFLS field \\citep{Condon03}, and 610-MHz Giant Metrewave Radio Telescope (GMRT) surveys of the xFLS \\citep{Garn07}, ELAIS-N1 \\citep{Garn08EN1}, Lockman Hole \\citep{Garn08LH} and ELAIS-N2 \\citep[][in prep.]{Garn09EN2} regions.  In Section~\\ref{sec:techniques} we describe the two stacking techniques which are used in the literature, and calculate the values of $q_{24}$, using 1.4-GHz data from within the xFLS field.  In Section~\\ref{sec:stack610} we extend this work to the four 610-MHz survey fields, and demonstrate that AGN contamination is not a significant problem for our infrared samples.  We compare the results from our four fields in Section~\\ref{sec:stackingdiscussion}, and find a field-to-field variation in the infrared / radio correlation, which is too great to be explained through cosmic variance.  We compare our results to those of \\citet{Boyle07} and \\citet{Beswick08}, and discuss potential explanations for the systematic variation. ",
        "conclusions": "We have compared the two methods commonly used in stacking experiments, and demonstrated that they give comparable measurements of the average radio flux density for binned sources.  We have shown that the median flux density is a much more reliable estimator than the mean, and that the noise level of stacked images decreases as $1/\\sqrt{N}$, at least until a depth of $\\sim0.5$~$\\umu$Jy~beam$^{-1}$ at 1.4~GHz.  Creating stacked images leads to less information being retained on the flux density distribution of sources, and we argue that future stacking experiments should not make images, but instead work directly from the individual flux density measurements that can be calculated at each source position.  We have calculated $q_{24}$ and $q'_{24}$ from sources within the xFLS field, and demonstrated that they can be related to each other with the simple assumption that all radio sources have the spectral index of $\\alpha=0.4$, over the 24-$\\umu$m flux density range of 150~$\\umu$Jy\\ -- 10~mJy.  Using this conversion, we have compared the stacking results of \\citet{Boyle07} to the results obtained from our three SWIRE field images, and demonstrated that they are consistent, within the errors that can be placed on the median stacked values of $q_{24}$.  There appears to be some field-to-field variation seen between the xFLS and SWIRE fields, with stacking results from the xFLS field leading to a value of $q_{24}$ that is $\\sim0.3$ lower than the value seen in the SWIRE fields.  This variation is also seen when stacking radio sources using the 70-$\\umu$m catalogues of each field. Through comparison with the results of \\citet{Appleton04}, who look at individual sources which are detected in the xFLS field, we conclude that our radio flux density measurements are not being significantly biased by the stacking procedure.  We have considered several potential explanations for the difference in the median values of $q_{24}$, and shown that effects such as a flux calibration error cannot be responsible for the difference in results.  While we can not explain the systematic offset which is being seen, we believe that it originates in some combination of a `stacking bias' and the varying noise levels of the different radio surveys.  The systematic effect leads to an uncertainty in the absolute values of $q_{24}$ and $q_{70}$, but not in their variation with infrared flux density.  We find no evidence for a variation in the median value of $q_{24}$ from any of the survey fields down to a 24-$\\umu$m flux density of 150~$\\umu$Jy.  This is in agreement with the results of \\citet{Boyle07}, although is in contrast to the tentative findings of \\citet{Beswick08}.  We find a similar result at 70~$\\umu$m, although over the brighter flux density range of 10 -- 100~mJy.  We conclude that there is no evidence for any significant variation in the infrared / radio correlation with infrared flux density, down to the limits set by the extremely sensitive {\\it Spitzer} observations of the xFLS and SWIRE fields."
    },
    "0812/0812.0909_arXiv.txt": {
        "abstract": "The double humped SED (Spectral Energy Distribution) of blazars, and their flaring phenomena can be explained by various leptonic and hadronic models. However, accurate modeling of the high frequency component and clear identification of the correct emission mechanism would require simultaneous measurements in both the MeV-GeV band and the TeV band. % Due to the differences in the sensitivity and the field of view of the instruments required to do these measurements, it is essential to identify active states of blazars likely to be detected with TeV instruments.  Using a reasonable intergalactic attenuation model, various extrapolations of the EGRET spectra, as a proxy for GLAST (Gamma-ray Large Area Space Telescope) measurements, are made into TeV energies for selecting EGRET blazars expected to be VHE-bright. Furthermore, estimates of the threshold fluxes at GLAST energies are provided, at which sources are expected to be detectable at TeV energies, with Cherenkov telescopes like HESS, MAGIC or VERITAS.% ",
        "introduction": "Blazars have been detected in $\\gamma$-rays in a broad range of frequencies. The EGRET (Energetic Gamma-Ray Experiment Telescope) instrument on the CGRO (Compton Gamma Ray Observatory) detected more than $60$ AGN (Active Galactic Nuclei) at energies $>100$\\,MeV (with the highest energy photons at $\\sim$\\,few GeV). A number of these sources were blazars. On the other hand in the VHE-band (Very High Energy, defined as \\mbox{E\\,$>100$\\,GeV)} the remarkable advances in the field of ground based $\\gamma$-ray astronomy made in the last two decades, has resulted in the detection of more than 20 VHE blazars, and still counting. Yet, no simultaneous measurements of blazar SED in $\\gamma$-rays at MeV to TeV energies has been made before 2008. Therefore the emission component spanning these energies has not been unambiguously explained. Both leptonic (electron/positron synchrotron emission followed by and inverse-Compton scattering on a photon field) as well as hadronic models (proton-synchrotron and particle decay) are still in contention. The recent launch of the GLAST (now renamed as Fermi Gamma-ray Space Telescope) satellite, has opened up the opportunity for doing simultaneous multi-wavelength measurements of blazars overlapping the MeV-TeV bands.  The LAT (Large Area Telescope) instrument on the GLAST satellite is assumed  to measure blazar spectra from $20$\\,MeV to $300$\\,GeV, with a sensitivity of \\mbox{$\\sim10^{-9}$\\,/cm$^2$/s}. The typical sensitivity of current generation of ACTs (Atmospheric Cherenkov Telescope) is \\mbox{$\\approx10^{-12}$\\,/cm$^2$/s} above $\\approx100$ GeV, which is much better than GLAST sensitivity at $> 100$\\,MeV energies. However since the SED of blazars is falling from the MeV-GeV band to the VHE band, to a large extent due to the EBL (Extragalactic Background Light) attenuation (also might be due to the intrinsic spectral characteristics), it is more challenging to catch extragalactic sources which have fluxes above the ACT sensitivity threshold; the high (\\mbox{$\\sim 3$} orders of magnitude higher) level of background also adds to the complexity of measuring a signal. The LAT sensitivity is much better than that of the EGRET instrument and it is expected that GLAST will be able to detect fainter blazars in shorter exposure times. In this work, firstly, a list of blazars likely to be bright in TeV energies have been presented from a sample of EGRET detected blazars. These were obtained by ranking a redshift-limited and photon-index constrained EGRET blazar sample according to their predicted TeV flux, after correcting for the attenuation due to the EBL in the IR-UV frequencies. Secondly, GLAST threshold-flux levels have been provided for a range of redshifts and  photon indices ($\\Gamma$, the intrinsic power law photon index assumed to be same over the MeV-TeV band, $\\frac{dN}{dE}$\\,$\\propto$\\,E$^{-\\Gamma}$), at which follow up observations at VHE energies should be triggered. ",
        "conclusions": "With the combination of the recently launched Fermi Gamma-ray Space Telescope and the existing Cherenkov telescopes like HESS, MAGIC and VERITAS as well as the soon to come HESS II and MAGIC II telescopes, it would be possible to observe the spectra of blazars over the entire MeV to TeV energy range. Such wide coverage will give unprecedented measurements of the high energy component of blazars.  Due to the low flux levels of most blazars at TeV energies it would be advantageous to know a priori when a blazar is in a high state. To this end a ranked list of suitable EGRET blazars that could be covered in this entire energy range are listed. Furthermore, threshold flux levels for the Fermi Gamma-ray Space Telescope (at specific measured spectral indices for sources at particular range of redshifts) are provided as trigger conditions, for follow up observations in VHE with modern Cherenkov telescopes.    \\begin{theacknowledgments} This work is supported by the DFG through SFB 439. \\end{theacknowledgments}"
    },
    "0812/0812.2417_arXiv.txt": {
        "abstract": "We investigate the role of neutron star superfluidity for magnetar oscillations. Using a plane-wave analysis we estimate the effects of a neutron superfluid in the elastic crust region. We demonstrate that the superfluid imprint is likely to be more significant than the effects of the crustal magnetic field. We also consider the region immediately beneath the crust, where superfluid neutrons are thought to coexist with a type~II proton superconductor. Since the magnetic field in the latter is carried by an array of fluxtubes, the dynamics of this region differs from standard magnetohydrodynamics. We show that the presence of the neutron superfluid (again) leaves a clear imprint on the oscillations of the system. Taken together, our estimates  show that the superfluid components cannot be ignored in efforts to carry out ``magnetar seismology''. This increases the level of complexity of the modelling problem, but also points to the possibility of using observations to probe the superfluid nature of supranuclear matter. ",
        "introduction": "Anomalous X-ray pulsars (AXPs) and Soft Gamma-ray Repeaters (SGRs) are widely believed to be magnetars; neutron stars powered by an ultrastrong magnetic field (see \\citet{woods} for a review). Observations (mainly conducted by X-ray satellites) have established basic parameters like the magnetic field intensity, $ B \\sim 10^{15}\\, {\\rm G} $, and spin period, $P \\sim 10 {\\rm s} $, for this class of objects. They have also revealed a complex emission pattern with alternating periods of burst activity and quiescence. SGRs, which are typically more active than AXPs,  are the only ones exhibiting giant flares. These flares, which are believed to be triggered by some sort of instability in the magnetic field \\citep{DT,TD},  are by far the most energetic events associated with magnetars.  An exciting contribution to magnetar phenomenology was provided by the recent discovery of quasi-periodic oscillations (QPOs) in the late tail spectrum of the two giant flares  \\citep{israel,anna1,anna2}. There may also be evidence for a single QPO in the data of the third known flare, observed back in the 1970s \\citep{barat}. The frequencies of the most prominent QPOs lie in the range $30-100\\,{\\rm Hz} $, exactly where one would expect to find the seismic oscillation modes of the magnetar's crust \\citep{hansen}. This is consistent with the theoretical expectation that the energy released in a giant flare  is sufficient to fracture the crust and excite its normal modes \\citep{duncan}. If this interpretation of the QPOs is correct then we may have the opportunity to carry out magnetar \"asteroseismology\"; a comparison between theoretically predicted mode frequencies and the QPO data, with the ultimate goal of constraining the properties of neutron star matter \\citep{lars}.  Indeed, several recent papers have attempted to constrain the bulk equation of state of neutron star matter, assuming that the observed QPO frequencies are identified with the first few toroidal seismic modes of the crust  (see \\citet{stroh_review} for a recent review and  references). This is natural as a first step, but in reality the situation is likely to be more complicated. As suggested by \\citet{levin06} and \\citet{toypaper} the strong magnetic field would likely couple an oscillating crust to the liquid core on a very short timescale. Then the observed QPOs would be the manifestation of the coupled crust-core dynamics rather than the dynamics of the crust alone. Possible evidence that the magnetic core plays an active role is given by the presence of a low frequency QPO in the data of the December 2004 flare in SGR 1806-20. This QPO is difficult to reconcile with the seismic mode interpretation \\citep{israel}. It is therefore conceivable that magnetar ``seismology''  also probes the (much less well known) properties of the interior magnetic field. This obviously comes at a price. We now have to model global crust-core oscillations, a problem that is considerably more challenging than that of pure seismic crust modes.  Another aspect of neutron star physics, which is directly relevant to the mode interpretation of the QPOs, has received almost no attention so far. Young and mature neutron stars (older than a month or so) are sufficiently cold that the bulk of their interior liquid matter is in a superfluid state. In the crust, for densities above the nuclear drip density $ \\rho \\approx 4 \\times 10^{11}\\,{\\rm g}/{\\rm cm}^3 $, the  \"dripped\" neutrons are expected to form a superfluid below a temperature $\\sim  5 \\times 10^{9}\\, {\\rm K} $. The dynamical role of these \"free\" neutrons could be important. After all, they account for $\\sim 80\\%$ of the total mass in the crust. Similarly, in the liquid core  we  expect to find both neutrons and protons in a superfluid state (below a similar threshold temperature). In the outer core, the protons most likely form a type II superconductor, provided that the interior magnetic field does not exceed a critical value $\\approx 3 \\times 10^{16}\\, {\\rm G} $. As a consequence, any magnetic field that penetrates the proton plasma will form a large number of quantised magnetic fluxtubes \\citep{baym}.  It is clearly relevant to ask to what extent  the physical components (the crust and the magnetic field) that play the leading role in the magnetar QPO problem are sensitive to neutron and proton superfluidity/superconductivity. The aim of this investigation is to provide an insight into this issue, and improve our understanding of the relative importance of the multifluid aspects of the magnetar oscillation problem. By carrying out a local analysis, i.e. considering uniform parameter models, we learn how the shear waves in the neutron star crust are affected by the presence of a superfluid neutron component. Similarly, a local analysis in the core tells us how the Alfv\\'en waves are altered by the presence of the neutron superfluid (which provides the bulk of the core mass). Not surprisingly, the entrainment between neutrons and protons turns out to be the key parameter in these problems. This initial (order of magnitude) analysis  serves as a useful guideline for future (more detailed) work for realistic neutron star models. ",
        "conclusions": "In this paper we have investigated the role of neutron star superfluidity for magnetar oscillations. The results impact on attempts to use data from observed quasiperiodic oscillations in the tails of magnetar flares to place contraints on neutron star parameters, see for example \\citet{lars}. Using a plane-wave analysis we  estimated the effects of the neutron superfluid in the elastic crust region. This, the first ever, analysis of the combined magnetic-elastic-superfluid crust problem demonstrated that  the superfluid imprint is likely to  be more significant than the effects of the crustal magnetic field. This is, of course, assuming that the SGR flare mechanism does not deposit sufficient heat in the crust to raise the system above the superfluid transition temperature. Available estimates, e.g. \\citet{lyb}, suggest that this is unlikely. We also considered the region immediately beneath the crust, where superfluid neutrons are thought to coexist with a type~II proton superconductor. Since the magnetic field in the latter is carried by an array of fluxtubes, see \\citet{supercon} for a recent discussion, the dynamics of this region differs from standard magnetohydrodynamics. We showed that the presence of the neutron superfluid (again) affects the oscillations of the system. This accords well with previous results of, in particular, \\citet{mendell}.  Our estimates show that the superfluid components cannot be ignored in efforts to carry out magnetar seismology. This increases the level of complexity of the modelling problem, but also points to the exciting possibility of using observations to probe the superfluid nature of supranuclear matter. Future work needs to extend our analysis to consider the global oscillations of magnetic-superfluid-elastic neutron stars. This is a very interesting problem because, in addition to enabling a more detailed seimsology analysis, it may also provide insight into rotational glitches in magnetars \\citep{dib}. It is generally believed that superfluidity plays a key role in radio pulsar glitches. Recent developments in modelling these events is showing some promise \\citep{glitch}, and it would obviously be highly relevant to extend this analysis to strongly magnetised systems."
    },
    "0812/0812.2635_arXiv.txt": {
        "abstract": "The X-ray Pulsar-based Autonomous Navigation(XNAV) was tested which use the Crab pulsar (PSR B0531+21) in the USA Experiment on flown by the Navy on the Air Force Advanced Research and Global Observation Satellite (ARGOS) under the Space Test Program recently. It provide the way that the spacecraft could autonomously determine its position with respect to an inertial origin. Now I analysis the sensitivity of the exist instrument and the signal process to use radio pulsar navigation and discuss the integrated navigation use pulsar, then give the different navigation mission analysis and design process basically which include the space, the airborne, the ship and the land of the planet or the lunar. So the pulsar navigation can give the continuous position in deep space, that means we can freedom fly successfully in the solar system use celestial navigation that include pulsar and traditional star sensor. It also can less or abolish the dependence of Global Navigation Satellite System (GNSS) which include GPS, GRONSS, Galileo and BeiDou et al. ",
        "introduction": "\\IEEEPARstart{T}{HE} gigantic success of celestial navigation is lead to Christopher Columbus awareness of the American continents in the Western Hemisphere in 1492, Even if it first been discovered in 1421 by ZhengHe or Viking, they all use the same way in navigation. The history of navigation use celestial origin from the ancient activity which include hunting and back to home in the night. The base principle of celestial navigation is used the moon, stars, and planets as celestial guides assuming the sky was clear in that times.  With the period of the Age of Discovery or Age of Exploration (the early 15th century and continuing into the early 17th century), many new technique be invented which include water clock, quadrant Sextant(1731,John Hadley), Chronometer(1761,John Harrison), the Sumner Line or line of position(1837,Thomas Hubbard Sumner), The Intercept Method or Marcq St Hilaire method(1875,Marcq St Hilaire) et al. The Sumner Line and Marcq St Hilaire method construct the foundation of modern nautical navigation. The lighthouse are used to mark dangerous coastlines, hazardous shoals and reefs, safe entries to harbors and can also assist in aerial navigation. For the physics and chemistry development, Gustaf Dal\u00e9n in recognition of his remarkable invention of automatic valves designed to be used in combination with gas accumulators in lighthouses and light-buoys(the noble prize in 1912).  After Guglielmo Marconi achieve the radio communication in 1895, navigation by radio as an aid has been practiced in Germany since 1907. Scheller invented the complimentary dot-dash guiding path, which can be seen as a 'landmark' for several decades of navigational aids. The first practical VHF radar system was installed on French ship in 1935. Radio navigation grow fast for the defense technique need in the World War II, the Long Range Navigation (Loran) system was invented. In 19 century, it was proposed that use artificial stellar navigation. Until 1960, the first satellite navigation system, Transit, was first successfully tested, used by the United States Navy. Then it evolved to Global Positioning System(GPS). The soviet also built the similar system - GLONASS for the same reason. Now the Global Navigation Satellite system(GNSS) in the realism and dreams still have Chinese BeiDou, Indian IRNSS(Indian Regional Navigational Satellite System), French DORIS(Doppler Orbitography and Radio-positioning Integrated System) and Japanese QZSS(The Quasi-Zenith Satellite System).  An Inertial Navigation System(INS) is a navigation aid that uses the computing and motion sensors to continuously track the position, orientation, and velocity (direction and speed of movement) of a moving object without the need for external references. The gyroscopic compass (or gyro compass) was introduced in 1907. In 1942,the first INS be applied in V2 missile, then it be used in aeronautics, nautical and space widely.  When GPS and INS still do not grown up, celestial navigation system(CNS) spread to aeronautics by US(B-52,FB111, B-1B,B-2A,C-141A,SR-71,F22 et al.) and Soviet(TU-16,TU-95,TU-160 et al.)\\cite{pdl+01,w07}. Then the star tracker (i.e. track one star or planet or angle between it)\\cite{mf63} be used to determine the attitude of the spacecraft in help orient the Apollo spacecraft enroute to and from the Moon. Now the advanced star sensor(i.e. sense many star simultaneous) has been developed for the optical CCD technique improved\\cite{khh92}.  Although GPS and INS almost can finish any job in this planet now, someone still think it is important for it can be used independently of ground aids and has global coverage, it cannot be jammed (except by clouds) and does not give off any signals that could be detected by an the others. The traditional maritime state which include US,Russia,UK,French, all spend many money in CNS for its unique advantage.  After radio pulsar been discovered by Hewish,A. and Bell,J. in 1967, Downs give a advice that use radio pulsars for interplanetary navigation in 1974\\cite{d74}. But in that time, both the radio and optical signatures from pulsars have limitations that reduce their effectiveness for spacecraft navigation. At the radio frequencies that pulsars emit (from 100 MHz to a few gigahertz) and with their faint emissions, radio-based systems would require large antennas (on the order of 25 m in diameter or larger) to detect sources, which would be impractical for most spacecraft. Also, neighboring celestial objects including the sun, moon, Jupiter, and close stars, as well as distance objects such as radio galaxies, quasars, and the galactic diffuse emissions, are broadband radio sources that could obscure weak pulsar signals\\cite{rwp06,s05,spr+06}. So Chester and Butman described spacecraft navigation using X-ray pulsars in 1981\\cite{cb81}. Then in the USA Experiment on the ARGOS, the X-ray Pulsar-based Autonomous Navigation(XNAV) were first tested which use the Crab pulsar. Dr Sheikh construct the complete frame about XNAV which based modern navigation technique that include Kalman filter et al. ",
        "conclusions": ""
    },
    "0812/0812.0367.txt": {
        "abstract": "We examine a wide class of multi-field inflationary models based on fields that decay or stabilize during inflation in a staggered fashion. The fields driving assisted inflation are on flat, short stretches, before they encounter a sharp drop; whenever a field encounters such a drop due to its slow roll evolution, its energy is transferred to other degrees of freedom, i.e. radiation. The rate at which fields decay is determined dynamically and it is not a free parameter in this class of models. To compute observables, we generalize the analytic framework of staggered inflation, allowing for more general initial conditions and varying potentials.  By searching for generic situations arising on the landscape, we arrive at a setup involving linear or hilltop potentials and evenly spread out initial field values. This scenario is not more fine tuned than large-field models, despite the fact that many more degrees of freedom are involved. Further, the $\\eta$-problem can be alleviated.  The additional decrease of the potential energy caused by the decay of fields provides leading order contribution to observables, such as the scalar and tensor spectral index or the tensor to scalar ratio, for which we derive general expressions. We compare the predictions with WMAP5 constraints and find that hilltop potentials are borderline ruled out at the $2\\sigma$-level, while linear potentials are in excellent agreement with observations. We further comment on additional sources of gravitational waves and non-Gaussianities that could serve as a smoking gun for staggered inflation. ",
        "introduction": "The embedding of single-field inflation within string theory is an active field of study, see \\cite{McAllister:2007bg,Cline:2006hu,Burgess:2007pz,Kallosh:2007ig} for recent reviews. One of the most successful setups yet is the KKLMMT proposal \\cite{Kachru:2003sx}. A complete implementation is not easily achieved since many conceptual challenges need to be overcome. Among these challenges is the requirement that all moduli fields need to be stabilized, that the inflaton potential ought to be extremely flat and that corrections to the potential must be understood well enough in order to guarantee  that the flatness is not spoiled once a field traverses a super-Planckian distance. Remarakably, it seems possible to addresses these challenges within the KKLMMT proposal, but at the cost of fine-tuning and a construction requiring delicately arranged ingredients \\cite{Baumann:2008kq,Baumann:2007ah} (see however \\cite{Hoi:2008gc}).  But if we turn our attention to a generic location on the landscape, single-field inflation appears unlikely: moduli-fields can be dynamical, potentials are generically too steep and they change significantly if the fields roll far enough.  A way to circumvent some of these problems is to implement assisted inflation \\cite{Liddle:1998jc,Malik:1998gy,Kanti:1999vt,Kanti:1999ie,Calcagni:2007sb}, a type of multi-field inflation, whereby many fields are dynamical and assist each other in driving the inflationary phase. In this scenario, the potentials can be steeper without spoiling slow roll, because a given field sees the driving force of its own potential, but it is slowed down by the combined Hubble friction of all fields. In addition, if the number of fields is large, none of them needs to traverse a super-Planckian stretch in fields space, thus alleviating the $\\eta$-problem to some extent \\cite{Liddle:1998jc,Kanti:1999vt,Kanti:1999ie}. From this point of view, multi-field models are more desirable and natural than single-field ones. Some recent realizations of assisted inflation are $\\mathcal{N}$-flation \\cite{Dimopoulos:2005ac,Easther:2005zr} (fields are identified with axions arising from some KKLT compactification \\cite{Kachru:2003aw} of type IIB string theory, see also \\cite{Misra:2007cq,Misra:2008tx}), M5-brane inflation \\cite{Becker:2005sg} (inflatons are identified with distances between branes that are spread out over an orbifold) or inflation from tachyons \\cite{Piao:2002vf,Majumdar:2003kd}.  Of course, all moduli-fields need to be stabilized once inflation is over, that is moduli stabilization has to occur during or towards the end of assisted inflation.  Current models of multi-field inflation are often described by an effective single-field model, see i.e. \\cite{Wands:2007bd} for a review; it is further assumed that all fields stabilize coherently, leading to a sudden end of inflation and (p)reheating just as in the single-field case \\footnote{Note that reheating predominantly hidden sectors can be a problem for multi-filed models \\cite{Green:2007gs}.}. However, the use of an effective single field is problematic for several reasons; first of all, there is no a priori reason to expect that fields decay (stabilize) all at once; indeed, there are cases, such as in inflation from multiple M5-Branes \\cite{Becker:2005sg,Krause:2007jr,Ashoorioon:2006wc,Ashoorioon:2008qr}, where a coherent decay (stabilization) is impossible \\footnote{During inflation the M5-branes separate from each other until they encounter a boundary brane into which they dissolve, one after the other.}. A proper computational framework has to incorporate this feature, such as in \\cite{Battefeld:2008py,Battefeld:2008ur} where analytic tools for staggered inflation were developed (see also the numerical treatment in \\cite{Ashoorioon:2006wc}, where this effect was named cascade inflation). Another problem of an effective single-field description is its inapplicability during reheating: even if inflation ends coherently for all fields, they generically de-phase during reheating, resulting in the suppression of many resonances \\cite{Battefeld:2008bu,Battefeld:2008rd}, albeit tachyonic preheating \\cite{Dufaux:2006ee} might still work \\cite{Battefeld:2008rd,Braden,inprep2}. Further, the production of isocurvature perturbations and non-Gaussianities cannot be computed correctly.  To incorporate that fields decay or stabilize during inflation, we extend the formalism of \\cite{Battefeld:2008py}, employing again the coarse grained description proposed therein: by smoothing out the number of fields $\\mathcal{N}(t)$ we can introduce a continuous decay rate $\\Gamma(t)\\equiv -\\dot{\\mathcal{N}}/\\mathcal{N}$. This rate causes energy transfer from the inflaton sector to an additional component of the energy momentum tensor $T^{\\mu\\nu}_r$, i.e. radiation. To avoid spoiling inflation, the ratio $\\bar{\\varepsilon}=3(1+w_r)\\rho_r/(2\\rho_{total})\\simeq \\Gamma/(2H)$ needs to be small. The resulting setup shares some similarities to warm inflation \\cite{Berera:1995ie,Berera:2008ar} or dissipative inflation \\cite{Hall:2007qw}, without many of its shortcomings (see i.e. \\cite{Yokoyama:1998ju} for some problems within warm inflation related to the additional frition term, which we avoid). A study of scalar perturbations within staggered inflation reveals that $\\bar{\\varepsilon}$ appears alongside the usual slow roll parameters $\\varepsilon$ and $\\eta$ in observables \\cite{Battefeld:2008py}.  In \\cite{Battefeld:2008py} several simplifying assumptions were made, such as identical initial field values and potentials for all inflatons. In this paper, we relax these conditions to allow for more general cases, but we still assume flat potentials. To be precise,  we assume that the change of the potential energy in a given field over its slow roll phase is small compared to its potential energy. Within this setup, we compute the spectrum of gravitational waves $\\mathcal{P}_T$, the spectral index $n_{T}$ and the tensor to scalar ratio $r$.  Equipped with this improved formalism, we construct new models of multi-field inflation motivated by generic potentials on the landscape \\cite{Susskind:2003kw}, with $\\mathcal{N}\\sim 10^{3}$ fields and $\\bar{\\varepsilon}\\gg \\varepsilon$ \\footnote{ $\\varepsilon\\sim \\tilde{\\varepsilon}\\bar{\\varepsilon}$ where $\\tilde{\\varepsilon}=\\Delta V_i/V_i$; thus $\\bar{\\varepsilon}\\gg \\varepsilon$ is identical with our assumption of flat potentials.}: we arrive at models where fields can participate in assisted inflation until they encounter a sharp drop, at which point they decay or stabilize quickly. Thus, the decay rate is determined by the dynamics and it is not a free parameter. We assume similar potentials (hilltop or linear) for all fields, but allow for a narrow spread of masses or slopes. These types of potentials are motivated by expanding around either flat stretches or maxima on the landscape. Regarding initial conditions, we spread the fields evenly over the allowed interval and argue that this choice is the least fine tuned one.  The scenario is similar in spirit to chain inflation \\cite{Freese:2004vs,Feldstein:2006hm,Freese:2006fk,Huang:2007ek,Chialva:2008zw,Ashoorioon:2008pj,Ashoorioon:2008nh}, which is supposed to operate via successive, rapid tunnelling between different vacua on the landscape. However, instead of a rapid succession of first order phase transitions, which is difficult to achieve and can cause fatal problems in chain inflation related to bubble nucleation \\cite{Ashoorioon:2008pj}, fields follow a smooth path through the landscape. The presence of such a path becomes more probable the higher the dimensionality of the moduli space is, but the danger of getting trapped in a meta-stable minimum remains.  Within staggered inflation, we compute observables, such as $n_s$ and $r$, that are dominated by $\\bar{\\varepsilon}$; a comparison with the WMAP5 data \\cite{Komatsu:2008hk} reveals that hilltop potentials are already ruled out at the $2\\sigma$-level, while linear potentials are in excellent agreement with observations. We also re-examine briefly inflation driven by tachyons, as proposed in \\cite{Majumdar:2003kd}, and conclude that a constant decay rate is already ruled out in this setup while a serial condensation is still viable.  We further comment on expected unique signatures of staggered inflation, such as additional gravitational waves and non-Gaussianities, that should arise whenever a field decays or stabilizes. To compute these signals, one has to go beyond the analytic formalism of this paper and study the decay of fields in detail. This would also provide a numerical check of the analytic framework.  We conclude that staggered inflation offers a novel, compelling possibility to implement inflation within string theory, enabling the use of potentials that would be considered unsuitable otherwise.  The concrete outline of this paper is as follows: we start by reviewing the background evolution of staggered inflaton in Sec.~\\ref{sec:bgr}, which is used in our discussion of scalar and tensor perturbations in Sec. \\ref{scalarperturbations} and \\ref{sec:tensor}. We follow with an examination of concrete models, linear potentials in Sec.~\\ref{sec:linear} and hilltop potentials in Sec.~\\ref{sec:hilltop}, allowing for different slopes and masses in Sec.~\\ref{sec:linearv} and \\ref{sec:quadraticv}. The predictions are compared with WMAP5 constraints in Sec.~ \\ref{sec:discussion}, where we also re-examine inflation driven by multiple tachyons. We elaborate on the theoretical motivation for these models in Sec.~\\ref{sec:landscape} and point out open questions within staggered inflation in Sec.~\\ref{sec:opensissues}, before concluding in Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "} Staggered multi-field inflation incorporates the notion of decaying/stabilizing fields during inflation, a feature that has been largely ignored in the literature. We extended the formalism to incorporate more general initial conditions and potentials, enabling the discussion of generic multi-field models as expected in the landscape picture. The potentials we considered are flat over short intervals, before a sharp drop is encountered. Such potentials are less fine tuned than the one in single-field models, which require flatness over long distances in field space. After spreading out the fields evenly over the allowed interval and using either linear or hilltop potentials with some spread of slopes or masses, we computed observable parameters such as the scalar spectral index, the tensor to scalar ratio or the tensor spectral index. These parameters are dominated by the contributions due to staggered inflation $\\bar{\\varepsilon} \\simeq \\Gamma / (2H) \\propto \\rho_{radiation}/\\rho_{total} > \\varepsilon_{slowroll}$; $\\bar{\\varepsilon}$ is not a free parameter but follows from the slow roll evolution.  We found that hilltop potentials are already ruled out at the $2\\sigma$ level by the WMAP5 data, whereas linear potentials, which are better motivated, are in excellent agreement with observations. We also reconsidered staggered inflation driven by multiple tachyons \\cite{Majumdar:2003kd,Battefeld:2008py}, where $\\Gamma(t)$ is an unknown function. We find that constant decay rates are already ruled out at the $2\\sigma$ level, but a serial condensation of tachyons is still viable.  One drawback of multi-field inflation is the apparent flexibility to choose initial values and potentials at will, increasing the tunability compared to single-field scenarios. As a result, multi-field models are in danger of loosing their predictability. To minimize fine tuning, we investigated which potentials and initial conditions appear to be generic on the landscape and gave arguments in favor of our choice. However, the absence of a canonical measure on the space of initial conditions and the landscape renders these arguments heuristic and open to debate. This uncertainty can be alleviated in certain models, for instance if the inflatons were axions with a finite domain.  We concluded with a summary of open questions, ranging from a numerical validation of the analytic framework, to potentially new sources of gravitational waves and non-Gaussianities; the latter ones may very well serve as a smoking gun for staggered inflation, but require a better understanding of the fields' decay/stabilization mechanism before they can be computed. In this light, a more concrete implementation within string theory of this new type of inflation on the landscape is of great interest.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0812/0812.2259_arXiv.txt": {
        "abstract": "We use measurements of the growth of cosmic structure, as inferred from the observed evolution of the X-ray luminosity function (XLF) of galaxy clusters, to constrain departures from General Relativity (GR) on cosmological scales. We employ the popular growth rate parameterization, $\\Omega_{\\rm m}(z)^{\\gamma}$, for which GR predicts a growth index $\\gamma\\sim 0.55$. We use observations of the cosmic microwave background (CMB), type Ia supernovae (SNIa), and X-ray cluster gas-mass fractions ($f_{\\rm gas}$), to simultaneously constrain the expansion history and energy content of the Universe, as described by the background model parameters: $\\Omega_{\\rm m}$, $w$, and $\\Omega_{\\rm k}$, i.e., the mean matter density, the dark energy equation of state parameter, and the mean curvature, respectively. Using conservative allowances for systematic uncertainties, in particular for the evolution of the mass--luminosity scaling relation in the XLF analysis, we find $\\gamma=0.51^{+0.16}_{-0.15}$ and $\\Omega_{\\rm m}=0.27\\pm 0.02$ (68.3 per cent confidence limits), for a flat cosmological constant ($\\Lambda$CDM) background model. Allowing $w$ to be a free parameter, we find $\\gamma=0.44^{+0.17}_{-0.15}$. Relaxing the flatness prior in the $\\Lambda$CDM model, we obtain $\\gamma=0.51^{+0.19}_{-0.16}$. When in addition to the XLF data we use the CMB data to constrain $\\gamma$ through the ISW effect, we obtain a combined constraint of $\\gamma=0.45^{+0.14}_{-0.12}$ for the flat $\\Lambda$CDM model. Our analysis provides the tightest constraints to date on the growth index. We find no evidence for departures from General Relativity on cosmological scales. ",
        "introduction": "\\label{introduction}  Recently, using the observed evolution of the X-ray luminosity function (XLF) of massive galaxy clusters, \\citet[][hereafter M08]{Mantz:08} presented new constraints on dark energy from measurements of the growth of cosmic structure \\cite[see also][]{Vikhlinin:09}. These results are consistent with and complementary to those based on measurements of the cosmic expansion history, as deduced from distances to type Ia supernovae \\citep{Riess:98,Perlmutter:98, Davis:07, Kowalski:08}, the gas mass fraction in galaxy clusters \\citep{Allen:04, Allen:08, Rapetti:05, Rapetti:07}, and Baryon Acoustic Oscillations in the distribution of galaxies \\citep{Cole:05,Eisenstein:05,Percival:07}. In combination with independent measurements of the cosmic microwave background \\citep{Spergel:03,Spergel:06,Komatsu:09}, such measurements argue for a universe that is spatially flat, with most matter being cold and dark, and with the energy density being currently dominated by a cosmological constant, $\\Lambda$.  Due to severe theoretical problems associated with the cosmological constant, a plethora of other dark energy models have been proposed \\citep[for a review, see][]{Frieman:08}. Recent works \\citep[][and references therein]{Carroll:06, Caldwell:07, Hu:07, Amin:08, Bertschinger:08} have studied the possibility that late--time cosmic acceleration is not driven by dark energy but rather by gravity. These authors propose various frameworks to test for both time and scale-dependent modifications to GR at late times and on large scales. Although, currently, modifications to GR may be argued to be theoretically disfavoured, it is essential to test for such deviations using powerful, new data that are now becoming available. To measure departures from GR on cosmological scales, experiments sensitive to the dynamical effects of GR on these scales are required. Such experiments include, for example, the cross-correlation of the integrated Sachs-Wolfe (ISW) effect with other matter tracers; galaxy clustering; weak gravitational lensing; or cluster number counts using, e.g., X-ray selected samples.  General Relativity has been thoroughly tested from laboratory to Solar System scales. However, GR has only just begun to be tested on cosmological scales. Several authors have recently investigated a simple parameterization of the growth rate, $\\Omega_{\\rm m}(z)^{\\rm \\gamma}$ \\cite[first introduced by][]{Peebles:80}, to test for time--dependent modifications to GR \\citep[see e.g.][]{Linder:05, Sapone:07, Polarski:07, Gannouji:08, Acquaviva:08, Ballesteros:08,Mortonson:09, Thomas:09}. For a growth index, $\\gamma$, of approximately $0.55$, this parameterization accurately models the growth rate of GR. Some authors \\citep{Nesseris:07,DiPorto:07,Wei:08,Gong:08} have recently estimated constraints on $\\gamma$ by combining results from measurements of redshift space distortions and evolution in the galaxy power spectrum, as well as measurements of the normalization of the matter power spectrum, $\\sigma_{8}(z)$, from Lyman-$\\alpha$ forest data. Using current cosmic shear and galaxy clustering data at low redshift, \\cite{Dore:07} placed constraints on scale-dependent modifications to GR. These authors constrained two phenomenological, although physically motivated, modifications of the Poisson equation on megaparsec scales (from $0.04$ to $10\\Mpc$).  In this paper, we use the XLF experiment developed by M08, and data from the {\\it ROSAT} Brightest Cluster Sample \\cite[BCS;][]{Ebeling:98}, the {\\it ROSAT}-ESO Flux Limited X-ray cluster sample \\cite[REFLEX;][]{Bohringer:04}, the MAssive Cluster Survey \\cite[MACS;][]{Ebeling:01} and the 400 square degrees {\\it ROSAT} PSPC cluster survey \\cite[400sd;][]{Burenin:07}, to constrain departures from GR on scales of tens of megaparsecs over the redshift range $z\\lt 0.9$. We use CMB \\citep{Komatsu:09}, SNIa \\citep{Kowalski:08}, and cluster $f_{\\rm gas}$ data \\citep{Allen:08} to simultaneously constrain the background evolution of the Universe. We examine three background models: flat $\\Lambda$CDM, flat $w$CDM and non-flat $\\Lambda$CDM. We employ a Markov Chain Monte Carlo (MCMC) analysis, accounting for systematic uncertainties in the experiments. Our results represent the first constraints on $\\gamma$ from the observed growth of cosmic structure in galaxy clusters. ",
        "conclusions": "\\label{sec:conclusions}  Combining XLF, $f_{\\rm gas}$, SNIa and CMB data, we have simultaneously constrained the background evolution of the Universe and the growth of matter density fluctuations on cosmological scales, allowing us to search for departures from GR. We parameterize the expansion history with simple models that include late-time cosmic acceleration, taking flat $\\Lambda$CDM as our default model, but also investigating flat, constant $w$ models, and non--flat $\\Lambda$CDM models. We parameterize the growth of cosmic structure using the growth index, $\\gamma$, which assumes the same scale dependence as GR, but allows for time--dependent deviations from it.  We have performed MCMC analyses with seven (or eight) interesting cosmological parameters, and an additional fourteen parameters to encompass conservative allowances for systematic uncertainties. Marginalizing over these allowances, we obtain the tightest constraints to date on the growth index. For the flat $\\Lambda$CDM background model, we measure $\\gamma=0.51^{+0.16}_{-0.15}$. Allowing $w$ or $\\Omega_{\\rm k}$ to be free, we obtain similar constraints: $\\gamma=0.44^{+0.17}_{-0.15}$ and $\\gamma=0.51^{+0.19}_{-0.16}$, respectively. Including also constraints on $\\gamma$ from the ISW effect, we obtain $\\gamma=0.45^{+0.14}_{-0.12}$ for flat $\\Lambda$CDM. Currently, we find no evidence for departures from General Relativity.  In the near future, improved XLF analyses should provide significantly tighter constraints on both $\\gamma$ and $\\zeta$, allowing us to explore more complex modified gravity models (Rapetti et al., in preparation). Our results highlight the importance of X-ray cluster data, and the potential of combined expansion history plus growth of structure studies, for testing dark energy and modified gravity models for the acceleration of the Universe."
    },
    "0812/0812.4649_arXiv.txt": {
        "abstract": "We present a detailed systematics for comparing warped brane inflation with the observations, incorporating the effects of both moduli stabilization and ultraviolet bulk physics. We explicitly construct an example of the inflaton potential governing the motion of a mobile D3 brane in the entire warped deformed conifold. This allows us to precisely identify the corresponding scales of the cosmic microwave background. The effects due to bulk fluxes or localized sources are parametrized using gauge/string duality. We next perform some sample scannings to explore the parameter space of the complete potential, and first demonstrate that without the bulk effects there can be large degenerate sets of parameters with observationally consistent predictions. When the bulk perturbations are included, however, the observational predictions are generally spoiled. For them to remain consistent, the magnitudes of the additional bulk effects need to be highly suppressed. ",
        "introduction": "The inflationary paradigm~\\cite{inflation} addresses a number of fine tuning problems of the standard hot big bang cosmology, such as the horizon and the flatness problems. It also predicts a nearly scale invariant power spectrum of the curvature perturbation, which has been verified to high accuracies by the observation of the thermal fluctuations in the cosmic microwave background (CMB) and the large scale structure of the universe~\\cite{Tegmark:2006az,Komatsu:2008hk}. Numerous models of inflation based on effective field theory have been proposed, however, distinct predictions of a given model crucially depend on its ultraviolet completion. To construct a truly predictive inflationary model, it is clearly important to embed it into a consistent microscopic theory of quantum gravity such as string theory.   During the past few years, our understanding of the various ingredients for obtaining string inflation has been significantly expanded, and many models with increasing sophistication and striking signatures have been proposed (For recent developments, see Ref.~\\cite{Reviews} and references therein). In the coming decade, {beyond the ongoing Sloan Digital Sky Survey~\\cite{Tegmark:2006az} and the Wilkinson Microwave Anisotropy Probe~\\cite{Komatsu:2008hk}, vastly improved cosmological data will become available from the advanced CMB observations~\\cite{Planck}, the CMB polarization experiments~\\cite{cmbpolarization}, the dark energy surveys~\\cite{jdem} as well as the map of large scale structure~\\cite{euclid}.} They will allow us to constrain the parameter spaces of these models, and possibly to even rule out some of them. It is therefore of timely interest to perform a thoroughly updated and complete case study in such a direction\\footnote{What we mean by ``complete'' should become clear momentarily.}. In this paper we shall focus on one of the most developed string inflation models in the literature, usually referred to as ``brane inflation''.   The original setup of brane inflation first introduced in Ref.~\\cite{DvaliTye} is to consider a pair of spacetime-filling D3-$\\overline{\\rm D3}$ branes, separated at some distance greater than the local string length on a compact six manifold. As D3 and $\\overline{\\rm D3}$ move towards each other under the Coulombic attraction, the canonical inflaton is then identified as the separation between them. Unfortunately in such a simple setup, the Coulombic attraction is too strong for the slow-roll inflation to persist. To overcome this obstacle, the authors of Ref.~\\cite{KKLMMT} considered instead placing the $\\rm D3$-$\\overline{\\rm D3}$ pair in a locally warped deformed conifold throat developed in a compact Calabi-Yau orientifold by background fluxes. The $\\overline{\\rm D3}$ is then stabilized at the tip of deformed conifold, and D3 is attracted weakly by the warped down D3-$\\overline{\\rm D3}$ potential given by \\begin{equation}\\label{VD3D3bar} V_{D3\\overline{D3}}(\\tau,\\sigma) = \\frac{D_0}{U^2} \\left( 1 - \\frac{3D_0}{16\\pi^2 T_3^2{|\\bf{y}-\\bar{\\bf{y}}|}^4} \\right) \\, . \\end{equation} Here $D_0=2T_3a_0^4$, $T_3=1/[(2\\pi)^3 g_s(\\alpha')^2]$ is the D3 brane tension, $a_0=\\exp[{-2\\pi K/(3g_sM)}]$ is the warp factor at the tip of deformed conifold with $-K$ and $M$ being the quanta of NSNS and RR three form fluxes, and $|\\bf{y}-\\bar{\\bf{y}}|$ is the D3-$\\overline{\\rm D3}$ separation. Furthermore, the factor $U$ is the universal K\\\"ahler modulus, whose role we shall discuss in detail later. Notice that the first term in (\\ref{VD3D3bar}) gives the positive contribution and uplifts the total potential energy, whereas the Coulombic attraction is highly suppressed by $a_0^8$ and only becomes dominating near the tip of deformed conifold. Eventually when $|\\bf{y}-\\bar{\\bf{y}}|$ becomes comparable to warped string length $\\sim a_0 l_s$, D3 and $\\overline{\\rm D3}$ annihilate through open string tachyon condensation, rapidly terminates the brane inflation.   There are two important further ingredients that have so far been missing in our discussion, and they are crucial for obtaining the inflationary phase and making detailed comparisons with observational data. The first ingredient is the stabilization of both closed and open sting moduli. They are usually stabilized by the perturbative flux potential~\\cite{GVW} and the non-perturbative superpotential generated by wrapped branes~\\cite{npW}. The second ingredient is the ultraviolet corrections arising from embedding the warped throat into a compact Calabi-Yau orientifold. The bulk fluxes and the distant branes, or additional supersymmetry breaking and moduli stabilization sources can give significant perturbations to the inflaton potential derived from the local sources and geometry.   The warped deformed conifold offers us an ideal venue to analyze these two ingredients. Its explicit metric~\\cite{Conifold} allows for studying the moduli stabilization and the construction of the inflaton potential valid for the {\\em entire} evolution, including precise identification of where inflation ends. Furthermore while bulk physics is largely unknown, the spectrum of supergravity states in singular conifold has been tabulated in Ref.~\\cite{T11}. The gauge/string duality then allows us to parametrize these symmetry breaking bulk perturbations to the inflaton potential by coupling dual (approximately) conformal field theory to these bulk modes~\\cite{Baumann2}. Combining these with the D3-$\\overline{\\rm D3}$ interaction, we can schematically parametrize the total potential of the inflaton $\\phi$ including both local and bulk effects, experienced by a mobile D3 brane in the warped deformed conifold as \\begin{equation}\\label{TotalV} {\\mathbb{V}}(\\phi) = V_{D3\\overline{D3}}(\\phi) + V_{\\rm stab.}(\\phi) + V_{\\rm bulk}(\\phi) \\, . \\end{equation} Here $V_{\\rm stab.}(\\phi)$ arises from moduli stabilization and $V_{\\rm bulk}(\\phi)$ encodes all other possible perturbations from bulk physics\\footnote{Although we follow similar scheme as in Ref.~\\cite{Baumann2}, we do not partition $V_{\\rm stab.}(\\phi)$ such that the inflaton mass term $\\sim H^2\\phi^2$ is singled out. As $H^2$ is usually a combination of microscopic parameters, for the purpose of full parameter scanning we shall calculate $V_{\\rm stab.}(\\phi)$ in full detail and express it explicitly in terms of the microscopic parameters while treating $V_{\\rm bulk}(\\phi)$ as further perturbations.}. Most of the quantities specifying $V_{\\rm stab.}(\\phi)$ are exclusively related to the local geometry of the throat, e. g. the warp factor $a_0$. However $V_{\\rm stab.}(\\phi)$ typically also depends on other quantities controlled by bulk physics, such as the one loop determinant of the non-perturbative superpotential. The quantities controlling $V_{\\rm stab.}(\\phi)$ are usually treated as free parameters and yield a landscape of possible inflaton potentials.   The current observational data~\\cite{Tegmark:2006az,Komatsu:2008hk} enable us to make comparisons with the predictions yielded by different parameter sets and constrain their allowed values. In order for this exercise to be instructive, it is crucial to include all the significant contributions to the inflaton potential. The existing literature in this direction~\\cite{branescan,ClineScan} has mostly focused on the first two contributions in (\\ref{TotalV}) without taking into account $V_{\\rm bulk}(\\phi)$. However in light of recent results in Ref.~\\cite{Baumann2} indicating that $V_{\\rm bulk}(\\phi)$ can be generically comparable to $V_{\\rm stab.}(\\phi)$, it is clearly necessary to apply such general results and scan the enlarged parameter space to include the bulk corrections.   Here we aim to provide some initial steps toward a complete systematic parameter scanning for warped brane inflation. We shall first consider a specific brane configuration and stabilize explicitly both universal K\\\"ahler modulus and some of the angular moduli of D3. We then construct an example of $V_{D3\\overline{D3}}(\\phi)+V_{\\rm stab.}(\\phi)$ valid for {\\em entire} warped deformed conifold throat. This potential should be regarded as the infrared completion to the model obtained in Ref.~\\cite{Baumann1} under the singular conifold limit (See also \\cite{McGill1,KP} for related work). It allows us to identify the end point of inflation, hence extrapolate precisely to the CMB scale. Next we shall briefly review the parametrization of the bulk effects $V_{\\rm bulk}(\\phi)$ given in Ref.~\\cite{Baumann2}, and discuss the microscopic and observational constraints on the inflationary parameter scanning. Finally we shall first present different degenerate parameter sets such that the resultant $V_{D3\\overline{D3}}(\\phi)+V_{\\rm stab.}(\\phi)$ yields observationally consistent curvature power spectrum $\\mathcal{P}_\\mathcal{R}$ and the corresponding spectral index $n_\\mathcal{R}$. We next demonstrate that the perturbations due to $V_{\\rm bulk}(\\phi)$ can have significant impact and need to be small to preserve these seemingly optimal parameter sets. These sample scannings aim to highlight the possible degeneracies and the important role of bulk effects. ",
        "conclusions": "In this paper, we have discussed in detail the inflaton potential governing the motion of a mobile D3 in the entire warped deformed conifold. In particular, we have included both the effects of moduli stabilization and other bulk physics. We then have performed some sample scannings to demonstrate that without the bulk perturbations, there can be significant degeneracies in the conifold deformation parameter $\\epsilon$ and the D7 embedding parameter $\\mu$ for producing observationally consistent predictions. However, as the bulk perturbations are included, we have explicitly shown that their magnitudes need to be $10^{-8}$-$10^{-9}$ to preserve the observationally consistent parameter sets\\footnote{An obviously interesting question would be whether the smallness of bulk perturbation coefficients $c_{\\Delta}$ really constitutes a significant fine-tuning, or they are just tied to the choice of having inflection point inflation in the throat. To answer this question fully, we believe it requires better than our current understanding of UV physics and beyond the scope of investigations here.}. The results presented here provide the beginning systematic steps towards a complete brane scanning in the warped throat, and, in particular, highlight the importance of the bulk effects.   It would be very interesting to follow the steps outlined here and perform a full scanning over the parameters listed in (\\ref{localindepparameters}). This clearly requires intensive computational undertakings. However given the rich parameter space and the degeneracies we have shown in the sample scannings, barring the observation of the primordial gravitational waves, it is likely that there remain significant regions in the parameter space for the warped brane inflation to match the future data. Moreover, a variant of the inflation model presented here is recently proposed in Ref.~\\cite{CHS}. In such a construction the gravitino mass $m_{3/2}$ can be made smaller than the Hubble scale $H$, hence circumventing the phenomenological bound given in Ref.~\\cite{KLbound}. It would clearly be interesting to generalize the analysis here and scan the parameter space for such variant, and search for an explicit example of a parameter set that gives TeV scale gravitino mass and observationally consistent cosmological predictions.     \\subsection*{Acknowledgement}   We thank Gary Shiu for collaboration and discussions at the early stage of this project. We are also grateful to Ana Ach\\'ucarro, Daniel Baumann, James Cline, Shamit Kachru, Gonzalo Palma, Fernando Quevedo, Koenraad Schalm and Bret Underwood for comments and suggestions. HYC appreciates the hospitality of Stanford Institute for Theoretical Physics, where part of this work was conducted. The work of HYC is supported in part by NSF CAREER Award No. PHY-0348093, DOE grant DE-FG-02-95ER40896, a Research Innovation Award and a Cottrell Scholar Award from Research Corporation, and a Vilas Associate Award from the University of Wisconsin. JG is partly supported by the Korea Research Foundation Grant KRF-2007-357-C00014 funded by the Korean Government at the early stage of this work, and is currently supported in part by a VIDI and a VICI Innovative Research Incentive Grant from the Netherlands Organisation for Scientific Research (NWO)."
    },
    "0812/0812.3405_arXiv.txt": {
        "abstract": "We present results of combined N-body and three-dimensional reionization calculations to determine the relationship between reionization history and local environment in a volume 1 Gpc $h^{-1}$ across and a resolution of about 1 Mpc.  We resolve the formation of about $2 \\times 10^6$ halos of mass greater than $\\sim 10^{12} M_\\odot$ at $z=0$, allowing us to determine the relationship between halo mass and reionization epoch for galaxies and clusters. For our fiducial reionization model, in which reionization begins at $z\\sim 15$ and ends by $z\\sim 6$, we find a strong bias for cluster-size halos to be in the regions which reionized first, at redshifts $10<z<15$. Consequently, material in clusters was reionized within relatively small regions, on the order of a few Mpc, implying that all clusters in our calculation were reionized by their own progenitors.  Milky Way mass halos were on average reionized later and by larger regions, with a distribution most similar to the global one, indicating that low mass halos are nearly uncorrelated with reionization when only their mass is taken as a prior.  On average, we find that most halos with mass less than $10^{13} M_\\odot$ were reionized internally, while almost all halos with masses greater than $10^{14} M_\\odot$ were reionized by their own progenitors.  We briefly discuss the implications of this work in light of the ``missing satellites'' problem and how this new approach may be extended further. ",
        "introduction": "The universe we observe at $z=0$ must bear the marks of reionization. Reionization began when the first stars polluted the intergalactic medium and created individual H~II regions \\citep{alvarez/etal:2006, abel/etal:2007,yoshida/etal:2007,wise/abel:2008}.  As the first galaxies grew in abundance, the H~II regions became longer lived, eventually containing perhaps tens of thousands of dwarf galaxies, growing and merging until they overlapped, marking the end of reionization \\citep{shapiro/giroux:1987,miralda-escude/etal:2000,gnedin:2000a,sokasian/etal:2001,nakamoto/etal:2001,ciardi/etal:2003,furlanetto/etal:2004,iliev/etal:2006,zahn/etal:2007,trac/cen:2007}. Observations of high-redshift quasars imply that this process was complete by redshift $z\\sim 6$ \\citep{becker/etal:2001,fan/etal:2002,white/etal:2003, willott/etal:2007}, while large-angle polarization measurements of the cosmic microwave background constrain the duration of reionization \\citep{spergel/etal:2003,komatsu/etal:2008}.  During this time, the temperature of the intergalactic medium increased from a few to tens of thousands of degrees, dramatically changing the evolution of gas as it responded to the highly dynamic underlying dark matter potential.  Low mass halos in ionized regions are less able cool, collapse, and form stars than those in neutral regions, due to the increase in the cosmological Jeans mass when gas is ionized and photo-heated, sometimes called ``Jeans mass filtering'' \\citep[e.g.,][]{shapiro/etal:1994, thoul/weinberg:1996,gnedin:2000b,dijkstra/etal:2004,shapiro/etal:2004}. This suppression of structure is one of the fundamental ways that reionization can leave its imprint on subsequent structure formation, even up until the present day.  Correlating reionization with the present-day environment may be the key to the so-called ``missing satellite problem'' --- many more satellite halos are predicted to form in CDM than are actually observed as galaxies \\citep{klypin/etal:1999}.  The leading explanation --- an alternative to more exotic possibilities like modifying dark matter or the amplitude of small-scale primordial density fluctuations --- is that the UV background maintains the intergalactic gas in a photo-heated state, preventing it from falling into the shallow potential wells of the progenitors of the satellite halos \\citep[e.g.,][]{bullock/etal:2000,benson/etal:2002}. For example, one might expect that regions that were reionized earlier will have fewer luminous satellites than regions that were ionized later.  However, biased regions, which are rich in early low-mass galaxy formation, would have reionized {\\em first}.  The latter effect implies that early reionization would lead to more satellite galaxies, while the former implies just the opposite.  Detailed three-dimensional models are necessary in order to disentangle these competing effects and quantify their dependence on the inevitable assumptions that must be made when modeling reionization on such large scales.  In this {\\em Letter}, we present our first calculations to address the correlation between reionization and local environment.  We take a novel approach, combining N-body simulations with a ``semi-numerical'' algorithm \\citep{zahn/etal:2007,mesinger/furlanetto:2007} for calculating the reionization history for the simulation, allowing us to achieve a higher dynamic range in resolving the scales of reionization than has been possible until now. We then report on the statistical correlations between halo properties and their reionization epoch and environment.  We present our hybrid N-body/semi-numerical method in \\S 2, our results in \\S 3, and end with a discussion in \\S 4.  Throughout, we assume a flat universe with $\\Omega_m=0.25$, $\\sigma_8=0.8$, $n_s=1$, $\\Omega_b=0.04$, and $h=0.7$.   \\begin{figure*} \\begin{center} \\includegraphics[width=1.0\\textwidth]{4panelsl.eps} \\end{center} \\caption{Visualization of the progress of reionization in our 1 Gpc$h^{-1}$ calculation. Redshifts $z=14$, $z=10$, $z=8$, and $z=6$ are shown from right to left. Ionized regions are blue and translucent, ionization fronts are red and white, and neutral regions are dark and opaque. A random sampling of 5 per cent (about 40,000) of all the halos at $z=0$ are shown in yellow.  Reionization is still quite inhomogeneous on these large scales, with large regions ionizing long before others. \\label{panels} } \\end{figure*} ",
        "conclusions": "Using large-volume and high-resolution coupled simulations of reionization and halo formation, we have developed a new method for connecting the $z=0$ distribution of halos to the reionization epoch. We have found that, when only their mass is known, galaxy scale halos are nearly uncorrelated with respect to reionization, with a distribution of H~II region bubble sizes and reionization epochs that are roughly consistent with having a random spatial distribution.  Higher mass halos, however, show a much stronger correlation, with none of the cluster scale objects having $z_r<8$ or $R_{\\rm HII}>30$~Mpc.  An important distinction is between {\\em internal} and {\\em external} reionization.  Figure~\\ref{schematic} describes these two possibilities.  In the external reionization case, the halo material was ionized by sources in a region with $R_{HII}\\gg R_{\\rm Lag}$, where $R_{\\rm Lag}$ is the comoving volume occupied by the mass of the halo at the cosmic mean density.  In this case, most of the sources that ionized the material were not progenitors of the halo, and the ionization front swept over the halo's progenitors quickly, leaving the halo with a relatively uniform reionization epoch.  For internal reionization, $R_{HII}<<R_{\\rm Lag}$, and the halo's reionization history is likely to be much more complex.  In general, more massive halos were internally reionized, while less massive ones were externally reionized.  Our definition is somewhat different from previous definitions \\citep[e.g.,][]{weinmann/etal:2007}, but we believe our definition is best suited for the method used here.  In our definition, halos are considered to be externally ionized if their Lagrangian radius, defined by $M_{\\rm halo}=4\\pi\\overline{\\rho}R_{\\rm Lag}^3/3$, is smaller than their H~II region radius.  For galaxy scale objects, with Lagrangian radii of order 2 Mpc, it is clear from Figure~\\ref{binr_m} that most of these objects were externally ionized.  Because our model does not resolve scales below about a Mpc, however, it is difficult to determine how few galaxies were internally ionized.  More detailed modeling of the small-scale structure on galactic scales, while still retaining the large volume presented here, is therefore necessary.   Our predictions for clusters are more robust. For these objects, with Lagrangian radii of order 20~Mpc, all but a tiny handful are internally ionized.  Our results may have important implications for galaxy formation, and in particular for the missing satellites problem.  Because we find such a large spread in reionization epochs for Milky Way mass halos, more information, such as larger-scale environment and accretion history, will be necessary to determine the reionization epoch of our own halo, even after the global reionization history is well constrained. If the abundance of Galactic satellites is strongly dependent on the reionization epoch of the galaxy, our results indicate that Milky Way mass halos would have a large spread in the number of observable satellites. We explore this issue in a companion paper (Busha et al., in preparation).  \\begin{figure}[t!] \\begin{center} \\includegraphics[width=0.45\\textwidth]{schematicm.eps} \\end{center} \\caption{Schematic diagram indicating the difference between external and internal reionization. $R_{HII}$ indicates the size of the reionizing bubble; $R_{Lag}$ indicates the Lagrangian radius of the halo. \\label{schematic} } \\end{figure}  Our results may also have implications for the issue of galaxy ``assembly bias'', the idea that galaxy clustering may be dependent on properties other than the mass of their host halos \\citep{wechsler/etal:2006, gao/white:2007, croton/etal:2007}. If the reionization epoch of halos at a given mass is correlated with halo formation time, and if the reionization epoch affects any aspects of the galaxy population, then assembly bias could be more important for such galaxies than it is for their host halos.  Further study will be required to investigate such effects.  The approach we have presented here will serve as the foundation for such more detailed future studies. These studies will investigate the statistical correlations between present-day structure and reionization, and will also incorporate detailed galaxy formation modeling. Such improvements will allow for an investigation of the detailed coupling between star formation histories and the local reionization history."
    },
    "0812/0812.1400_arXiv.txt": {
        "abstract": "{} { The initial cluster mass function (ICMF) in spiral discs is constrained and compared with data for old globular clusters and young clusters in starbursts. } { For a given absolute magnitude, the cluster age distribution depends on the ICMF. Here, the behaviour of the median age-magnitude relation is analysed in detail for Schechter ICMFs with various cut-off masses, \\mbox{$M_c$}. The calculated relations are compared with observations of the brightest clusters in spiral galaxies. Schechter functions are also fitted directly to observed mass functions (MFs). } { A single Schechter ICMF with an $\\mbox{$M_c$}$ of a few times 10$^5\\mbox{$M_{\\odot}$}$ can reproduce the observed ages and luminosities of the brightest (and \\mbox{$5^{\\rm th}$} brightest) clusters in the spirals if disruption of optically visible clusters is dominated by relatively slow secular evolution. A Schechter function fit to the combined cluster MF for all spirals in the sample yields $\\mbox{$M_c$} = (2.1\\pm0.4)\\times10^5\\mbox{$M_{\\odot}$}$. The MFs in cluster-poor and cluster-rich spirals are statistically indistinguishable. An $\\mbox{$M_c$}=2.1\\times10^5\\mbox{$M_{\\odot}$}$ Schechter function also fits the MF of young clusters in the Large Magellanic Cloud. If the same ICMF applies in the Milky Way, a bound cluster with $M>10^5\\mbox{$M_{\\odot}$}$ will form about once every $10^7$ years, while an $M>10^6\\mbox{$M_{\\odot}$}$ cluster will form only once every 50 Gyr. Luminosity functions (LFs) of model cluster populations drawn from an $\\mbox{$M_c$}=2.1\\times10^5\\mbox{$M_{\\odot}$}$ Schechter ICMF generally agree with LFs observed in spiral galaxies. } { The ICMF in present-day spiral discs can be modelled as a Schechter function with $\\mbox{$M_c$}\\approx2\\times10^5\\mbox{$M_{\\odot}$}$.  However, the presence of significant numbers of $M>10^6\\mbox{$M_{\\odot}$}$ (and even $M>10^7\\mbox{$M_{\\odot}$}$) clusters in some starburst galaxies makes it unlikely that the \\mbox{$M_c$} value derived for spirals is universal. In high-pressure environments, such as those created by complex gas kinematics and feedback in mergers, \\mbox{$M_c$} can shift to higher masses than in quiescent discs. } ",
        "introduction": "Star clusters connect many branches of astrophysics. Classically, they have been important test benches for models of stellar evolution \\citep{mm81}. A large fraction of all star formation occurs in clusters \\citep{ll03}, and massive, long-lived globular clusters (GCs) are potentially useful tracers of stellar populations in their parent galaxies \\citep{bs06}. It is therefore of considerable interest to understand how the properties of star clusters are related to those of their parent galaxy, and in particular whether the formation (and survival) of GC-like objects requires special conditions. Three important properties of cluster populations which bear on this issue can, to some extent, be studied separately: 1) the initial cluster mass function (ICMF), 2) the formation efficiency of (bound) star clusters \\citep{lr00,bastian08,pf08}, and 3) dynamical evolution and cluster disruption, addressed both theoretically and observationally in many studies \\citep{vz03,bm03,lam05,mf08,gb08,whit07}. All three may, in principle, depend on time and location within a galaxy. This paper is concerned with the mass function, primarily its high-mass end and the possible existence of an upper cut-off mass.  Highly luminous, compact star clusters are observed to form in on-going starbursts such as those associated with major gas-rich mergers and, in smaller numbers, in many other external galaxies with elevated star formation rates \\citep[e.g.,][]{lar04}. However, conversion of observed cluster luminosities to masses requires knowledge of the mass-to-light ratios which, in turn, are strongly age-dependent. Dynamical mass estimates via integrated-light measurements of velocity dispersions have been obtained for only a limited number of clusters, but have confirmed that young clusters with masses above $10^5 M_\\odot$ (and, in some cases, $10^6 \\mbox{$M_{\\odot}$}$) are quite common in external star-forming galaxies \\citep{hf96,lar01,lbh04,lr04,mg07,sg01,dp07}. On the other hand, most young star clusters known in the Milky Way are relatively sparse with masses of $10^2$ - $10^4 M_\\odot$. This apparent contrast between the properties of Milky Way open clusters and the luminous star clusters encountered in some external galaxies raises important questions about the universality of the cluster formation process. Do the most massive clusters form preferentially in starburst environments for some physical reason, or do they just happen to be statistically more likely to form there because the total numbers of stars and clusters formed are very large?  Most studies of the statistical properties of star cluster populations have concentrated on \\emph{luminosity} functions (LFs), which are more readily derived for large samples of clusters than mass functions (MFs). There is a roughly linear relation between the luminosity of the brightest cluster in a galaxy and the total number of clusters, as  expected if luminosities are drawn at random from a power-law distribution with slope $\\approx-2$ \\citep{bhe02,lar02,weidner04,whit03,whit07}. This implies that if the LF has a physical upper limit, no cluster system studied to date is sufficiently rich to sample the LF up to that limit, and the (young) cluster populations in most galaxies might in fact be drawn from very similar underlying LFs.  Clearly, the MF would be a physically more interesting property to study than the LF.  The LF generally depends on both the mass- and age distribution of the cluster sample in question \\citep[e.g.][]{fall06}, and since the latter is often unknown, the observed LFs offer only limited insight into the underlying MFs.  If the goal is to obtain information about the \\emph{initial} MF, disruption effects must also be taken into account.  In the small number of cases where MFs have been determined, they are generally found to be consistent with power-laws $dN/dM \\propto M^\\alpha$ with a slope $\\alpha \\approx -2$ \\citep{bik03,deg03,mg07,zf99,dow08}. However, the mass ranges differ from one study to another (e.g.\\ $2.5\\times10^3 < M/\\mbox{$M_{\\odot}$} < 5\\times10^4$ for the Bik et al.\\ study of \\object{NGC~5194}, and $10^4 < M/\\mbox{$M_{\\odot}$} < 10^6$ for \\object{the Antennae} clusters studied by Zhang \\& Fall), so it is not clear that the similarity of the slopes derived for different galaxies actually implies a universal ICMF.  Even if the MF can typically be fit by a power-law, this might only apply over a limited mass range. One potential difficulty is the large number of very massive clusters predicted in the \\object{Milky Way} for a uniform power-law MF. For a star formation rate in bound star clusters of $5.2\\times10^{-10} \\mbox{$M_{\\odot}$}$ yr$^{-1}$ pc$^{-2}$ in the Solar neighbourhood \\citep[][hereafter LAM05]{lam05}, the Milky Way disc would have formed about 100 clusters with $M > 10^6 \\mbox{$M_{\\odot}$}$ over a 10 Gyr lifetime, assuming that the cluster formation rate was constant in the past (probably a conservative assumption; e.g.\\ \\citealt{fuchs08}). These are not observed. Since the completeness of existing surveys of young Milky Way clusters remains extremely poorly quantified it cannot be excluded that some remain hidden behind massive extinction, and others may have dissolved (Sect.~\\ref{sec:disrupt}). Another possibility is that they simply never formed, or at least in smaller numbers than predicted by extrapolation of a pure power-law MF.  An upper limit to the ICMF need not be a strict cut-off, but could also take the form of an exponential decline above some characteristic mass \\mbox{$M_c$}, i.e.\\ a \\citet{sch76} function: \\begin{equation} \\frac{dN}{dM} \\propto (M/\\mbox{$M_c$})^\\alpha \\exp \\left(-M/\\mbox{$M_c$}\\right), \\label{eq:schechter} \\end{equation} According to \\citet{bastian08}, the relation between the luminosity of the brightest cluster in a galaxy and the star formation rate (SFR) can be reproduced if clusters in all galaxies are drawn from a Schechter function with $\\mbox{$M_c$} = (1-5)\\times10^6\\mbox{$M_{\\odot}$}$. More indirect evidence of a truncation or steepening of the MF in young cluster systems comes from the detection of a bend in the LF of clusters in \\object{NGC~5194} (M51), \\object{NGC~6946} and the Antennae galaxies (NGC 4038/39), which is consistent with a power-law MF truncated at a mass of $5\\times10^5 - 10^6\\mbox{$M_{\\odot}$}$ \\citep{zf99,gieles06}.  The aim of this paper is to examine, for a larger sample of ``normal'' spiral galaxies, the behaviour of the cluster MF at its high-mass end and constrain the upper mass limit, $\\mbox{$M_c$}$. The most direct way to do this would perhaps be to simply obtain luminosities and ages for a large sample of clusters in each galaxy and converting these to masses, ideally using \\emph{Hubble Space Telescope} (HST) imaging to obtain clean cluster samples. Although a large amount of data for nearby galaxies have by now accumulated in the HST archive, many datasets suffer from two basic shortcomings: 1) they lack (deep) imaging in a (roughly) $U$-band equivalent filter, which is essential for age-dating of the clusters, and 2) they cover only a small fraction of the galaxy. The latter limitation is particularly problematic when studying the most massive clusters, which tend to be rare. The only spiral which is sufficiently nearby to take advantage of HST's imaging resolution for cluster identification and is fully covered by multi-passband HST imaging including the ultraviolet is NGC~5194 \\citep{scheep07,hl08}, which is arguably not a ``normal'' spiral due to its on-going interaction with \\object{NGC~5195}.  Instead of studying the complete MF in detail for each individual galaxy, we begin by taking an approach similar to that adopted in several previous studies of the LF: from the correlation between the magnitude of the \\emph{brightest} cluster and some estimate of the \\emph{total} cluster population, one can test whether the data are consistent with random sampling from a single, universal LF. Here, we go one step further and also attempt to draw conclusions about the MF from observations of the brightest (and \\mbox{$5^{\\rm th}$} brightest) clusters.  As will become clear below, the crucial extra information that makes this possible are the cluster \\emph{ages}, but because these are needed only for the brightest  clusters, ground-based data are generally adequate.  The general outline of the paper is as follows: In Sect.\\ \\ref{sec:montecarlo} we determine the relation between the median age and luminosity of the brightest and \\mbox{$5^{\\rm th}$} brightest clusters (or, in fact, any sample with magnitudes in a fixed, narrow range) for a given sample size, cut-off mass $\\mbox{$M_c$}$, and age distribution.  This statistical age-luminosity relation should not be confused with the general fading  of an individual star cluster that is the result of stellar evolution and cluster dissolution. The (log(age), $\\mathcal{M}_V$) predictions are compared (in Sect.\\ \\ref{subsec:data}) with previously published ground-based data for clusters in a sample of nearby spiral galaxies in order to put constraints on \\mbox{$M_c$}.  In Sect.\\ \\ref{sec:mfs}, a direct Schechter function fit to observed MFs is performed and yields $\\mbox{$M_c$}=(2.1\\pm0.4)\\times10^5\\mbox{$M_{\\odot}$}$. In Sect.~\\ref{sec:lfs} the properties of the LF for a cluster sample drawn from an $\\mbox{$M_c$}=2.1\\times10^5\\mbox{$M_{\\odot}$}$ Schechter MF are compared with observed LFs. In Sect. \\ref{sec:elsewhere} the implications of an $\\mbox{$M_c$}=2.1\\times10^5\\mbox{$M_{\\odot}$}$ Schechter MF are discussed in the context of other cluster systems, and it is concluded that while such a MF is consistent with existing data for Milky Way and \\object{Large Magellanic Cloud} clusters, it is unlikely to be universal. Finally, in Sect.\\ \\ref{sec:discussion} the results are discussed in the context of host galaxy properties and current views on star/cluster formation. ",
        "conclusions": "This paper provides three main lines of evidence suggesting that the mass function of young star clusters in spiral galaxies can be approximated by a Schechter function with $\\mbox{$M_c$}=2.1\\times10^5 \\mbox{$M_{\\odot}$}$ and low-mass slope $\\alpha=-2$: \\begin{itemize} \\item The observed age-luminosity relations for the brightest and \\mbox{$5^{\\rm th}$} brightest cluster in a sample of mostly late-type spirals are consistent with those expected for cluster populations drawn from a Schechter MF with an \\mbox{$M_c$} of a few times $10^5\\mbox{$M_{\\odot}$}$. This result is not significantly affected by secular cluster disruption. \\item A maximum-likelihood fit of a Schechter function to the combined MF for clusters more massive than $10^5\\mbox{$M_{\\odot}$}$ in these spirals yields $\\mbox{$M_c$}=(2.1\\pm0.4)\\times10^5\\mbox{$M_{\\odot}$}$ for fixed low-mass slope $\\alpha=-2$. The cluster MFs for cluster-poor and cluster-rich subsamples of the galaxies are statistically indistinguishable. \\item The luminosity function for a model cluster sample with masses drawn from an $\\mbox{$M_c$}=2.1\\times10^5\\mbox{$M_{\\odot}$}$ Schechter function is consistent with existing data for the LFs of clusters in spiral galaxies, including the well studied M51 cluster system. \\end{itemize} An $\\mbox{$M_c$}=2.1\\times10^5\\mbox{$M_{\\odot}$}$ (or $\\mbox{$M_c$}\\approx1.4\\times10^5\\mbox{$M_{\\odot}$}$ for a Kroupa- or Chabrier IMF) Schechter function also fits the MF of clusters in the Large Magellanic Cloud, and it is argued that it is probably consistent with current surveys of young clusters in the Milky Way. It is, however, unlikely to be universal. In merger galaxies and other starbursts, a Schechter function fit requires a significantly higher \\mbox{$M_c$} than derived for the spirals. Since some spirals are not very different from mergers like the Antennae in terms of overall star formation rate or gas content, differences in the upper end of the cluster MFs are more likely related to the different gas kinematics in these systems. The fact that \\mbox{$M_c$} values are found to be greater for old GC systems than for young cluster populations in discs suggests that star (cluster) formation in present-day spiral discs generally proceeds in a relatively more orderly way than when most old GCs formed. In the future, observations with ALMA should provide crucial insight into the properties of molecular gas on scales that are directly relevant to formation of individual clusters in galaxies well beyond the Local Group."
    },
    "0812/0812.1636_arXiv.txt": {
        "abstract": "{} { Our goal is to investigate the properties of the circumstellar disks around intermediate mass stars to determine their occurrence, lifetime and evolution. } {We completed a search for circumstellar disks around Herbig Be stars using the NRAO Very Large Array (VLA) and the IRAM Plateau de Bure (PdB) interferometers. Thus far, we have observed 6 objects with 4 successful detections. The results towards 3 of these stars (R~Mon, MWC~1080, MWC~137) were presented elsewhere. We present our new VLA and PdBI data for the three objects MWC~297, Z~CMa, and LKH$\\alpha$~215. We constructed the SED from near-IR to centimeter wavelengths by adding our millimeter and centimeter data to the available data at other wavelengths, mainly Spitzer images. The entire SED was fitted using a disk+envelope model. In addition, we compiled all the disk millimeter observations in the literature and completed a statistical analysis of all the data. } { We show that the disk mass is usually only a small percentage (less than 10$\\%$) of the mass of the entire envelope in HBe stars. For the disks, there are large source-to-source variations. Two disks in our sample, R~Mon and Z~CMa, have similar sizes and masses to those found in T Tauri and Herbig Ae stars. The disks around MWC~1080 and MWC~297 are, however, smaller (r$_{out}$$<$100~AU). We did not detect the disks towards MWC~137 and LkH$\\alpha$~215 at millimeter wavelengths, which limits the mass and the size of the possible circumstellar disks. } { A comparison between our data and previous results for T Tauri and Herbig Ae stars indicates that although massive disks ($\\sim$0.1~M$_\\odot$) are found in young objects ($\\sim$10$^4$ yr), the masses of the disks around Herbig Be stars are usually 5-10 times lower than those around lower mass stars. We propose that disk photoevaporation is responsible for this behavior. In Herbig Be stars, the UV radiation disperses the gas in the outer disk on a timescale of a few 10$^5$ yr. Once the outer part of the disk has vanished, the entire gaseous disk is photoevaporated on a very short timescale ($\\sim$10$^5$ yr) and only a small, dusty disk consisting of large grains remains. } ",
        "introduction": "Herbig Ae/Be (HAeBe) stars are intermediate-mass (M$\\sim$2--8~M$_\\odot$) pre-main-sequence objects. Since these objects share many characteristics with high-mass stars (clustering, PDRs), but are far closer to us and less embedded, the detection of circumstellar disks around these stars is crucial to the understanding of the massive-star-formation process. Herbig Ae (HAe) stars are the precursors of Vega-type systems: determining the frequency and timescales of the disks around HAe stars is therefore also important for planet formation studies.  Substantial theoretical and observational efforts have been carried out recently for the understanding of the disk occurrence and evolution in HAeBe stars \\citep{mee01, vin02, mil01, ack05}. On the basis of the different mid-infrared excesses, \\cite{mee01} proposed that disks around Herbig Ae/Be stars can be classified into two Groups that they interpreted in terms of different geometries: Group I sources with strong mid-infrared (20--100 $\\mu$m) flux excesses and Group II sources with modest infrared excesses. Group I sources have Chiang \\& Goldreich-like flared disks and Group II sources host flatter, self-shadowed disks. \\cite{lei04} confirmed this hypothesis on the basis of 10~$\\mu$m interferometry data. While most HAe stars belong to Group I and are surrounded by disks similar to those in T Tauri (TT) stars, Herbig Be (HBe) stars are mostly found in Group II and seem to be surrounded by disks with a flatter geometry \\citep{ack05}.  Although the existence of different types of disks is well accepted, whether this difference is due to the different characteristics of the central star (hotter stars produce flatter disks), or to different characteristics of the circumstellar matter (different grain opacities), or to an evolutionary link between flared, self-shadowed, and debris disks is not established. If the evolutionary scenario stands, the different geometry between disks in HAe and HBe stars is the consequence of a more rapid grain growth and a shorter dissipation timescale in the latter. The grain growth causes the optical depth of the disk to decrease and allows the UV radiation to penetrate deep into the circumstellar disk and photoevaporate the disk external layers \\citep{dul04}. Disks around HBe would flatten and lose a large fraction of their mass before the pre-main-sequence phase ($<$0.1~Myr).  One problem in discerning between the different scenarios is that most studies are based on the optical-NIR and mid-IR observations. Optical-NIR and mid-IR observations probe the region within a few AU around the star. In addition, since the disk is optically thick at these wavelengths, they only provide information about the disk surface and cannot be used to derive the disk mass. An obvious step towards the formation of planetesimals in disks is coagulation and growth of the submicron sized grains accreted from the proto-stellar envelope \\citep{dom07}. Although the end results are plainly visible in our own planetary system, the details and mechanisms of dust evolution are not understood. The main problem is that the cooler, large grains do not emit significantly at infrared wavelengths. Only by observing the submm-mm range of the spectrum are we able to determine the sizes and properties of these large grains. % Millimeter observations must be used to trace this more evolved grain component. Thus, imaging disks at mm wavelengths is required to trace the outer part of the disk and determine its mass and grain properties.  We have carried out a search for circumstellar disks around HBe stars using the NRAO Very Large Array (VLA) and the IRAM Plateau de Bure (PdB) interferometers to investigate the properties of the circumstellar disks around intermediate-mass stars and to determine eventually their occurence, lifetime, and evolution. Some results have already been published elsewhere \\citep{fue03,fue06,alo07,alo08}. In this Paper, we present new observational data (LkH$\\alpha$~215, Z~CMa and MWC~297) that complete our survey and we provide a comprehensive analysis of all the data obtained thus far. ",
        "conclusions": "We briefly summarize the main conclusions of our study:  \\begin{itemize} \\item We have carried out a search for circumstellar disks around HBe stars using the NRAO Very Large Array (VLA) and IRAM Plateau de Bure (PdB) interferometers. Interferometric observations are necessary to separate the disk from the surrounding envelope, although the disks themselves usually remain unresolved. Our goal is to investigate the properties of the circumstellar disks around intermediate mass stars to determine eventually their occurrence, lifetime, and evolution. Thus far, we have observed 6 objects with 4 successful detections.  \\item Our first result is that the disk mass is usually only a small percentage (less than 10\\%) of the mass of the entire envelope. There are significant variations in the disk mass from source to source. Two disks of our sample, R~Mon and Z~CMa, have similar sizes to those found in TT and HAe stars. Z~CMa is a FU Orionis object and cannot be directly compared with the others. The mass of the other large disk, R~Mon, is a factor of 5 below that of an average HAe star. The disk around MWC~1080 is smaller (r$_ {out}$$<$100~AU) and less massive. Around MWC~297, we find an inner compact disk (r$_{out}$$\\sim$~29~AU) and an outer ring at a distance of 200~AU. The 1.3mm/2.7mm spectral index indicates that grain growth has proceeded in all disks in our sample.  \\item A comparison between our data and previous results in TT and HAe stars show that the masses of the circumstellar disks around HBe stars are on average 5-10 times lower than those around Herbig Ae stars. We propose that disk photoevaporation is responsible for this behavior. In HBe stars, the UV radiation disperses the gas in the outer disk on a timescale of a few 10$^5$ yr. Once the outer part of the disk dissapears, the entire gaseous disk is photoevaporated on a short timescale ($\\sim$10$^5$ yr), and only a dusty disk composed of large grains remains. \\end{itemize}   \\begin{acknowledgement} We thank the IRAM staff in Grenoble for their help and support during the observations and data reduction. This work was partly supported by INAF PRIN 2006 grant $``$From disks to planetary systems$\"$. We thank the referee for his/her fruitful comments. \\end{acknowledgement}"
    },
    "0812/0812.0010_arXiv.txt": {
        "abstract": "We revisit Lyman-$\\alpha$ bounds on the dark matter mass in $\\Lambda$ Warm Dark Matter ($\\Lambda$WDM) models, and derive new bounds in the case of mixed Cold plus Warm models ($\\Lambda$CWDM), using a set up which is a good approximation for several theoretically well-motivated dark matter models. We combine WMAP5 results with two different Lyman-$\\alpha$ data sets, including observations from the Sloan Digital Sky Survey. We pay a special attention to systematics, test various possible sources of error, and compare the results of different statistical approaches. Expressed in terms of the mass of a non-resonantly produced sterile neutrino, our bounds read $m_\\dw \\ge 8$~keV (frequentist 99.7\\% confidence limit) or $m_\\dw \\ge 12.1$~keV (Bayesian 95\\% credible interval) in the pure $\\Lambda$WDM limit. For the mixed model, we obtain limits on the mass as a function of the warm dark matter fraction $\\Fwdm$. Within the mass range studied here ($5~\\mathrm{keV} < m_\\dw < \\infty$), we find that any mass value is allowed when $\\Fwdm < 0.6$ (frequentist 99.7\\% confidence limit); similarly, the Bayesian joint probability on $(\\Fwdm, 1/m_\\dw)$ allows any value of the mass at the 95\\% confidence level, provided that $\\Fwdm < 0.35$. ",
        "introduction": "\\label{sec:intro}  The nature of Dark Matter (DM) is one of the most intriguing questions of particle astrophysics. Its resolution would have a profound impact on the development of particle physics beyond its Standard Model (SM). Although the possibility of having massive compact halo objects (MACHOs) as a dominant form of DM is still under debates (see recent discussion in~\\cite{Calchi:07} and references therein), it is widely believed that DM is made of non-baryonic particles.  Yet the SM of elementary particles does not contain a viable DM particle candidate -- a massive, neutral and stable (or at least cosmologically long-lived) particle.  Active neutrinos, which are both neutral and stable, form structures in a top-down fashion~~\\cite{Zeldovich:70,Bisnovatyi:80,Bond:80,Doroshkevich:81,Bond:83}, and thus cannot produce the observed large scale structure. Therefore, the DM particle hypothesis implies an extension of the SM. Thus, constraining the properties of DM helps to distinguish between various DM candidates and may help to differentiate among different models Beyond the Standard Model (BSM).  What is known about the properties of DM particles?  DM particle candidates may have very different masses (for reviews of DM candidates see e.g.~\\cite{Bergstrom:00,Bertone:05,Carr:06,Taoso:07}): massive gravitons with mass $\\sim 10^{-19}\\ev$~\\cite{Dubovsky:04}, axions with mass $\\sim 10^{-6}\\ev$~\\cite{Wilczek:77,Weinberg:77,Holman:83}, sterile neutrinos with mass in the keV range~\\cite{Dodelson:93}, supersymmetric (SUSY) particles (gravitinos~\\cite{Pagels:82}, neutralinos~\\cite{Lee:77,Haber:85}, axinos~\\cite{Covi:99,Covi:01} with mass ranging from a few eV to hundreds GeV), supersymmetric Q-balls~\\cite{Kusenko:97b}, WIMPZILLAs with mass $\\sim 10^{13}\\gev$~\\cite{Kuzmin:98,Chung:99}, and many others. Thus, the mass of DM particles becomes an important characteristics which may help to distinguish between various DM candidates and, more importantly, may help to differentiate among different BSMs.  A quite robust and model-independent \\emph{lower bound} on the mass of spin-one-half DM particles was suggested in~\\cite{Tremaine:79}. The idea was based on the fact that for any fermionic DM the average phase-space density (in a given DM-dominated, gravitationally bound object) cannot exceed the phase-space density of the degenerate Fermi gas.  This argument, applied to the most DM-dominated dwarf Spheroidal (dSph) satellites of the Milky Way leads to the bound $m_\\dm > 0.41\\kev$~\\cite{Boyarsky:08a}. For particular DM models (with known primordial velocity dispersion) and under certain assumptions about the evolution of the system during structure formation, this limit can be strengthened.  This idea was developed in a number of works~\\cite{Madsen:84,Madsen:91,Madsen:90,Dalcanton:00,Hogan:00,Madsen:00,Madsen:01,Boyanovsky:08,Boyarsky:08a,Gorbunov:08b} (for recent results and a critical discussion see~\\cite{Boyarsky:08a}).   For any DM candidate there should exist a production mechanism in the early Universe. Although it is possible that the DM is produced entirely through interactions with non-SM particles (e.g. from the inflaton decay) and is inert with respect to all SM interactions, many popular DM candidates are produced via interactions with the SM sector.  One possible class of DM candidates are \\emph{weakly interacting massive particles} -- WIMPs~\\cite{Lee:77}. They have interactions with SM particles of roughly electroweak strength.  The WIMPs are completely stable particles, as their interaction strength would otherwise imply a lifetime incompatible with a sizable relic density today. Therefore, they are produced in the early Universe via annihilation into SM particles with roughly weak interaction cross-section (see e.g.~\\cite{Kolb:90}).  To have a correct abundance, these particles should have a mass in the GeV to TeV range. Such a mass and interaction strength means that they decouple from the primeval plasma deeply in the radiation dominated epoch, while being non-relativistic.  Such particles have a Boltzmann velocity distribution function at the moment of freeze-out and are called ``cold DM'' (CDM). In the CDM scenario, structure formation goes in the bottom-up fashion: the smallest structures (with supposed masses as small as that of the Earth) form first, and then combined into larger structures.  Another wide class of DM candidates can be generically called \\emph{superweakly} interacting (see e.g.~\\cite{Feng:2003xh}), meaning that their cross-section of interaction with SM particles is much smaller than the weak one.  There are many examples of \\emph{super-WIMP} DM models: sterile neutrinos~\\cite{Dodelson:93}, gravitinos in theories with broken R-parity~\\cite{Takayama:00,Buchmuller:07}, light volume moduli~\\cite{Conlon:07} or Majorons~\\cite{Lattanzi:07}. Many of these candidates (e.g. sterile neutrino, gravitino, Majoron) possess a two-body decay channel: $\\dm \\to \\gamma + \\nu,\\gamma+\\gamma$.  These scenarios are very interesting, since apart from laboratory detection such particles can be searched in space. Looking for a monochromatic decay line in the spectra of DM-dominated objects provides a way of \\emph{indirect detection} of DM, and could help to \\emph{constrain} its interaction strength with SM particles. The astrophysical search for \\emph{decaying} DM is in fact more promising than for stable (annihilating) DM. Indeed, the decay signal is proportional to the \\emph{column density} $\\int\\rho_\\dm(r)dr$, i.e.  the density integrated along the line of sight, and not the squared density $\\int\\rho^2_\\dm(r)dr$ (as it is the case for annihilating DM). As a consequence, \\emph{(i)} a vast variety of astrophysical objects of different nature would produce roughly the same decay signal~\\cite{Boyarsky:06c,Bertone:07}; \\emph{(ii)} this gives more freedom for choosing the observational targets, and avoid complicated astrophysical backgrounds (e.g. one does not need to look at the Galactic center, expecting a comparable signal from dark outskirts of galaxies, clusters and dark dSph's); \\emph{(iii)} if a candidate line is identified, its surface brightness profile may be measured (since it does not decay quickly away from the centers of objects), in order to distinguish it from astrophysical lines (which usually decay in outskirts) and to compare with the signal from several other objects. For all these reasons, astrophysical search for decaying DM can almost be viewed as another type of \\emph{direct detection}.  A search for DM decay signals was conducted both in the keV -- MeV range~\\cite{Boyarsky:05,Boyarsky:06b,Boyarsky:06c,Boyarsky:06d,Boyarsky:06e,Boyarsky:06f} and GeV range~\\cite{Bertone:07}.  The small strength of interaction of light super-WIMP particles usually implies that: \\begin{inparaenum}[\\em (i)] \\item they were never in thermal equilibrium in the early Universe; and \\item they were produced deep in the radiation dominated epoch, while still being relativistic (unlike CDM). \\end{inparaenum} The latter property means that the density perturbation of these particles are suppressed for scales below their \\emph{free-streaming length},\\footnote{The free-streaming length is the equivalent of the Jeans length for a collisionless fluid. It is given by the ratio $\\lambda_\\fs \\sim \\langle v \\rangle /H$, where $\\langle v \\rangle$ is the velocity dispersion of the particles.} which affects structure formation on those scales. If the maximum free-streaming length roughly coincides with galaxy scales, the particles are called Warm Dark Matter (WDM), see e.g.~\\cite{Bode:00}.  Currently, the best way to distinguish between WDM and CDM models is the analysis of Lyman-$\\alpha$ (Ly-$\\alpha$) forest data~(for an introduction see e.g.~\\cite{Hui:97,Gnedin:01,Weinberg:03}).  This method requires particular care. Part of the physics entering into the \\lya analysis is not fully understood (for a recent review see \\cite{Meiksin:07}), and can be significantly influenced by DM particles~(see e.g.~\\cite{Viel:2003fx,Kim:07a,Bolton:07a,Gao:07,Biermann:06,Stasielak:06,Faucher:07}). The results can also be affected by approximations related to computational difficulties.  Indeed, at the redshifts probed by Ly-$\\alpha$ data, structures are already undergoing non-linear evolution. In order to relate the measured flux spectrum with the parameters of each cosmological model, one would have to perform a prohibitively large number of numerical simulations.  Therefore, various simplifying approximations have to be realized. Different approaches are presented in~\\cite{Viel:04,McDonald:05,Viel:05c}.  Previous works analyzing \\lya constrains on WDM~\\cite{Hansen:01,Viel:05,Abazajian:05b,Seljak:06,Viel:06,Viel:07} assumed the simplest $\\Lambda$WDM model with a cut-off in the linear power spectrum (PS) of matter density fluctuations. Such a PS arises when WDM particles are \\emph{thermal relics}.  However, in many super-weakly interacting DM models, the PS has a complicated non-universal form due to the non-thermal primordial velocity distribution~(see e.g.~\\cite{Feng:2004zu,Roszkowski:04,Cerdeno:2005eu,Choi:2008zq}).  For example, in gravitino models with broken R-parity, gravitino production occurs in two stages: thermally at high temperatures (see e.g.~\\cite{Bolz:00,Pradler:06,Rychkov:07}), and then non-thermally through the late decay of the next-to-lightest supersymmetric particle~(see e.g.~\\cite{Borgani:96,Feng:2004zu}). Therefore the primordial velocity distribution results from the superimposition of a colder (resonant) and a warmer (thermal) component, and the PS is characterized by a step and a plateau on small scales (rather than a sharp cut-off).  Another interesting super-WIMP WDM model with a non-trivial PS is the sterile neutrino DM of the \\numsm -- an extension of the SM by three sterile (right-handed, gauge singlet) neutrinos~\\cite{Asaka:05a,Asaka:05b}.  Having masses of all new particles below the electro-weak scale, this model is able to explain simultaneously three BSM phenomena: \\begin{inparaenum}[\\em (i)] \\item the transition between the neutrino of different flavours (\\emph{neutrino oscillations}); \\item the production of a non-zero baryon number in the early Universe; and \\item the existence of DM candidate with correct relic density.\\footnote{This model also eases the \\emph{hierarchy problem}, as the \\numsm can be thought of as a consistent quantum field theory, valid all the way to the Planck scale. The hierarchy problem is then shifted into the realm of quantum gravity, broadly understood as a fundamental theory, superseeding at the Planck scale~\\cite{Shaposhnikov:07b}. } \\end{inparaenum} Sterile neutrino DM in the \\numsm is produced via its mixing with active neutrinos in the early Universe~\\cite{Dodelson:93,Shi:98,Abazajian:01a,Asaka:06c,Laine:08a}. In the absence of lepton asymmetry, the production rate is strongly suppressed at temperatures above few hundreds of MeV and peaks roughly at $ T_{peak} \\sim 130\\left(\\frac{M_\\dm}{1\\kev}\\right)^{1/3}$~\\mbox{MeV}~\\cite{Dodelson:93}. The resulting momentum distribution of sterile neutrino DM is approximately that of a rescaled Fermi-Dirac~\\cite{Dodelson:93,Dolgov:00}, with its temperature equal to that of active neutrinos (see~(\\ref{eq:fp}) below). We will call this production mechanism \\textbf{NRP} -- \\emph{non-resonant production}.  Notice that this production channel is always present (although it may be subdominant, see below).  The sterile neutrino DM production changes in the presence of a lepton asymmetry~\\cite{Shi:98}.  In this case the dispersion relations for active and sterile neutrinos may cross each other. This results in an effective transfer of an excess of active neutrinos (or antineutrinos) to the population of DM sterile neutrinos. The lepton asymmetry necessary for this \\emph{resonant production} (\\textbf{RP}) is generated via the decay of heavier sterile neutrinos~\\cite{Shaposhnikov:08a}. In the RP scenario, for each mass and lepton asymmetry, the velocity dispersion has a colder (resonant) component and a warmer (non-resonantly produced) tail. For a recent review see e.g.~\\cite{Boyarsky:09a}.  Sterile neutrinos having colder and warmer velocity components can also be produced through the decay of a gauge-singlet scalar, followed by a subsequent non-resonant production~% \\cite{Shaposhnikov:06,Kusenko:06a,Petraki:07,Petraki:08,Boyanovsky:08b} (see also \\cite{Khalil:08,Wu:09} for other mechanisms of production of sterile neutrino).  The decay of this scalar field should occur at temperatures $\\gtrsim 100$~GeV and the produced sterile neutrinos do not share subsequent entropy releases~\\cite{Shaposhnikov:06,Petraki:07}. As a result the characteristic momentum of sterile neutrinos, produced in that way, is several times colder than the average momentum of the NRP component (see e.g.~\\cite{Shaposhnikov:06}). Notice that in this case the production is determined not only by the mixing between active and sterile neutrinos, but also by the coupling of the sterile neutrino with the scalar.  We see that in several well-motivated models, the primordial momentum distribution of DM particles contains a warm (Fermi-Dirac-like component) and a colder one. % Therefore, in this work, apart from updating existing bounds on $\\Lambda$WDM models, we consider \\lya bounds in the \\emph{$\\mathit{\\Lambda}$CWDM} scenario, including a mixture of a CDM and WDM. This model is characterized by two independent parameters beyond the usual dark matter fraction $\\Omega_{\\rm DM}$: the velocity dispersion (or free-streaming length) of the WDM component, and the fraction of WDM density. We will overview and provide a critical analysis of uncertainties in the \\lya method for this model, and derive constraints on the above parameters.  Finally, we will discuss the consequences of these results for some physically relevant DM models of this kind.  The first qualitative analysis of \\lya bounds in the $\\Lambda$CWDM scenario was performed in~\\cite{Palazzo:07}. In this work, the authors computed the \\emph{linear} $\\Lambda$CWDM power spectra projected along the line of sight. The PS suppression (as compared to the pure $\\Lambda$CDM case) at some fiducial scale was compared with that quoted in~\\cite{Seljak:06} for several WDM masses. A $\\Lambda$CWDM model was considered as ruled out at some confidence level if it produced the same suppression in the 1D projected linear PS as a pure $\\Lambda$WDM model ruled out at the same level in~\\cite{Seljak:06}.  This qualitative method should be improved in two directions. First of all, the \\lya method provides a measurement of the 1D \\emph{non-linear} flux spectrum.  The non-linear evolution ``mixes'' different scales. Therefore, two linear PS passing through the same point at a given wave number but having a different behavior otherwise would fit \\lya data differently. Secondly, the effects of CWDM suppression on the matter PS can be somewhat compensated by the change of cosmological parameters (DM abundances, $\\sigma_8$, etc.). Therefore, a proper \\lya analysis should include a simultaneous fitting of all cosmological plus CWDM parameters to the \\lya (and possibly other cosmological) data. Such an analysis is performed in this work.  The paper is organized as follows. In section \\ref{sec:plateau-cwdm-models}, we describe the effect of a mixture of cold and warm dark matter on the linear matter power spectrum, and present some analytic approximations. In section \\ref{sec:flux-power-spectrum}, we introduce the two \\lya data sets used in this analysis: the Viel, Haenelt \\& Springel (VHS) data compiled in Ref.~\\cite{Viel:2004bf}, and the Sloan Digital Sky Survey (SDSS) data complied in Ref.~\\cite{McDonald:04b,McDonald:05}. In section \\ref{sec:veloc-init-cond}, we discuss the issue of including thermal velocities in hydrodynamical simulations in presence of warm dark matter. The reader only interested in our results could go directly to section \\ref{sec:results-vhs}, where we present the bounds based on the conservative VHS data set, and to section \\ref{sec:sdss-results}, which contains our analysis for the more constraining SDSS data. In section \\ref{sec:syst-uncert}, we describe the systematic errors included in our analysis. Finally, in section \\ref{sec:conclusion}, we discuss some implications of our results. ",
        "conclusions": "\\label{sec:conclusion}  \\begin{figure}[t] \\centering \\includegraphics[width=.8\\linewidth]{lya-xray} \\caption{Comparison of X-ray bounds with the results of our \\lya analysis of $\\Lambda$CWDM, for the case in which NRP sterile neutrinos contribute to 60\\% (upper panel) or 20\\% (lower panel) of the dark matter density.  The wide gray band (labeled ``NRP production'') shows the results of the computation of the NRP sterile neutrino abundance, taking into account uncertainties from QCD (see~\\cite{Asaka:06b,Asaka:06c} for details). \u00c2\u00a0The X-ray bounds are from~\\cite{Boyarsky:06d} (label ``MW''), from \\cite{Boyarsky:07a} (label ``M31'') and from~\\cite{Boyarsky:06c} (yellow region in the lower right corner). To account for possible uncertainties in the determination of the DM amount in various object, the X-ray bounds have been rescaled by a factor of 2 (c.f.~\\cite{Boyarsky:07a}). The vertical dashed line marks the $5\\kev$ lower bound on the mass found in the present analysis, assuming 60\\% of WDM. For 20\\% of DM, we found that the mass remains unconstrained by \\lya data.} \\label{fig:lya-xray} \\end{figure}  In this work we have performed a thorough analysis of \\lya constraints on warm and cold plus warm DM models using WMAP5 data, combined with VHS or SDSS \\lya data.  As expected from the quality of the data, the results based on SDSS \\lya data are much more constraining, and below we only summarize this case.   Our results apply to any model in which the phase space distribution of the warm component can be approximated by a Fermi-Dirac or rescaled Fermi-Dirac distribution. The latter case (assuming a temperature $T=T_\\nu$) is a first-order approximation for the case of non-resonantly produced sterile neutrinos. In most of the paper, we expressed our mass bounds in terms of the mass $m_\\dw$ defined in this case.   Within a conservative frequentist approach, for pure $\\Lambda$WDM models, masses $m_{\\dw}$ below $8\\kev$ are excluded at 99.7\\%\\ CL ($m_\\dw \\ge 9.5\\kev$ at 95\\%\\ CL).  In the case of CWDM models, our analysis shows that for $m_\\dw = 5\\kev$ (the smallest WDM mass probed in this investigation) as much as $60\\%$ of WDM is allowed at 99.7\\%\\ CL ($40\\%$ at 95\\%\\ CL). We believe that these results are robust within the current state-of-the-art. Of course, as the \\lya method is still being developed, these results could be affected by still unknown systematic uncertainties (see discussion in the Section~\\ref{sec:syst-uncert}).  From the same data we have also obtained Bayesian credible intervals on the same parameters. The bound on $m_\\dw$ (at 95\\% CL), defined as in previous works~\\cite{Viel:05,Viel:06,Seljak:06}), are $m_\\dw\\ge 12.1\\kev$ for pure $\\Lambda$WDM; for $\\Lambda$CWDM, the joint probability of $(\\Fwdm, 1/m_\\dw)$ is displayed in figure~\\ref{fig:mwdm_fwdm}: this shows that for the smallest probed mass $5\\kev$, the data is compatible with $\\Fwdm \\le 35\\%$.  Our results can be easily translated for the case thermal relics, the corresponding limits are the following. In the frequentist framework, for pure $\\Lambda$WDM, masses $m_\\textsc{tr}$ below $1.5\\kev$ are excluded at 99.7\\%\\ CL ($m_\\textsc{tr} \\ge 1.7\\kev$ at 95\\%\\ CL).  In the case of CWDM models, our analysis shows that for $m_\\textsc{tr} = 1.1\\kev$ (the smallest WDM mass probed in this investigation) as much as $60\\%$ of WDM is allowed at 99.7\\%\\ CL ($40\\%$ at 95\\%\\ CL). In the Bayesian approach, for pure $\\Lambda$WDM, the 95\\% CL bound defined in the same way as before reads $m_\\textsc{tr} \\ge 2.1 \\kev$. For the $\\Lambda$CWDM, the joint constraints on the mass and on the warm fraction can be easily obtained from the left plot in figure~\\ref{fig:mwdm_fwdm}, translated in terms of $m_\\textsc{tr}$ using equation~(\\ref{eq:vtr}).  Let us now discuss the implications of our results for models in which the dark matter (or part of it) consists of non-resonantly produced sterile neutrinos.  The lower mass bound for the pure $\\Lambda$WDM case should be compared with X-ray bounds~\\cite{Boyarsky:05,Boyarsky:06b,Boyarsky:06c,Riemer:06,Watson:06,Boyarsky:06e,% Abazajian:06b,Boyarsky:06d,Boyarsky:06f,Boyarsky:07a}, that restrict the mass of NRP sterile neutrinos from above. Confronting the upper bound $m_\\dw \\le 4\\kev$ (see e.g. Ref.~\\cite{Boyarsky:07a}) with the bound $m_\\dw\\ge 8\\kev$ obtained in this work (both at 99.7\\% CL) excludes the NRP scenario.  Again, this conclusion should be taken with care in view of the above discussion of systematic errors.  Next, we analyze the case of mixture of NRP sterile neutrino and some other, cold DM particle (similar to~\\cite{Palazzo:07}). To this end, one should rescale the X-ray constraints of~\\cite{Boyarsky:05,Boyarsky:06b,Boyarsky:06c,Riemer:06,Watson:06,Boyarsky:06e,% Abazajian:06b,Boyarsky:06d,Boyarsky:06f,Boyarsky:07a} and the results of the NRP DM abundance computation~\\cite{Asaka:06b,Asaka:06c} by a factor $\\Fwdm<1$. We can compare the upper bounds rescaled in this way with the results of our CWDM runs (see Fig.~\\ref{fig:lya-xray}). We see that for $\\Fwdm \\ge 0.6$, SDSS \\lya bounds are in conflict with X-ray constraints at $99.7\\%$~CL. For an NRP fraction $\\Fwdm = 0.4 - 0.6$, an allowed window for the mass appears around $m_\\dw \\approx 5\\kev$ (c.f. Fig.~\\ref{fig:lya-xray}, upper panel). For smaller \\Fwdm, the upper mass bound from X-rays quickly increases (as the lower panel of Fig.~\\ref{fig:lya-xray} demonstrates). At the same time, given that 1 and 2$\\sigma$ \\lya contours for $\\Fwdm$ become horizontal as $m_\\dw$ approaches $5\\kev$ (c.f.  Fig.~\\ref{fig:mwdm_fwdm}), we expect that masses smaller than 5~keV with $\\Fwdm \\lesssim 0.4$ should be allowed (although a precise analysis should properly take into account thermal velocities, c.f. Sec.~\\ref{sec:veloc-init-cond}). Thus, we conclude that for $\\Fwdm < 0.6$ a window of allowed NRP masses opens up, and for $\\Fwdm \\sim 0.2$ we see that all masses (above the phase-space density bound~\\cite{Boyarsky:08a}) are allowed up to $\\sim 15\\kev$.   Another DM model which has similarities with the $\\Lambda$CWDM scenario discussed here is that of resonantly produced sterile neutrinos~\\cite{Shi:98,Shaposhnikov:08a,Laine:08a}.  In this case, the DM velocity dispersion has a colder (resonant) and a warmer (NRP produced) component. In this case, a specific analysis is needed in order to apply the results of this work. We will present these results elsewhere~\\cite{Boyarsky:08d}.  Further improvements in the results presented here would require performing simulations for a larger grid of points $(m_\\dw,\\Fwdm)$, with a better control of systematics as well as better understanding of the physics of the IGM. Ideally, one would like to fit all the IGM statistics at once (flux power spectrum, flux probability distribution function, flux bispectrum and all the line-based statistics) in the spirit of \\cite{Tytler:04,Bolton:07a} to find a consistent IGM model from low to high redshifts.   High redshift QSOs are important for WDM (and any matter power spectrum related) investigations, since at higher redshifts the structures are more linear and should more closely resemble the underlying matter power spectrum. However, at $z>5$ the IGM becomes more neutral and the mean flux level lowers. Thus, interpreting the almost saturated spectra would require hydrodynamical simulations at much higher resolution in order to account for very small structures that contribute to absorption. Furthermore, the number density of QSOs drops significantly at high redshifts and thereby the progress will be somewhat limited also in a statistical sense.  In principle, QSO spectra (especially those at high resolution and with high signal-to-noise) carry information down to scales comparable to the resolution element (for VLT, Very Large Telescope, UVES-spectra $\\sim 0.02$ \\AA\\ which could be roughly translated into 200 $h^{-1}$ comoving kpc at $z=3$), thereby starting to probe the sub-Mpc scales. Although these numbers are very promising, it is important to stress that at these scales the flux power spectrum becomes very steep, so that a small change in the astrophysical and cosmological parameters could have a large impact on it. To interpret it with hydrodynamical simulations could become challenging both for resolution and physical issues (all the physical ingredients that could have an influence on galactic scales need to be modelled and their effect re-computed). It is probably more promising to tackle some physical effects such as the IGM thermal evolution with independent probes (like line statistics) and use them as strong priors in the interpretation of the flux power at small scales, in order to restrict the parameter space.  A large number of intermediate resolution QSO spectra, between that of SDSS and VLT spectra, will be also very important in order to cross-check the effects of systematics on the data and improve the limits obtained in this work. This effort could be possible quite soon with the X-shooter spectrograph \\cite{Kaper:08}."
    },
    "0812/0812.2015_arXiv.txt": {
        "abstract": "{% The 21$\\mum$ and 30$\\mum$ bands are the strongest dust emission features detected in evolved low- and intermediate-mass C-rich stars (i.e. asymptotic giant branch [AGB] stars, proto-planetary nebulae [PPN], and planetary nebulae [PN]). While the 21$\\mum$ feature is rare and exists only in the transient PPN phase, the 30$\\mum$ feature is more common and seen in the entire late stage of stellar evolution, from AGB to PPN and PN phases, as well as in the low-metallicity galaxies: the Large Magellanic Cloud (LMC) and the Small Magellanic Cloud (SMC). The carriers of these features remain unidentified. Eleven of the twelve well-identified 21$\\mum$ sources also emit in the 30$\\mum$ band, suggesting that their carriers may be somewhat related. % } ",
        "introduction": "The circumstellar chemistry of the low- and intermediate-mass stars from AGB on is dichotomized into O-rich and C-rich classes according to the C/O abundance ratio, since C and O would combine at relatively high temperature to form the very stable molecule CO. The dominant dust species in O-rich circumstellar envelope (CSE) is silicate, abundant in both amorphous and crystalline forms. In C-rich objects, the C-bearing molecules (e.g. C$_{2}$, C$_{3}$, CN, and some hydrocarbon compounds) dominate, while the most common dust feature should be assigned to the SiC 11.3$\\mum$ band detected in more than 700 stars most of which are C-rich objects. However, the strongest and most enigmatic spectral features of C-rich objects are the 21$\\mum$ and 30$\\mum$ emission bands. The combined contribution of these two bands can amount to more than 20\\% of the total infrared flux, so they greatly influence the CSE evolution. Furthermore, the dispersion of the CSE dust to the interstellar medium (ISM) brings the dust for these features enter into the recycling of cosmic dust. ",
        "conclusions": ""
    },
    "0812/0812.0821_arXiv.txt": {
        "abstract": "We study the dynamical evolution of spiral structure in the stellar disks of isolated galaxies using high resolution Smoothed Particle Hydrodynamics (SPH) simulations that treat the evolution of gas, stars, and dark matter self-consistently.  We focus this study on the question of self-excited spiral structure in the stellar disk and investigate the dynamical coupling between the cold, dissipative gaseous component and the stellar component.  We find that angular momentum transport from the gas to the stars inside of corotation leads to a roughly time-steady spiral structure in the stellar disk.  To make this point clear, we contrast these results with otherwise identical  simulations that do not include a cold gaseous component that is able to cool radiatively and dissipate energy, and find that spiral structure, when it is incipient, dies out more rapidly in simulations that do not include gas.  We also employ a standard star formation prescription to convert gas into stars and find that our results hold for typical gas consumption time scales that are in accord with the Kennicutt-Schmidt relation.  We therefore attribute the long-lived roughly time steady spiral structure in the stellar disk to the dynamical coupling between the gas and the stars and the resultant torques that the self-gravitating gaseous disk is able to exert on the stars due to an azimuthal phase shift between the collisionless and dissipative components. ",
        "introduction": "Spiral galaxies present two faces to observers.  When seen in blue light, these galaxies look disordered, with spurs, branches and feather-like structures seen even in grand design galaxies, while K-band images show that these galaxies possess smooth, usually two-armed spiral patterns.  We study here the dynamical interaction between the Population I component seen in blue light and the older Population II component revealed in the longer wavelengths, specifically with the intent of understanding how the interaction between the gas and the stars affects the time evolution of spiral structure in the stellar disk.  The idea that the stellar spiral structure is not a material phenomenon, but a roughly time-steady density wave rotating at a pattern speed was developed in the seminal work of Lin \\& Shu (1964; henceforth LS64).  Toomre (1964) gave the minimum radial stellar velocity disperion required to suppress the onset of local instabilities in a razor-thin stellar disk stabilized by rotation and the velocity dispersion of stars on large and small scales respectively.  Much of the early numerical work of the study of spiral structure utilized N-body simulations to study the growth of spiral instabilities in the collionsionless components of galaxies.  Transient spiral structure was evinced in the early N-body simulations of Hohl (1970), as well as in the later works of Sellwood \\& Carlberg (1984; henceforth SC84).  Toomre (1981) showed that a shearing disk with a supply of leading waves will swing amplify these leading disturbances to trailing disturbances, and was one of the earliest works to emphasize the role of cosmological accretion of gas.  Tidal disturbances or minor mergers can torque the stellar disk and create transient spiral structure (Hernquist \\& Mihos 1995).  Some of these early influential works noted that the dissipative nature of gas may provide a self-regulatory influence in the combined gas-disk system (Lin \\& Shu 1964; Bertin \\& Lin 1996), with dissipation in the gas layer serving to continually regenerate the shearing spiral instability (Goldreich \\& Lynden-Bell 1965).  Contemporary numerical investigations of spiral structure have focused mainly on the response of a gas disk when driven by a spiral potential, to varying levels of physical detail and numerical sophistication.  Kim \\& Ostriker (2002) carried out local magnetohydrodynamic (MHD) simulations to study the development of spiral substructure in a gas disk forced by a spiral potential.  Chakrabarti, Laughlin \\& Shu (2003) (henceforth CLS) studied the development of spiral substructure due to ultraharmonic resonances using global simulations which incorporate the self-gravity of the gas and found that inter-arm features of long azimuthal extent occurred at resonance locations, which they called ``branches''.  CLS also found that when the gas disk is forced by a time-steady, but nonlinear spiral potential (which is composed of many Fourier components), overlap of resonances produced a highly time-dependent response in the gas which they speculated may signal the onset of turbulence in the purely hydrodynamic regime.  Shetty \\& Ostriker (2006; henceforth SO06) carried out global MHD simulations and found that the magneto-Jeans instability is effective at forming spiral-substructure in low-shear (i.e., near the spiral arms) regions when the magnetic field is large.  The stabilizing effect of magnetic pressure allowed these simulations to be carried out for longer times than those of CLS.  Yanez et al.'s (2007) hydrodynamic simulations also evince the role of ultraharmonic resonances in the development of spiral sub-structure in the gaseous component.  A significant work on the analysis of spiral structure in self-gravitating disks is that of Laughlin et al. (1998) - these authors rigorously show that the primary unstable mode is endemic to the disk and not due to the numerical implementation.   One of the essential points evinced by these simulations of the response of a self-gravitating gas disk to a time-steady stellar potential is that low $Q$ gas disks display a very time-dependent response and continue to do so over many dynamical times (where we refer to the Toomre Q parameter, i.e., $Q=c_{i}\\kappa/\\pi G \\Sigma$, where $c_{i}$ is the sound speed of the relevant component; Toomre 1964).  This is in marked contrast to purely stellar disks which may exhibit transient spiral structure (if they are low $Q$ disks) but the spiral patterns disappear after a few dynamical times.  Block et al. (1994; 1996) found that many spiral galaxies have a grand design structure in infrared light (Elmegreen \\& Elmegreen 1984; Rix \\& Zaritsky 1995) while appearing quite flocculent in blue light.  They interpret this discrepancy in morphology to mean that the Population I and Population II components are dynamically decoupled.  Recent multi-wavelength observations have begun to address in detail the long term evolution of spiral structure in nearby galaxies.  Notably, Kendall et al. (2008) infer from analysis of high spatial resolution multi-wavelength photometry of M81 (from the B-band to IRAC bands) a phase shift between the gas and the stars that is suggestive of a long-lived spiral structure in the stellar disk.  Global torques may have a role in fueling star formation over many Gyrs if they are able to transport the large amounts of mass in extended HI disks (Wong \\& Blitz 2002) to the inner regions.  These issues require a ${\\it global}$ analysis that treats the interplay between gas and stars rather than treating the response of the gas to an imposed stellar potential as has been studied previously (CLS; SO06, among other previous papers).   In this paper, we employ high resolution three-dimensional simulations performed with the Smoothed Particle Hydrodynamics (SPH) code GADGET-2 (Springel 2005) to simulate the evolution of $\\it{self-excited~spiral~structure}$ in the stellar disks of galaxies that are evolving in relative isolation. The recent analysis of the chemical and star formation history of the Milky Way by Colavitti et al. (2008) suggests that the Milky Way has not undergone recent (i.e., after $z \\sim 1$) significant merger events.  We study here the problem of the generation and evolution of spiral structure in an isolated galaxy with gas, stars, and a dark matter halo, where these components are free to interact dynamically with each other.  Specifically, by including not only the stars with their gravitational forces and velocity dispersions, but also the cold interstellar gas with its self-gravity, pressure, its ability to cool radiatively and dissipate energy through shocks, and form new stars, we focus our study on the dynamical coupling between the gas and the stars and the torques that the gas is able to exert on the stars, an effect that has not been analyzed in previous studies which have focused on the effect of the stars (whether live or modeled through a potential) on the gas (Hernquist \\& Mihos 1995; CLS; S006; Kim \\& Ostriker 2007).  We also comment on the coupling with the live halo and its role as a sink of angular momentum, as has been noted in previous studies of bar instabilities (Athanassoula 2003; Martinez-Valpuesta et al. 2006; Berentzen et al. 2007).  We find that the dissipative nature of the gas leads to $\\it{sharp}$, $\\it{asymmetric}$ azimuthal features, with the gas response leading the stellar response, which allows for a phase shift between these components such that the gas is able torque the stars.  It is this angular momentum transport interior to co-rotation that allows for the stellar disk to be unstable to spiral structure for a longer time (many revolutions) compared to simulations that do not include the dissipative gaseous component.  The paper is organized as follows.  In \\S 2, we review the simulation methodology, specifically the initial equilibrium structure of the disks and implementation of radiative cooling, viscosity, and star formation prescriptions for turning gas into stars.  In \\S 3, we give our results for simulations that treat the collionsionless components only, and in \\S 4, we describe results from simulations that include the time evolution of gas, stars, and a dark matter halo. We discuss our results in the context of observations in \\S 5, and we conclude in \\S 6.  The Appendices are devoted to a detailed exploration of simulation parameters and assumptions that can affect the gas-star angular momentum transport, specifically the artificial bulk viscosity and equation of state for the gas. ",
        "conclusions": "$\\bullet$ We find that angular momentum exchange between the self-gravitating gas disk and the stars leads to relatively long-lived spiral structure in the stellar disk, with spiral structure lifetimes of order a few Gyr.  Spiral structure is washed out, when it is incipient, in simulations that do not include a cold, dissipative gaseous component.  $\\bullet$ The gas disk is able torque the stellar disk and donate angular momentum interior to co-rotation.   The gas is able to cool radiatively and dissipate energy - leading to a time lag (or azimuthal phase shift on the order of $\\sim 1~\\rm radian$ close to co-rotation) between the response of the gas and the response of the stars.  Even though the gas mass is less than the stellar mass, the gas can exert appreciable torques (and donate non-negligible angular momentum) as the azimuthal response which is highly asymmetric (or azimuthal gradient of the potential) is much larger than that of the stars.  Over timescales of many Gyr, there is radial inflow of gas which ultimately limits the donation from the gas (if it is not replenished due to cosmological infall); star formation (depending on gas consumption time-scales) will convert the cold, dissipative component to stars which heat up and have no cooling vent.  $\\bullet$ The primary (specifically the $m=2$) Fourier components of the stellar surface density are roughly time-steady in simulations that include a gaseous component, with variation in surface density of $\\sim 3$ over a Gyr.  The pattern speed of the $m=2$ component is more time steady in simulations that include gas than those that do not.  $\\bullet$ The halo is a net sink of angular momentum, as has been seen in previous studies of bar instabilities with live halos.  \\bigskip \\bigskip"
    },
    "0812/0812.0360_arXiv.txt": {
        "abstract": "We demonstrate that the Sommerfeld correction to cold dark matter (CDM) annihilations can be appreciable if even a small component of  the dark matter is extremely cold. Subhalo substructure provides such a possibility given that the smallest clumps are relatively cold and contain even colder substructure  due to incomplete phase space mixing. Leptonic channels can be enhanced for plausible models and the solar neighbourhood boost required to account for PAMELA/ATIC data is plausibly obtained, especially in the case of  a few TeV mass neutralino for which the Sommerfeld-corrected boost is found to be $\\sim10^4-10^5.$ Saturation of the Sommerfeld effect   is shown to occur below $\\beta\\sim 10^{-4},$ thereby making this result largely independent on the presence of substructures below $\\sim 10^5\\rm M_\\odot$. We find that  the associated diffuse gamma ray signal from annihilations  would  exceed EGRET constraints unless the channels annihilating to heavy quarks or to gauge bosons are suppressed. The lepton channel gamma rays are potentially detectable by the FERMI satellite, not from the inner galaxy where substructures are tidally disrupted, but rather as a quasi-isotropic background from the outer halo, unless  the outer substructures are  much less concentrated than the inner substructures  and/or the CDM density profile  out to  the virial radius steepens significantly. ",
        "introduction": "The motivation for studying dark matter annihilation signatures (see e.g. \\cite{Bertone:2004pz}) has received considerable recent attention following reports of a 100 GeV excess in the PAMELA data on the ratio of the fluxes of cosmic ray positrons to electrons \\cite{Adriani:2008zr}. In the absence of any compelling astrophysical explanation, the signature is  reminiscent of the original prediction of a unique dark matter annihilation signal \\cite{Silk:1984zy}, although there are several  problems that demand attention before any definitive statements can be made.  By far the most serious of these is the required annihilation boost factor. The remaining difficulties with a dark matter interpretation, including most notably the gamma ray signals from the Galactic Centre and the inferred leptonic branching ratio, are,  as we argue below, plausibly circumvented or at least alleviated. Recent data from the ATIC balloon experiment  provides evidence for a cut-off in the positron flux near 500 GeV that supports a Kaluza-Klein-like candidate for the annihilating particle \\cite{chang:2008zzr} or a neutralino with incorporation of suitable radiative corrections \\cite{berg}.   In a pioneering paper,  it was  noted \\cite{Profumo:2005xd} that the annihilation signal can be boosted by a combination of coannihilations and Sommerfeld corrrection. We remark  first that the inclusion of coannihilations to boost the annihilation cross-section modifies the relic density, and opens the 1-10 TeV  neutralino mass window to the observed (WMAP5-normalised) dark matter density. As found by  \\cite{Lavalle et al.}, the outstanding problem now becomes  that of normalisation.  A boost factor of around  100 is required to explain the HEAT data in the context of a 100 GeV neutralino.  The flux is suppressed by between one and two powers of neutralino mass, and the problem becomes far more severe with the 1-10 TeV neutralino  required by the PAMELA/ATIC data \\cite{Cirelli:2008pk}, a boost of $10^4$ or more being required. These latter authors included a Sommerfeld correction appropriate to our $\\beta\\equiv v/c=0.001$ dark halo and incorporated channel-dependent boost factors to fit the data, but  the required boosts still fell short of plausible values by  at least an order of magnitude.  Here we propose a solution to the boost problem  via Sommerfeld correction in the presence of a model of substructure that incorporates a plausible phase space structure for cold dark matter (CDM). We reassess the difficulty with the  leptonic branching ratio and show that it is not insurmountable  for supersymmetric candidates. Finally, we evaluate  the possibility of independent confirmation via photon channels.  Substructure survival means that as much as 10\\% of the dark matter is at  much lower $\\beta$. This is likely in the solar neighbourhood and beyond, but not in the inner galaxy where clump destruction is prevalent due to tidal interactions.  Possible annihilation signatures from the innermost galaxy such as the WMAP haze  of synchrotron emission  and the EGRET flux of diffuse gamma rays  are likely to be much less affected by clumpy substructure than the positron flux in the solar neighbourhood.  We show in the following section that incorporation of the Sommerfeld correction means that clumps dominate the annihilation signal, to the extent that the initial clumpiness  of the dark halo survives. ",
        "conclusions": "In the previous section we have shown how it is possible to get a boost factor of order $10^4-10^5$ for a dark matter particle mass of order $4.5\\,\\TeV$. This is tantalizing because this is roughly the value one needs to explain the PAMELA data for a dark matter candidate with this given mass, as can be inferred by analysis of Fig. 9 of Ref \\cite{Cirelli:2008pk}. Although we have made several approximations concerning the clump distribution and velocity, it should be noted that our results still hold as long as the majority of the clumps are very cold ($\\beta\\lesssim 10^{-4}$) because this is the regime in which the enhancement becomes constant. The saturation of the Sommerfeld effect also plays a crucial role in showing that the very coldest  clumps are unable to contribute significantly  to the required boost factor if the dark matter mass is not close to one of the Sommerfeld resonances. Because of saturation below $\\beta\\sim 10^{-4}$, the Sommerfeld boost is insensitive to extrapolations beyond the currently resolved  scales in simulations. Note however that the precise value for the dark matter particle mass is uncertain because of such model-dependent assumptions as the adopted mass-splitting, the multiplet nature of the supersymmetric particles, and the possibility of different couplings, weaker than weak.  The model presented here does not pose any problem from the point of view of the high energy gamma-ray emission from the centre of the galaxy, since very few clumps are presents in the inner core and thus there is no Sommerfeld enhancement. Thus there is no possibility of violating the EGRET or HESS observations of the galactic center or ridge, contrary to what is argued in Ref. \\cite{Bertone08}. There is a potential problem however with gamma ray production beyond the solar radius out to the outer halo. From \\cite{springela}, the simulations are seen to yield an additional enhancement due to clumpiness alone  above  $10^5\\rm M_\\odot$ of around 80\\% at $r_{200}$ in the annihilation luminosity. Extrapolating to earth mass clumps, the enhancement is 230 in the annihilation luminosity at the same radius.  This is what a distant observer would see. The incorporation of the Sommerfeld factor would greatly amplify this signal by $S\\sim 10^4-10^5.$  The expected flux that would be observed by looking in a direction far from the galactic center can be readily estimated. Assuming an effective cross section $\\sigma v= 3\\times 10^{-22}$ cm$^3$ s$^{-1}$, corresponding to a Sommerfeld boost of $10^4$ on top of the canonical value of the cross section times velocity, the number of annihilations on the line of sight is roughly  $4\\times 10^{-9} (\\mdm/\\TeV)^{-2}$ cm$^{-2}$ s$^{-1}$. We have assumed a Navarro-Frenk-White (NFW) profile. The effect of the clumpiness is still not included in this estimate. Following the results of the simulation in Ref. \\cite{springela}, this value should be multiplied by a factor $\\sim 200.$ Convolving with the single annihilation spectrum of a 5 TeV dark matter particle yields  the flux shown in Fig. \\ref{fig:diff}. There we show the spectrum that would be produced if the dark matter particle would annihilate exclusively either to $W$ bosons, b quarks or $\\tau$ leptons (blue, red and green curves, respectively). We also consider a candidate that annihilates to $\\tau$ leptons 90\\% of the time and to $W$s the remaining 10\\% of the time (model ``Hyb1'') and a candidate that annihilates only to quarks and leptons, with the same cross section apart from color factors (model ``Hyb2'').  The gamma ray signal  mostly originates from the outer halo and should be detectable as an almost isotropic hard gamma-ray background. Candidates annihilating to heavy quarks or to gauge bosons seem to be excluded by EGRET. On the other hand,  a dark matter particle annihilating to $\\tau$ leptons is compatible with the measurements of EGRET at these energies \\cite{Strong:2004ry}, and within the reach of FERMI.   \\begin{figure} \\begin{center} \\includegraphics[width=0.45\\textwidth,  keepaspectratio]{gal-diffuse} \\caption{Contribution to the diffuse galactic photon background from the annihilation of a 5 TeV dark matter particle, for different channels, when both clumpiness and the Sommerfeld enhancement in cold clumps are taken into account, compared with the measurements of the diffuse gamma background from EGRET \\cite{Strong:2004ry}.  The label ``Hyb1'' (solid black line) stands for a hybrid model in which the dark matter annihilates to $\\tau$ leptons 90\\% of the time and to $W$ pairs the rest of the time. The label ``Hyb2'' (dashed black line) stands for a model in which the dark matter annihilates to leptons and quarks only, with the same cross-section apart from color factors. The latter could be realized through the circumvention of helicity suppression.} \\label{fig:diff} \\end{center} \\end{figure}  There are however at least two reasons that induce  significant uncertainty into any estimates. Firstly,  the halo density profile in the outer galaxy may be  substantially steeper than is inferred from an NFW profile, as current models are  best fit by an Einasto profile \\cite{gao}, $\\rho(r)\\propto \\exp[(-2/\\alpha ((r/r_s)^\\alpha - 1)],$ as opposed to the asymptotic NFW profile $\\rho(r)\\propto r^{-3}$. Using the Einasto profile  yields at least a  10\\% reduction. Another possibility is to use a Burkert profile \\cite{Burkert:1995yz}, that gives a better phenomenological description of the dark matter distribution inside the halo, as it is inferred by the rotation curves of galaxies \\cite{Gentile:2004tb,Salucci:2007tm}. Using a Burkert profile, the flux is reduced by a factor 3. Secondly, and  more importantly,  the subhalos are much less concentrated at greater distances from the Galactic Centre \\cite{diemand}. These effects should substantially reduce the gamma ray contribution from the outer halo.  A future application will be to evaluate the extragalactic diffuse gamma ray background where the evolution of clumpiness with redshift should play an interesting role in producing a possible spectral feature in the isotropic component. Note that the annihilation rate originating from very high redshift subhalo substructure and clumpiness near the neutralino free-streaming scale \\cite{Kamionkowski:2008gj} is mostly suppressed due to the saturation of the Sommerfeld effect that we described above.  Because of the  saturation of the Sommerfeld boost, it should be possible to focus future simulations on improved modelling of the radial profiles and concentrations of substructures in the outer halo. It is these that contribute significantly to the expected diffuse gamma if our interpretation of the PAMELA and  the ATIC data,  and in particular the required normalisation and hence boost, is correct. Of course, there are other possible explanations of the  high energy positron data, most notably the flux from a local pulsar \\cite{ahar, yuks,Hooper:2008kg} that has recently been detected as a TeV gamma ray source.  An interesting consequence of the model proposed here is the production of synchrotron radiation emitted by the electrons and positrons produced in the dark matter annihilations, similar to the one that is possibly the cause of the observed ``WMAP haze'' \\cite{Hooper:2007kb,Cumberbatch:2009ji}. For a TeV candidate, this synchrotron emission would be visible in the $\\nu \\gtrsim 100$~GHz frequency region. This region will be probed by the Planck mission; the synchrotron radiation would then give rise to a  galactic foreground ``Planck haze'' in the microwave/far infrared part of the spectrum. This quasi-isotropic high frequency synchrotron component will be an additional source of B-mode foregrounds that will need to be incorporated into proposed attempts to disentangle any primordial B-mode component in the cosmic microwave background. Another interesting application would be to look at the gamma-ray emission from specific objects, like the Andromeda Galaxy (M31). M31 has been observed in the relevant energy range by the CELESTE and HEGRA atmospheric Cherenkov telescopes, and limits on the partial cross section to photons, in the absence of boost, were obtained in Ref.   \\cite{Mack:2008wu}.   Finally, we note that  in Sec. \\ref{sec:lept} we have described a mechanism that can enhance the production of leptons (especially light leptons) in neutralino dark matter annihilations, making the leptonic channel as important as the gauge boson channel. A dark matter candidate annihilating mainly into leptons can simultaneously fit the PAMELA positron and antiproton data, owing to the fact that no antiproton excess is produced. The enhancement of the lepton branching ratio can possibly alleviate the problem of antiproton production following neutralino annihilation into a pair of gauge bosons. It should however be noted that the mechanism in question also enhances the quark channel in a similar way, thus introducing an additional source of antiprotons. It would thus be desirable to suppress in some way the quark annihilation channel. This could be realised in a variation of the above mentioned mechanism, if the lightest neutralino is quasi-degenerate in mass with the lightest slepton $\\tilde l$; this is what happens for example in the $\\tilde \\tau$ coannihilation region. In this case, the Sommerfeld enhancement would proceed through the creation of an intermediate $\\tilde l^+\\tilde l^-$ bound state that would subsequently annihilate to the corresponding standard model lepton pair, without producing any (tree-level) quark. This points to the necessity of further investigating different models in order to assess if the boost in the leptonic branching ratio is indeed compatible with the PAMELA data."
    },
    "0812/0812.0683_arXiv.txt": {
        "abstract": "We analyzed the statistics of subhalo abundance of galaxy-sized and giant-galaxy-sized halos formed in a high-resolution cosmological simulation of a 46.5Mpc cube with the uniform mass resolution of $10^6 M_{\\odot}$. We analyzed all halos with mass more than $1.5 \\times 10^{12}M_{\\odot}$ formed in this simulation box. The total number of halos was 125. We found that the subhalo abundance, measured by the number of subhalos with maximum rotation velocity larger than 10\\% of that of the parent halo, show large halo-to-halo variations. The results of recent ultra-high-resolution runs fall within the variation of our samples. We found that the concentration parameter and the radius at the moment of the maximum expansion shows fairly tight correlation with the subhalo abundance. This correlation suggests that the variation of the subhalo abundance is at least partly due to the difference in the formation history. Halos formed earlier have smaller number of subhalos at present. ",
        "introduction": "The cold dark matter (CDM) model is widely accepted as the standard theory of the formation and evolution of the universe.  According to this CDM model, the structure formation in the universe proceeds hierarchically,  from small-scale structures to larger-scale ones.   The CDM model successfully reproduced large-scale structures \\citep[e.g.,][]{Davis1985, Springel2005}.  However, it has been claimed that there are serious discrepancies between the structure predicted by the CDM model and that observed, in the scale of galaxy-sized objects and below.  \\citet{Moore1999a} and \\citet{Klypin1999} calculated the evolution of galaxy-scale dark mater halos using high-resolution cosmological $N$-body simulations.  They found that dark matter halos of the mass comparable to the Local Group contained far too many subhalos compared to the number of known dwarf galaxies in the Local Group. In the case of dark matter halos of the size of typical clusters of galaxies, the theoretical prediction and observation agreed pretty well. The theoretical prediction for galaxy-sized and cluster-sized halos are similar, and observations of clusters of galaxies and that of Local Group are very different.  This discrepancy, now called ``missing dwarf problem'', was studied by many researchers. Follow-up simulation studies \\citep[e.g.,][]{Diemand2004, Reed2005, Kase2007} gave essentially the same results as those of \\citet{Moore1999a} and \\citet{Klypin1999}. Accordingly, it is now regarded as one of the most serious problems of the CDM model.  A number of solutions for this problem have been proposed.  For example, we can reduce the number of subhalos in small scales by changing the nature of dark matter in small scales. Models in this direction are warm dark matter \\citep{Kamionkowski2000} and self-interacting dark matter \\citep{Spergel2000}.  Another direction is to explain the discrepancy as the difference between the number of nonobserved dark matter subhalos and that of observed satellite galaxies. Maybe not all subhalos host dwarf galaxies, or the measured velocity dispersion of dwarf galaxies might be much smaller than the rotation velocity of the parent subhalos. Models in this direction includes the suppression of star formation by early reionization \\citep{Susa2004} and self-regulation of star formation in small halos \\citep{Stoehr2002, Kravtsov2004}. However, none of these explanations are widely accepted as the clear-cut solution for the missing dwarf problem.   In almost all previous studies of substructures in dark-matter halos, the so-called re-simulation method has been used.  In this method, one simulation is done in the following two steps. In the first step, a relatively large volume (for galaxy-sized halos typically a 50-100Mpc cube) is simulated with low-mass resolution, and the candidate regions for high-resolution simulations are identified. In the second step, the selected regions are simulated with higher mass resolution. In practically all recent high-resolution simulations, this re-simulation method is used \\citep{Diemand2007, Diemand2008, Stadel2008, Springel2008}.  This re-simulation method allows us to resolve the structure of a selected halo with a very high resolution.  On the other hand, it is not clear what kind of selection bias is introduced by the use of this method. Typically, the authors claimed that the halo which looks similar to the halo of the Local Group was selected to compare with the Local Group. However, in most cases the selection criteria were not described in detail, and thus it is difficult to see if any selection bias affected the results or not.  An obvious way to avoid the selection bias, which might be introduced by the use of the re-simulation method, is simply not to use it and do the high-resolution simulation of the entire initial simulation box. In \\citet[][hereafter Paper I]{Ishiyama2008}, we simulated a relatively small region (the size of the simulation box is 21.4 Mpc) with high-mass resolution (mass of particles = $3 \\times 10^6M_{\\odot}$).  We identified 21 dark halos with maximum rotation velocity larger than 200$\\rm kms^{-1}$ and investigated the environmental effect on the subhalo abundance. The main conclusion of Paper I is that the abundance of subhalos has large variation. The most subhalo-rich halos contain the comparable number of subhalos as reported in \\citet{Moore1999a}, while least subhalo-rich ones contain around 1/10 subhalos. They also found that only the galaxy-sized halos (mass less than $3\\times 10^{12}M_{\\odot}$) show large variation in subhalo abundance, and more massive ones show much smaller variation. The subhalo abundance shows correlation with the concentration parameter $c$ (larger $c$, smaller abundance) and measures of the local density such as the smoothed local density and distance to the nearest massive halo.  The accuracy of the result of Paper I is, however, limited by both the mass resolution (particle mass $3 \\times 10^6M_{\\odot}$) and small number statistics. The total number of halos we analyzed was 21, and only 10 had the mass less than $3 \\times 10^{12}M_{\\odot}$.  In this paper, we report the result of much larger, and thus much more accurate, simulation than that in Paper I.  We simulated a box of the size 46.48 Mpc with $1600^3$ particles of mass of $1\\times 10^{6}M_{\\odot}$. We also reduced the softening parameter from 1.8kpc to 700 pc. Thus, compared to the calculation in Paper I, we used three times better mass resolution, two times better spatial resolution, and 10 times larger simulation volume.  We identified all halos with the maximum rotation velocity larger than 160$\\rm kms^{-1}$ and the virial mass $M$ in the range of $1.5\\times 10^{12}M_{\\odot} \\le M < 1\\times 10^{13}M_{\\odot}$.  We found 125 such halos. In all halos, the number of particles in the virial radius is more than $1.5\\times 10^6$, which is large enough to give reliable results for the statistics of subhalos with rotation velocity more than 1/10 of that of the parent halo \\citep{Kase2007}.  In Section 2, we describe the initial model and the simulation method. In Section 3, we present the result. Section 4 presents summary and discussion. ",
        "conclusions": "\\subsection{Comparison with Aquarius simulation}  \\citet{Springel2008} claimed that they found very small halo-to-halo variations in the subhalo abundance. They used the definition of the subhalo count different from what was used in our work and most previous works. In this section, we compare our result with theirs.  Figure \\ref{fig12} shows the cumulative numbers of subhalos as a function of their maximum rotation velocities $V_{\\rm c}$ normalized by the circular velocity of the parent halos at $r_{50}$.  We define the radius $r_{50}$ as inside which the averaged spherical overdensity is 50 times the critical value and count the number of subhalos within $r_{50}$. This definition is the same as that used for Figure 10 in \\citet{Springel2008}. We plot the result of the 44 galaxy-sized halos with $M_{50} < 3.0\\times10^{12}M_{\\odot}$, where $M_{50}$ is the mass within $r_{50}$.  \\begin{figure} \\centering \\includegraphics[width=8cm,trim=80 0 100 0,clip]{fig12.eps} \\caption{ Cumulative numbers of subhalos as a function of their maximum rotation velocities $V_{\\rm c}$ normalized by the circular velocity of the parent halos at $r_{50}$. Three thick solid curves show the average (middle) and $\\pm 1\\sigma$ values (top and bottom). The thick dashed curves show the results of \\citet[][simulation Aq-A-1, Aq-C-2, Aq-E-2]{Springel2008}. } \\label{fig12} \\end{figure}  If we measure the subhalo abundance again using $N_{>0.1}$, the number of subhalos with normalized rotation velocity more than 0.1, we can conclude that our result and \\citet{Springel2008} result are practically identical. However, our results for $N_{>0.1}$ may be affected by the resolution, because $V_{50}$ is typically 20\\% smaller than the maximum rotation velocity $V_{\\rm p}$.  Thus, our halos are probably slightly more subhalo-rich than those of \\citet{Springel2008}, but the halo-to-halo variation, calculated using the definition of \\citet{Springel2008}, is largely similar, though it is difficult to tell much on the variation in \\citet{Springel2008} result because of small number of runs. We show three results here, and the total number of runs performed was six.  Compared with Figure \\ref{fig2}, the variation of the subhalo abundance in Figure \\ref{fig12} is significantly smaller. This result is quite natural because of the two differences in the definition of the subhalo count. First, the use of $r_{50}$ means that many subhalos are at the outermost region which are not virialized. Though subhalos in this outer region are gravitationally bound to the main halo, their evolution is not affected by the structure of the main halo simply because they are far from the central high-density region. Second, the use of $V_{50}$ instead of $V_{\\rm p}$ would also reduce the halo-to-halo variation, since centrally concentrated halos have higher $V_{\\rm p}$.  \\citet{Springel2008} claimed that the small halo-to-halo variation they found contradicted with the result of Paper I.  Their rms halo-to-halo scatter was about 8\\%, while ours (in Figure 2) is about 41\\%. From the comparison above, it is clear that this ``contradiction'' they found is partly due to the difference in the definition of the subhalo abundance.  The halo-to-halo scatter as seen in Figure \\ref{fig12} is smaller than one in Figure \\ref{fig2}.  However, there is still some difference between our results and their ones.  The remaining difference might be caused by their way to select parent halos. They selected the halos which host normal spirals in their semianalytic modelling.  The systematic difference between \\citet{Springel2008} results and our result may be due to the way \\citet{Springel2008} selected their halos for high-resolution calculations. They selected halos without massive neighbors at $z=0$, and they also put a criterion that in the semi-analytic calculation the halo should host a late-type galaxy. This would mean the selected halos should not experience major merger events at low-$z$, and probably implies that their formation epoch is early. Thus, it is natural that their halos have systematically smaller numbers of subhalos than ours have. They found systematic difference between their result and Via Lactea \\citep{Diemand2007,Diemand2008} results. However, our result agrees with both.  \\subsection{Is the missing-dwarf problem solved?}  The difference with the observed number of dwarf galaxies in Local Group or our galaxy is around a factor of 2 for the most subhalo-poor ones in our simulation, at the rotation velocity 10\\% of that of the parent halo. For smaller velocities, the difference is larger, but recent SDSS results indicated that there are still many dwarf galaxies not identified. Thus, the missing dwarf problem does not seem to be a very serious problem. A factor-of-two difference may be easily explained by the effect of baryons. Also, \\citet{Springel2008} found the subhalo abundance even lower than the lowest number we found. Thus, if our galaxy corresponds to most subhalo-poor halos formed in $\\Lambda$CDM universe, there is little discrepancy between the observed number of dwarf galaxies and the expected number of dark matter subhalos.  Obvious question is why our galaxy belongs to most subhalo-poor halos, but as we found the halos with high central concentration and early formation epoch have generally smaller number of subhalos. Also, since our galaxy is a spiral galaxy, it is unlikely that it experienced recent major merger event. If we look at the Local Group as a whole, it will experience a major merger in some future, but in this case we should look at two main subhalos, our galaxy and M31, since their structure determine the total number of subhalos in the Local Group.   \\subsection{Summary}  We simulated $1600^3$ particles with mass of $1.0\\times10^6M_{\\odot}$ in a 46.48Mpc cubic box to create an unbiased sample of halos with resolution high enough to determine the abundance of subhalos. We confirmed the large variation of subhalo abundance we found in Paper I, but could not confirm the dependence on environmental parameters such as average density or distance to the nearest massive halo. However, we found clear dependence on the concentration parameter, which means the number of subhalos is determined primarily by the initial condition and formation time.  The ``missing-dwarf problem'' pointed out by \\citet{Moore1999a} turned out to be much less serious, since now the discrepancy is around a factor of 2. A detailed study of baryon evolution of dwarf galaxies might be necessary to close up this remaining discrepancy."
    },
    "0812/0812.2479_arXiv.txt": {
        "abstract": "UV radiation from early astrophysical sources could have a large impact on subsequent star formation in nearby protogalaxies, and in general on the progress of cosmological reionization.  Theoretical arguments based on the absence of metals in the early Universe suggest that the first stars were likely massive, bright, yet short-lived, with lifetimes of a few million years.  Here we study the radiative feedback arising from such stars using hydrodynamical simulations with transient UV backgrounds (UVBs) and persistent Lyman-Werner backgrounds (LWBs) of varying intensity.  We extend our prior work in Mesinger et al. (2006), by studying a more typical region whose proto-galaxies form at lower redshifts, $z\\sim$ 13--20, in the epoch likely preceding the bulk of reionization. We confirm our previous results that feedback in the relic HII regions resulting from such transient radiation, is itself transient.  Feedback effects dwindle away after $\\sim30\\%$ of the Hubble time, and the same critical specific intensity of $J_{\\rm UV} \\sim 0.1 \\times 10^{-21}{\\rm ergs~s^{-1}~cm^{-2}~Hz^{-1}~sr^{-1}}$ separates positive and negative feedback regimes. This suggests that overall feedback is fairly  insensitive to the large-scale environment, overdensity, and redshift-dependent halo parameters, and can accurately be modeled in this regime with just the intensity of the impinging UVB. Additionally, we discover a second episode of eventual positive feedback in halos which have not yet collapsed when their progenitor regions were exposed to the transient UVB.  When exposed to the transient UVB, this gas suffers relatively little density depletion but a significant enhancement of the molecular hydrogen abundance, thus resulting in net positive feedback.  This eventual positive feedback appears in all runs, regardless of the strength of the UVB. However, this feedback regime is very sensitive to the presence of Lyman-Werner radiation, and notable effects disappear under fairly modest background intensities of $\\JLW \\gsim 10^{-3} \\times 10^{-21}{\\rm ergs~s^{-1}~cm^{-2}~Hz^{-1}~sr^{-1}}$, assuming the region is optically thin for LW photons. Nevertheless, when exposed to the same LWB, halos inside relic HII regions always have a higher H$_2$ abundance and shorter cooling times than halos outside relic HII regions, allowing gas to cool faster once it finally begins to collapse onto the halo. We conclude that UV radiative feedback in relic HII regions, although a complicated process, seems unlikely to have a major impact on the progress of cosmological reionization, provided that present estimates of the lifetime and luminosity of a Pop III star are accurate.  More likely is that the build-up of the LWB ultimately governs the feedback strength until a persistent UV background can be established. ",
        "introduction": "\\label{sec:intro}  Semi-analytic models and numerical simulations predict that the first astrophysical objects formed at redshifts $z\\gsim$ 20.   They form out of metal-free gas with inefficient fragmentation, thus numerical simulations predict that these first objects could be very massive ($M_\\ast\\gsim 100 \\Msun$), so-called 'Population III' (Pop III) stars \\citep{ABN02, BCL02, YOH08}.  These stars would be short-lived, with lifetimes of $\\sim 3$ Myr, but could produce ten times more ionizing photons per baryon than regular Population II stars \\citep{Schaerer02, Schaerer03}.  Also, such star formation could have been somewhat synchronized, given that these early objects formed in highly biased and clustered environments.  Hence, Pop III stars could have a significant impact on subsequent generations of objects.  Indeed Pop III stars can strongly affect the progress of cosmological reionization \\citep{HH03, Cen03a, Cen03b, WL03_postWMAP}, although their contribution is virtually unknown from first principles and is generally applied to simulations with a ``tuning knob\" approach.  Thus ionizing photons from the first stars and lack thereof have been used to explain both early (e.g. \\citealt{Cen03b}) and fairly late (e.g. \\citealt{HB06}) reionization, depending on which scenario was favored when the works were published.  Therefore, it is quite important to gain physical insight into how the radiation of the first stars impacted their surroundings.  Radiative feedback can be either positive or negative, in that it can enhance or suppress subsequent star--formation.  Positive feedback can result when the enhanced free-electron fraction from ionizing photons (e.g. \\citealt{OH02}) or hydrodynamical shocks \\citep{SK87} catalyzes the formation of molecular hydrogen (H$_2$). If the ionizing background ``turns-off\" (resulting in so-called ``relic\" HII regions), H$_2$ cooling can provide the dominant cooling channel at high-densities and low temperatures. Conversely, negative feedback can result from heating by ionizing radiation which can photo-evaporate gas in low-mass halos \\citep{Efstathiou92, BL99, Gnedin00filter, SIR04, Dijkstra04, MD08}.  Also, an active background of Lyman-Werner (LW) radiation (with photon energies in the 11.18--13.6eV range) can dissociate H$_2$, thus decreasing the gas's cooling capabilities (e.g., \\citealt{HRL97, HAR00, CFA00, MBA01, WA07, Oshea08}).  Indeed, simulations find radiative feedback to be nuanced at very high redshifts.  Positive feedback dominates in flash ionized gas \\citep{OShea05}, or gas exposed to a weak transient ultraviolet background (UVB) (\\citealt{MBH06}, hereafter MBH06) such as might be present close to the edges of HII regions \\citep{RGS02b, KM05}.  On the other hand, negative feedback in relic HII regions can occur in regions closer to the Pop III star (MBH06; \\citealt{SU06, AS07, Yoshida07}).  Furthermore, radiative transfer effects can also impact the strength and sign of feedback (e.g. \\citealt{ISR05, SU06, AWB07, Whalen08, WA08}).  Because of variations in halo mass, collapse redshift, central baryon density, stellar spectrum, and distance to the ionizing source, it is difficult to accurately deal with the large parameter space governing photoevaporation.  In MBH06, we attempted to tackle this issue with a statistical approach. By giving up on simulating ``realistic\" (as much as this is possible given present uncertainties) relic HII regions left behind by deceased Pop III stars, we were able to explore larger swaths of parameter space and have a fairly large sample of second-generation objects.  We did this by applying various uniform ultraviolet (UV) and LW backgrounds on our entire simulation box, and statistically comparing the resulting evolution against the fiducial run without radiation.  Although this approach clearly does not realistically treat the photoionization around a particular source, it might provide a good statistical description of a ensemble of reionization studies.  Clearly, at the onset of early reionization not all halos lie within HII regions, as they do in our models. Our primary purpose is not to simulate early reionization or the photoevaporation of well-formed Pop III halos in detail (which would require radiative transfer) but to evaluate the impact of relic HII regions on the formation of halos at later times. Our uniformly illuminated boxes provide an ensemble of evaporated halos at high redshift that statistically sample the effects of early photoionization on the assembly of halos and formation of cold dense gas at lower redshifts.  However, since our analysis in MBH06 focused on a highly biased region which went non-linear at a high redshift, we could not follow the evolution to redshifts much lower than $z\\sim20$. This meant that we could only extrapolate the observed trends suggesting that the feedback was transient.  \\citet{OH03} claim that, even in the presence of a weak LWB, excess entropy will eventually suppress star formation in protogalaxies forming inside relic HII regions. In MBH06, we were unable to conclusively verify or refute this claim.  In this work, we extend the study by MBH06 by analyzing radiative feedback in a less-biased region, whose typical proto-galaxies form at lower redshifts, $z\\sim$ 13--20.  Given that the 5-yr {\\it WMAP} polarization data place a constraint on the reionization redshift (assuming instantaneous reionization) of $z=11.0 \\pm 1.4$ \\citep{Dunkley08}, such cosmological regions are likely to host the majority of feedback effects from Pop III stars, before a persistent UVB and/or a strong LWB become entrenched (e.g. \\citealt{HRL96}).  We apply the same statistical approach as in MBH06, and study various combinations of transient UVBs and persistent LWBs.  As such, we attempt to provide a framework for analytically incorporating such radiative feedback into semi-analytic, and large-scale numerical and semi-numerical studies.  In \\S \\ref{sec:sims} we enumerate and describe the simulations which are used in this work. In \\S \\ref{sec:trans}, we study the initial radiative feedback caused by a transient UVB, adding a persistent LWB in \\S \\ref{sec:LW}.  Then in \\S \\ref{sec:pos}, we describe a related mechanism that results in eventual positive feedback, regardless of the strength of the UVB. In \\S \\ref{sec:ss}, we discuss the impact of neglecting self-shielding.  Finally, we offer our conclusions in \\S \\ref{sec:conc}.  Throughout this paper, we adopt the background cosmological parameters ($\\Omega_\\Lambda$, $\\Omega_{\\rm M}$, $\\Omega_b$, n, $\\sigma_8$, $H_0$) = (0.7, 0.3, 0.047, 1, 0.92, 70 km s$^{-1}$ Mpc$^{-1}$), consistent with the measurements of the power spectrum of CMB temperature anisotropies by the first year of data from the {\\it WMAP} satellite \\citep{Spergel03}.  Although the 5-yr data prefers a slightly lower value of $\\sigma_8=0.82$ \\citep{Dunkley08, Komatsu08} which somewhat delays structure formation, we keep the cosmology the same as in MBH06, to facilitate direct comparison.  Unless stated otherwise, we quote all quantities in comoving units. ",
        "conclusions": "\\label{sec:conc}  UV radiation from early astrophysical sources could have a large impact on subsequent star formation in nearby protogalaxies, and in general on the progress of cosmological reionization.  Theoretical arguments based on the absence of metals in the early Universe suggest that the first stars were likely massive, bright, yet short-lived, with lifetimes of a few million years.  Here we study the impact of the transient radiation arising from such stars on early protogalaxies.  We apply the same statistical approach as in MBH06, studying various combinations of transient UVBs and persistent LWBs, using the hydrodynamical simulation code, Enzo.  We also study a more typical and relevant region whose proto-galaxies form at lower redshifts, $z\\sim$ 13--20. This allows us to trace feedback effects on longer time-scales and lower redshifts.  We confirm the results of MBH06 that feedback in the relic HII regions resulting from such transient radiation, is itself transient.  Feedback effects dwindle away after $\\sim30\\%$ of the Hubble time, and the same critical specific intensity of $J_{\\rm UV} \\sim 0.1 \\times 10^{-21}{\\rm ergs~s^{-1}~cm^{-2}~Hz^{-1}~sr^{-1}}$ separates positive and negative feedback regimes. A weaker UVB stimulates subsequent star formation inside the relic HII regions, by enhancing the molecular hydrogen abundance. A stronger UVB delays star-formation by reducing the gas density at the centers of collapsing halos. The fact that we have confirmed the results in MBH06 (which were mostly inferences obtained by extrapolating from higher redshifts since we were unable to continue those simulations to low enough redshifts) suggests that overall feedback is fairly  insensitive to the large-scale environment, overdensity, and redshift-dependent halo parameters, and can accurately be modeled in this regime with just the intensity of the impinging UVB.  We note that a value of $\\Jlwb\\sim 10^{-3}$--$10^{-2}$ separates feedback regimes dominated by the LW background from those dominated by the transient UVB.  As discussed in section~\\ref{sec:ss}, there is some uncertainty in these values due to our optically thin treatment.  In particular, accounting for LW self-shielding may increase the values of $\\Jlwb$ quoted here.  Additionally, we discover a second episode of eventual positive feedback in halos which have not yet collapsed when their progenitor regions were exposed to the transient UVB. Gas in such late-forming objects did not have time to set-up a steep density profile, which would translate into a strong pressure shock when photo-heated.  Instead they formed out of material which was of more moderate densities at $\\zuvbon$ where the lasting impact of the transient UVB was positive: seeding the gas with an enhanced abundance of molecular hydrogen, which aids in its eventual cooling. However, this feedback regime is very sensitive to the presence of Lyman-Werner radiation, and disappears under fairly modest background intensities of $\\JLW \\gsim 10^{-3} \\times 10^{-21}{\\rm ergs~s^{-1}~cm^{-2}~Hz^{-1}~sr^{-1}}$. Nevertheless, when exposed to the same LWB, halos inside relic HII regions always have a higher H$_2$ abundance and shorter cooling times than halos outside relic HII regions, allowing gas to cool faster once it finally begins to collapse onto the halo. Although it is difficult to accurately estimate the build-up of the LWB, it seems likely that such modest values were surpassed well before the bulk of reionization \\citep{HAR00}.  Thus this persistent eventual positive feedback might only be of academic interest, provided LW self-shielding is not efficient.  We conclude that radiative feedback from a short-lived UVB seems unlikely to have a major impact on the progress of global cosmological reionization, provided that present estimates of the luminosities and lifetimes of Pop III stars are accurate (e.g. \\citealt{Schaerer02}).  Thus it is likely that the LWB plays the dominant role in regulating feedback in relic HII regions.    \\vskip+0.5in  We thank the referee for comments which significantly improved the presentation of this paper.  Support for this work was provided by NASA through Hubble Fellowship grant \\#HF-01222.01 to AM, 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.  GB acknowledges support from NSF grants AST-05-07161, AST-05-47823, and AST-06-06959, as well as computational resources from the National Center for Supercomputing Applications. ZH thanks the Pol\\'{a}nyi Program of the Hungarian National Office of Technology."
    },
    "0812/0812.2962_arXiv.txt": {
        "abstract": "We present the complete galaxy cluster catalog from the Northern Sky Optical Cluster Survey, a new, objectively defined catalog of candidate galaxy clusters at $z\\lsim0.25$ drawn from the Digitized Second Palomar Observatory Sky Survey (DPOSS). The data presented here cover the Southern Galactic Cap, as well as the less-well calibrated regions of the Northern Galactic Cap. In addition, due to improvements in our cluster finder and measurement methods, we provide an updated catalog for the well-calibrated Northern Galactic Cap region previously published in Paper II. The complete survey covers 11,411 square degrees, with over 15,000 candidate clusters. We discuss improved photometric redshifts, richnesses and optical luminosities which are provided for each cluster. A variety of substructure measures are computed for a subset of over 11,000 clusters. We also discuss the derivation of dynamical radii $r_{200}$ and its relation to cluster richness. A number of consistency checks between the three areas of the survey are also presented, demonstrating the homogeneity of the catalog over disjoint sky areas. We perform extensive comparisons to existing optically and X-ray selected cluster catalogs, and derive new X-ray luminosities and temperatures for a subset of our clusters. We find that the optical and X-ray luminosities are well correlated, even using relatively shallow ROSAT All Sky Survey and DPOSS data. This survey provides a good comparison sample to the MaxBCG catalog based on Sloan Digital Sky Survey Data, and complements that survey at low redshifts $0.07<z<0.1$. ",
        "introduction": "The construction of large catalogs of galaxy clusters for use in studies of cosmology, large scale structure, and galaxy evolution has often proven to be a difficult task \\citep[see][for a review]{gal08}. Indeed, the last such catalog generated using optical data, and covering the entire high-galactic-latitude Northern sky was that of \\citet{abe58}, updated in 1989 \\citep{aco89}. More recently, the Northern ROSAT All-Sky (NORAS) Galaxy Cluster Survey \\citep{boh00} has provided an X-ray selected catalog covering a similar region, but with many fewer clusters, while the largest catalog using purely digital observations \\citep{koe07b} covers a smaller area, albeit with much better photometry.  Because much improved optical data has become available, with automated techniques to generate objective, well-characterized cluster samples, we undertook the generation of a modern, optically selected cluster catalog, the Northern Sky Optical Cluster Survey \\citep[NoSOCS][hereafter Papers I and II]{gal00a,gal03}, and its deeper extension \\citep[hereafter Paper IV]{lop04}. The need for a modern cluster survey covering a significant portion of the sky is striking. The catalog from Paper II has already been used to suggest a connection between short-duration gamma-ray bursts and clusters \\citep{geh05,blo06}, to search for giant arcs \\citep{hen08}, to associate compact groups and large-scale structure \\citep{dec05,and06}, and to examine X-ray and optical cluster properties \\citep{lop06}.  Here we present the second and final installment of this catalog, including photometric redshifts, richnesses, optical luminosities and substructure measures. Although superior imaging data is now available from the Sloan Digital Sky Survey (SDSS, \\citealt{yor00}), the imaged area covers less than half of the Northern sky, and is not expected to ever reach the area coverage of our catalog. The currently published cluster catalog from SDSS \\citep{koe07b} covers $\\sim7500$ deg$^2$, containing nearly 14,000 clusters. That survey, using the MaxBCG technique, has a lower redshift cutoff of $z=0.1$; our survey extends down to $z=0.07$. This provides a sample of more local clusters whose properties can be examined in detail \\citep{lop08}, especially using the extensive SDSS spectroscopic database.   The regions covered in this paper are shown in Figure~\\ref{plates}. Three separate areas are covered: (1) The well-calibrated North Galactic Pole (NGP) region already described in Paper II (dotted lines); (2), the portion of the NGP with less well calibrated plates, not covered in Paper II (solid lines); and (3), the Southern Galactic Pole region (SGP, dashed lines). The area covered by the NGP-poor region is 2813$\\Box^{\\circ}$, while the SGP covers 2917$\\Box^{\\circ}$. Together with the 5681$\\Box^{\\circ}$ surveyed in Paper II, the final NoSOCS catalog covers 11411 square degrees. The distribution of clusters in the survey is shown in Figures~\\ref{alleq} and~\\ref{allgal}, in equatorial and galactic coordinates, respectively. Only clusters with richness $N_{gals}>20$ are shown for clarity.  \\begin{figure*} \\plotone{fig1.eps} \\caption{The distribution of DPOSS plates used in NoSOCS, showing the NGP well-calibrated (solid lines), poorly calibrated (dotted lines) and SGP (dashed lines) regions.\\label{plates}} \\end{figure*}   \\begin{figure*} \\plotone{fig2.eps} \\caption{Aitoff projection of the complete NoSOCS cluster catalog in equatorial coordinates.\\label{alleq}} \\end{figure*}  \\begin{figure*} \\plotone{fig3.eps} \\caption{Aitoff projection of the complete NoSOCS cluster catalog in galactic coordinates.\\label{allgal}} \\end{figure*}  The survey methodology is described in \\S2. Although the overall detection technique is very similar to that discussed in Paper II, we have modified our definition of bad areas due to very bright objects, reducing contamination by spurious detections. Changes to the photometric redshift estimation yield more robust (and realistic) errors. The richness estimator has also been improved, providing robust error estimates, as well as total $r$-band optical luminosities. Then, in \\S3, we discuss two new sets of parameters computed for our cluster sample: estimates of cluster substructure and the dynamical radius $r_{200}$.  In \\S4 we describe the general characteristics of our cluster sample, and present consistency tests for the three sky regions utilized. The selection functions describing the completeness as a function of richness and redshift are presented, as is an estimate of the contamination by projection effects. Our complete and final cluster catalog, including an updated version covering the area from Paper II, is presented. This cluster sample is compared to other optically-selected catalogs in \\S5. In particular, we compare our catalog to the SDSS MaxBCG catalog of \\citet{koe07b}, the only modern optical cluster catalog covering a similar area and redshift range. We then examine the correlation between our redshifts and richnesses and the X-ray measurements from the Northern ROSAT All-Sky Galaxy Cluster Survey (NORAS, \\citealt{boh00}). Because NORAS consists of many fewer clusters than our catalog, we use our optical positions to measure X-ray fluxes and upper limits from RASS for a significant subsample of NoSOCS. ",
        "conclusions": "We have presented NoSOCS, a new cluster catalog based on the $|b|>30^{\\circ}$ plate scans from the Digitized Second Palomar Observatory Survey. Spanning over $\\pi$ steradians, this is the largest area optical cluster catalog created since those of \\citet{abe58} and \\citet{aco89}. In terms of area coverage, it will only be superseded by new sky surveys such as Pan-STARRS \\citep{kai04} and LSST \\citep{tys06}, both of which have cosmology through clusters as important science drivers. We show consistency among the three regions covered by NoSOCS, and with the SDSS MaxBCG cluster catalog of \\citet{koe07b}. However, interesting discrepancies between these two large surveys remain. These include large numbers of poor clusters missed by one survey but found in another, suggesting lower completeness, higher contamination, or some combination of the two for such systems, in either or both surveys. Even for supposedly rich clusters there are sufficient discrepancies to call into question our ability to use such surveys for high-precision cosmological constraints. Understanding the sources of these disagreements requires further investigation of both systematic errors and individual cluster candidates.  We have also derived X-ray luminosities for a large subset of our cluster sample from ROSAT all-sky X-ray survey data. We demonstrate that the optical richness and X-ray luminosities are well correlated, albeit with moderate scatter. We find that, despite the poor photometric data and low X-ray luminosities of most NoSOCS clusters, the correlation between $L_{x}$ and $N_{gals}$ is in good agreement with literature results using better data and stacking analyses. Refinements to both the optical richnesses and especially deeper X-ray survey data will be necessary to improve this relation and truly understand the utility of optical richnesses for mass estimation. Nevertheless, our results show promise for using large surveys for such measurements in cases where the data quality is less than superb, as may be expected for the highest redshift clusters even in upcoming deep surveys such as Pan-STARRS and LSST. Furthermore, our ability to measure X-ray luminosities for hundreds of clusters not originally detected in the RASS argues for improved multi-wavelength detection methods that leverage multiple surveys (optical, infrared, X-ray, S-Z) to find distant and/or poor clusters which would otherwise fall below the significance cutoff in a single passband."
    },
    "0812/0812.1150_arXiv.txt": {
        "abstract": "The ZEPLIN-III experiment in the Palmer Underground Laboratory at Boulby uses a 12\\,kg two-phase xenon time projection chamber to search for the weakly interacting massive particles (WIMPs) that may account for the dark matter of our Galaxy. The detector measures both scintillation and ionisation produced by radiation interacting in the liquid to differentiate between the nuclear recoils expected from WIMPs and the electron recoil background signals down to $\\sim$10~keV nuclear recoil energy.  An analysis of 847\\,kg$\\cdot$days of data acquired between February $27^{th}$ 2008 and May $20^{th}$ 2008 has excluded a WIMP-nucleon elastic scattering spin-independent cross-section above $8.1\\times 10^{-8}\\,$pb at 60\\,GeVc$^{-2}$ with a 90\\% confidence limit. It has also demonstrated that the two-phase xenon technique is capable of better discrimination between electron and nuclear recoils at low-energy than previously achieved by other xenon-based experiments. ",
        "introduction": "\\subsection{Motivation} Searches for weakly interacting massive particles (WIMPs) are motivated by the coming together of unification schemes, such as supersymmetry, which predict new particle species, and extensive observational evidence which demonstrates the need for additional non-baryonic gravitational mass within the Universe.  That the WIMPs of supersymmetry naturally fulfill this need is remarkably persuasive.  Indeed, WIMPs occur in other frameworks too. As a generic class of particle they are assumed to only interact non-gravitationally with baryonic matter via the weak interaction. Whilst this offers a mechanism for energy transfer and hence detection, it also implies rather low event rates and energy deposits: $<$0.1~events/day/kg and $<$50~keV respectively. This dictates the use of sensitive underground experiments capable of specifically identifying energy deposits due to elastic scattering of incoming particles from target nuclei.  ZEPLIN-III is the latest in a progressive series of instruments designed to push steadily the sensitivity limits by exploring alternative approaches using xenon-based targets~\\cite{alner05,alner07}.  \\subsection{\\label{1B}ZEPLIN-III} ZEPLIN-III is a two-phase (liquid/gas) xenon time-projection chamber specifically designed to search for dark matter WIMPs. Its design and performance details have already been presented elsewhere~\\cite{araujo06, akimov07} and only a brief reminder is given here. The experiment is operating 1100~m underground.  The active volume is a disc of 35~mm thickness and $\\sim$190~mm diameter which contains $\\sim$12~kg of liquid xenon  above an array of 31 2-inch diameter photomultipliers (PMTs). The PMTs employed during this first science run were ETL D730/9829Q~\\cite{araujo04}, and they were used to record both the rapid scintillation signal, S1, and a delayed second signal, S2, produced by proportional electroluminescence in the gas phase above the liquid~\\cite{dolg70}. The PMT array was immersed in the liquid viewing upwards.  The electric field in the target volume was defined by a cathode wire grid 36~mm below the liquid surface and an anode plate 4~mm above the surface in the gas phase. These two electrodes alone produce the drift field in the liquid (3.9~kV/cm), the field for extraction of the charge from the surface, and the electroluminescence field in the gas (7.8~kV/cm). A fiducial volume for WIMP searches was defined by using a time window for delays between S1 and S2, which selected a depth slice within the liquid, and by 2-D position reconstruction from the PMT signals to select a radial boundary at 150~mm. The time window was set between 500~ns and 13,000~ns which selected depths between 1.29~mm and 33.43~mm.  These together defined a fiducial volume containing 6.5~kg of xenon. \\\\ The PMT signals were digitised at 2\\,ns sampling over a time segment of $\\pm18\\,\\mu$s either side of the trigger point. Each PMT signal was fed into two 8-bit digitisers (ACQIRIS DC265) with a $\\times10$ gain difference between them provided by fast amplifiers (Phillips Scientific 770), to obtain both high and low sensitivity read-out covering a wide dynamic range.  The PMT array was operated from a common HV supply with attenuators (Phillips Scientific 804) used to normalise their individual gains.  The trigger was created from the shaped sum signal of all the PMTs.  For nuclear recoil interactions the trigger was always caused by an S2 signal for energies up to S1=40~keVee, where keVee is an energy unit referenced to the equivalent S1 signal produced by 122~keV $\\gamma$-rays from $^{57}$Co. The trigger threshold was $\\sim$11 ionisation electrons and this corresponded to $\\sim$0.2~keV for electron recoils (for nuclear recoils see Section~\\ref{ec}).  This S2 threshold was set to avoid excessive triggers from single electron emission events and from electron and nuclear recoils whose primaries would otherwise have been undetectable as they fall below the S1 detection threshold. \\\\ The xenon target was contained within a vessel itself located within a vacuum jacket both made from low-background oxygen-free copper. Cooling was provided by a $40\\,$litre liquid nitrogen reservoir, also made from copper, inside the vacuum jacket. Thermal stability to $<$0.5~$\\rm{^o}$C was achieved over the entire run by controlling the flow of cold nitrogen boil-off gas through the base-flange of the xenon vessel. Pressure stability to 2\\% was maintained.  The ZEPLIN-III detector was completely surrounded by a shield of $30\\,$cm thick polypropylene and $20\\,$cm thick lead, giving $10^5$ attenuation factors for both $\\gamma$-rays and neutrons from the cavern walls.  Dedicated access through the shield was provided for the radioactive calibration source delivery, instrument levelling screws and pipe-work to the external gas purification system.  \\subsection{\\label{1C}Science Data}  WIMP-search data were collected over 83~days of continuous operation in the Boulby Laboratory starting on $27^{th}$ February 2008. An 84\\% live time was achieved during the science run and some $847\\,$kg$\\cdot$days of raw data were collected from the $12\\,$kg target volume.  $^{57}$Co calibration measurements were made every day. Nuclear recoil calibrations were made with an AmBe neutron source at the beginning and end of the 83~day period (5~hrs each). A typical event, from a neutron elastic scattering interaction in the liquid with S1=5~keVee, is shown in Figure~\\ref{event} as recorded through the high-sensitivity sum channel.  A short Compton calibration was performed using a $^{137}$Cs source at the beginning of the run with a much longer run at the end (122~hrs). Ten percent of the science data (every $10^{th}$ file) were used to develop initial data analysis and selection cuts, to establish the level of the electron-recoil background, and to define the boundaries for the WIMP-search box and its acceptance. At first, the remaining 90\\% of the science data were retained unopened to carry out a `blind' analysis, but these data were eventually used for perfecting some data-selection cuts as detailed below, making the final analysis non-blind.\\\\ \\\\ \\begin{figure} \\includegraphics[width=3.5in,clip=]{event_paper_1.eps} \\caption{\\label{event}Segment of the high-sensitivity summed waveform for a neutron elastic scattering event with S1$= 5~$keVee, showing a small primary pulse (S1) preceding a large secondary pulse (S2).  Some PMT after-pulsing and, possibly, single electron emission can be seen following S2.  Note that only excursions $>$3~{\\it rms} on individual channels are added into the summed waveform.  See later text for more detailed discussion of some of these points.} \\end{figure} Pulse-finding algorithms were used to identify signals in the 62 waveforms (independently for each PMT and for high and low sensitivity channels).  These were then categorised as S1 or S2 candidates based on a pulse width parameter (charge mean arrival time, $\\tau$): scintillation pulses are much shorter ($\\tau$$\\lesssim$40~ns) than electroluminescence pulses, with durations corresponding to the drift time across the gas gap ($\\tau$$\\sim$550~ns).  Viable S1 and S2 candidates were then subject to software thresholds ($\\ge$3 channels recording signals above $1/3$~photoelectron (p.e.) for S1 and a minimum area of $\\sim$5~ionisation electrons for S2).  Only events with one S1 and one S2  were considered for further analysis. Of particular note here, $\\chi^2$ goodness of fit indicators within the position reconstruction of both S1 and S2 were used to remove multiple-scatter events, and this was particularly effective for those with one vertex in a `dead' region of the xenon, which would otherwise have been a troublesome background.  Such `dead' regions include the reverse-field volume between the cathode wire and the PMT grid wire \\cite{akimov07} and the thin (0.5~mm) layer of xenon surrounding the PMT bodies.  Double-Compton interactions with at least one vertex in these regions, referred to as `multiple-scintillation single-ionisation' (MSSI) events, fulfil the previous selection criteria since there is no S2 pulse from the dead region and the coincident scintillation pulses are added together in a single S1. Unfortunately, perfecting this selection eventually required use of the full data-set as will be described in more detail below. ",
        "conclusions": "An analysis of 847~kg$\\cdot$days of data from the first science run of ZEPLIN-III has resulted in a signal lower limit consistent with zero, and an upper limit on the spin-independent WIMP-nucleon elastic scattering cross-section of $8.1 \\times 10^{-8}\\,$pb, at 90\\% confidence level.  In reaching this result it was necessary to confront an unexpected mismatch between the nuclear recoil spectrum shown in the AmBe calibration data and the Monte Carlo simulation. A careful and thorough analysis of efficiency factors and threshold effects (including the use of alternative data-sets with different thresholds, systematic changes to software cuts and thresholds, visual scanning and manual analysis of large samples of data and modelling and direct verification of the performance of the DAQ) did not resolve this mismatch.  As a more credible alternative explanation the possibility of a non-linearity in the nuclear recoil energy scale has been studied. Non-linearity as such is not unexpected and, indeed it would be surprising if it did not exist at low energy, and a similar approach has been used by others for xenon~\\cite{sore08}. Using this analysis it has been possible to reconcile the data with a non-linearity setting in at the same energy as in \\cite{sore08} but with a more significant effect at lower energies. In itself this may not be surprising given the very different operating conditions within ZEPLIN-III and XENON10: the most obvious being that the electric field in the liquid is 6 times stronger in the former. Indeed, there are other clear differences in the performances of the two instruments. However, it is clear that the physics underlying the low-energy performance is poorly understood. This is true of both the response to electron recoils~\\cite{shutt07} and to nuclear recoils~\\cite{sore08}. As a point of reference, if the mismatch between the AmBe simulation and the data were interpreted solely as an instrument efficiency, the effect on the upper limit would not have been dramatic ($<$40\\% increase) as this approach has a better effective threshold for nuclear recoils but a poorer efficiency.  The analysis presented is not blind as one of the analysis routines was changed after opening of the full data-set as was the limit setting procedure.  In applying the limit analysis no use was made of any background estimates (neither electron-recoil or neutron scattering) and this was done deliberately to avoid underestimating the upper limit."
    },
    "0812/0812.3649_arXiv.txt": {
        "abstract": "One of the goals of the AIRFLY (AIR FLuorescence Yield) experiment is to measure the absolute fluorescence yield induced by electrons in air to better than 10\\% precision.  We introduce a new technique for measurement of the absolute fluorescence yield of the \\unit[337]{nm} line that has the advantage of reducing the systematic uncertainty due to the detector calibration. The principle is to compare the measured fluorescence yield to a well known process -- the \\v Cerenkov emission. Preliminary measurements taken in the BFT (Beam Test Facility) in Frascati, Italy with 350 MeV electrons are presented. Beam tests in the Argonne Wakefield Accelerator at the Argonne National Laboratory, USA with 14 MeV electrons have also shown that this technique can be applied at lower energies. ",
        "introduction": "\\label{Sec:Intro}  The detection of ultra high energy ($\\gtrsim 10^{18}$eV) cosmic rays using nitrogen fluorescence emission induced by extensive air showers (EAS) is a well established technique \\cite{flyseye}. Atmospheric nitrogen molecules, excited by EAS charged particles (mainly $e^{\\pm}$), emit fluorescence light in the $\\approx$ 300--400 nm range. The fluorescence detection of UHECR is based on the assumption that the number of fluorescence photons of wavelength $\\lambda$ emitted at a given stage of a cosmic ray shower development, {\\it i.e.} at a given altitude $h$ in the atmosphere, is proportional to the energy $E_{dep}^{shower}(h)$ deposited by the shower particles in the air volume. Since a typical cosmic ray shower extends up to about 15 km altitude, the fluorescence yield must be known over a wide range of air pressure and temperature. Measurements by AIRFLY of the fluorescence yield dependence on atmospheric parameters (pressure, temperature and humidity), together with the spectral distribution between 280 nm and 430 nm are presented in two separate contributions \\cite{Andreas,Hal}.  It should be noted that  $E_{dep}^{shower}(h)$ is the sum of the energies deposited by EAS particles with a spectrum spanning from keV to GeV.  It is thus important to verify the proportionality of the fluorescence emission to the energy deposit over a wide range of electron energies. In \\cite{EnergyScan}, the proportionality of the fluorescence light to the energy deposit at a few \\% level was tested over the energy ranges 0.5 to 15 MeV, 50 to 420 MeV and 6 to 30 keV. However, only relative measurements within each range were performed, and absolute measurements of the fluorescence yield are in principle needed to verify that the proportionality constant is the same in the three measured energy ranges.  The absolute fluorescence yield is currently one of the main systematics on the cosmic ray energy determination by EAS experiments which employ the fluorescence technique. It is only known at the level of \\unit[15]\\% and for a few electron energies \\cite{others}. In this work, we will report preliminary results of the measurements of the absolute fluorescence yield of the most prominent line -- 2P(0,0) 337~nm by a technique intended to keep the systematics below  \\unit[10]\\%.  The data were taken in the BTF (Beam Test Facility) of the INFN Laboratori Nazionali di Frascati, which can deliver 50--800 MeV electrons. Additionally, we have performed several tests at the the Argonne Wakefield Accelerator (AWA), located at the Argonne National Laboratory, which can deliver 3--15 MeV electrons. The results presented here are preliminary and the intention of the authors is to show that the methodology is appropriate to achieve accuracies below the \\unit[10]\\% level.  The paper is organized as follows: in Section \\ref{Sec:BTF} the technique proposed to measure the absolute fluorescence yield is presented and applied to the measurements in the BTF; in Section \\ref{Sec:Argonne} recent measurements in the AWA are presented and the additional systematic uncertainties due to the smaller electron energies discussed; in Section \\ref{Sec:Concl} we conclude and discuss future work.  \\begin{figure*}[hbt] \\begin{equation} \\begin{split} \\underbrace{N_{fl}(337)}\\ &=\\quad \\underline{\\fly} \\quad\\times\\ \\ \\underbrace{G_{fl}}\\ \\times\\ \\underbrace{T\\ \\times\\ \\ Q(337)}\\ \\ \\times\\ \\ \\underbrace{N_{e^-}}\\cr \\overbrace{N_{c}(337)}^{measured}\\ &=\\ \\overbrace{\\mind Y_{c}}^{known}\\ \\times\\ \\overbrace{G_{c}}^{MC}\\ \\times\\ \\ \\overbrace{T\\ \\times\\ Q(337)}^{\\sim cancel}\\ \\times\\overbrace{N_{e^-}}^{relative}\\times\\ \\overbrace{R_{m}}^{measured} \\end{split} \\label{eq:cer/fl} \\end{equation}  \\end{figure*} ",
        "conclusions": ""
    },
    "0812/0812.1827_arXiv.txt": {
        "abstract": "Dark energy and dark matter are only indirectly measured via their gravitational effects. It is possible that there is an exchange of energy within the dark sector, and this offers an interesting alternative approach to the coincidence problem. We consider two broad classes of interacting models where the energy exchange is a linear combination of the dark sector densities. The first class has been previously investigated, but we define new variables and find a new exact solution, which allows for a more direct, transparent and comprehensive analysis. The second class has not been investigated in general form before. We give general conditions on the parameters in both classes to avoid unphysical behavior (such as negative energy densities). ",
        "introduction": "}  Cosmological observations point to the existence of non-baryonic cold matter and of a late-time acceleration of the universe (see, e.g.~\\cite{Dunkley:2008ie,Tegmark:2006az,Percival:2007yw}). If gravity is modelled on cosmological scales by general relativity, and if we assume that the universe is homogeneous and isotropic on these scales, then the late-time acceleration is sourced by a dark energy component, and the universe is dominated by the ``dark sector\". The quest to uncover the true nature of dark matter (DM) and dark energy (DE) is one of the most pressing topics of modern cosmology.  One of the fundamental puzzles within this quest is the ``coincidence problem\": how is it that we seem to live in a time when the densities of DM and DE are of the same order of magnitude, given that they evolve very differently with redshift? An interesting proposal is that interaction between the dark fields could perhaps alleviate the coincidence problem. Various interaction models have been put forward and studied (see, e.g.~\\cite{Wetterich:1994bg,Amendola:1999qq,Billyard:2000bh, Zimdahl:2001ar,Dalal:2001dt,Farrar:2003uw,Chimento:2003iea,Cai:2004dk,Majerotto:2004ji,Nunes:2004wn,Olivares:2005tb,Sadjadi:2006qp,Guo:2007zk,GarciaCompean:2007vh,Setare:2007we,Boehmer:2008av,Quartin:2008px,Quercellini:2008vh,Valiviita:2008iv,He:2008si,Chen:2008ft}).  For a flat FRW universe, the background dynamics after recombination are governed by the equations of energy balance and the Raychaudhuri field equation: \\begin{subequations} \\label{eq:motion} \\begin{eqnarray} \\dot{\\rho}_b &=& - 3H \\rho_b \\, , \\label{eq:motiond} \\\\ \\dot{\\rho}_c &=& - 3H \\rho_c + Q \\, , \\label{eq:motiona} \\\\ \\dot{\\rho}_x &=& - 3(1+w_x) H \\rho_x - Q \\, , \\label{eq:motionb} \\\\ \\dot{H} &=& - 4\\pi G \\left[ \\rho_b + \\rho_c + (1+w_x) \\rho_x \\right] , \\label{eq:motionc} \\end{eqnarray} \\end{subequations} where $H = \\dot{a}/a$ is the Hubble parameter, $\\rho_c$ is the cold DM density, $\\rho_b$ is the baryonic density, $\\rho_x$ is the density of DE and $w_x <0$ is its constant equation of state. The baryons only interact gravitationally with the dark sector, and $Q$ is the rate of energy transfer in the dark sector. The Friedmann constraint equation is \\begin{equation} H^2 = \\frac{8\\pi G}{3} \\left( \\rho_b + \\rho_c + \\rho_x \\right) \\,. \\label{eq:friedmann} \\end{equation} Note that the field equations~(\\ref{eq:motionc}) and (\\ref{eq:friedmann}) are independent of $Q$, because of total energy conservation. A positive $Q > 0$ represents transfer of energy from DE to DM; a negative $Q < 0$ represents transfer of energy from DM to DE.  In this paper, we consider interactions that are linear combinations of the dark sector densities: \\be Q= A_c\\rho_c + A_x \\rho_x\\,. \\label{Q} \\ee Here the rate factors $A_I$ are either proportional to $H$ or constants, leading to two classes of interaction model: \\be \\mbox{ Model I}~ && A_I= 3\\alpha_I H\\,, \\\\ \\mbox{ Model II} && A_I= 3\\Gamma_I \\,, \\ee where $\\alpha_I$ are dimensionless constants and $\\Gamma_I$ are constant transfer rates. Observations impose the general constraint that the interaction should be sub-dominant today, so that \\be |\\alpha_I| \\ll 1\\,, ~~~~ |\\Gamma_I| \\ll H_0\\,. \\label{con} \\ee  Model I has been recently analysed by~\\cite{Quartin:2008px} (and earlier work considered the special cases $\\alpha_c=\\alpha_x$~\\cite{Chimento:2003iea} and $\\alpha_x=0$~\\cite{Billyard:2000bh}). We use an alternative approach, defining new variables to simplify the parameter space, and finding the general exact solution. We are able to recover previous results more directly and simply, and to provide some new insights into the model.  The $H$ in the $Q$-term for Model~I is motivated purely by mathematical simplicity. By contrast, the energy exchange in Model~II is motivated by similar models that have been used in reheating~\\cite{Turner:1983he}, curvaton decay~\\cite{Malik:2002jb} and decay of DM to radiation~\\cite{Cen:2000xv}. As far as we know, Model~II for the dark sector interaction has has not been treated in the general case before. The special case $\\Gamma_x=0$ has been analyzed by~\\cite{Valiviita:2008iv} (and by~\\cite{Boehmer:2008av} in the case where DE is modelled by exponential quintessence).  A summary of the paper is as follows. Some properties of the general case ($Q$ not specified) are presented in Sec.~\\ref{sec:math-backgr}. Model~I is studied in Sec.~\\ref{sec:model-i.-q}. In Sec.~\\ref{sec:model-ii.-q}, we analyze Model~II. Finally, we conclude in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "}  We have made a careful analysis of two simple models of interaction between DM and DE. The first model, Model~I, has an interaction term proportional to the Hubble parameter times a linear combination of the dark sector densities, and is one of the most studied in the literature.  We developed different mathematical techniques to investigate the properties of the model under simple but general enough assumptions. Our results recover those of previous studies, and we found new analytic expressions that clarify the limitations of Model~I.  To begin with, we absorbed the (constant) DE equation of state $w_x$ into the equations of motion, so that the parameter-space is truly 2-dimensional; this very much simplified the study of the allowed values of the interaction parameters (compare our Fig.~\\ref{fig:parameters} with the corresponding 3-dimensional figures in~\\cite{Quartin:2008px}).  The absorption of the equation of state is not a mere mathematical trick. The original equations of motion would have been \\begin{subequations} \\begin{eqnarray} x^\\prime &=& -3 w_x x \\left( 1 - x \\right) - z \\, , \\\\ y^\\prime &=& 3w_x xy + z \\, . \\end{eqnarray} \\end{subequations} We can see that Eqs.~(\\ref{eq:dyna-sys}) are recovered by setting $w_x=-1$ in the above equations. Thus, the cases discussed are in some sense isomorphic to the cosmological constant interacting case.  This fact also shows the degeneracy between interacting models with a constant DE equation of state; if there is a successful model with a cosmological constant, one can find another one with a different value of $w_x$. However, the models are distinguishable in their particular evolution, as the DE equation of state should appear explicitly in the final expression of the DE-DM ratio, see Eq.~(\\ref{eq:simple-ratio}), and other quantities.  Even though our general assumption was $w_x < 0$, it is clear that in practice we have to restrict ourselves to values of the DE equation of state that allow an accelerating expansion at the present time, i.e. $w_{\\rm tot} < -1/3$. It is also possible to allow \\emph{phantom} values $w_x<-1$, although it is not clear whether any physically consistent models exist in this regime. A comparison with observations would then be required to distinguish among the different possibilities, but this is beyond the purposes of the present study.  We presented for the first time the surprisingly simple equation~(\\ref{DEtoDM}) for the DE-DM ratio, and its simple analytic solution~(\\ref{eq:simple-ratio}). We stress that the exact solution~(\\ref{eq:simple-ratio}) allows us to easily and directly uncover the limitations of this interaction model. The exact solution has only one free integration constant, and thus we need a careful choice of initial conditions ensure an evolution that remains close to that of the standard non-interacting cosmology.  Surprisingly restrictive limitations are needed on the interaction parameters, $\\alpha_c$ and $\\alpha_x$. Our simplified study of the parameter space showed that the parameters should take very small values. This is also in agreement with other studies in which one can find a careful comparison of the model with cosmological observations (see for instance~\\cite{Quartin:2008px,delCampo:2008sr,delCampo:2008jx,Olivares:2007rt,Guo:2007zk}). We showed that the model that seems to work better is the simple case $\\alpha_c =0$. This model is almost indistinguishable from the standard $\\Lambda$CDM at early times, but provides a finite DE-DM ratio at late times so that the coincidence problem is alleviated.  Other interesting properties arise from the interacting Model~II of Sec.~\\ref{sec:model-ii.-q}. Even though it can be thought of as a time-dependent version of Model I, the properties of the critical points differ significantly.  First of all, its vanishing interaction variables at early times allow Model~II to have a true matter dominated epoch, which in fact corresponds to a (unstable) critical point of the dynamical system. However, a positive DM interaction $\\Gamma_c > 0$ is necessary in order to keep the DE component positive at early times.  Also, the DE-DM ratio can be finite and positive at late times, but this requires the dark interactions to have opposite signs and to comply with the condition $\\Gamma_x > \\Gamma_c$. It is then apparent that we cannot, in general, have a realization of Model~II in which the dark components are well behaved for all times and the coincidence problem is addressed.  There is though a simple model that may be realistic and corresponds to $\\Gamma_c =0$. This model can satisfy all the constraints and the DE interaction $\\Gamma_x$ can be adjusted so that the Universe can have appropriate matter and DE eras. The only difference is that DE domination is a transient event: the Universe would eventually go back to a DM dominated epoch at late times.  Finally, we note that the problem of negative dark sector densities is not the only problem with Models~I and II: the curvature perturbation also has a non-adiabatic instability on large scales~\\cite{Valiviita:2008iv}. As pointed out in~\\cite{Valiviita:2008iv}, both of these pathologies in the interaction models are related to the way in which we treat DE, i.e. as a fluid with constant $w_x$. If $w_x$ is allowed to vary in the early universe (as happens with scalar field DE), then the pathologies can be avoided."
    },
    "0812/0812.1220_arXiv.txt": {
        "abstract": "We have examined the relationship between rotation and activity in 14 late-type (M6-M7) M dwarfs, using high resolution spectra taken at the W.M. Keck Observatory and flux-calibrated spectra from the Sloan Digital Sky Survey. Most were selected to be inactive at a spectral type where strong H$\\alpha$ emission is quite common. We used the cross-correlation technique to quantify the rotational broadening; six of the stars in our sample have $v$sin$i$$\\geq$3.5 km\\,s$^{-1}$.  Our most significant and perplexing result is that three of these stars do not exhibit H$\\alpha$ emission, despite rotating at velocities where previous work has observed strong levels of magnetic field and stellar activity.  Our results suggest that rotation and activity in late-type M dwarfs may not always be linked, and open several additional possibilities including a rotationally-dependent activity threshold, or a possible dependence on stellar parameters of the Rossby number at which magnetic/activity ``saturation'' takes place in fully convective stars. ",
        "introduction": "Many M dwarfs, which are the most abundant stars in the Milky Way, have strong magnetic dynamos that give rise to chromospheric and coronal heating, producing emission from the X-ray to the radio. Although this magnetic heating (or activity) has been observed for decades, the exact mechanisms that control magnetic activity in M dwarfs are still not well-understood.  In the Sun, magnetic fields are generated at the boundary between the convective and the radiative zones (known as the tachocline), where the differential rotation in the convective zone creates a rotational shear \\citep{Parker93,Ossendrijver03,Thompson03} that allows magnetic fields to be generated, stored and ultimately rise to the surface. These fields drive the heating of the stellar chromosphere and corona, resulting in both flares and lower level quiescent magnetic activity. A common diagnostic of this heating in optical spectra (of M dwarfs) is the H$\\alpha$ emission line.  Rotation in solar-type stars slows with time due to angular momentum loss from magnetized stellar winds; as a result, magnetic activity decreases.  \\citet{Skumanich72} found that both activity (as measured by Ca II emission) and projected rotational velocity decrease over time as a power law (t$^{-0.5}$).  Subsequent studies confirmed the Skumanich results and demonstrated that there is a strong link between age, rotation and activity in Solar-type stars \\citep{Barry88,Soderblom91,Pizzolato03, MH08}.  There is strong evidence that the rotation-activity relation extends from stars more massive than the Sun to smaller dwarfs \\citep{Pizzolato03, MB03, Kiraga07}.  However, at a spectral type of $\\sim$M3 \\citep[0.35 M$_{\\odot}$;][]{NLDS, Chabrier1997}\\footnote{Although mass and spectral type are closely linked in evolved low-mass dwarfs, young, low-mass objects cool rapidly with age and migrate across a range of spectral types while at the same stellar mass; mass and spectral type are only coupled in stars that have settled on the main sequence ($>$ 500 Myr for most late-type M dwarfs).}, stars become fully convective and the tachocline presumably disappears.  This transition marks an important change in the stellar interior that has been thought to affect the production and storage of internal magnetic fields. Despite this change in the stellar interiors, magnetic activity persists in the late-type M dwarfs; the fraction of active stars peaks around a spectral type of M7 before decreasing into the brown dwarf regime \\citep{Hawley96, Gizis2000, W04}.  A few previous studies of H$\\alpha$ emission have uncovered evidence of a possible rotation-activity relation extending past the M3 convective transition and into the brown dwarf regime \\citep{D98, MB03, RB07}.  In addition, recent simulations of magnetic dynamo generation in fully convective stars find that rotation may play a role in magnetic field generation \\citep{Dobler06,Browning08}.  However, the lack of an unbiased sample of high resolution spectra of late-type M dwarfs complicates the situation.  Using over 30000 spectra from the Sloan Digital Sky Survey \\citep[SDSS;][]{DR6}, \\citet{W06,W08} showed that the activity fraction of M dwarfs varies as a function of stellar age (using Galactic height as a proxy for age) and that the H$\\alpha$ activity lifetime for M6-M7.5 stars is 7-8 Gyr.  Nearby samples of late-type M dwarfs are therefore biased towards young populations with high levels of activity; until recently every known M7 dwarf was observed to be magnetically active \\citep{Hawley96, Gizis2000, W04}.  Due to the intrinsic faintness of late-type M dwarfs, previous rotation studies were limited (and therefore biased) toward, bright, nearby, and therefore active M dwarfs; the 21 M6-M7.5 dwarfs that have previous rotation measurements are all active \\citep{D98, MB03, RB07}. The advent of large spectroscopic surveys of M dwarfs (e.g.\\,SDSS) allows for improved sample selection that can minimize observational bias in the study of late-type M dwarfs.  In this paper we present results from our study of the $v$sin$i$ rotation velocities for a small sample of M6-M7 dwarfs, most of which were selected to be inactive or weakly active from the SDSS low-mass star spectroscopic sample.  We describe the sample and observations in \\S2, our analysis in \\S3, and the observational results in \\S4.  We then discuss our results in \\S5. ",
        "conclusions": "We conducted high resolution spectral observations of 14 M6-M7 dwarfs and found $v$sin$i$ rotation velocities for 6 for the stars.  Three of the stars showed both activity and rotation, 6 of the stars showed neither rotation nor activity, 2 of the stars showed activity but no rotation and 3 stars showed rotation but no activity.  These results are in contrast with previous studies that found a strong connection between rotation and activity in all (active) M6-M7 dwarfs \\citep{D98,MB03, Reiners07}.  As can be seen in Figure \\ref{figure:vsini}, those studies produced one out of 20 stars which was active but not rotationally broadened (very reasonably explained by the inclination effect; there is a 94\\% chance of selecting one out of 20 with the inclination needed to take it from 4 km\\,s$^{-1}$ to 2 km\\,s$^{-1}$), and one which showed rotation and low (but not zero) activity.  Our sample is the first rotation study to include by design M6-M7 dwarfs that are inactive.  Whereas previous studies sampled a range of dynamic populations and therefore stellar ages \\citep{D98, MB03}, our sample was selected from a more distant, mostly inactive and potentially older population \\citep[M7 dwarfs have activity lifetimes of $\\sim$8 Gyr;][]{W08}. Although our sample is not representative of the entire M6-M7 population, it does examine a different part of parameter space than previously sampled and observes something not previously seen: rotation in inactive M7 dwarfs.  Currently, there is no clear explanation for the rotation-activity patterns seen in our sample. However, we present a number of possible explanations that may account for this unusual dataset.  One possible explanation is that rotation and activity are not strongly linked in the fully convective envelopes of M6-M7 dwarfs. Previous studies only examined active M6-M7 dwarfs, many of which have activity levels with no notable correlation with rotation beyond a general tendancy for rotators to lie in a broadly spread ``saturated'' activity regime.  In fact, several active dwarfs have no detectable rotation (see Figure \\ref{figure:vsini}).  Besides saturation, another explanation for the lack of a strong relationship between activity and rotation is inclination; the rotation axes of the stars may be inclined and therefore have reduced line-of-sight components to their rotation velocities.  On the other hand, the 3 inactive stars with detected rotations certainly cannot be resolved by simply imposing inclination effects, which can only serve to make their true rotations even higher.  Another possibility is that there is a rotationally-induced threshold for activity at V$_{rot}>$6 km\\,s$^{-1}$ \\citep{Hawley96}. This would explain the inactive rotators (they are below the threshold) and require that active stars with $v$sin$i$ rotation velocities less than the threshold have inclined rotation axes. Combining all of the previous rotation studies for M6-M7 (Figure \\ref{figure:vsini}), we see that 7 of the 26 active dwarfs have $v$sin$i\\leq$ 6 km\\,s$^{-1}$; the rest of the stars are rotating several km\\,s$^{-1}$ above the detection thresholds. To quantify this possibility, we follow the statistical analysis of \\citet{BBMW08}, where we compute the probability (from the binomial distribution) of observing n objects with inclination $i$ $<$ $i_{crit}$.  We find that the probability of randomly drawing a sample of 7 stars with the inclinations necessary to take the slow but active stars from a V$_{rot}$=6.5 km\\,s$^{-1}$ to the observed $v$sin$i$ is 85.7\\%.  If we instead calculate the more conservative probability of drawing inclinations that take the 7 stars from 8 km\\,s$^{-1}$ (the median $v$sin$i$) to their respective $v$sin$i$ values, we obtain 55.9\\%.  In this scenario, the inactive stars, which are presumably older, have spun down enough to fall below the rotation-activity threshold.  In earlier M dwarfs, the mere detection of rotation implies that the star is in the ``saturation'' regime; an effect that should have become even more pronounced for these even smaller objects (assuming the convective overturn time is similar).  All of our detected rotators have estimated Rossby numbers that are beyond the ``saturation'' threshold for early-type M dwarfs \\citep{RBB08}.  One possibility is that in these fully convective late-type M dwarfs, the Rossby number ``saturation'' threshold might be substantially smaller. However, this model would imply that stars go from fully active to completely inactive over a small range of Rossby numbers - something that contradicts what is seen in early-type M dwarfs, where the activity falls smoothly with increasing Rossby number \\citep{Kiraga07, RBB08}. It is clear that the velocity thresholds proposed here would have to be lowered (as a function of spectral type) to be compatible with the results of \\citet{RBB08}; too many of their stars are active below $v$sin$i=$ 6 km\\,s$^{-1}$ (including several in the M5-6 spectral range).   Of course one additional possibility is that M6-M7 dwarfs might have activity cycles and that we are observing the inactive stars at a minimum in their cycles.  Although this might be observed for a small number of stars, it is unlikely that an activity minimum can explain the majority of inactive M dwarfs seen in the Milky Way.  \\citet{W08} showed that the activity fractions of M dwarfs are strongly correlated with the height above the Galactic plane (and therefore age), a trend that cannont be explained by an activity cycle.  The dearth of a large sample of nearby, inactive late-type dwarfs precludes large-scale follow-up that would further test the age-rotation-activity relation of fully convective low-mass dwarfs. However, future efforts to extend our analysis to both inactive M5 and M8 dwarfs will help elucidate the problem.  The multi-epoch photometric observations produced by all sky time-resolved surveys (e.g. Pan-STARRS and LSST) will produce rotation periods for a large number of late-type M dwarfs that will eliminate the inclination effects.  Of course, inactive M dwarfs present a problem to this method in that they may not exhibit photometric variability (if they don't produce starspots).  The combination of these and other future observations will help advance the understanding of what role rotation plays in the production of magnetic fields in fully convective objects, and provide crucial constraints for theoretical models of fully convective dynamos.  They will also be a key component to understanding the feedback between magnetic fields and stellar angular momentum evolution through stellar activity."
    },
    "0812/0812.3549_arXiv.txt": {
        "abstract": "We have started high precision photometric monitoring observations at the AIU Jena observatory in Gro{\\ss}schwabhausen near Jena in fall 2006. We used a 25\\,cm Cassegrain telescope equipped with a CCD-camera mounted picky-pack on a 90\\,cm telescope. To test the obtainable photometric precision, we observed stars with known transiting planets. We could recover all planetary transits observed by us.\\\\ We observed the parent star of the transiting planet TrES-2 over a longer period in Gro{\\ss}schwabhausen. Between March and November 2007 seven different transits and almost a complete orbital period were analyzed. Overall, in 31 nights of observation 3423 exposures (in total 57.05\\,h of observation) of the TrES-2 parent star were taken. Here, we present our methods and the resulting light curves. Using our observations we could improve the orbital parameters of the system. ",
        "introduction": "The search for extrasolar planets is one of the most important research field in Astronomy today. Up to now the most successful method to detect exoplanet candidates is the radial velocity technique. But in the last ten years another indirect detection method has established itself as a highly successful technique in finding planets and confirming candidates - the transit method. \\\\ The transit event is a strictly periodic phenomenon. In a systen where a known planet transits its host star, a second planet in that system will cause the time between transits to vary. It is becoming increasingly popular because even small ground-based observatories have already obtained the photometric precision necessary to detect sub-Earth mass planets by the transit timing variation method \\cite[(Steffen et al. 2007)]{Steffen_etal07}.\\\\ We have started high precision photometric monitoring observations at the AIU Jena observatory in Gro{\\ss}schwabhausen near Jena in fall 2006. In this work we use the transit method to observe the known transiting planet TrES-2. The aim is to test procedures with our telescope and determine the obtainable photometric precision with the currently existing camera. We paid special attention to the accurate determination of transit times in order to identify precise transit timing variations that would be indicative of perturbations from additional bodies and to refine the orbital parameters of the systems. ",
        "conclusions": "During the observations of the transiting extrasolar planet TrES-2 at our university observatory in Gro{\\ss}schwabhausen with the 25\\,cm Cassegrain telescope equipped with the optical CCD camera CTK we obtained a timing accuracy of $\\sim$2\\,min. The timing residuals  are not consistent with zero within the measurement errors. The second last data point is 3$\\sigma$ deviant. The deviations to both sides of zero could be a first indication of timing anomalies caused by additional planets or moons. We will continue observing TrES-2 to confirm these transit time variations. Therefore, we work on methods to improve the accuracy of our transit times.\\\\ This year we get a new CCD camera for the Schmidt focus of the 90\\,cm reflector. This camera will have a smaller pixel scale and a higher sensitivity. Part of the preparation for the new camera is the improvement of the software for relative photometry. Our transit observations will benefit strongly from the new camera.\\\\ The transit observations with the AIU Jena telescope in Gro{\\ss}schwabhausen provide anchors for future searches for transit time variations."
    },
    "0812/0812.0583_arXiv.txt": {
        "abstract": "This paper is a first step towards developing a formalism to optimally extract dark energy information from number counts using multiple cluster observation techniques. We use a Fisher matrix analysis to study the improvements in the joint dark energy and cluster mass-observables constraints resulting from combining cluster counts and clustering abundances measured with different techniques. We use our formalism to forecast the constraints in $\\Omega_{\\rm DE}$ and $w$ from combining optical and SZ cluster counting on a 4000 sq. degree patch of sky. We find that this cross-calibration approach yields $\\sim 2$ times better constraints on $\\Omega_{\\rm DE}$ and $w$ compared to simply adding the Fisher matrices of the individually self-calibrated counts. The cross-calibrated constraints are less sensitive to variations in the mass threshold or maximum redshift range. A by-product of our technique is that the correlation between different mass-observables is well constrained without the need of additional  priors on its value. ",
        "introduction": "The evolution of the number of clusters of galaxies provides a powerful tool to study the nature of dark energy. Clusters are sensitive probes of the growth of structure because cluster abundances are exponentially dependent on the linear density perturbation field. In addition, cluster surveys are sensitive to the evolution of the volume element with redshift so that cluster surveys also probe the background cosmology.  Planned and ongoing cluster surveys will detect millions of clusters using a variety of techniques such as counts of optically detected galaxies (e.g. DES \\cite{des05}, LSST, \\cite{tys02}), the Sunyaev-Zel'dovich (SZ) flux decrement (e.g. SPT \\cite{ruh04} and ACT \\cite{kos03}), X-ray temperature and surface brightness (e.g. eRosita, \\cite{pre07}), and weak lensing shear. Because different cluster techniques suffer from different sources of errors, combining the information from different surveys is essential to reduce random errors and control the systematics.  One of the major challenges in extracting dark energy information from clusters is that cluster masses are not directly observable. One must rely on observable proxies for mass which only correlate statistically with the true mass. The inherent uncertainties in the observable-mass relation will degrade cosmological constraints if not well understood. Methods have been developed to use additional cluster properties such as the cluster power spectrum \\citep{maj04}, sample covariance from counts in cells \\citep{lim04}, or the shape of the observed mass function \\citep{lim05,hu03,roz07,lev02} to ``self-calibrate'' the mass-observable relation by simultaneously solving for the cosmological and mass-observable parameters.  Other works have investigated combining different cluster techniques to cross-calibrate the mass-observable relations of each \\citep{maj04,maj03,you06}. In \\cite{maj04,maj03}, the cross-calibration is between an SZ or X-ray survey and a detailed mass follow-up to calibrate the mass-observable relation, whereas \\cite{you06} combine SZ and X-ray surveys. However, these studies have assumed that the two surveys were independent, so that the joint constraints were estimated by adding the Fisher matrices of both experiments. But if two surveys observe the same patch of sky, the measurements are {\\bf not} independent. The goal of this paper is to show how to exploit the interdependence of cluster surveys over the same patch of sky to improve constraints on dark energy and mass-nuisance parameters.  The paper is organized as follows. In \\S \\ref{sec:selfcal} we describe the Fisher matrix formalism to forecast cosmological constraints from cluster counts and clustering using a single and multiple observables. We describe the major cluster mass determination techniques in \\S \\ref{sec:techniques} and explain our parametrization of the errors in the observables, i.e. the mass-observable distributions. Results are presented in \\S \\ref{sec:res} and our conclusions and prospects for future work are given in \\S \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}   We developed a formalism to derive joint cosmological and cluster mass-observable constraints from cluster number counts and clustering sample variance of multiple cluster finding techniques. The improvement we find relative to previous works arises from our use of the interdependence of cluster measurements performed over the same patch of sky to cross-calibrate the mass-observable relations of the different techniques. When combining an SPT-like and DES-like survey, the full cross-calibration method yields $\\sim 2$ times smaller constraints on $\\DE$ and $w$ compared to simply adding the Fisher matrices of the individual experiments. Furthermore, constraints from the full cross-calibration are less sensitive to $\\Mth$ and $z_{\\rm max}$ than the single mass-observable constraints.  The cross-calibration places tight constraints on the correlation between the observables without the need of additional priors. Conversely, priors on the mass-variance and bias can significantly improve the dark energy constraints. Constraints on $\\DE$ are most sensitive to priors on the mass biases. On the other hand, constraints on $w$ are more sensitive to priors on the redshift-dependent part of the scatters. Priors on the optical nuisance parameters are more relevant than priors on SZ nuisance parameters for both $\\DE$ and $w$ constraints.  Our technique can still be improved. Combining more than two techniques at a time should further improve constraints. But we can only combine multiple techniques if we use a more efficient binning strategy, to minimizes the number of mass bins needed to extract the useful information. It is possible that a more efficient binning may improve even the two observable case, particularly in cosmologies with low $\\sigma_8$.   Work still needs to be done before the self-calibration or full cross-calibration can be applied to real data. The cross-calibration estimates presented here are sensitive to the parametrization of the mass errors. Simulations are needed to determine what parametrizations are robust to theoretical and experimental uncertainties. Our results assumed a perfect selection, but selection effects may bias the cosmological constraints. \\cite{wu08} have shown that if the halo selection depends on halo concentration, and if the halo bias depends on the assembly history, the sample variance due to clustering will deviate from that of a random selection of halos with the same mass distribution. If the clustering sample variance is modeled incorrectly, the self-calibration may bias the recovered dark energy parameters. Since the different cluster surveys are expected to have selections with different dependence on the halo concentration, cross-calibration should mitigate selection effects, though we are yet to test this hypothesis. Finally, we must still account for the relation between photo-z and mass-observable errors. Regardless of the simplifications adopted here, we conclude that having overlap between surveys is very important to maximize the effectiveness of cross-calibration techniques."
    },
    "0812/0812.2065_arXiv.txt": {
        "abstract": "We present our analysis of photometric data in the Johnson $B$ and $V$ filter of the southern Blazhko star SS For. In parallel, we analyzed the $V$ observations obtained with the ASAS-3 photometry of the star gathered between 2000 and 2008. In the frequency spectra resulting from a Fourier analysis of our data, the triplet structure is detectable up to high order, both in the $B$ and $V$ data. Moreover, we find evidence for quintuplet components. We confirm from our data that the modulation components decrease less steeply than the harmonics of the main frequency. We derived the variations of the Fourier parameters quantifying the light curve shape over the Blazhko cycle. There is good agreement between the spectroscopic abundance and the metallicity determined from the Fourier parameters of the average light curve. SS For is peculiar as a Blazhko star because of its strong variations around minimum light. ",
        "introduction": "A large fraction of the RR Lyrae stars shows a periodic amplitude and/or phase modulation with a period of typically ten to hundreds of times the pulsation period. This phenomenon is referred to as the Blazhko effect, after the Russian astronomer who first reported it (Blazhko 1907).  The origin of the Blazhko effect is still a matter of controversy.  The most widely discussed models attribute the effect to either the consequences of the interaction of a magnetic field with the main radial pulsation (Shibahashi \\& Takata 1995), or a resonance between the main mode and a non-radial mode of low degree (Van Hoolst, Dziembowski \\& Kawaler 1998; Nowakowski \\& Dziembowski 2001; Dziembowski \\& Mizerski 2004). However, also models that do not require nonradial modes have been proposed, e.g. by Stothers (2006). This scenario attributes the Blazhko variation to a variable turbulent convection due to transient magnetic fields in the star. All models presently proposed for the Blazhko effect have shortcomings in explaining the variety of features shown by Blazhko stars.  Therefore, theoretical efforts to revise or expand the models would be worthwhile, and even the exploration of alternative explanations.  Since its discovery almost a century ago, most studies of the Blazhko effect were carried out from the northern hemisphere (e.g., Szeidl 1988, and references therein; Smith 1995, and references therein). For a long time there were fewer well-established {\\it field} Blazhko stars at southern declinations.  The available data for southern field Blazhko stars (e.g., Hoffmeister 1956; Kinman 1961; Clube et al. 1969; Lub 1977) are often insufficient to determine the Blazhko periods with the required accuracy. Accurate and complete photometric data sets of southern field Blazhko stars are still lacking. Extensive multi-target surveys such as MACHO (Alcock et al. 2000, 2003) and OGLE (Moskalik \\& Poretti 2003), and ASAS (All Sky Automated Survey, Pojmanski 2000) have significantly contributed to our knowledge of the Blazhko effect, and expanded the list of known Blazhko stars. Wils \\& Sodor (2005) published a list of new and confirmed Blazhko targets, and Szczygiel \\& Fabrycky (2007) obtained interesting new results, among which is a star with multiple Blazhko periods.  Nevertheless, long-term campaigns dedicated to particular Blazhko stars and yielding complete light curves at different phases in the Blazhko cycle remain of great scientific value.  Only these can reveal changes in the light curve shape, as well as long-term changes in the characteristics of the Blazhko effect such as amplitudes, phases, and periods. In this way they can give crucial information for deciding among the different hypotheses for the Blazhko effect.  Photometry with a good spread over both the pulsation and Blazhko cycles allows us to perform the first rigorous frequency analysis of a southern field Blazhko star. The first accurate photometric data set covering both the pulsation and the Blazhko cycle was published by Jurcsik et al. (2005a) for the northern short-period Blazhko star RR Gem.  SS For (SAO 167572; $\\alpha(J2000)$: 02$^{\\rm h}$ 07$^{\\rm m}$ 52$^{\\rm s}$, $\\delta(J2000)$: -26$^{\\circ}$ 51' 54'') is a relatively bright southern Blazhko star. According to the General Catalogue of Variable Stars (Kukarkin 2003) its $V$ brightness changes within the range 9.45-10.60. The General Catalogue of Variable Stars (GCVS) lists a period of $P$ = 0.495432 d (Kholopov et al. 1998) or about 11 h 53 min, corresponding to a frequency of $f_0 = 2.0184405$ cd$^{\\rm -1}$.  We selected this target on the basis of references to its variable light curve in the literature (Lub 1977). Our analysis of both the HIPPARCOS and the ASAS-3 photometry demonstrated the existence of a Blazhko effect with strong variations around minimum.  The deduced Blazhko period of SS For (about 35 d) turned out to be short enough to cover several modulation cycles during one observing season. All our data of SS For are available online at http://www.univie.ac.at/tops/blazhko/data/.  In Section\\,2 we present the new observations as well as the ASAS-3 data used for comparison.  The different steps in the data analysis and its results are described in Section\\,3. Section\\,4 is dedicated to a discussion of the results obtained. Finally, some concluding remarks are given in Section\\,5. ",
        "conclusions": "  \\begin{itemize}  \\item  From a Fourier analysis of the combined SAAO+SSO data set we obtain $f_0 = 2.018433 \\pm 0.000001 $ cd$^{\\rm -1}$ for the radial pulsation frequency and $f_N = 2.04705 \\pm 0.00004$ cd$^{\\rm -1}$ for the right-side peak frequency.  \\vspace{2mm} \\item {\\it From our data} we find a value for the Blazhko period $P_B=34.94 \\pm 0.05$ d.  \\vspace{2mm} \\item The amplitudes in the $B$ data are a factor of about 1.2 larger than in the $V$ data. The harmonics of the main frequency as well as the triplet structure around the main frequency and its harmonics are detectable up to high order, especially in the $B$ data (15th order). Also the Blazhko modulation is a factor of 1.2 larger in the $B$ than in the $V$ data.  \\vspace{2mm} \\item After subtracting the fit including triplet components, we find clear evidence of quintuplet frequencies in the data.  We found several clearly significant peaks located at $kf_0 \\pm 2f_B$ frequency values.  These frequencies emerged most clearly from the $B$ data set in which the modulation components have higher amplitudes, but were also retrieved from the $V$ data. Our data thus clearly show the existence of quintuplet components in SS For, as were also found by Hurta et al. (2008) in RV UMa and by Jurcsik et al. (2008) in MW Lyr. Quintuplet components are predicted by the magnetic model for explaining the Blazhko effect (Shibahashi \\& Takata 1995), but so far there has been no convincing, unambiguous positive detection of a magnetic field in RR Lyrae stars.  \\vspace{2mm} \\item We also find evidence for a so-called septuplet component in our data (at $6f_0+3f_B$).  Multiplet components of order $j \\ge 3$ in $kf_0 \\pm jf_B$ have also been found by Jurcsik et al. (2008) in their extensive data set of MW Lyr.  They have yet to be explained by the models for the Blazhko effect. Uninterrupted satellite data of Blazhko stars, covering several Blazhko cycles, such as delivered by the COROT satellite (Baglin 2007), will undoubtedly shed new light upon the existence of higher-order multiplet structures in Blazhko stars. We may expect to find many more in high-quality data sets.    \\vspace{2mm} \\item A subdivision of the data into 10 Blazhko phase intervals shows the variations of the light curve over the Blazhko cycle.  As all the Blazhko phase intervals in $V$ have sufficient coverage, we calculated the Fourier parameters and plotted their variations.  Observed values at different phases of the Blazhko cycle may be useful for confrontation with non-linear convective models such as those developed by Feuchtinger (1999). It would be worthwhile to find out whether the observed variations in the Fourier parameters can be reproduced theoretically for a star with constant mass and metallicity.  \\vspace{2mm} \\item Application of the empirical $P-\\phi_{31}-$[Fe/H] relation developed by Jurcsik \\& Kov\\'acs yields a good result for the average light curve derived from all our data. This strengthens the assumption that the average light curve of Blazhko stars can be employed in the empirical formula calibrated by Jurcsik \\& Kov\\'acs on stricly periodic stars.  However, the average light curve of a Blazhko star can only be determined accurately if the Blazhko cycle is sufficiently covered by the available data.  \\vspace{2mm} \\item From a study of the light curve shape at different Blazhko phases we clearly see that the strong variations around minimum light in SS For are related to the position and strength of the bump. A detailed study of the bump behaviour over the Blazhko cycle, in photometric as well as spectroscopic data, may shed new light upon the understanding of the Blazhko effect.  \\vspace{2mm} Theoretical efforts to revise or expand the existing models for the Blazhko effect would be worthwhile, as well as the exploration of alternative explanations.  \\end{itemize}  \\begin{figure*} \\includegraphics[width=12cm, angle=0]{ssfor_figures/ssfor_phasebins_B.jpg} \\caption{$B$ light data (crosses) in different 0.1 phase intervals of the Blazhko cycle (34.94 $\\pm$ 0.05 d from this data set).  The mean light curve over all Blazhko phase intervals is also shown as a solid line.  Note the strong variations at the bump phase before minimum light. } \\end{figure*}"
    },
    "0812/0812.2586_arXiv.txt": {
        "abstract": "{} {This is the first paper of a series aimed to understand the formation and evolution of bars in early-type spirals and their influence in the evolution of the galaxy.} {Optical long-slit spectra along the major axis of the bar of a sample of 20 galaxies are analyzed. velocity and velocity dispersion profiles along the bar are presented.  Line-strength indices in the bar region are also measured to derive stellar mean-age and metallicity distributions along the bars using stellar population models.} {We obtain mean ages, metallicities and chemical abundances along the bar of 20  galaxies with morphological types from SB0 to SBbc. The main result is that we find large variation in age and metallicity along the bar in $~$45\\% of our sample. We find three different types of bars according to their metallicity and age distribution along the radius:  1) Bars with negative metallicity gradients. They show mean young/intermediate population ($<$~2~Gyr),  and have amongst the lowest stellar maximum central velocity dispersion of the sample. 2) Bars with null metallicity gradients. These galaxies that do not show any gradient in their metallicity distribution along the bar and have negative age gradients (i.e younger populations at the bar end).  3)  Bars with positive metallicity gradients, i.e. more metal rich at the bar ends. These galaxies are predominantly those with higher velocity dispersion and older mean population.  We found no significant correlation between the age and metallicity distribution, and bar/galaxy parameters such as the AGN presence, size  or the bar strength. From the kinematics, we find that all the galaxies show a disk--like central component.} {The results from the metallicity and age gradients indicate that most galaxies with high central stellar velocity dispersion host bars that could have been formed more than 3~Gyrs ago, while galaxies with lower central velocity dispersions show a wider distribution in their population and age gradients.  A few bars show characteristics compatible with having been formed less than $<$2~Gy ago. However, we do not have a definite answer to explain the observed gradients and these results place strong constrains to models of bar formation and evolution. The distribution of mean stellar population parameters in the bar with respect to $\\sigma$ is similar to that found in bulges, indicating a close link in the evolution of both components. The disk-like central components also show the important role played by bars in the secular evolution of the central structure.} ",
        "introduction": "Bars are known to be efficient mechanisms to redistribute angular momentum and matter in the galaxies up to large distances and, therefore, they are expected to play an important role in the secular evolution of disk galaxies. The bar gravity torques make the gas lose angular momentum which, in turn, provokes an inflow of gas towards the central parts. This inflow of gas accumulates in central mass concentrations (CMC), fueling the central black hole (e.g Sholsman et al. 1989) and, possibly, forming a stellar bulge. Bars are found in $\\sim$60\\% of spiral galaxies in the local Universe (e.g. Eskridge et al.\\ 2000; Whyte et al.\\ 2002; Men\\'endez-Delmestre et al. 2007; Marinova \\& Jogee 2007)\\nocite{menendezdelmestre2007,marinova2007}\\nocite{eskridge2000,whyte2002}, or 30\\% if only strong bars are considered. Despite their obvious importance in disk galaxy evolution, their own fate is still a matter of debate. Several mechanisms are able to destroy the bar: for example; 1) the bar may dissolve in the presence of a sufficiently massive central component (e.g. a black hole) (Friedli \\& Pfenniger 1991; Friedli \\& Benz 1993, Norman, Sellwood \\& Hasan 1996; Berentzen et al. 1998; Fukuda et al.\\ 2000; Bornaud \\& Combes 2002; Shen \\& Sellwood 2004; Athanassoula, Lambert \\& Dehnen 2005; Bornaud, Combes \\& Semelin 2005)\\nocite{shen,friedli91,friedli93,norman96,berentzen98}. However, the masses of the CMC needed to destroy the bar are much larger than those of the most super massive black holes found so far in disk galaxies (Shen \\& Selwood 2004; Athanassoula, Lambert \\& Dehnen 2005; Ferrarese \\& Ford 2005)\\nocite{ferrarese}. 2) Transfer of angular momentum from infalling gas to the bar has also been proposed as a mechanism that could strongly weaken the bar. The combined effect of this angular momentum transfer and the CMC can destroy the bar in Sb-Sc galaxies in 1-2 Gyr  (Bournaud \\& Combes 2002)\\nocite{bournaud}. Then, if gas is present, a new bar could form with characteristics different from those of the parent bar. However, simulations of the effect of gas on the stellar evolution of a bar embedded in a live halo (Berentzen et al. 2007)\\nocite{berentzen} have shown that, althogh there are some structural  differences in bars evolving with and without gas, in both cases the stellar bar can survives, at least, 5 Gyr (the whole computation time of these simulations  Do all bars go trough these processes? Do some bars die while others are robust over many Hubble times? Previous studies have tried to put some limits to the age of the bars by analyzing the evolution with redshift of the bar's fraction in galaxies. These studies have also  obtained  discrepant results: Some authors have found a fraction of barred galaxies at z$>$0.5 lower than the local fraction (Abraham et al.\\ 1999)\\nocite{abraham}, although some authors claim that this may be the consequence of  selection effects, due to the high angular resolution needed to find bars (Elmegreen, Elmegreen \\& Hirst.\\ 2004; although see van den Bergh et~al.\\ 2002)\\nocite{elmegreen2004,vandenbergh}. In order to avoid the problem of the angular resolution, several studies have carried out this analysis using the ACS camera on the HST: Elmegreen et al. (2004) and Jogee et al.\\ (2004)\\nocite{jogee2004} found the same bar fraction ($\\sim$0.4) at redshift z=1.1 than in the local Universe, suggesting that the bar dissolution cannot be common during a Hubble time unless of course, the bar formation rate is comparable to the bar destruction rate. On the contrary, Sheth et al.\\ (2008)\\nocite{sheth2008}, in a recent study, using images from the COSMOS survey (Scoville et al.\\ 2007)\\nocite{scoville2007} and using a larger sample than previous studies, found that the bar fraction at z=0.84 is one-third of the present-day value. They also found a much stronger evolution for low mass galaxies and late-type spectral types. Part of the discrepant results may be due to the selection effects and other systematic effects that need to be still further investigated.  Detailed analysis of the stellar populations in the bar region of local galaxies can shed some light to the formation and evolution of bars. Gadotti \\& de Souza (2006)\\nocite{gadotti} obtained the color and color gradients in the bar region of a sample of 18 barred galaxies. They interpreted the color differences as differences in stellar ages. They concluded that younger bars were hosted by galaxies of later types.  However, the conclusions are hampered by the assumption that the color traces the age, and that the metallicity effect or dust extinction are not important. Clearly, a study attempting to break the age--metallicity degeneracy and taking into account the effect of dust is needed. In addition, the influence of dust in the line-strenght indices have been shown to be minimum (MacArthur 2005)\\nocite{macarthur} while ages and metallicities from broad--band colors are heavily affected by dust as well as by the age--metallicity degeneracy.  To date, it has been difficult to obtain stellar population parameters along the bar because a high signal-to-noise is required for this analysis. Although bars are high surface brightness structures, this still implies long integration times on medium-size telescopes. Recent work has presented the radial distribution of line-strength indices along bars (P\\'erez, S\\'anchez-Bl\\'azquez \\& Zurita 2007) \\nocite{perez07} for 6 galaxies.  The theoretical model of abundance gradients in bars developed by Friedli (1994) \\nocite{friedli94} predicted, using N-body simulations of bars with pre-existing exponential abundances, a null evolution of the stellar abundance profile along the bar (although with a decrease in the mean metallicity) while the gas abundance profile flattened rapidly.  Studies of the radial gas-phase abundance distribution have shown that there is, indeed, a flattening in the gas abundance  along the bar (Martin \\& Friedli 1997, 1999)\\nocite{martin97,martin99}.  However, recent work (P\\'erez, S\\'anchez-Bl\\'azquez \\& Zurita 2007) has shown, from stellar absorption line-strengths,  that  some bars present stellar metallicity gradients opposite to those found in the disks (i.e. more metal rich population at the end of the bar than in the regions of the bar closer to the center) , contradicting the numerical results. This surprising result needed a more systematic study of the stellar population along the bars, properly comparing the line-strength indices of a larger sample of galaxies with the state-of-the art stellar populaton models to derive and quantify the variations of age and metallicity along the bar.  In this paper, we present a  detailed stellar population analysis of the bar region of  a sample of 20 early-type spirals (from S0 to Sbc) with and without nuclear activity in order to derive ages, metallicities and relative chemical abundances ratios.  The sample selection is described in Sect.~\\ref{sample}. The observations and data reduction are presented in Sect.~\\ref{observation}. The velocity dispersion measurements and the emission correction applied to the data are explained in Section~\\ref{velocity}. The line-strength derivation is described in Sect.~\\ref{line}.  The morphological characterization of the sample is described in Sect.~\\ref{morphology}. Section~\\ref{deri.para} gives a description of the methodology followed to derived the stellar population parameters.The results of the line strength-indices and the kinematics are presented in Sect.~\\ref{results}. In this work, we concentrate only in the stellar populations of the bar region. The results of the central indices, bulge and bulge gradients will be presented in the second series of this paper. ",
        "conclusions": "}   We have derived the kinematics and the distribution of ages and metallicities along the bar in a sample of 20 barred galaxies to study the question whether bars are long--lived or weather they are destroyed and reformed in a few Gyrs. We also want to explore the relation between the age, and stellar populations of the bar and other global physical properties of the host galaxies.  From the study of the kinematics along the bar, we have found that all the galaxies in the sample show a disk--like component in their centers, showing as dips or plateaus in their central stellar velocity dispersion and as stellar rotating disks. Therefore, this implies that bars, as expected, have a strong influence in the building up and later evolution of the central component. This aspect will be studied in more detail in S\\'anchez-Bl\\'azquez, P\\'erez and Zurita (2009, in prep) where we compare the bulges ages and metallicities of barred and unbarred bulges of otherwise similar galaxies.  We have derived ages and metallicities from Lick/IDS line-strength indices of single stellar population (SSP). It is clear that the true star formation history in these objects will be more complicated than a simple SSP. However, a lot of information can be derived from the SSP analysis. The SSP-equivalent age of a composite stellar population (CSP) is dominated primarily by the age of the youngest component and the mass fraction of the different populations. The SSP-equivalent metallicity reflects a $V-band$  luminosity weighted chemical composition (Serra and Trager 2007)\\nocite{serra}.  As bars form from stars already in the disc, we have to distinguish between stellar population gradients  that may have already been present when the bar formed and those  that formed during the bar formation or afterwards. There are not many studies analysing the stellar population of discs outside the MW. In particular, for early-type spirals, only broadband colours have been traditionally used to estimate the ages and metallicities of the stellar discs. However, stellar parameters are very difficult to obtain with broad band colours. Recently, Yoachim \\& Dalcanton (2008)\\nocite{yoachim08} have measured Lick indices in the thin and thick disc of edge-on galaxies. They found a  very strong age gradient and a null or negative metallicity gradient dependending on the index they were using. In the MW a negative age and metallicity gradient are measured. We will assume, for our discussion,  that  this is the general trend on the discs before the formation of the bar.  We have found three different types of bars according to their metallicity and age distribution along the bar:  1) Bar with negative metallicity gradient. These bars show a mean young/intermediate population ($<$~2Gyr) and have amongst the lowest central stellar velocity dispersion of the sample. They also tend to have positive age gradients (i.e. younger populations at the beginning of the bar). In fact, all the  galaxies  showing these trends posses circumnuclear regions with young population ($~$1.5~Gyr). The negative gradient can be understood as the original gradient of the disc. This is because a small fraction of young stars can bias the mean-age measured with SSP towards very small values. However, the mean metallicity is more sensitive to the more massive component. In this picture, the circumnuclear region would be strongly affecting the mean age at the beggining of the bar (lowering it)  but its influence on the mean metallicity would not be so strong. Since one would expect the original metallicity gradients  to flatten by the  velocity dispersion of the bar (Friedli et al. 1994)\\nocite{friedli94}, the presence of metallicity gradients may be indicating that they formed recently.  2) Bars with null metallicity gradients. The galaxies that do not show any gradient in their metallicity distribution along the bar have negative age gradients, that is, their ends of the bar are younger. These galaxies present gradients in both metallicity and age compatible with the models that studied mixing of metals and the distribution of stellar population ages in bars using numerical simulations (e.g Wozniak 2007\\nocite{wozniak07}; Friedli et al. 1994). They predict low age regions at the end of bars due to the accumulation of young stellar populations trapped in orbits near the ultra-harmonic resonance. During the first Gyr, in the simulations, most of the star formation happens. We observe that, although, the population at the end of the bar is younger than the average bar value, it is still larger than 3 ~Gyr. This fact might be indicating that the bar in these galaxies is at least 3~Gyrs old.  3) Bars with positive metallicity gradients. These galaxies are the ones with older mean ages in their stellar populations. They are also the ones with more massive bulges (larger $\\sigma$). In principle, this positive metallicity gradients are difficult to understand in light of our current knowledge of formation and evolution of bars. Hydrodynamical simulations indicate that, even if metal rich population can be formed at the end of the bar, the diffusion timescale in this region is very short, of the order of 50 Myr (see figure 3 in Wozniak 2007). The metallicity gradient in disc galaxies is believed to be negative (see references above) and, therefore, it is not likely that this gradient is a relic of the original metallicity gradient on the disc. We do not have a definitive answer to the origin of this metallicity gradients. One posibility is the presence of an inner ring at the beggining of the bar with enriched SF that could be populating the extremes of the bar in the same fashion as we explained for the presence of young stellar populations in the outer parts of the bar (case 2). However, we do not see this effect in all the galaxies hosting a nuclear young ring. Furthermore, the luminosity weighted mean age of the stars at the beggining of the bar is older than that in the external parts which is not expected if  the bar ends  are being populated with an enriched population coming from the inner ring. Projection effect of the old metal-rich population of the bulge on the bar region is not a likely scenario, since, at it can see from the kinematics, the outer part of the bars are dominated by a cooler stellar component.    It has been claimed that the star formation in bars follows a sequential pattern, where younger bars ($<$~1Gyr) would show star formation along the entire bar and later the star formation would be concentrated on the bar--ends and the nuclear region (Phillips 1993; Martin \\& Friedli 1997; Martin \\& Roy 1995; Friedli \\& Benz 1995; Verley et al. 2007)\\nocite{phillips,martin95,friedli95,martin97,verley}. We have checked the galaxies for emission lines coming from H{\\small II} regions to see if the ones showing negative metallicity gradients present systematically H{\\small II} along their bars.  NGC~2665 and NGC~1530 show nebular emission along their bars belonging to H{\\small II} regions, which reflects very recent star formation events. Unfortunately, we do not have high S/N information on the stellar indices for the whole bar in NGC~1530 that would allow us to establish the stellar population age and metallicity. NGC~2665, on the other hand is one of the galaxies that show a negative metallicity and a  positive age gradient along its bar. This is compatible with the idea of this sequential star formation pattern if we assume that these bar have formed less than 3 Gy ago. However, other bars such as, NGC~4245 and NGC~2681, although belonging to the same case as NGC~2665, do not present any  H{\\small II} regions along their bars.  The lowest average bar value of the $\\alpha/$Fe, here expressed as [E/Fe] (Trager et al. 2000),  tends to have lower central velocity dispersion values (Fig.~\\ref{mean.age.sig}), these galaxies have also in general a lower average age (Fig.~\\ref{mean.age.sig}).  The galaxies with higher central velocity dispersion stopped forming star earlier than galaxies with lower velocity dispersion. However, this is not an indication that the galaxies with lower central velocity dispersion host a younger (i.e. a more recently formed) bar."
    },
    "0812/0812.2253_arXiv.txt": {
        "abstract": "We use the spherical evolution approximation to investigate nonlinear evolution from the non-Gaussian initial conditions characteristic of the local $f_{nl}$ model. We provide an analytic formula for the nonlinearly evolved probability distribution function of the dark matter which shows that the underdense tail of the nonlinear PDF in the $f_{nl}$ model should differ significantly from that for Gaussian initial conditions. Measurements of the underdense tail in numerical simulations may be affected by discreteness effects, and we use a Poisson counting model to describe this effect. Once this has been accounted for, our model is in good quantitative agreement with the simulations. ",
        "introduction": "The most common inflation model (the single scalar field, slow-roll inflation) predicts the primordial perturbations to be approximately described by the Gaussian statistics. Detections of non-gaussianity can discriminate between different inflation models, so the study of various cosmological probes of primordial non-gaussianity has attracted much recent attention.  For example, the CMB has been used to constrain non-gaussianity \\citep{cmb5yr,hikageetal08,yw08,mhlm08}. In addition, various aspects of large scale structure have also been examined as probes of non-gaussianity.  These include halo abundances and bias \\citep{lmv88,rb00,mvj00,fnlverde,dalaletal08,mv08,cvm08,at08, shshp08,fnlvincent}; higher order effects on the galaxy power spectrum \\citep{mcdonald08,tkm08}; and the galaxy bispectrum \\citep{ssz04,sk07}. And, \\citet{slosar08} has discussed how to optimize the constraint on $f_{nl}$ by applying weightings on subsamples from a single tracer.  Recently, \\citet{grossi08} used N-body simulations to study the effect of primordial non-gaussianity on the nonlinear probability distribution function (PDF) of the dark matter field.  They found that departures from the Gaussian prediction are strongest in underdense regions and on small scales.  The main goal of the present study is to provide some analytic understanding of their results.  We do this by studying the nonlinear evolution of the dark matter PDF. When the initial conditions are Gaussian, then the nonlinear evolution can be modelled by the spherical and ellipsoidal collapse models \\citep{lamshethreal,lamshethred}. In these models, the evolution of the PDF depends on two ingredients: a model for the dynamical evolution from an initial state to a final one, and the appropriate average over the initial states. In particular, the spherical and ellipsoidal collapse models approximate the nonlinear dynamics as 1$\\to$1 or 3$\\to$1 mappings from the initial to the final state that do not depend on the Gaussianity of the initial conditions.  Primordial non-Gaussianity enters only because it determines the initial set of states. As a result, essentially all of the machinery developed for the Gaussian case can be simply carried over to the present study.  We recap the definitions of the local non-gaussian $f_{nl}$ model in Section~\\ref{section:nosmooth}. The effect of smoothing the initial distribution is described in Section~\\ref{section:initC}. Section~\\ref{section:nlpdf} describes the nonlinear PDF associated with the spherical evolution model and compares the results with the Gaussian case.  Section~\\ref{section:sims} compares our predictions with measurements in simulations. We summarise our results in Section~\\ref{section:discussion}. ",
        "conclusions": "\\label{section:discussion} We used the spherical evolution model to study the nonlinear evolved probability distribution function of the dark matter density field in the local primordial non-Gaussian $f_{nl}$ model.  (The spherical model is able to provide a good description of the nonlinear PDF when $f_{nl}=0$.) For currently acceptable values of $f_{nl}$, our approach shows that the signatures of primordial non-Gaussianity are small, but are most evident in underdense regions (equation~\\ref{eqn:pdfsph} and related discussion). The PDF measured in simulations can be affected by discreteness effects, especially in small underdense cells.  The effect of this can be approximated by using a Poisson counting model (equation~\\ref{eqn:poisson}). Once this is done, our model is in very good agreement with measurements in numerical simulations (Figures~\\ref{fig:sphPTr8} and~\\ref{fig:sphPTr4}), so we hope our equation~(\\ref{eqn:pdfsph}) will be useful in studies which require knowledge of the nonlinear evolved PDF.  Recently, following the same logic for why cluster abundances should be good probes of primordian non-Gaussianity, \\cite{kvj08} have suggested that void abundances should also be good probes.  Our results provide further motivation for studying underdense regions. We are in the process of developing a more complete model of voids and void shapes \\cite[following][]{sw04}.  Finally, note that the parameter which describes the non-Gaussianity in the smoothed initial and final fields, $\\sigma\\,S_3$, is only weakly scale-dependent.  Therefore, our analysis indicates that the non-gaussian distribution of the resulting nonlinear PDF (equation~\\ref{eqn:pdfsph}) can be written in terms of the scaled variable $\\nu=\\delta_l/\\sigma$.  So one could formulate a reconstruction of the initial $f_{nl}$ field analogously to how this is done for the Gaussian case \\citep{lamshethreal}.  We have not pursued this further."
    },
    "0812/0812.0917_arXiv.txt": {
        "abstract": "\\noindent We present observations of the Lynds' dark nebula LDN\\,1111 made at microwave frequencies between 14.6 and 17.2\\,GHz with the Arcminute Microkelvin Imager (AMI). We find emission in this frequency band in excess of a thermal free--free spectrum extrapolated from data at 1.4\\,GHz with matched {\\it uv}-coverage. This excess is $> 15\\sigma$ above the predicted emission. We fit the measured spectrum using the spinning dust model of Drain \\& Lazarian (1998a) and find the best fitting model parameters agree well with those derived from Scuba data for this object by Visser et~al. (2001). ",
        "introduction": "Recent pointed observations (Finkbeiner et~al. 2002, 2004; Casassus et~al. 2004, 2006; Watson et al. 2005; Scaife et~al. 2007; Dickinson et~al. 2007) have provided some evidence for the anomalous microwave emission commonly ascribed to spinning dust (Drain \\& Lazarian 1998a,b). Although this emission was originally seen as a large scale phenomenon in CMB observations (see e.g. Kogut et~al. 1996) it has been suggested that the emission occurs in a number of distinct astronomical objects, such as dark clouds, and {\\sc Hii} and photo-dissociation regions. It often appears to be correlated with thermal dust emission as supported by the pointed observations mentioned previously but it must be stated that this is not always the case, see for example Casassus et~al. (2008).  In spite of these predictions the evidence for anomalous microwave emission in compact objects is often contradictory. Early observations  below 10\\,GHz of the molecular cloud LPH\\,96 (Finkbeiner et~al. 2002), which showed a rising spectrum were later contradicted by observations at 31 and 33\\,GHz which found emission consistent with an optically thin free--free spectrum extrapolated from lower frequencies (Dickinson et~al. 2006; Scaife et~al. 2007). Although some evidence for an excess was found in a sample of Southern {\\sc Hii} regions (Dickinson et~al. 2007) and more significantly in RCW175 (Dickinson et~al. 2008), no emission inconsistent with free--free was found in a sample of Northern {\\sc Hii} regions (Scaife et~al. 2008). Casassus et~al. (2004) proposed a flux density of approximately 1\\,Jy from the Helix planetary nebula at 31\\,GHz to be in excess of a free--free spectrum extrapolated from lower frequencies. However, based on flux densities from the literature at 1.4, 2.7, 6.63\\,GHz (Ehman, Dixon, \\& Kraus 1970; Higgs 1971; Wall, Wright, \\& Bolton 1976) and the recent WMAP 5-year densities at 23--94\\,GHz (Wright et~al. 2008), which are all also $\\approx 1$\\,Jy, we suggest that there is no evidence for this reported excess in the flux density spectrum, although the nebula may be anomalous in other ways.  In this letter we present observations of the Lynds' dark nebula LDN\\,1111, taken from the AMI sample of compact Galactic star formation regions (Scaife et~al., in prep). This sample was selected from the SCUBA sample of compact Lynds' clouds (Visser, Richer \\& Chandler 2001).  The spectra of dark clouds at gigahertz frequencies is poorly documented in all but a few cases (Finkbeiner et~al 2004; Casassus et~al. 2006; Casassus et~al. 2008). In those cases where cm-wave data is available (Casassus et~al. 2006; Casassus et~al. 2008) the behaviour of these objects has been found to be anomalous in a number of ways and in the case of LDN\\,1622 (Casassus et~al. 2006) to show a distinct excess of microwave emission.    LDN\\,1111 ($\\alpha = 21^{\\rm{h}} 40^{\\rm{m}} 30^{\\rm{s}}, \\delta = +57^{\\circ} 48' 00''$ J2000) lies on the inside edge of the heel of the large ($\\approx 3$\\,degree) horseshoe shaped {\\sc Hii} region IC\\,1396 (Sh 2-131; Sharpless 1959). It has no known IRAS association (Parker 1988) but has been studied in the submillimetre at 850\\,$\\mu$m with the SCUBA instrument (Visser, Richer \\& Chandler 2001). An opacity class 6 object ($A_{\\nu} \\ge 5$\\,mag; Lynds 1962), it is one of the most opaque dark nebulae.   \\begin{section}{Observations}  \\begin{subsection}{Calibration and Data Reduction}  The AMI Small Array (SA) is a radio interferometer which observes in eight frequency channels in the band 12--18\\,GHz at the Mullard Radio Astronomy Observatory, Lord's Bridge, Cambridge, UK. In practice, the lowest two frequency channels are generally unused due to a low response in this frequency range, and interference from geostationary satellites. \\citet{0807.2469} discusses the telescope in more detail.  Observations of LDN\\,1111 were made in 13 hours over two days in 2007 November. Data reduction was performed using the local software tool \\textsc{reduce}. This applies both automatic and manual flags for interference and shadowing and hardware errors. It also applies phase and amplitude calibrations; it then Fourier transforms the correlator data to synthesize the frequency channels before output to disk in uv FITS format suitable for imaging in \\textsc{aips}.  Flux calibration was performed using short observations of 3C286 near the beginning and end of each run. We assumed I+Q flux densities for this source in the AMI SA channels consistent with Baars et al. (1997) $\\simeq 3.3$\\,Jy at 16\\,GHz. As \\citet{1977A+A....61...99B} measure I and AMI SA measures I+Q, these flux densities include corrections for the polarisation of the source derived by interpolating from VLA 5, 8 and 22\\,GHz observations. A correction is also made for the changing intervening air mass over the observation. From other measurements, we find the flux calibration is accurate to better than 5 per cent (Scaife et~al. 2008; Hurley--Walker et~al., in press).  The phase was calibrated using hourly interleaved observations of the point source J2201+508 ($\\alpha = 22^h 01^m 43^s.5,  \\delta = 50^{\\circ} 48' 56''.4$), which has a flux density of 0.3\\,Jy at 16\\,GHz. It was selected from the Jodrell Bank VLA Survey (JVAS; \\citealt{1992MNRAS.254..655P}). After calibration, the phase is generally stable to $5^{\\circ}$ for channels 4--7, and $10^{\\circ}$ for channels 3 and 8. In this work we use only channels 4--7 due to their superior phase stability.   The FWHM of the primary beam of the AMI SA is $\\approx 20$\\arcmin at 16\\,GHz. The FWHM of the synthesised beam of the combined channel map towards LDN\\,1111 is $2.4'\\times2.1'$ using natural weighting.  \\end{subsection}  \\begin{subsection}{Imaging}  \\begin{figure} \\centerline{\\includegraphics[height=8cm,width=9cm,angle=0]{./L1111ami.ps}} \\caption{Cleaned combined channel map of LDN\\,1111 at 16\\,GHz. Contours are increments of 1\\,mJy beam$^{-1}$, with the first contour at 1\\,mJy beam$^{-1}$. The primary beam FWHM of the AMI SAis shown as a black circle and the synthesized beam as a grey ellipse in the bottom left-hand corner. \\label{fig:amiplot}} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[height=7.5cm,width=8.5cm,angle=0]{./L1111csamp100.ps}} \\caption{Cleaned map of the $uv$ sampled CGPS data at 1.4\\,GHz. Contours are increments of 1\\,mJy beam$^{-1}$, with the first contour at 1\\,mJy beam$^{-1}$. The primary beam FWHM of the AMI SA is shown as a black circle and the synthesized beam as a grey ellipse in the bottom left hand corner.\\label{fig:l1111cgps}} \\end{figure}  The reduced visibilities were imaged using the AIPS data package. Dirty images were deconvolved using {\\sc imagr} which applies a differential primary beam correction to the {\\sc clean} components to account for the different frequency channels of the AMI instrument. Maps were made from both the combined channel set, shown here, and for individual channels.  Two objects are labelled in both Fig.~\\ref{fig:amiplot} and Fig.~\\ref{fig:l1111cgps}: object $A$, which is the dark cloud; and object $B$, which is the NVSS radio source J213955+574859.  \\end{subsection}  \\begin{subsection}{Radio Spectra}  \\begin{table} \\centering \\caption{Flux densities for  J213955+574859\\vspace{0.2cm}\\label{tab:l1111src}} \\begin{tabular}{ccc} \\hline\\hline Frequency & Flux density & Reference\\\\ (GHz) & (mJy) & \\\\ \\hline 0.327 & 20.0$\\pm$4.8 & WENSS; Rengelink et~al. 1997\\\\ 0.408 & 13.4$\\pm$1.3 & this work; data from \\\\ && Taylor et~al. 2003\\\\ 1.420 & 12.2$\\pm$1.25& this work; data from \\\\ && Taylor et~al. 2003\\\\ 1.420 & 11.1$\\pm$0.5 & NVSS; Condon et~al. 1996\\\\ 14.631& 3.48$\\pm$0.56& this work\\\\ 15.381& 2.88$\\pm$0.60& this work\\\\ 16.130& 3.02$\\pm$0.55& this work\\\\ 17.150& 2.93$\\pm$0.57& this work\\\\ \\hline \\end{tabular} \\end{table}   Since the frequency coverage of AMI tells us nothing about the spectral behaviour at longer radio wavelengths we must combine our new data with existing archival data. Ideally total power measurements are required, with the same or better angular resolution than AMI. Data of this type allow sampling in the $uv$ plane to match exactly the measured angular scales of AMI. Here we use data from the Canadian Galactic Plane Survey (hereinafter CGPS; Taylor et~al. 2003) at 1.42\\,GHz. These data are a combination of single-dish and synthesis data, and provide a total power measurement of this region with a resolution of $\\approx 1'$. In order to mimic an AMI observation, these data are modulated by the primary beam of AMI before sampling in the $uv$ plane to match the AMI $uv$ coverage. These data are then processed in the same way as the true AMI visibilities to give the final sampled image. Fig.~\\ref{fig:l1111cgps} shows the sampled data which result from matching the visibility coverage to AMI channel 4.  Spectra of both LDN\\,1111 and the radio source J213955+574859 are presented, identified as objects $A$ and $B$ respectively in Fig.~\\ref{fig:amiplot}. Errors on the data points in these spectra are calculated using a contribution for the rms noise in each channel, a conservative 5\\% flux calibration error and a contribution for the uncertainty in flux extraction. The morphology of LDN\\,1111 and its surroundings makes this extraction subject to the fitting area used. To account for this in the errors 10 independent `fits' are made. The recorded flux density is the mean of these measurements and the variance is referred to as $\\sigma^2_{\\rm{fit}}$. These errors are combined in quadrature: $\\sigma_{\\rm{S}} = \\sqrt{\\sigma_{\\rm{rms}}^2 + (0.05S_{\\rm{I}})^2 + \\sigma^2_{\\rm{fit}}}$. The flux extraction is performed using the {\\sc{fitflux}} program (Green 2007), which takes a user defined polygonal aperture and performs aperture photometry after subtracting a twisted plane background level. These errors are dominated by the 5\\% calibration uncertainty, with the map r.m.s. being $\\approx$0.2\\,mJy per channel.  The spectrum of J213955+574859 is presented in order to demonstrate the flux calibration of AMI, see Fig.~\\ref{fig:l1111src}. The data points for this plot are given in Table~\\ref{tab:l1111src}. In addition to data points from the literature, flux densities at 408\\,MHz and 1.42\\,GHz from archival CGPS data are also plotted. These data have been extracted using the {\\sc{fitflux}} program as described above. Spectral indices are calculated using firstly the AMI data points plus the literature data, $\\alpha=0.52$; and secondly using all the data points, $\\alpha=0.46$. These fits are shown as dot-dash and dashed lines respectively. Both are consistent with $\\alpha_{\\rm{\\sc{wenss}}}^{\\rm{\\sc{nvss}}}=0.43\\pm0.18$.  For J213955+574859, which is approximately point-like, we expect a negligible amount of flux loss from extended structure across the AMI bandwidth. In the case of LDN\\,1111 the same is not true. Using the CGPS map at 1.42\\,GHz we find $>$ 50\\% flux loss: the integrated flux from the original unsampled map, calculated as above, is $S_{1.42,{\\rm{total}}} = 25.4\\pm2.7$\\,mJy. However, since the morphology of the source varies between 1.4 and 16\\,GHz we have not used the CGPS data as a model to calculate flux loss.  Flux densities for LDN\\,1111 are listed in Table~\\ref{tab:floss}. It is immediately obvious that the flux densities for LDN\\,1111 between 14 and 18\\,GHz are inconsistent with an optically thin free--free spectrum extrapolated from 1.42\\,GHz. At the high frequency end of the AMI spectrum there is an excess of almost 50\\,mJy, $> 15\\sigma_{\\rm{S}}\\, (\\approx 250 \\sigma_{\\rm{rms}})$.  Although an absolute measure of the flux loss is not feasible, it is possible to make an estimate of how the {\\it relative} flux loss will affect the shape of the microwave spectrum. This is done using a multi-variate Gaussian model for the source with dimensions of the deconvolved source at 14.63\\,GHz (AMI channel 4) as given by the {\\sc{aips}} task {\\sc{jmfit}}, $7'\\times 4'$. This model is then sampled in $uv$ using the exact coverage of the individual channels, and their flux loss relative to the lowest channel is recorded. We assume an uncertainty of 10\\% on the percentage flux loss and propagate the errors. The results of this calculation are shown in Fig.~\\ref{fig:l1111fullspec} and tabulated in Table~\\ref{tab:floss}. It should be noted that this is a conservative uncertainty and the relative flux loss will only be in error should the morphology change significantly between 14 and 17\\,GHz.  Fig.~\\ref{fig:l1111fullspec} includes a data point at 353\\,GHz (850\\,$\\mu$m) from the Scuba instrument (Visser, Richer \\& Chandler 2001) for illustartive purposes. This data point is uncorrected for flux loss. These Scuba data at 850\\,$\\mu$m are chopped at 2.5\\,arcmin and consequently we would assume that information on scales larger than this is lost. To correctly estimate the flux loss at this frequency we would require information on the morphology of the vibrational dust emission. Parker (1988), from whose paper the Scuba sample were selected, lists the dimensions of LDN 1111 as $1.7 \\times 1.1$\\,arcmin. A multivariate Gaussian model based on this data would suggest a flux loss of approximately 73\\%.  Source extraction methods applied to the WMAP data at 23--91\\,GHz (L{\\'o}pez-Caniego et~al. 2007) are only able to give us upper limits on any compact emission towards LDN\\,1111, with the lowest limit being at 41\\,GHz, $S_{41} \\le 0.30$\\,Jy. This value is consistent with the noise level at that frequency.   \\begin{figure} \\centerline{\\includegraphics[height=8cm,width=5cm,angle=-90]{./L1111srcspec.ps}} \\caption{Spectrum of the source J213955+574859. Data points from the literature are shown as crosses and are described in the text. Data points extracted from the CGPS data are shown as un-filled triangles. Data points from the AMI are shown as filled circles. \\label{fig:l1111src}} \\end{figure}  \\begin{table} \\centering \\caption{Flux densities for LDN\\,1111 \\vspace{0.2cm}\\label{tab:floss}} \\begin{tabular}{cccc} \\hline\\hline Frequency & Flux density & Relative & Corrected \\\\ & & flux loss & flux density \\\\ (GHz) & (mJy) & (\\%) & (mJy)\\\\ \\hline 1.420$^a$  & 7.2 $\\pm$0.7 & -  & 7.2$\\pm$0.7  \\\\ 14.631 & 50.7$\\pm$2.9 & -  & 50.7$\\pm$2.9 \\\\ 15.381 & 53.9$\\pm$3.2 & 10 & 59.9 $\\pm$6.8\\\\ 16.130 & 54.5$\\pm$3.3 & 17 & 65.7 $\\pm$7.6\\\\ 17.150 & 55.5$\\pm$3.3 & 24 & 73.0 $\\pm$8.4\\\\ \\hline \\end{tabular} \\begin{minipage}{7cm} {\\small\\vspace{0.1cm} $^{a}$ $uv$ sampled CGPS data point} \\end{minipage} \\end{table}  \\end{subsection} \\end{section}  \\begin{section}{Analysis}  The emission seen between 14.63 and 17.15\\,GHz by AMI is clearly in excess of a simple free--free spectrum extrapolated from the CGPS data at 1.4\\,GHz. A number of possibilities may throw some light on the nature of this excess. The thermal (vibrational) dust spectrum of proto-planetary disks around T-Tauri stars is expected to extend into the cm-regime. In these circumstances grain growth has increased the population of cm-sized grains, or pebbles, causing $\\beta$ to approach zero and the greybody spectrum to fall off with an index approaching 2. However the size of these disks is small and would be unlikely to be resolved by either AMI or Scuba. LDN\\,1111 is quite obviously elongated in the AMI data, as it is in the Scuba map which has a resolution of 14\\,arcsec. Although we do not deny the possibility of there being such a disk embedded within LDN 1111 or projected along the line of sight through this object, in the absence of further evidence we choose to pursue an alternative explanation.  A further possibility for a superposition of LDN\\,1111 and another object is that there may be an {\\sc Hii} region, with a density such that it is still optically thick at 16\\,GHz, whose flux is contributing to the spectrum. An {\\sc Hii} region with a turn-over frequency $>30$\\,GHz, as would be required here, would possess an emission measure in excess of $4\\times 10^9$\\,pc cm$^{-6}$. This would put it in the regime of ultra-compact and hyper-compact {\\sc Hii} regions. The spectrum of these objects is inverted with an index of $\\simeq -2$. The AMI data points alone have a spectral index of $\\alpha = -2.86\\pm0.35$, rather steeper than would be expected. Indeed hyper-compact {\\sc Hii} regions exhibit slightly shallower spectra (Franco et~al. 2000) which extends even to mm-wavelengths. The lack of a visible source in the WMAP data towards LDN\\,1111 would suggest that this is not the case here. In addition hyper-compact {\\sc Hii} regions tend to have very broad radio recombination line (RRL) widths, typically in excess of 50 km s$^{-1}$ (Kurtz 2005). Ultra-compact {\\sc Hii} regions also show broadened line widths, of typically 30--40 km s$^{-1}$. The RRL measurement of Heiles, Reach \\& Koo (1996) which encompasses LDN\\,1111 has a width of only 21.4 km s$^{-1}$, although we note that this measurement is not directed at LDN\\,1111 specifically. An emission measure of this magnitude would also imply an {\\sc Hii} region mass of around 400\\,M$_{\\odot}$. The total mass of the cloud from the Scuba data (Visser et~al. 2002) is 0.3\\,M$_{\\odot}$ only, which given the calculated flux loss in this data implies an upper limit on the mass of 1\\,M$_{\\odot}$.  A third possibility for the excess emission seen in the AMI data is that of dipole emission from rapidly rotating small dust grains (Drain \\& Lazarian 1998a,b). The observational evidence for this emission mechanism has been explored in the introduction to this letter. It is a relatively new emission mechanism and the evidence for its existence is not conclusive. We assess the possibility that the emission seen in LDN\\,1111 arises as a consequence of spinning dust by comparing the data to the model of Draine \\& Lazarian (1998a;b),which has been used extensively in the past. To make this comparison we use the MCMC based software {\\sc metro} (Hobson \\& Baldwin 2004) to find the best--fit parameters.  We fitted a model which has a free--free component normalized to give the flux density at 1.4\\,GHz from the sampled CGPS data with a spectral index of $\\alpha = 0.1$. To this we add a spinning dust component scaled from the DL98 molecular cloud model.  The model is parameterized by the column density, $N({\\rm{H_2}})$, and the angular size of the object. Visser et~al. (2001) calculate the average and peak column densities for LDN\\,1111 as 5 and 13$\\times10^{21} \\rm{cm}^{-2}$, respectively. The results of our model fitting are $N({\\rm{H_2}})$ = 6.72$\\pm$0.58$\\times 10^{21} {\\rm{cm}}^{-2}$, and $\\Delta \\theta$ = 5.38$\\pm$0.26\\,arcmin. These values provide a $\\chi^2$ of 1.03 (79\\% probability). However, given the small number of degrees of freedom in this case we do not propose this statistic as being conclusive. The resulting model is shown in Fig.~\\ref{fig:l1111fullspec}. The column densities agree well with those of Visser et~al., although as might be expected there is a degeneracy between the two parameters.   \\begin{figure}  \\centerline{\\includegraphics[height=8cm,width=5cm,angle=-90]{./L1111dl98.ps}}  \\caption{LDN\\,1111: Data points are listed in Table~\\ref{tab:floss}, with the addition of the SCUBA flux density at 850\\,$\\mu$m. A free--free spectrum as described in the text and a greybody spectrum scaled to the Scuba data point are shown as dashed lines. The best-fit spinning dust model of Draine \\& Lazarian (1998) is shown as a dot--dash line and the total spectrum as a solid line. \\label{fig:l1111fullspec}} \\end{figure}  \\end{section}  \\begin{section}{Discussion and Conclusions}  Given the available evidence we propose the excess of emission we see towards the dark cloud LDN\\,1111 in the microwave band to be a result of emission from small spinning dust grains. This excess is quantified relative to a flux density derived from lower frequency data sampled to give the same $uv$ coverage and consequently measuring the same angular scales on the sky. Lacking further data at lower frequencies with suitable angular resolution this one point is the basis for the assumption that there is an excess in the AMI frequency band. We therefore qualify this excess by using the original CGPS data to measure the {\\it total power}, i.e. power on all scales, flux density at 1.42\\,GHz to be $S_{1.42,{\\rm{total}}} = 25.4\\pm2.7$\\,mJy. Since the object is extended this value may be regarded as an upper limit on the measureable flux at 1.42\\,GHz. In addition the total power CGPS data at 408\\,MHz may then be used to constrain better the lower frequency spectrum: $S_{408,{\\rm{total}}} = 21.6\\pm11.8$\\,mJy. The large error on this flux density is dominated by the fitting uncertainty. In combination these two flux densities yield a spectral index of $\\alpha = -0.16\\pm0.12$, which predicts a {\\it total power} flux at 16\\,GHz of $S_{16,{\\rm{total}}} = 36.9^{+13.8}_{-10.1}$\\,mJy. This suggests that even in terms of total power there is a clear excess at 16\\,GHz.  Relative to a canonical free--free spectrum extrapolated from our sampled CGPS data at 1.42\\,GHz we see an excess of $S_{\\rm{excess}} = 60.1\\pm7.6$\\,mJy, $\\approx 8\\sigma$, assuming a fixed spectral index of $\\alpha = 0.1$. Using the spectral index derived from the CGPS total power maps at 408\\,MHz and 1.42\\,GHz then there is an excess of $S_{\\rm{excess}'} = 55.10$\\,mJy, which is still $> 7\\sigma$.  In conclusion, we have presented observations of the Lynds' dark nebula LDN\\,1111 at frequencies of 14.6--17.2\\,GHz. These measurements show an excess of emission towards this object relative to an extrapolated free--free spectrum at a significance of $\\approx 15\\sigma$. We have proposed that this excess may be due to emission from small spinning dust grains and find that it is well-described by the model of Draine \\& Lazarian (1998a;b).   \\end{section} ",
        "conclusions": ""
    },
    "0812/0812.0126_arXiv.txt": {
        "abstract": "{A bright feature 100 pc away from the core in the powerful jet of M\\,87 shows mysterious properties. Earlier radio, optical and X-ray observations have shown that this feature, labelled HST-1, is superluminal, and is possibly connected with the TeV flare detected by HESS in 2005. To examine the possible blazar-like nature of HST-1, we analyzed $\\lambda$2\\,cm VLBA data from dedicated full-track observations and the 2\\,cm survey/MOJAVE VLBI monitoring programs observed from 2000 to 2008. Applying wide-field imaging techniques, the HST-1 region was imaged at milliarcsecond resolutions. Here we present the first 15\\,GHz VLBI detection of this feature and discuss the connection between our radio findings and the TeV detection.}  \\FullConference{The 9th European VLBI Network Symposium on The role of VLBI in the Golden Age for Radio Astronomy and EVN Users Meeting\\\\ September 23-26, 2008\\\\ Bologna, Italy}  \\begin{document} ",
        "introduction": "Active Galactic Nuclei (AGN) are among the most interesting phenomena in the Universe, and they have been studied since the first discovery in 1943 \\cite{seyfert43}. Until now, although there are clues that imply a super-massive black hole (SMBH) is the engine launching the powerful jet \\cite{urry95}, the exact mechanism remains unknown.  M\\,87 (also known as Virgo\\,A) is a nearby elliptical galaxy located in the Virgo cluster. It hosts a very powerful one-sided jet emerging from the central region. Observations show that M\\,87 contains $2.4\\times10^{9}$ $M_{\\odot}$ within a 0.25$^{\\prime\\prime}$ radius, which give hints of a SMBH in the center \\cite{harms94}. Due to its closeness (16\\,Mpc), M\\,87 is a major candidate for studying AGN phenomena, and has been monitored at different wavelengths over the last decades. Superluminal motion was reported from {\\it Hubble Space Telescope} ({\\it HST\\/}) observations within 6$^{\\prime\\prime}$ of the jet with speeds of 4\\,$c$ to 6\\,$c$ \\cite{biretta99}. Discrepant speeds were reported with values between 0.25\\,$c$ to 0.4\\,$c$ \\cite{ly07}, and a value of 2\\,$c$ \\cite{walker08} from VLBA 43\\,GHz observations. VLBA $\\lambda$2\\,cm observations found apparent speeds <\\,0.4\\,$c$ during 1994 to 2001 \\cite{kellermann04,kovalev07}. Therefore, the kinematical properties of the jet in M\\,87 are still under discussion.  From the \\textit{HST} observation in 1999, a bright knot along the jet was discovered, and was named HST-1. This feature is located 0.8$^{\\prime\\prime}$ away from the core, which corresponds to a projected distance of 0.1\\,kpc, and is active in the radio, optical and X-ray regimes. VLBA $\\lambda$20 cm observations show that HST-1 has sub-structure and contains superluminal moving components up to 4\\,$c$ \\cite{cheung07}. These observations interpret HST-1 as a collimated shock in an AGN jet.  However, recent multi-wavelength observations show that HST-1 could be related to the origin of the TeV emission of M\\,87. In 2005, the HESS telescope detected a TeV flare from M\\,87 \\cite{aharonian06}. Comparing with the results in soft X-rays (\\textit{Chandra}), and VLA $\\lambda$2\\,cm observations \\cite{cheung07, harris06}, the light curves of HST-1 reach the maximum in 2005, while the resolved core shows no correlation with the TeV flare. Therefore, the TeV emission from M\\,87 might be associated to HST-1 \\cite{harris08}. The AGN standard model considers the blazar behavior to originate at the vicinity of the SMBH. However, HST-1 is 100\\,pc away from the core. If the hypothesis is true, the present AGN model would face a challenge. Here we examined this hypothesis with $\\lambda$2\\,cm VLBI wide-field imaging of the HST-1 feature. ",
        "conclusions": "With tapering and natural weighting to the data, the most extended inner-jet structures were traced until 100\\,mas (8\\,pc), and the noise levels of HST-1 \\textsc{clean} images are listed in Table \\ref{tab:epoch_list}. The inner-jet structure had a total flux density of 2.5\\,Jy, with overall flux density changes up to 0.5\\,Jy during 2000 to 2008. Out of 13 epochs, HST-1 was detected in 4 epochs from 2004.61 to 2005.85 (see Table \\ref{tab:epoch_list} and Figure \\ref{fig:HST-1 overlay}). The feature had a peak surface brightness of 3 to 5\\,mJy\\,beam$^{-1}$, and the total flux density was 14 to 22\\,mJy, with an extended structure of 10\\,mas. These detections show that HST-1 has a weak peak at $\\lambda$2\\,cm.  Due to its low flux density and extended structure, only the brightest feature could be detected. The peak position changes show an outward motion from the core. By linear-fitting the distance of the peak position to the central core as a function of time, a proper motion of 0.44\\,$\\pm$\\,2.24 mas\\,yr$^{-1}$ was derived, which corresponds to a projected linear speed of 0.11\\,$\\pm$\\,0.56\\,$c$. Previous observations report higher speeds up to 6\\,$c$ from the HST-1 region \\cite{biretta99, cheung07}. Figure \\ref{fig:HST-1 overlay} shows two overlaid images of HST-1 at $\\lambda$20\\,cm (Cheung, priv.\\ comm.) and $\\lambda$2\\,cm. Note that the feature detected at $\\lambda$2\\,cm was located at the same position as the HST-1d component by Cheung et al \\cite{cheung07}. The HST-1d component shows a stationary behavior with a speed of <\\,0.25\\,$c$, which is comparable with our measurements.  The emission intensity of HST-1 has reached the maximum in different wavebands at 2005, and the light curves from the VLA $\\lambda$2\\,cm, VLBA $\\lambda$20\\,cm, and X-ray observations share the same tendency \\cite{cheung07, harris06}. Our results were consistent with those observations. The HESS team reports that a TeV flare was detected in 2005 \\cite{aharonian06}. It is believed that the very high energy emissions come from the vicinity of the SMBH, however, {\\it Chandra\\/} observations show that the resolved core has a flat light curve during this time, while HST-1 has a significant peak in 2005. Our results also show that the total flux density of the inner-jet does not have a peak in 2005, while HST-1 was only detected during 2004 to 2005. Those findings support that HST-1, a feature 100\\,pc away from the SMBH, could be the source of the TeV emission. However, according to our detection, HST-1  appeared to be extended, which does not support the blazar type activity. Moreover, no emerging or rapidly moving features were found among the detections. Therefore, our results does not support that HST-1 was the origin of the TeV emission, however, one could not exclude the possibility.   \\begin{figure}% \\centering \\includegraphics[angle=270, width=0.75\\textwidth]{2005.3_overlay_mas_final.ps} \\includegraphics[angle=270, width=0.75\\textwidth]{2005.9_overlay_mas_final.ps} \\caption{VLBA images of the HST-1 region in M\\,87 at $\\lambda$\\,2\\,cm (black contour, beam size at half power level\\,=\\,2$\\times$1\\,mas, P.A.\\,=\\,$-16^{\\circ}$, plotted bottom left) and $\\lambda$\\,20\\,cm (grey contour, beam size\\,=\\,8$\\times$3\\,mas, P.A.\\,=\\,0$^\\circ$, plotted bottom right). Upper panel: epoch 2005.30 ($\\lambda$\\,2\\,cm, peak surface brightness: 3.4\\,mJy\\,beam$^{-1}$) and 2005.35 ($\\lambda$\\,20\\,cm, peak: 45\\,mJy\\,beam$^{-1}$); lower panel: epoch 2005.85 ($\\lambda$\\,2\\,cm, peak: 3.3\\,mJy\\,beam$^{-1}$) and 2005.82 ($\\lambda$\\,20\\,cm, peak: 42\\,mJy\\,beam$^{-1}$). The coordinates are relative to the core location. The lowest contour is 0.7\\,mJy\\,beam$^{-1}$, and the contour levels are separated by a factor of $\\sqrt2$.} \\label{fig:HST-1 overlay} \\end{figure}  \\paragraph*"
    },
    "0812/0812.2912_arXiv.txt": {
        "abstract": "A significant fraction of stars in globular clusters (about 70\\%-85\\%) exhibit peculiar chemical patterns with strong abundance variations in light elements along with constant abundances in heavy elements. These abundance anomalies can be created in the H-burning core of a first generation of fast rotating massive stars and the corresponding elements are convoyed to the stellar surface thanks to rotational induced mixing. If the rotation of the stars is fast enough this matter is ejected at low velocity through a mechanical wind at the equator. It then pollutes the ISM from which a second generation of chemically anomalous stars can be formed. The proportion of anomalous to normal star observed today depends on at least two quantities : (1) the number of polluter stars; (2) the dynamical history of the cluster which may lose during its lifetime first and second generation stars in different proportions.  Here we estimate these proportions based on dynamical models for globular clusters.  When internal dynamical evolution and dissolution due to tidal forces are accounted for, starting from an initial fraction of anomalous stars of 10\\% produces a present day fraction of about 25\\%, still too small with respect to the observed 70-85\\%.  In case gas expulsion by supernovae is accounted for, much higher fraction is expected to be produced.  In this paper we also address the question of the evolution of the second generation stars that are He-rich, and deduce consequences for the age determination of globular clusters. ",
        "introduction": "\\begin{figure} \\centering \\includegraphics[width=0.6\\textwidth]{scenario2} \\caption{Schematic evolution of a globular cluster: a first generation of stars is born from a giant molecular cloud. Massive stars of the first generation evolve and give birth to a second generation of low-mass stars (dashed symbols in middle panels). Then the cluster evolves and today (lower panels) a mixture of first and second generation low-mass stars is present. In the right panel the cluster is initially mass segregated and massive stars (and hence stars of second generation) are concentrated toward the cluster centre.} \\label{fig:selfenrich} \\end{figure}  It has long been known that globular cluster stars present some striking anomalies in their content in light elements whereas their heavy elements (i.e., Fe-group, $\\alpha$-elements) remain fairly constant from star to star (with the notable exception of $\\omega$~Cen).  While in all the Galactic globular clusters studied so far one finds ``normal'' stars with detailed chemical compositions similar to those of field stars of same metallicity (i.e., same [Fe/H]), one also observes numerous ``anomalous'' main sequence and red giant stars that are simultaneously deficient (to various degrees) in C, O, and Mg, as well as enriched in N, Na, and Al (for reviews see \\citealp{GrattonSneden2004,Charbonnel2005}).   These abundance variations are expected to result from H-burning nucleosynthesis at high temperatures around $75\\times 10^6$~K \\citep{DenisenkovDenisenkova1989,DenisenkovDenisenkova1990,LangerHoffman1995,PrantzosCharbonnel2007}. Such temperatures are not reached in the low-mass main sequence and RGB stars that are chemically peculiar, meaning that the stars inherited their abundance anomalies at stellar birth.  Here we follow the work of \\citet{PrantzosCharbonnel2006} and \\citet{DecressinMeynet2007} who propose that abundance anomalies are build up by fast rotating, fast evolving massive stars. During their main sequence evolution, rotationally-induced mixing transports elements synthesised in the convective H-burning core to the stellar surface. For stars heavier than 20~\\Ms{}, the surface reaches break-up at the equator (i.e., the centrifugal equatorial force balances gravity), providing their initial rotational velocity is high enough. In such a situation, a slow mechanical wind develops at the equator and forms a disc around the stars similar to what happens to Be stars \\citep{TownsendOwocki2004,EkstromMeynet2008}. Matter in discs is strongly enriched in H-burning products and has a slow escape velocity that allows it to stay in the potential well of the cluster. On the contrary, matter released later through radiativelly-driven winds during most of the He-burning phase and then through SN explosion has a very high velocity and is lost by the cluster. Therefore, new stars can form only from the matter available in discs with the abundance patterns we observe today. Thus globular clusters can contain two populations of low-mass stars: a first generation which has the chemical composition of the material out of which the cluster formed (similar to field stars with similar metallicity); and a second generation of stars harbouring the abundance anomalies born from the ejecta of fast rotating massive stars. This scenario is sketched in Fig.~\\ref{fig:selfenrich}. ",
        "conclusions": ""
    },
    "0812/0812.0310_arXiv.txt": {
        "abstract": "We report on low resolution ($R\\approx 3000$) spectropolarimetry of the A0 supergiant star HD~92207.  This star is well-known for significant spectral variability.  The source was observed on seven different nights spanning approximately 3 months in time.  With a rotation period of approximately 1 year, our data covers approximately a quarter of the star's rotational phase.  Variability in the continuum polarization level is observed over this period of time.  The polarization across the H$_\\alpha$ line on any given night is typically different from the degree and position angle of the polarization in the continuum.  Interestingly, H$\\beta$ is not in emission and does not show polarimetric variability. We associate the changes at H$_\\alpha$ as arising in the wind, which is in accord with the observed changes in the profile shape and equivalent width of H$_\\alpha$ along with the polarimetric variability.  For the continuum polarization, we explore a spiral shaped wind density enhancement in the equatorial plane of the star, in keeping with the suggestion of Kaufer \\etal\\ (1997).  Variable polarization signatures across H$_\\alpha$ are too complex to be explained by this simple model and will require a more intensive polarimetric follow-up study to interpret properly. ",
        "introduction": "Spectropolarimetry is a valuable tool for studying a number of important properties of stars.  Circular polarization is one of the few direct means of measuring stellar magnetic fields (e.g., Babcock 1958).  Linear polarization, such as from Thomson scattering in circumstellar envelopes of hot stars, has been important for ascertaining deviations from sphericity in circumstellar media (e.g., the Be stars [Poeckert \\& Marlborough 1976; McLean \\& Brown 1978; Wood, Bjorkman, \\& Bjorkman 1997] and supernovae [Wang \\etal\\ 1996; Leonard \\etal\\ 2005]).  A challenge for this area is that net polarizations tend to be small so that large telescopes are required to obtain sufficiently high quality data for sources that are only moderately faint.  Fortunately, astronomers now have access to several 8--10~m class or larger telescopes that are outfitted with polarimeters.  With such instrumentation, studies of polarizations across spectral lines -- combining good spectral resolution with large telescope collecting areas -- have become increasingly popular and important for understanding structure in circumstellar media.  Here we present an analysis of variable linear polarization across the H$_\\alpha$ line and in the adjacent continuum for the A0 supergiant star HD~92207 at seven epochs over a period of three months.  This particular star was chosen for several different reasons. It has been studied spectroscopically by Kaufer \\etal\\ (1996) who found substantial line profile and line equivalent width changes at H$_\\alpha$.  These variations appear to be related to the well-known discrete absorption components (or ``DACs'') that are commonly seen in UV lines of early-type stars (e.g., Massa \\etal\\ 1995; Kaper \\etal\\ 1996). HD~92207 is not far removed from the Luminous Blue Variable stars and O supergiants that show variable continuum polarizations (Lupie \\& Nordsieck 1987; Taylor \\etal\\ 1991; Harries, Howarth, \\& Evans 2002; Davies, Vink, \\& Oudmaijer 2007).  Moreover, the relatively late spectral type of HD~92207 within the early-type class indicates that H$_\\alpha$ acts more nearly as a scattering line than a recombination line (Puls \\etal\\ 1998), thus making the star a prime candidate for exploring variable polarization across a line dominated by wind emission.  A description of the observations is presented in the following section.  An analysis of the data is provided in section~3.  A discussion of the implications of the results is given in section~4. ",
        "conclusions": "We have presented relatively medium resolution spectropolarimetric data of the highly variable A supergiant HD~92207.  It has been suggested by Kaufer \\etal\\ (1997) that the observed photometric and H$_\\alpha$ line variations are the result of a corotating structure in the wind, which they consider to be in the star's equatorial plane.  To explore that possibility, we begin with a consideration of the variable {\\em continuum} polarization.  We have developed a phenomenological model for Thomson scattering from an enhanced scattering region in a spherical wind in the form of a spiral pattern, similar in spirit to the work of Brown \\etal\\ (1994) who sought to model DACs in emission lines.  Our objective is to obtain a reference model for the global wind morphology that we can use for interpreting the H$_\\alpha$ polarization.  We consider a corotating stucture as a simple spiral that is top-bottom symmetric about the plane of the star's rotational equator.  Hence, the spiral structure has a guiding center that is always in the equatorial plane.  This center obeys an equation of motion for the radial wind velocity law and conservation of angular momentum in the rotating frame, so it follows a ``streak line'' (e.g., see Ignace, Bjorkman, Cassinelli 1998).  We adopt a standard ``$\\beta$-law'' for the radial flow:  \\begin{equation} v_{\\rm r} = v_\\infty\\,\\left(1-bu\\right)^\\beta, \\end{equation}  \\noindent where $u=R_\\ast/r$ and $b<1$ determines the radial speed of the wind at its base, with $v_0 = v_\\infty (1-b)^\\beta$.  \\begin{table}[t] \\centering \\caption[]{ Model Parameters \\label{tab2} } \\centering \\begin{tabular}{@{}cc@{}} \\hline\\hline Factor & Value \\\\ \\hline $R_\\ast^a$  & $223 R_\\odot$ \\\\ $v_{\\rm rot}\\sin i^a$  & 30 km s$^{-1}$ \\\\ $\\dot{M}^b$ & $1.3\\times 10^{-6}\\, M_\\odot$ yr$^{-1}$ \\\\ $v_\\infty^a$  & 275 km s$^{-1}$ \\\\ $\\beta^c$  & 1 \\\\ $b^c$  & 0.98 \\\\ $\\eta^c$  & 25 \\\\ $\\delta^c$  & $55^\\circ$ \\\\ $\\psi_0^c$  & $8^\\circ$ \\\\ $i_0^c$  & $70^\\circ$ \\\\ \\hline \\end{tabular} \\begin{center} {\\small $^a$ Przybilla \\etal\\ (2006); $^b$ Kudritzki \\etal\\ (1996); $^c$ best fit model parameters (see text).} \\end{center} \\end{table}  Velocity laws with $\\beta=1$ and $\\beta=2$ were considered, and the case of $\\beta=1$ produced a better match to the data.  In this case the location of the guiding center is given analytically with azimuth $\\varphi$ by  \\begin{eqnarray} \\varphi(u) & = & \\varphi_0(t) -\\frac{\\Omega\\,R_\\ast}{v_\\infty}\\,\\left\\{ \\frac{1-u}{u} +b\\ln\\left[\\frac{w(u)}{uw_0}\\right]\\right. \\nonumber \\\\ &  & \\left. -\\frac{1}{b}\\,\\ln\\left[\\frac{w(u)}{w_0}\\right]\\right\\}, \\end{eqnarray}  \\noindent where $w_0=v_0/v_\\infty$ and $\\Omega = 2\\pi/P$.  The density for the spiral-shaped perturbed region is treated as an {\\em excess} of density above the otherwise spherically symmetric wind. This density excess in the spiral pattern is taken to scale with the spherical wind density in our ``toy'' model, thus we conveniently parametrize the excess by a constant factor $\\eta = n_{\\rm excess}/ n_{\\rm sph}$, for $n_{\\rm sph}$ the spherical wind density.  In addition to the solution for the guiding center, we also need the cross-section of the spiral.  The cross-section (i.e., the intersection of the spiral pattern with a spherical shell) is treated as circular.  This spherical ``cap'' is axisymmetric, and so assuming the electron scattering is optically thin, the polarization from any given slice of the spiral is given by Brown \\& McLean (1978), along with the finite star depolarization correction factor of Cassinelli, Nordsieck, \\& Murison (1987).  Summing up contributions from all the caps yields the polarization from the structure as a function of rotational phase and viewing inclination for the rotation axis of the star.  Note that our model accounts for occultation of scattered light by the intervening star, but only in an approximate way.  We consider a slice as entirely occulted if its guiding center lies behind the star, and unocculted otherwise.  The principal model parameters are the density excess $\\eta$, the half-opening angle of the spiral $\\delta$, an orientation angle between the observer $Q-U$ axes and those of the star system $\\psi_0$, and finally the inclination of the rotation axis of the star $i_0$. Using a reduced chi-square evaluation for a grid of model polarization light curves, Table~\\ref{tab2} lists the model parameters that provide the best fit to the observed {\\em continuum} data.  The star and wind properties of HD~92207 are taken from Przybilla \\etal\\ (2006), except that the mass-loss rate is taken from Kudritzki \\etal\\ (1999) and does not account for clumping.  The most reasonable match to the observed continuum polarizations in the neighborhood of H$_\\alpha$ is shown in Figure~\\ref{fig3}.  The upper panel is for $P_Q$ and the lower one for $P_U$, displayed as percentage polarizations.  The rotational phases depend on the star's rotation period.  Given the radius and minimum rotation speed from Table~\\ref{tab2}, the maximum period is $P_{\\rm max} \\approx 376$\\,d.  The true period is $P_{\\rm rot} = P_{\\rm max}/\\sin i_0$, where $i_0$ is constrained from our model fitting.  As the ephemeris is not known, we assigned the rotational phase ``0'' to the date of our first observation, and the phases appearing in Figure~\\ref{fig3} represent values for our best fit model at $i_0 = 70^\\circ$.  For the model fitting, there are seven free parameters: the four listed above plus three offsets -- one for rotational phase, a vertical offset in $Q$, and an independent one in $U$.  With fourteen data points in total, the reduced chi-square for our best simultaneous fit to the $P_Q$ and $P_U$ lightcurves is 1.6.  That value is primarily a reflection of one discrepant data point, the first one in the observed $P_U$ ligthcurve. Our model is inherently smooth, whereas the line variability indicates the presence of variable wind structure.  For the continuum polarization, this wind structure can act as a source of ``noise''.  Better temporal and phase coverage is needed to interpret the polarimetric lightcurve more fully.  \\begin{figure}[t] \\centering{\\epsfig{figure=f3.eps, height=8.0cm}} \\caption{\\small Model $P_Q$ (top) and $P_U$ (bottom) light curves plotted for two rotational cycles based on the model described in the text with parameters listed in Tab.~\\ref{tab2}.  Observed data are plotted as points with error bars. The data points have been shifted horizontally and vertically to provide the best fits. \\label{fig3}} \\end{figure}  The non-zero variation of $P_U$ indicates that $i_0$ is not edge-on to the system, since no $\\Delta P_U$ could result for such a perspective when the circumstellar scattering environment is top-bottom symmetric. However, $\\Delta P_U$ is much less than $\\Delta P_Q$, suggestive that $i_0$ is not close to pole-on.  Since the interstellar polarization is not known, there is freedom to shift the observed points with respect to the model in both phase (i.e., $t_0$) and vertical offset to best match a model curve.  The vertical offsets for $P_Q$ and $P_U$ are independent, but the phase shift must be the same for both.  There is also freedom to rotate the model system relative to the observer $Q$-$U$ system, but we found the angle to be small, with $\\psi_0 = 5^\\circ$.  \\begin{figure*}[t] \\centering{\\epsfig{figure=f4.eps, height=11.0cm}} \\caption{\\small An overplot of the P Cygni lines against the $P_Q$ (dotted curve, or red for the online version) and $P_U$ (dashed curve, or blue for the online version) spectra using the format of Fig.~\\ref{fig2} to highlight the velocity locations of features. The H$_\\alpha$ line is normalized to the local continuum, and plotted as the solid curve using the scale on leftside axes; polarizations are shown as percent using the scale on the rightside axes. \\label{fig4}} \\end{figure*}  The match between the model and the continuum data appears to be best around an inclination of $i_0=70^\\circ$ with a half-opening angle of $\\delta=55^\\circ$.  A $55^\\circ$ opening angle corresponds to the Van Vleck angle and yields the maximum polarization when other parameters are held fixed, thus allowing us to minimize the value of the excess density. Because the spherical wind is relatively thin in electron scattering ($\\tau_{\\rm e} \\approx 0.014$ for solar abundances with ionized H and He neutral), a fairly large value of $\\eta=25$ is needed to match the observed $\\Delta P_Q$ and $\\Delta P_U$.  Such a large excess would need to be justified through physical modeling of the stellar wind and atmosphere.  Overall, the match is reasonably good except for the first data point in the lower panel for $P_U$ which we cannot reconcile to the model after exploring a range of opening angles, viewing inclinations, and $\\eta$ values.  Based on the work of Davies \\etal\\ (2005) for luminous blue variables, wind inhomogeneities are likely contributing a random contribution of unknown amplitude to the observed polarizations from HD~92207, which makes the interpretation of any globally coherent variation with phase challenging for such sparse sampling.  Even though our spiral model is simplistic, it does capture the essential points inherent to any corotating structure in terms of the generic trends of $P_Q$ and $P_U$ with phase, viewing inclination, and density.  It is certainly clear that much better sampling of the polarimetric light curve in the continuum is needed to critically assess the existence of a corotating region in the wind.  Turning to the H$_\\alpha$ line, the P Cygni morphology indicates that the line is dominated by the stellar wind.  In contrast, the H$_\\beta$ line is in absorption and shows no variable polarization.  At the spectral type of HD~97702, the $n=2$ level of hydrogen can act as the effective ground state (Puls \\etal\\ 1998) explaining why it shows such a strong P~Cygni shape.  It also means that line scattering polarization (e.g., Hamilton 1947) may influence the observed variable polarization in H$_\\alpha$.  For example, such effects were explored by Jeffery (1987, 1989) in the case of SN1987A.  Figure~\\ref{fig4} displays the polarized $P_Q$ and $P_U$ spectra in velocity along with the P Cygni profiles to emphasize where polarization changes occur within the line in relation to the absorption trough and emission peak.  The vertical scale for the measured percent polarizations are shown on the right axes, whereas the scale for the continuum normalized P Cygni lines is displayed on the left axes. For the first 5 panels, it is typically the case that either $P_Q$, $P_U$, or both vary most around the line core over an interval that is about a third of the wind terminal speed.  One clear exception is for March 26 that shows the strongest change in polarization across the line in $Q$ and aligns quite well with the redshifted emission peak, although it appears that $P_U$ actually increases in absolute value.  The observed polarizations are negative, so zero polarization is toward the top of each panel.  Curiously, there is a depolarization at the same location the night prior, except in $U$ instead of $Q$.  The overall depolarizations across the emission peak in the line relative to the continuum would seem to be a classic ``line effect''.  Polarization is normalized at each wavelength to the total emission.  If the line formation leads to unpolarized radiation that is largely unscattered by free electrons before emerging from the wind, then the additional line emission tends to reduce the polarization at those wavelengths relative to values outside the line.  In fact, this can be an excellent way of placing an upper limit to the interstellar polarization to the star.  However, the H$_\\alpha$ line of HD~92207 is perplexing because its behavior is not consistent.  Moreover, night-to-night variations in the H$_\\alpha$ polarization are hard to understand in terms of the spiral structure that we have considered for the continuum polarization given that the star's rotation period is of order a year.  It is worth pointing out that Harries (2000) modeled wavelength-dependent polarization across H$_\\alpha$ for the O supergiant $\\zeta$~Pup. His models did not include contributions to the polarization from line scattering; rather, he obtained changes of polarization across the line owing to the influence of line opacity for the polarization produced by Thomson scattering.  The application of Harries' models to HD~92207 is unclear.  There is the likelihood of a significant line scattering contribution to H$_\\alpha$ because of HD~92207's much later spectral type as compared to $\\zeta$~Pup. Other effects such as dynamic changes in the wind or variable stellar illumination (c.f., Al-Malki \\etal\\ 1999) will be needed to interpret the polarization across H$_\\alpha$.  For example, it would be useful to combine the methods of Li \\etal\\ 2000 and Davies \\etal\\ 2007 for polarimetric variability from wind clumping and electron scattering with the resonance scattering polarization effects explored by Ignace, Nordsieck, \\& Cassinelli (2004) for stationary winds.  With limited coverage of the rotational phase and only modest spectral resolution, the existing dataset is too poor to undertake a detailed calculation of the H$_\\alpha$ line to interpret the observed polarizations.  With future data, such an effort would be worthwhile because the continuum polarization constrains the global wind morphology whereas the H$_\\alpha$ line is sensitive to vector velocity flow."
    },
    "0812/0812.0304_arXiv.txt": {
        "abstract": "A new means of incorporating radiative transfer into smoothed particle hydrodynamics (SPH) is introduced, which builds on the success of two previous methods - the polytropic cooling approximation as devised by \\citet{Stam_2007}, and flux limited diffusion (e.g. \\citealt{Mayer}).  This hybrid method preserves the strengths of its individual components, while removing the need for atmosphere matching or other boundary conditions to marry optically thick and optically thin regions.  The code uses a non-trivial equation of state to calculate temperatures and opacities of SPH particles, which captures the effects of H$_{2}$ dissociation, H$^{0}$ ionisation, He$^{0}$ and He$^{+}$ ionisation, ice evaporation, dust sublimation, molecular absorption, bound-free and free-free transitions and electron scattering.  The method is tested in several scenarios, including: (1) the evolution of a \\(0.07\\,M_{\\odot}\\) protoplanetary disc surrounding a \\(0.5\\,M_{\\odot}\\) star; (2) the collapse of a \\(1\\,M_{\\odot}\\) protostellar cloud, and (3) the thermal relaxation of temperature fluctuations in a static homogeneous sphere. ",
        "introduction": "\\noindent Smoothed Particle Hydrodynamics (SPH) \\citep{Lucy,Gingold_Monaghan,Monaghan_92} is a Lagrangian method which represents a fluid by a distribution of particles.  Each particle is assigned a mass, position, internal energy and velocity: state variables such as density and pressure can then be calculated by interpolation - see reviews by \\citet{Monaghan_92, Monaghan_05}.  The effects of gravitation are also included as standard in most SPH codes.  Recently, many authors have constructed  versions of SPH with effects of radiative transfer \\citep{WB_1,Stam_2005_1,Stam_2005_2,Mayer}.  A full description of polychromatic 3D radiative transfer is currently not possible within SPH (at least while current computational limitations prevent it).  Even describing a snapshot from a simulation using full polychromatic radiative transfer is quite expensive \\citep{Stam_2005_1,Stam_2005_2}.  In the past, approximations to individual features of radiative transfer were used - for example the cooling time formalism: \\(\\dot{U} = \\frac{-U}{t_{cool}}\\), \\citep{Ken_1} which only describes energy loss from the system, and does not model transport of energy between neighbouring particles.  \\\\  \\noindent An example of where radiative transfer plays a fundamental role is in the physics of gravitational instabilities (GIs) in protoplanetary discs. The simple parametrisation of the cooling time method allowed these GIs to be probed and characterised effectively.  Gravitational fragmentation of protoplanetary discs is key to the disc instability theory of giant planet formation \\citep{Boss_science}.  However, it is disputed whether fragmentation can indeed occur, with strong debate between different groups using different methods of simulation and different formalisms for radiative transfer.  At the time of writing there is no strong consensus as to whether gravitational instability and fragmentation in protoplanetary discs can be a realistic mechanism for giant planet formation.  \\\\  \\noindent For fragmentation to occur, the disc must become gravitationally unstable, so that gravity can overcome pressure support and rotational support.  The disc can become gravitationally unstable to axisymmetric instabilities if \\citep{Toomre_1964}:  \\begin{equation} Q = \\frac{c_s \\kappa}{\\pi G \\Sigma} < 1 \\end{equation}  \\noindent where \\(c_s\\) is the local sound speed, \\(\\kappa\\) is the epicyclic frequency, and \\(\\Sigma\\) is the surface density of the disc.  In a Keplerian disc, the epicyclic frequency is replaced by the angular frequency, \\(\\Omega\\).  If the perturbation is nonaxisymmetric, the condition becomes \\(Q<1.5-1.7\\)  \\citep{Durisen_review}.  Further to this condition, the cooling of the gas must be efficient enough to radiate away energy gained by compression during the contraction.  Use of the cooling time formalism allowed a quantitative statement of this second condition: \\(t_{cool} \\leq  3\\Omega^{-1}\\) \\citep{Gammie,Ken_1,Mejia_2}. \\\\  \\noindent More recent efforts have used more sophisticated approximations to capture the equation of state (EoS) of the gas, and model realistic radiative cooling and radiation transport, using both SPH \\citep{WB_2,Mayer} and grid codes \\citep{Mejia_4, BDNL}, but these are becoming increasingly complex, with some methods requiring mapping of the photosphere (which is often of non-trivial geometric shape), and extra conditions to be applied there (matching atmospheres as in \\citet{Mejia_4}, or specifying cooling at the photosphere as in \\citet{Mayer}).  Also, identifying the photosphere often requires extra free parameters, the changing of which will affect the final results.  The latest radiative transfer approximations (such as \\citet{BDNL}, which solves the full radiative transfer equation explicitly in the vertical direction) are now attempting to remove this parametrisation. \\\\  \\noindent This paper presents a new radiative transfer approximation, which relies on two separate methods working in tandem.  The first method is the polytropic approximation devised by \\citet{Stam_2007}, which models the cooling of particles over a range of optical depths (\\(0 < \\tau \\lesssim  10^{11}\\)).  Its formulation ensures that cooling will be at its most efficient where the optical depth is around unity, in accordance with the definition of the photosphere.  This avoids the necessity of explicitly computing the location of the photosphere, or imposing boundary conditions upon it.  The second method is the flux-limited diffusion approximation, used by many authors \\citep{boden_lambda,CM_diffuse,WB_1,WB_2,BDNL,Mayer,Mejia_4,Boss_2008} to simulate radiation transport in optically thick regimes.  Although presented as an algorithm for particle-based codes, the hybrid algorithm can be applied to grid-based codes also.  This hybrid method captures all the physics of frequency-averaged radiative transfer, without relying on parametrisation.\\\\  \\noindent The paper is organised as follows. In section \\ref{sec:Method} the constituents of the new hybrid algorithm are presented; it is then shown how these methods are combined to create the new algorithm.  In section \\ref{sec:Tests}, the results from the following test scenarios are given: the evolution of a \\(0.07\\,M_{\\odot}\\) protoplanetary disc used as an example by \\citet{Mejia_1}, \\citet{Mejia_2}, \\citet{Mejia_3}, and \\citet{Mejia_4}; the collapse of a \\(1M_{\\odot}\\) molecular cloud \\citep{Masunaga_1}; and the thermal relaxation of a static sphere with seeded temperature fluctuations \\citep{Spiegel,Masu_98}.  Finally, in section \\ref{sec:Conclusions}, the method is summarised, and some indications of future work are given. ",
        "conclusions": "\\label{sec:Conclusions}  \\noindent This paper has presented a new means of modelling radiative transfer in SPH by fusing two well tested methods, polytropic cooling and flux-limited diffusion, in order that they may complement each other, and perform the functions that the other cannot.  By this fusion, the physics of three-dimensional frequency-averaged radiative transfer is captured without the need for complex boundary conditions, photosphere mapping or extra parameters.  Temperatures and opacities are obtained using a non-trivial equation of state which captures the effects of H$_{2}$ dissociation, H$^{0}$ ionisation, He$^{0}$ and He$^{+}$ ionisation, ice evaporation, dust sublimation, molecular absorption, bound-free and free-free transitions and electron scattering.  This data is tabulated pre-simulation for interpolation by the code. \\\\  \\noindent The algorithm is fast: only a 6\\% increase in CPU time is incurred in comparison to standard SPH simulations performed with a barotropic equation of state.  It has shown itself to be accurate in the tests outlined in the previous section: the evolution of a protoplanetary disc (with parameters proposed by Mej\\'{i}a, Boley, Cai et al \\citep{Mejia_1,Mejia_2}) from a uniform state through ring formation and contraction to instability; the complex thermal history of a collapsing molecular cloud (as studied by \\citet{Masunaga_1}); and the smoothing of temperature fluctuations in a homogeneous, static sphere \\citep{Spiegel,Masu_98}.  However, the scheme is still approximate, and can only partially describe radiative effects that occur over midrange distances (unlike the scheme proposed by \\citet{BDNL}, albeit in the vertical direction only).\\\\  \\noindent Comparisons with simulations using polytropic cooling alone have shown that the hybrid method is in effect only a small correction to the polytropic cooling method, which however can become important in some problems where temperature gradients and system geometries become complex (for example the protoplanetary disc simulations described in this paper).  \\\\  \\noindent Future work will see this algorithm applied to a variety of protostellar and protoplanetary environments, primarily a study of initial conditions for disc formation and evolution, as well as the effects of interactions between discs and binary companions."
    },
    "0812/0812.2906_arXiv.txt": {
        "abstract": "We have tested the two main theoretical models of bubbles around massive star clusters, Castor et al.~and Chevalier \\& Clegg, against observations of the well studied Carina Nebula.  The Castor et al.~theory over-predicts the X-ray luminosity in the Carina bubble by a factor of 60 and expands too rapidly, by a factor of 4; if the correct radius and age are used, the predicted X-ray luminosity is even larger.  In contrast, the Chevalier \\& Clegg model under-predicts the X-ray luminosity by a factor of 10.  We modify the Castor et al.~theory to take into account lower stellar wind mass loss rates, radiation pressure, gravity, and escape of or energy loss from the hot shocked gas.  We argue that energy is advected rather than radiated from the bubble. We undertake a parameter study for reduced stellar mass loss rates and for various leakage rates and are able to find viable models.  The X-ray surface brightness in Carina is highest close to the bubble wall, which is consistent with conductive evaporation from cold clouds. The picture that emerges is one in which the hot gas pressure is far below that found by dividing the time-integrated wind luminosity by the bubble volume; rather, the pressure in the hot gas is set by pressure equilibrium with the photoionized gas at $T=10^4\\K$. It follows that the shocked stellar winds are not dynamically important in forming the bubbles. ",
        "introduction": "Over a dynamical time only 2\\% of the gas in a typical disk galaxy is turned into stars \\citep{kennicutt98}, in spite of the fact that the disks are marginally gravitationally stable \\citep{kennicutt89}.  This puzzling result suggests that something other than gravity plays a (negative) role in star formation. Passive support against collapse and subsequent star formation can come from magnetic fields or turbulence. Examination of star forming galaxies and individual star forming regions commonly shows evidence for more active support against gravity, for example, expanding bubbles. Active mechanisms, which come under the rubric of `feedback', include energy or momentum input by active galactic nuclei, or from stars. Since most disk galaxies, including our own Milky Way, show little or no nuclear activity, and since bubbles in individual star forming regions are clearly not due to active galactic nuclei, in this paper we focus on feedback from stars.  Stars of type B2 and earlier supply several types of feedback, including protostellar jets, main-sequence stellar winds, radiation pressure, gas pressure associated with ionizing radiation, and supernovae. The mechanical energy imparted via a supernova is comparable to that injected by the star's stellar wind over its lifetime \\citep{cas75}. However, it is clear that many observed bubbles have formed before any of the stars in the central cluster have exploded in a supernova. Since we are interested in bubble formation, we will ignore supernovae, leaving open the question of whether supernovae provide the feedback necessary to limit the rate of star formation.  There are two competing theoretical models of stellar wind bubbles in the literature: that of Castor et al.~(i.e.~\\citet{cas75} and \\citet{wea77}) and that of \\citet{che85}, as implemented in, e.g., \\citet{ste03}.  In the theory explored by Castor, the stellar-wind shocked gas is confined by a cool shell of swept up interstellar medium (ISM).  In contrast, the Chevalier \\& Clegg theory ignores any surrounding material and simply has a steady-state wind.  The temperatures, pressures and predicted sizes of the X-ray emitting regions in the two theories are significantly different.  However, these differences can be masked if the stellar content of the bubble is not well constrained. Recent work by \\citet{smi06a} provides a detailed accounting of the number of early type stars in the well studied Carina nebula. Combined with the wealth of multi-wavelength observations available for Carina, this allows for a stringent test of the models.  Observations of star forming regions in the Milky Way show that the gas density is highly variable. In particular, the projected surface density is consistent with a log-normal distribution, e.g., \\citet{goodman,wong}. These observations are consistent with both analytic and numerical studies of supersonic turbulent flows, which predict log-normal density and column density distributions \\citep{passot,osg01}. These results call into question one of the fundamental assumptions of the Castor et al.~model, that the gas surrounding the central star or star cluster has a rather uniform density distribution. If the mass distribution is non-uniform, the shell swept up by the shocked stellar wind is likely to have holes. If so, this will lead to incomplete confinement of the hot gas and a consequent reduction in the pressure and associated X-ray emission of the hot gas. Such a situation would be intermediate between the Castor et al.~model of a completely confined wind and the Chevalier \\& Clegg picture of a free-flowing wind. We explore such intermediate models below.  This paper is organized as follows. In the next section we summarize some relevant observations of Carina, as well as some order of magnitude estimates of the pressures and forces acting on the surrounding cold, dense gas. In \\S \\ref{sec:methods} (and the appendices) we describe the Chevalier \\& Clegg model, the Castor et al. model, and our modifications to the Castor et al.~models. We also describe our numerical methods. In \\S \\ref{sec:results} we present the results of our modeling, including a parameter study of the (modified) Castor et al.~bubbles. In \\S \\ref{sec:globules} we discuss the X-ray emission. In \\S \\ref{sec:discussion} we discuss the results and compare to previous work. We give our conclusions in the final section. ",
        "conclusions": "We use both the Castor et al.~theory and the Chevalier and Clegg theory for bubble evolution around massive star clusters to try to understand observations of Milky Way and LMC bubbles.  The Castor theory over-predicts X-ray luminosity in the Carina bubble by a factor of 60 and expands too fast by a factor of 4.  In contrast, the Chevalier and Clegg model under predicts the X-ray luminosity by a factor of 20; since it is a steady state it has no expansion time.  These results suggest there is a partially confined hot, homogeneous interior but there are also holes in the shell through which the hot gas escapes like a free flowing wind.  We constructed a model in which the confining shell covered only a fraction $C_f$ of the sky as seen from the cluster, and allowed hot wind gas to escape at its sound speed from the uncovered portion of the shell. As an example one model consistent with observations of Carina is an isothermal sphere GMC model with $R_G = 60 \\pc$, $\\alpha = 0.2$, and $C_f = 0.5$, $R_b = 16.5 \\pc$, $L_x = 1.2 \\times 10^{34} \\ergs$, and $P_b = 1.1 \\times 10^{-10}\\,\\dynes \\cm^{-2}$.  This compromise is the best way to fit the many observations of the Carina nebula and is likely the best solution for all bubbles.  Our models suggest that the exact size, distribution and evolution of holes does not matter much, it is the weighted average of the fraction of total area that matters.  The bottom line is that the dynamics of X-ray bubbles are not strongly affected by the presence of luminous stellar winds; rather, they are controlled by a combination of the self-gravity of the surrounding GMC gas, radiation pressure, and H II gas pressure. This conclusion is in contrast to that of \\citet{cas75} and \\citet{shull}, but in agreement with \\citet{McKee84}. We argue that the X-ray fluxes of the bubbles are {\\em always} below those predicted by the Castor et al.~picture, properly interpreted, i.e., using the observed rather than predicted bubble radii. As noted by \\citet{Dorland86,DM87}, the X-ray observations show that the pressure in the hot gas is well below that predicted by \\citet{cas75}; unlike \\citet{Dorland86,DM87} and \\citet{McKee84}, we argue that this is a result of leakage of hot gas from the bubble, rather than leakage of energy in the form of radiation.  The results of the models in this paper are dependent upon the ISM model used, both the run of mean density, and the fluctuations in the surface density in different directions from the bubble center. According to both simulations and observations, the latter is distributed log-normally, meaning that in many directions the surface density will be far below the mean surface density. We suggest that hot gas escapes along these directions, which in our simplified models we treat as holes in the bubble wall.  Full hydrodynamical models are needed to check that the semi-analytical solutions shown in this paper are correct and to understand both the early evolution of bubbles and the development of holes in the shell."
    },
    "0812/0812.0903_arXiv.txt": {
        "abstract": "We report on a multi-epoch study of the Fornax dwarf spheroidal galaxy, made with the Infrared Survey Facility, over an area of about $42' \\times 42'$. The colour-magnitude diagram shows a broad well-populated giant branch with a tip that slopes down-wards from red to blue, as might be expected given Fornax's known range of age and metallicity. The extensive AGB includes seven Mira variables and ten periodic semi-regular variables. Five of the seven Miras are known to be carbon rich. Their pulsation periods range from 215 to 470 days, indicating a range of initial masses. Three of the Fornax Miras are redder than typical LMC Miras of similar period, probably indicating particularly heavy mass-loss rates. Many, but not all, of the characteristics of the AGB are reproduced by isochrones from Marigo et al. for a 2 Gyr population with a metallicity of Z=0.0025.  An application of the Mira period-luminosity relation to these stars yields a distance modulus for Fornax of 20.69 $\\pm 0.04$ (internal), $\\pm 0.08$ (total) (on a scale that puts the LMC at 18.39 mag) in good agreement with other determinations. Various estimates of the distance to Fornax are reviewed. ",
        "introduction": "The Fornax dwarf spheroidal galaxy ($\\l=237^{\\rm o}.24$, $b=-65^{\\rm o}.66$) is one of the most populous dwarf spheroidal companions to the Milky Way, second only to the disrupting Sagittarius dwarf. It shows evidence for a very extended history of star formation that ceased only a few million years ago (e.g. Gallart et al. 2005; Coleman \\& de Jong 2008) and a significant range of metallicity (e.g. Battaglia et al. 2006). It has its own system of five metal deficient globular clusters (Buonanno et al. 1999; Letarte et al. 2006), an extended asymptotic giant branch (AGB) with numerous carbon stars (e.g. Westerlund, Edvardsson \\& Lundgren 1987; Lundgren 1990) and evidence of substructure, possibly a consequence of merging with another galaxy (Coleman et al. 2004; Olszewski et al. 2006). The young stars are more concentrated towards the centre of the galaxy (Coleman \\& de Jong 2008) and there is an anti-correlation between the iron content and the velocity dispersion (Battaglia et al. 2006), suggestive of an age metallicity relation.  This paper is one of a series aimed at finding and characterizing luminous AGB variables within Local Group galaxies; it follows similar work on Leo I and Phoenix (Menzies et al. 2002, 2008). Here we report on multi-epoch $JHK_S$ photometry which enables us to identify the AGB variables within Fornax. The large-amplitude variables, generally known as Miras are of particular interest, first, because they tell us about the intermediate age population of which they are the most luminous representatives, secondly because they provide an independent distance calibration and thirdly because they will be major contributors to the processed material currently entering the interstellar medium of Fornax and are therefore an important source of enrichment.  Estimates for the interstellar extinction towards Fornax are mostly in the range $0.03<E(B-V)<0.1$ mag (e.g. Greco et al. 2007 and references therein). Here we use a value near the middle of this range, $E(B-V)=0.07$ mag, which amounts to $A_K<0.02$ and is therefore almost negligible in our work.  \\begin{figure} \\includegraphics[width=8.5cm]{pos.ps} \\caption{The fields covered by our survey and the distribution of the AGB variables and AGB candidates. Periodic variable stars are indicated as solid symbols: circles (Miras), squares (SRs with long term trends) and triangles (SRs). Variables with undetermined periods are shown as open triangles. Globular clusters (Hodge 1961) are shown as large crosses and the overdensity noted by Coleman et al. (2004) is shown as a bar.  The centroid of Fornax, from van den Bergh (2000), is at the centre of the image (RA: $\\rm 02^h\\ 39^m\\ 53^s$, Dec: $\\rm -34^o\\ 30'\\ 16''$ equinox 2000). The dashed lines mark the boundaries of our fields and the dotted line shows the region covered by the shallow survey of Gullieuszik et al. (2008).} \\label{fig_pos} \\end{figure} ",
        "conclusions": ""
    },
    "0812/0812.0242_arXiv.txt": {
        "abstract": "Stars can produce steady-state winds through radiative driving as long as the mechanical luminosity of the wind does not exceed the radiative luminosity at its base.  This upper bound on the mass loss rate is known as the photon-tiring limit.  Once above this limit, the radiation field is unable to lift all the material out of the gravitational potential of the star, such that only part of it can escape and reach infinity.  The rest stalls and falls back toward the stellar surface, making a steady-state wind impossible. Photon-tiring is not an issue for line-driven winds since they cannot achieve sufficiently high mass loss rates. It can however become important if the star exceeds the Eddington limit and continuum interaction becomes the dominant driving mechanism.  This paper investigates the time-dependent behaviour of stellar winds that exceed the photon-tiring limit, using 1-D numerical simulations of a porosity moderated, continuum-driven stellar wind.  We find that the regions close to the star show a hierarchical pattern of high density shells moving back and forth, unable to escape the gravitational potential of the star.  At larger distances, the flow eventually becomes uniformly outward, though still quite variable. Typically, these winds have a very high density but a terminal flow speed well below the escape speed at the stellar surface.  Since most of the radiative luminosity of the star is used to drive the stellar wind, such stars would appear much dimmer than expected from the super-Eddington energy generation at their core. The visible luminosity typically constitutes less then half of the total energy flow and can become as low as ten percent or less for those stars that exceed the photon-tiring limit by a large margin. ",
        "introduction": "The high luminosity of hot, massive stars leads to strong radiatively driven mass loss. In the more or less steady winds of O, B, and WR stars, the driving is through scattering of radiation by an ensemble of spectral lines (e.g., Castor, Abbott, \\& Klein 1975, hereafter \\cite{cak}). The tendency of these driving lines to become saturated with increasing density acts as an effective wind regulator, limiting the wind mass loss rates to ca.~$10^{-4} M_{\\odot}$/yr in even the most luminous stars. While potentially quite important for the star's evolution, this level of mass loss is energetically of only minor significance, with a total (kinetic and potential) mechanical luminosity of only a few percent of the ``photon tiring limit'' set by the total available stellar luminosity.   By contrast, the mass loss rates during the giant eruptive phase of Luminous Blue Variable (LBV) stars like $\\eta$~Carinae are inferred to be as high as $0.1-1 \\, M_{\\odot}$/yr, much higher than can be reasonably explained by models based on line-driving \\citep{so06}. In the case of $\\eta$~Carinae, the inferred average mechanical luminosity during its ca.\\ decade-long eruption in the 1840's is comparable to the estimated average radiative luminosity, $\\sim 2.5 \\times 10^{7}~L_{\\odot}$, during this time. This luminosity likely well exceeds the {\\em Eddington limit}, at which the radiative force associated with {\\em continuum} scattering by free electrons exceeds the inward force of gravity, implying then that the extreme mass loss can be driven by continuum, instead of line, opacity.  Unlike line driving, continuum driving does not suffer saturation at high densities, meaning then it can {\\em initiate} an {\\em arbitrarily large} mass flux, since all available photon-energy can be used to drive the matter. However, if the mechanical luminosity needed to lift this mass flux fully out of the gravitational potential exceeds the stellar luminosity, then the ``tiring'' or reduction in radiative energy flux will make it unable to sustain the outflow. This paper presents time-dependent hydrodynamical simulations of the complex combination of stagnated outflow and subsequent infall that occurs in continuum-driven models with initial, base flux that exceeds this photon-tiring limit.   Previous analyses of continuum-driven mass loss have focussed on developing steady-state wind outflow models with mass loss rates that approach, but do not exceed the photon-tiring limit. To provide the necessary driving regulation and limitation of the mass flux, such models invoke a  ``porosity'' moderation of the opacity associated with a clumped medium \\citep{s98,s01}. Such clumps can result from instabilities, such as the photon-bubble instability \\citep{s76,hak97, st99, b01, s01b} in the stellar atmosphere and the upper layers of the stellar envelope. These ``photon-bubbles'' are a phenomenon that occurs in super-Eddington atmospheres and can support large density inhomogeneities.  This reduces the continuum driving in the dense layers where clumps are optically thick, thus allowing the stellar base to be hydrostatically bound even though the luminosity is formally above the Eddington limit. But as the clumps become optically thin in the higher-level, lower-density layers, they are exposed to the full radiative driving of the assumed super-Eddington luminosity. The intermediate layer where the net outward force equals gravity then becomes the sonic point of a wind outflow, with the density at this point setting the wind mass loss rate. This identification of the sonic radius can be used to predict mass loss rates that are consistent with the observations of clear-cut super-Eddington objects \\citep{s01}.  For a given Eddington parameter $\\Gamma$, the mass loss initiated in such models depends on the nature of the assumed porosity. Specifically, within the ``power-law porosity'' formalism developed by  Owocki, Gayley \\& Shaviv (2004; hereafter \\cite{ogs04}), it depends on a power-index and the ratio of a characteristic ``porosity length'' to the local scale height (see \\cite{ogs04} and section~2.4 below). Analytic solutions derived by \\cite{ogs04} show that physically plausible porosity parameters combined with moderate super-Eddington parameters, $\\Gamma \\gtrsim 1$ can indeed lead to mass loss rates comparable to that inferred for the extreme LBV giant eruption of $\\eta$~Carinae, even approaching the ``photon-tiring'' limit.  A recent companion paper (van Marle, Owocki \\& Shaviv 2008; hereafter \\cite{mos08}) used numerical radiation hydrodynamics simulations to test these analytic, porosity-moderated wind solutions of \\cite{ogs04} for cases that do not exceed this photon-tiring limit. The results showed that the asymptotic steady-states of the numerical simulations are generally in quite close agreement with the analytic results. However, with only modest variations in the choice of the key porosity parameters, these analytic scaling laws also allow base mass fluxes that can {\\em exceed} the photon-tiring limit, with the consequence that the initial wind solution must necessarilly stagnate at some finite radius \\citep{mos08a},  thereby precluding a net, steady mass outflow.  The central thrust of the current paper is to explore the nature of the necessarily time-dependent flow that ensues from such a porosity model with a base mass flux above this tiring limit. In principle this should involve complex, three-dimensional (3-D) patterns of both outflow and infall. But given the computational and conceptual challenges in developing dynamically realistic 3-D description, we first explore here much simpler, albeit idealized one-dimensional (1-D) models. A key rationale for this approach is that the underlying phenomenon of photon-tiring is not inherently multi-dimensional, but rather depends largely on the {\\em radial} variation of driving vs. gravity. The insight gained in studying 1-D simulations of the complex, time-dependent combination of outflowing and inflowing layers will form a basis for future, more complete models that include multi-dimensional (2-D or even 3-D) flow.  As detailed in \\S 2 (and also the appendix), a key feature of our simplified method here is to maintain a global energy conservation that accounts both for the loss of radiative luminosity in regions of outward acceleration, as well as the re-energization of the luminosity from regions of infall and shock heating. The results in \\S 3 illustrate the complex pattern of inflow and outflow, including also the strong effect on the emergent radiative luminosity. We follow in \\S 4 with a discussion of what these results mean for understanding the nature of LBVs, and then conclude in \\S 5 with a summary of the work thus far and a brief discussion of directions for future work. ",
        "conclusions": "\\label{sec-concl}  It should be emphasized here that the photon-tiring effects that are the focus of this paper are only relevant for the most extreme forms of radiatively driven mass loss. For even the strongest, steady-state, line-driven winds from WR stars, the mass loss rates can range up to a few times $10^{-5} \\, M_{\\odot}$/yr and the flow speeds up to $v_{\\infty} \\approx 3000 \\,$km/s. This then implies a mechanical wind liminosity that is less than a few percent of a typical WR luminosity of ca. $10^{6} \\, L_{\\odot}$.  By contrast, the Homunculus nebulae that resulted from the 1840-1850 giant eruption of $\\eta$~Carinae is estimated to harbor at least $10 \\, M_{\\odot}$ with expansion speeds on order $600 \\,$km/s \\citep{s02}, implying a characteristic average mechanical luminosity that is quite comparable to the estimated radiative luminosity $2.5 \\times 10^{7} \\, L_{\\odot}$ during the epoch. This seems a strong indication that photon tiring may have played a key role in the driving of this extensive mass loss.  But there are some key differences between these properties of the $\\eta$~Carinae eruption and the 1-D photon-tired wind models computed here. For one, the  observed terminal flow speed of ca. 600~km/s is much higher than the ca. 50~km/s found for the tiring-limited models computed here. Moreover, in contrast to the rough parity between mechanical and radiative luminosity for $\\eta~$Carinae's giant eruption, these tiring models give mechanical luminosities that are several times higher than the emergent radiation, which itself is also quite variable.  However, as already noted, it seems quite likely that these general properties of 1-D photon-tiring could be quite different in full 3-D case. In particular, the likely lateral variations between regions of infall and outflow could allow a stronger, more regular escape of radiative flux. Indeed it might even resemble the extensive, complex 3-D clump structure envisioned in the parameterized porosity formalism. The net effect could be a self-regulation of porosity to keep the mass flux below the tiring limit, and thus allow the kind of quasi-steady, higher-speed wind outlfow obtained in both the analtyic models of \\cite{ogs04} and the numerical simulations of \\cite{mos08}.  The up-flow and down-flow nature of the mass flux close to the star is also reminiscent of convective motion, albeit driven by radiative luminosity rather than buoyancy. In the stellar interior, convection can effectively transport energy at such a rate that the star can locally exceed the Eddington limit \\citep{jso73}. But in surface layers it becomes inefficient due to the lower density \\citep{ogs04}. Similarly, the shock layers allow for mechanical energy advection to keep the photon-tired layer with a reduced luminosity. However, unlike convection where the flows are sub-sonic, the photon tired layer is composed of highly super-sonic motion.  Even in the 1-D framework, the present simulations assume an instantaneous nature for both the propagation and diffusion of radiation, and neglect optically thick backwarming effects in favor of a simple, constant temperature. Including an optically thick radiative diffusion with a finite speed would allow energy to be stored in the radiation field, with a likely smoothing effect on both luminosity variation and the resulting motions of mass shells. In 3-D or even 2-D such shells should break up into clumps. This could allow more photons to escape between the dense clumps, reducing the mass flux and increasing the radiative luminosity at the outer boundary. This higher luminosity at large radii could also accelerate the flow to higher terminal speeds."
    },
    "0812/0812.0597_arXiv.txt": {
        "abstract": "\\noindent Clusters of galaxies have not yet been detected at gamma-ray frequencies; however, the recently launched Fermi Gamma-ray Space Telescope, formerly known as GLAST, could provide the first detections in the near future. Clusters are expected to emit gamma rays as a result of (1) a population of high-energy cosmic rays fueled by accretion, merger shocks, active galactic nuclei and supernovae, and (2) particle dark matter annihilation. In this paper, we ask the question of whether the Fermi telescope will be able to discriminate between the two emission processes. We present data-driven predictions for the gamma-ray emission from cosmic rays and dark matter for a large X-ray flux limited sample of galaxy clusters and groups. We point out that the gamma-ray signals from cosmic rays and dark matter can be comparable. In particular, we find that poor clusters and groups are the systems predicted to have the highest dark matter to cosmic ray emission ratio at gamma-ray energies. Based on detailed Fermi simulations, we study observational handles that might enable us to distinguish the two emission mechanisms, including the gamma-ray spectra, the spatial distribution of the signal and the associated multi-wavelength emissions. We also propose optimal hardness ratios, which will help to understand the nature of the gamma-ray emission. Our study indicates that gamma rays from dark matter annihilation with a high particle mass can be distinguished from a cosmic ray spectrum even for fairly faint sources. Discriminating a cosmic ray spectrum from a light dark matter particle will be instead much more difficult, and will require long observations and/or a bright source. While the gamma-ray emission from our simulated clusters is extended, determining the spatial distribution with Fermi will be a challenging task requiring an optimal control of the backgrounds. ",
        "introduction": "Clusters and groups of galaxies are the largest gravitationally bound matter structures observed in the Universe. Although these objects are expected to host several high energy phenomena, the resulting electromagnetic non-thermal emission is far from being fully understood \\cite{brunetti,heclusters}. A hallmark of the occurrence of non-thermal phenomena in these large structures is the detection, in numerous clusters, of extended radio emission associated to the synchrotron losses of relativistic cosmic-ray electrons \\cite{brunetti,heclusters,Berrington:2002bh,Colafrancesco:1998us,Reimer:2004au}. The acceleration of cosmic rays in galaxy clusters can originate from a number of physical processes, including violent shocks produced in cluster-cluster mergers and the accretion of smaller structures \\cite{harris,sarazin,shocks,Keshet:2002sw}, the re-acceleration of cosmic rays injected by galactic sources like active galactic nuclei and supernovae \\cite{sarazin02}, and inelastic collisions of primary cosmic-ray protons  \\cite{dennison} producing showers of secondary particles, including relativistic electrons and positrons as well as gamma rays.  The radio emission from clusters, perhaps the most solid source of observational information on non-thermal phenomena in these  objects, broadly falls into two classes featuring different spatial distribution, polarization and emission location within the cluster. The first class,  {\\em radio relics}, is characterized by irregular radio morphologies and is typically located in external regions of the cluster \\cite{kempner}. The second class of diffuse cluster radio sources is that of {\\em radio halos}  \\cite{giovannini}, whose emission is typically centered on the cluster and follows a similar spatial distribution as e.g. the thermal X-ray emission from the intra-cluster medium (ICM) gas. Radio relics, unlike radio halos, exhibit prominent polarization, and are often associated to cluster regions thought to host shock activity. The origin of radio halos, instead, is far from being fully understood \\cite{brunetti,heclusters}. The upcoming generation of low-frequency radio arrays, including the Giant Metrewave Radio Telescope (GMRT), the LOw Frequency ARray for radio astronomy (LOFAR), the Murchison Widefield Array (MWA) and the Long Wavelength Array (LWA), will improve the current observational situation in the near future.  In the recent past, claims of hard X-ray emission from nearby cluster of galaxies have been reported \\cite{hardxray}. One possibility is that this X-ray emission originates from the inverse Compton (IC) scattering off background radiation of the same non-thermal high energy electron population responsible for the radio emission. Most of the clusters reportedly detected at hard X-ray frequencies are merging clusters \\cite{nevalainen2004}, and the detections themselves are still debated \\cite{rossetti04,Ajello:2008rd}. In addition, it is unclear if simple models for the non-thermal population fueling radio and hard X-ray emissions can self-consistently explain data from clusters (see e.g. the case of new radio data from the Ophiuchus cluster of galaxies \\cite{ophradio} in connection with the recent claim of a non-thermal hard X-ray detection with INTEGRAL \\cite{ophhxr}). Future hard X-ray missions, including the Nuclear Spectroscopic Telescope Array (NuSTAR) \\cite{nustar} and the International X-ray Observatory (IMX) \\cite{syx}, will play a key role in settling the mentioned controversies, as well as in fostering our understanding of processes underlying non-thermal activity in clusters \\cite{Ajello:2008rd}.  Gamma rays, covering the highest end of the electro-magnetic spectrum, can potentially provide essential information on high-energy phenomena in groups and clusters of galaxies. So far, these objects have not been detected in gamma-rays, and data from the EGRET telescope on board the Compton Gamma-Ray Observatory \\cite{egretref} have only produced upper limits \\cite{reimer} (see however \\cite{grdetect}). Statistically conclusive discoveries and a true revolution in our understanding of the highest energy phenomena in galaxy clusters are anticipated with the Fermi Gamma-ray Space Telescope, formerly known as GLAST \\cite{glastref}, which was successfully launched on June 11, 2008. The main instrument onboard Fermi, the Large Area Telescope (LAT), represents a remarkable jump in sensitivity compared to its predecessor EGRET \\cite{egretref}. The LAT also extends the EGRET energy range (20 MeV to 10 GeV) to much higher gamma-ray energies, up to about 300 GeV. Aside from probing high and ultra-high energy particle physics processes in astrophysical sources, the detection of gamma rays from clusters would also help to establish the properties of the primary proton population within clusters, and possibly clarify its role in various cluster phenomena \\cite{heclusters}. These include e.g. the question of the origin of radio halos \\cite{dennison}, the ``cooling-flow'' problem \\cite{fabian} and particle acceleration within cluster merger shocks \\cite{sarazin04}.  Clusters of galaxies are potentially powerful observational probes of cosmology (see e.g. \\cite{voit,HMH01, albrecht06}). In this context, the accurate understanding of non-thermal phenomena in clusters is crucial to their ultimate utility as cosmological probes. Specifically, the estimate of cluster masses, which is at the basis of all cosmological applications using clusters, frequently relies on the assumption of hydrostatic equilibrium between gravitational forces and the thermal pressure supplied by the ICM. The accuracy of hydrostatic mass determinations is therefore limited by our understanding of the non-thermal pressure provided by cosmic rays, turbulence and magnetic fields in the ICM \\cite{ensslin97,nagai2007}. The Fermi Gamma-Ray Space Telescope, as well as current and planned ground based atmospheric Cherenkov telescopes (ACTs) will be able to probe the energy density supplied by cosmic rays, and possibly the evolution with redshift of cosmic ray pressure, to an unprecedented level of accuracy \\cite{pfrommer2004,nagaiando}.  A more exotic possibility for non-thermal activity in galaxy clusters was first envisioned by Totani in Ref.~\\cite{totani}: the pair annihilation of weakly interacting massive particles (WIMP) constituting the dark matter halo. Following that seminal work, Colafrancesco, Profumo and Ullio calculated in \\cite{Colafrancesco:2005ji} the complete multi-frequency spectrum, for the case of the Coma cluster, resulting from dark matter annihilation. The emission spectrum extends from radio to gamma-ray frequencies, and includes the secondary emissions from the non-thermal electrons and positrons produced as final stable particles in dark matter annihilation events. In addition, Ref.~ \\cite{Colafrancesco:2005ji} also studied the heating of the ICM produced by the energy injected by dark matter annihilation, as well as the induced Sunyaev-Zeldovich signal. As far as indirect signals from particle dark matter annihilation, by far the best studied non-thermal radiative emission is the production of gamma-rays \\cite{Gunn:1978gr,Stecker:1978du}. Specifically, when two dark matter particles annihilate, gamma rays result both directly from loop-suppressed diagrams as well as from the subsequent decays or radiative emission (e.g. from final state radiation)  of standard model particles produced in the annihilation final state, like quarks, leptons and gauge and Higgs bosons. The resulting gamma rays have energies extending up to the kinematic limit set by the WIMP mass (the pair annihilation event occurs for highly non-relativistic dark matter particles), predicted to be in the 10-1000 GeV range in the best motivated models \\cite{dmreviews,kkdm} (see however \\cite{profumoheavy,profumolight}).  The LAT instrument onboard Fermi is a tremendous tool for the indirect search for particle dark matter with gamma rays \\cite{Baltz:2008wd,ourgc}. Of special relevance for the present study, the flux of gamma rays from WIMP dark matter annihilation in clusters of galaxies has been shown to be in principle large enough to be detectable by Fermi-LAT \\cite{Colafrancesco:2005ji,ophiuchus}. Dark matter annihilation does not only produce gamma rays, but also additional stable particle species, such as energetic electrons and positrons. These, in turn, produce synchrotron, IC and bremsstrahlung radiation, with unique spectral features. The multi-wavelength emission from dark matter annihilation was studied in detail in \\cite{gondolo,bertone0101134,aloisio0402588,baltzwai,coladraco,xrdwarf,ullioregis,haze1,haze2}, and specifically in clusters of galaxies in \\cite{Colafrancesco:2005ji,ophiuchus,colabullet}. Interestingly, it was demonstrated that if dark matter annihilation fuels to some appreciable degree either the radio emission in the Coma cluster \\cite{Colafrancesco:2005ji}, or the hard X-ray emission in the Ophiuchus cluster \\cite{ophiuchus}, Fermi is almost guaranteed to have the sensitivity to detect gamma rays from dark matter annihilation.   Upcoming gamma ray observations of clusters of galaxies have therefore profound implications for cosmology and, possibly, the discovery of New Physics. If Fermi detects gamma rays from clusters, a crucial point will be to conclusively assess the nature of the mechanism responsible for the emission. In the present study, we focus on how to tell apart gamma rays produced by standard, astrophysical mechanisms such as cosmic rays from those resulting from WIMP dark matter annihilation. On general grounds, we expect to have three handles to differentiate the two emission mechanisms: \\begin{itemize} \\item the {\\em gamma-ray spectrum}: models for cosmic ray production of gamma rays predict a flux as a function of energy which differs from what expected out of dark matter annihilation. The low photon statistics represents a challenge for meaningful discrimination based on the gamma-ray spectrum. We thus study the best angular and energy range, and propose a technique based on hardness ratios that will help to discriminate cosmic rays from dark matter \\item the {\\em spatial distribution}: depending upon assumptions on the dark matter substructure distribution and density profile, as well as on the primary cosmic ray source distribution, we predict that clusters can appear to be extended gamma-ray sources. We study whether this can be used to differentiate gamma rays emitted by dark matter from those produced by cosmic rays \\item the {\\em multi-wavelength emission}: comparing the results of hydrodynamical simulations of cosmic rays in clusters to our predictions for a dark matter scenario, we find that the ratio of the hard X-ray to gamma-ray emission is a potential diagnostic to understand the origin of non-thermal phenomena in clusters. \\end{itemize} In our simulations and analysis, we use the latest LAT instrumental response function and observation strategy and the Fermi Science Tools software package which is currently being used to analyze Fermi data.   As an application of our theoretical study, we present predictions for gamma-ray fluxes from a large X-ray limited sample of 130 nearby groups and clusters of galaxies. Specifically, our predictions are based on X-ray data, and on a fixed set of assumptions for both dark matter and cosmic rays. This also allows for a meaningful comparison of the dark matter to cosmic ray induced emission at gamma-ray frequencies. We present a ranking of plausible candidates where one might expect a bright gamma-ray signal, and of sources where the dark matter contribution is expected to be stronger compared to the one fueled by astrophysical cosmic rays. In particular, we discovered that low-redshift groups are the most promising class of objects to search for a dark matter signal from distant extra-galactic systems.  The organization of our paper is as follows. The following Section \\ref{sec:method} illustrates the model we use to compute the gamma-ray emission resulting frtom cosmic rays (\\ref{sec:cr}) and from dark matter annihilation (\\ref{sec:dm}), and gives details on the Fermi simulation setup (\\ref{sec:glastsim}). Section \\ref{sec:origin} discusses how to study the origin of gamma rays from clusters, including our analysis of the optimal angular region, the spectra, hardness ratios, spatial extension and multi-wavelength counterparts. We present in Section \\ref{sec:cat} our predictions and ranking of nearby clusters and groups according to their gamma-ray emission (the complete list is provided in the Appendix). Section \\ref{sec:concl} gives a discussion and summary of our results, and concludes. ",
        "conclusions": "\\label{sec:concl}  The detection of gamma rays from clusters of galaxies might be a milestone in the scientific legacy of the Fermi Gamma-ray Space Telescope. Clusters are known to host a variety of high-energy phenomena that could fuel cosmic rays, producing, in turn, gamma rays as a result of collisions with the ICM gas. Being the largest bound dark matter structures, it is also reasonable to expect that clusters  feature a significant gamma-ray emission from the annihilation of dark matter particles. The scope of the present theoretical study was to investigate how to tell apart these two potential mechanisms of gamma-ray production with data from Fermi.  One potential source of gamma-rays from cluster regions that we have not considered here are the bright, central AGN known to be present in some clusters (NGC 1275 in Perseus and M87 in Virgo), which could inhibit our ability to detect faint extended gamma-ray emission.  However, this source of contamination is not expected from some of the best candidate, nearby clusters, like the Coma cluster, where strong radio galaxies are not observed.  The best method to detect diffuse gamma-ray emission, from cosmic rays or dark matter annihilation, will be to concentrate on those clusters lacking bright AGN.  For clusters with central AGN, the expected extended nature of the cluster gamma-ray emission may reveal this component, and the well-known multiwavelength (radio and X-ray) AGN spectra will allow modeling of the contribution of these sources to the gamma-ray emission.  By making simple analytical assumptions on the cosmic ray spectra and source distribution and on the dark matter particle properties and density distribution, we proposed a set of various benchmark models for both gamma-ray production mechanisms under investigation. We believe the set of models we considered is representative of the variety of possibilities one can realistically expect to encounter in clusters of galaxies. We then proceeded to simulate the expected 1 and 5 years gamma-ray signals with the Fermi Science Tools, and analyzed our results.  We summarize below the main results of the present theoretical study. \\begin{itemize} \\item We find that for cosmic rays the absolute signal to noise peaks at around one degree, while for dark matter it can peak at larger regions of interest, possibly as large as 3 degrees. Since the Fermi instrument response function forces us to consider relatively large angular regions to include low energy photons, we conclude that the best region of interest for gamma-ray studies of cluster of galaxies is around 3 degrees. \\item The spectral analysis of the simulated signal for our benchmark models shows that gamma ray emission from dark matter annihilation with a relatively large dark matter particle mass ($m_{{\\rm WIMP}}>50$ GeV) can be distinguished from a cosmic ray spectrum even for fairly faint sources.  Distinguishing a cosmic ray spectrum from a low dark matter particle mass appears to be more challenging, and would require deep data and/or a bright source. \\item We defined optimal energy bands for a simple hardness ratio, (H -- S)/(H + S), estimate of the spectrum.  The energy bands we propose are S $= 0.15 \\leq E < 0.7$ GeV and H $= E \\geq 0.7$ GeV. The hardness ratio is correlated to the nature of the emission and is predicted to be negative for a cosmic ray emission and positive for a dark matter annihilation signal.  Similar to the full spectral analysis, the hardness ratios for a simulated low mass dark matter particle model and a hard cosmic ray model are similar within the statisitcal errors. However, it is possible to use this simple two energy band ratio to distinguish cosmic ray and high mass dark matter particle models. \\item The level of the gamma-ray emission from clusters produced by cosmic rays can be comparable to that from dark matter annihilation. In the case of a mix of the two emissions, we find that for bright enough sources, if the dark matter contribution to the cluster flux is significant and the particle mass is not very low ($m_{WIMP} > 40$ GeV) the presence of a dark matter component can be seen even in the presence of a significant gamma-ray flux from cosmic rays.  However, tight constraints on the model parameters ($m_{WIMP}$ and $\\alpha_p$) may be problematic unless the dark matter particle mass is known. \\item Our cluster gamma-ray simulated emissions appear, after data reduction, as extended rather than point sources, with extensions which depend on the spatial models and on the emission mechanism.  However, determining the spatial distribution with Fermi will be a challenging task requiring an optimal control of the backgrounds. \\item We showed that the ratio of the integrated gamma-ray to hard X-ray flux in galaxy clusters can be used as a diagnostics for the discrimination of the origin of non-thermal phenomena. The generic expectation is that a dark matter induced signal would produce a brighter hard X-ray emission as opposed to cosmic rays, for a given gamma-ray flux. \\item We presented X-ray data-driven predictions for the gamma-ray flux from 130 clusters and groups of galaxies in the HIGFLUCS and GEMS catalogs. We found that the clusters with the brightest gamma-ray emission from cosmic rays include the Perseus, Coma, Ophiuchus, Abell 3627 and Abell 3526 clusters; the most luminous clusters in dark matter emission are predicted to be the Fornax group, Ophiuchus, Coma, Abell 3526 and Abell 3627. \\item We discovered that the objects with the largest dark matter to cosmic ray gamma-ray luminosity in our sample are groups and poor clusters. In particular, the highest ratios are associated to the groups NGC 5846, 5813 and 499, M49 and HCG 22. Of these, M49 ranks overall 6th/130 in terms of predicted dark matter induced gamma-ray emission, NGC 5846 ranks 12th and NGC 5813 15th; Fornax also has a relatively high dark matter to cosmic ray gamma-ray flux, and is the brightest object in dark matter emission. All these objects are very promising candidates for a search for gamma-ray emission with the Fermi telescope that could potentially be related to particle dark matter: specifically, as shown in the Appendix, we predict, for the DM\\_HS setup, an integrated flux of photons above 0.1 GeV of $4\\times 10^{-9}\\ {\\rm cm}^{-2}{\\rm s}^{-1}$ for Fornax and of  $2\\times 10^{-9}\\ {\\rm cm}^{-2}{\\rm s}^{-1}$ for M49, both around the projected Fermi point-source sensitivity for one year of data. With the same setup, we predict a flux of $0.8\\times 10^{-9}\\ {\\rm cm}^{-2}{\\rm s}^{-1}$ and $0.5\\times 10^{-9}\\ {\\rm cm}^{-2}{\\rm s}^{-1}$ for NGC 5846 and 5813, respectively, which would make them detectable or marginally detectable sources for Fermi within the anticipated mission lifetime. \\end{itemize}"
    },
    "0812/0812.4182_arXiv.txt": {
        "abstract": "{In string gas cosmology, extra dimensions are stabilised by a gas of strings. In the matter-dominated era, competition between matter pushing the extra dimensions to expand and the string gas pulling them back can lead to oscillations of the extra dimensions and acceleration in the visible dimensions. We fit this model to supernova data, taking into account the Big Bang Nucleosynthesis constraint on the energy density of the string gas. The fit to the Union set of supernova data is acceptable, but the fit to the ESSENCE data is poor. } ",
        "introduction": "\\paragraph{String gas cosmology.}  String gas cosmology is a cosmological scenario motivated by string theory \\cite{Kripfganz:1988, Brandenberger:1989, Tseytlin:1991} (see \\cite{Battefeld:2005b, Brandenberger:2008} for reviews and \\cite{Mathur:2008} for another perspective). In the original formulation of string gas cosmology, all spatial dimensions are treated on an equal footing: they are all toroidal and start out at the string size. The aim is that dynamical processes in the early universe will allow only three dimensions to expand to macroscopic size, while the extra dimensions are stabilised at the string size by a gas of strings. Assuming that the dilaton is stabilised by some other mechanism, the string gas can stabilise the extra dimensions during the radiation-dominated era \\cite{Patil:2004, Patil:2005a} (see also \\cite{Kripfganz:1988, Watson:2003a, Berndsen:2004, Brandenberger:2005a, Patil:2005b, Berndsen:2005, Kanno:2005, Danos:2008, Sano:2008}). However, when the universe becomes matter-dominated, the matter will push the extra dimensions to open up \\cite{Kripfganz:1988, Patil:2004, Patil:2005a, Weiss:1986}. It was shown in \\cite{Ferrer:2005} that the gas of strings can still prevent the extra dimensions from growing too large, but they cannot be completely stabilised. There is a competition between the push of matter and the pull of strings. If the number density of the strings is too small, the extra dimensions will grow to macroscopic size. If the strings win, the size of the extra dimensions will undergo damped oscillations around the self-dual radius. The oscillations between expansion and contraction of the extra dimensions induce oscillations in the expansion rate of the large dimensions, which can involve alternating periods of acceleration and deceleration \\cite{Ferrer:2005}. (This kind of a mechanism has also been studied in \\cite{Perivolaropoulos}.)  Since the oscillations can start only after the universe becomes matter-dominated, they provide an in-built mechanism for late-time acceleration in string gas cosmology, one that alleviates the coincidence problem in a manner similar to scaling and tracker fields \\cite{scaling, tracker}. The mechanism is based on ingredients already present in string gas cosmology and does not require adding new degrees of freedom or turning on new interactions. However, the oscillating expansion history is quite different from the \\LCDM model which is known to be a good fit to the observations. (For comparison of some oscillating models to observations, see \\cite{osc, Dutta:2008}.)  We compare the model studied in \\cite{Ferrer:2005} to observations of type Ia supernovae (SNe Ia), using the Big Bang Nucleosynthesis (BBN) constraint on new radiation degrees of freedom. In section 2 we describe the string gas model, and in section 3 we fit the model to the Union and ESSENCE sets of SNIa data. The distances predicted by the model provide an acceptable fit to the Union data, but the fit to the ESSENCE data is poor. In section 4 we summarise our results and discuss how to make the model more realistic. ",
        "conclusions": "\\label{sec:disc}  \\paragraph{Conclusions.}  We have studied the late-time acceleration due to the oscillations of extra dimensions in string gas cosmology in the simple model discussed in \\cite{Ferrer:2005}. We have fitted the expansion history to the Union and ESSENCE sets of type Ia supernovae. The string gas model does not fit the SNIa data as well as the \\LCDM model. With the Union SNIa data, the difference in the goodness-of-fit is small, but the fit to the ESSENCE data is poor. Also, the best-fitting string gas models have a significant fraction of the energy density during BBN in the string gas, and taking into account the BBN constraint on new radiation degrees of freedom makes the fit worse. Further, we have considered a rather conservative BBN limit, allowing for neutrino chemical potential. Since the best-fit model is at the boundary of the region allowed by BBN, we would expect the fit to become worse as the BBN limit becomes more stringent. In the model, the matter density is also too high and the Hubble parameter today somewhat low, so taking further observational constraints into account would be likely to degrade the fit further. In any case, the model can still provide a stabilisation mechanism for the extra dimensions during the matter dominated era, for which it was originally introduced.  Leaving aside the physical origin of the oscillations and the constraint from BBN, the model demonstrates how an expansion history which is very different from the \\LCDM model, with strong oscillations of the Hubble parameter, can still provide a good fit to the supernova data. (In this context, it may be interesting that the Hubble parameter inferred from observations of the ages of passively evolving galaxies shows oscillations \\cite{ages}, though it is premature to draw strong conclusions from the data.) The fit only becomes poor when the change in the expansion rate-distance relationship due to the extra dimensions is taken into account. This in turn is a concrete example of a model where this FRW consistency condition, discussed in \\cite{Clarkson:2007}, is strongly violated.  \\paragraph{Improving the model.}  As discussed in \\cite{Ferrer:2005}, the energy-momentum tensor for the string gas is expected to be more complex than \\re{rhos}--\\re{Ps}. The energy density \\re{rhos} corresponds to a gas of strings which all have the same momentum $M l_s/a$ in the visible dimensions, while a realistic gas would have a distribution of strings with different momenta. The evolution of terms with different values of $M$ is qualitatively the same: they scale like radiation in the radiation-dominated era and start tracking the matter during the matter-dominated era until the onset of oscillations. However, the different terms will lead to quantitatively slightly different oscillations, and as we have seen, the evolution is very sensitive to the parameters of the string gas. In order to explore this possibility, we would have to know the distribution of string momenta, which depends on how the string gas was created in the early universe and whether it has thermalised.  \\ack  TM gratefully acknowledges support from the Academy of Finland, project no. 111953. \\\\  \\appendix"
    },
    "0812/0812.1591_arXiv.txt": {
        "abstract": "In the current strategy of microlensing planet searches focusing on high-magnification events, wide and close binaries pose important sources of contamination that imitates planetary signals.  For the purpose of finding systematic differences, we compare the patterns of central perturbations induced by a planet and a binary companion under severe finite-source effect.  We find that the most prominent difference shows up in the morphology of the edge features with negative excess that appear at the edge of the circle with its center located at the caustic center and a radius equivalent to the source radius.  It is found that the feature induced a binary companion forms a complete annulus, while the feature induced by a planet appears as several arc segments.  This difference provides a useful diagnostic for immediate discrimination of a planet-induced perturbation from that induced by a binary companion, where the absence of a well-developed dip in the residual from the single-lensing light curve at both or either of the moments of the caustic center's entrance into and exit from the source star surface indicates that the perturbation is produced by a planetary companion. We find that that this difference is basically caused by the difference between the shapes of the central caustics induced by the two different types of companions. ",
        "introduction": "Microlensing planets are being discovered at an accelerating rate and the total number of detections now reaches 8 \\citep{bond04, udalski05, beaulieu06, gould06, gaudi08, dong08, bennett08}. The microlensing signal of a planet is a brief perturbation to the smooth standard light curve of the primary-induced lensing event occurring on a background star. To achieve the monitoring frequency required to detect short-duration planetary signals, current planetary lensing searches are being conducted by using a combination of survey and follow-up observations, where alerts of ongoing lensing events are issued by the survey observations (OGLE: Udalski et al.\\ 1994, MOA: Bond et al.\\ 2001) and the alerted events are intensively monitored by  the follow-up observations (Micro-FUN: Dong et al.\\ 2006-, PLANET: Kubas et al.\\ 2008).  However, the number of telescopes available for follow-up observations is far less for intensive monitoring of all alerted events \\citep{dominik08} and thus observations are focused on events which can maximize the planet detection probability with follow-up observations. Currently, the prime targets of follow-up observations are high-magnification events for which the source trajectories always pass close to the perturbation region around the central caustic induced by a planet and thus planet detection efficiency is intrinsically high \\citep{griest98}.   Although high-magnification events yield high sensitivity to planets, interpretation of the observed perturbation often suffers from several potential degeneracies.  There are two major degeneracies causing this complication.  The first well-known wide/close degeneracy arises due to the fact that the perturbation induced by a planet with a projected separation in units of the Einstein ring, $s$, is very similar to the perturbation induced by a planet with a separation $1/s$ \\citep{dominik99, an05, chung05}.\\footnote{We note that the wide/close degeneracy occurs not only for planetary events but also for binary events in general \\citep{dominik96}.} The other degeneracy arises due to the fact that a central perturbation of a high-magnification event can also be produced by a very close ($s\\ll 1$) or a very wide ($s\\gg 1$) binary with roughly equal mass components. Hereafter, we refer the latter degeneracy as the `planet/binary' degeneracy. In the sense that the planet/binary degeneracy causes indeterminacy of both of the mass ratio and the separation between the lens components, while the wide/close degeneracy causes ambiguity only in the separation, the planet/binary degeneracy poses more serious problem in the interpretation of an observed perturbation. Fortunately, the perturbations induced by a planet and a binary are intrinsically different and thus it would be possible to determine whether the perturbation is caused by a planet or a binary. However, distinguishing between the two interpretations usually requires detailed modelling which demands time-consuming search for a solution in the vast space of many parameters. Therefore, a simple diagnostic that can resolve the planet/binary degeneracy would be very helpful not only for the prompt interpretation of an observed perturbation but also for the establishment of the observational setup  optimizing the coverage of the characteristic features helping to resolve the degeneracy.  \\citet{han08a} provided such a diagnostic but this diagnostic applies to a subset of light curves with double-peak features.    Recently, \\cite{dong08} reported a planet detected  from the analysis of a new type of high-magnification event where the angular extent of the perturbation region induced by the planet is significantly smaller than the angular size of the source and thus finite-source effect is very severe.  The main perturbation features of this event are double spikes in the residuals, which are approximately centered at the times when the lens enters and exits the source.  In 2008 season, several more such events were detected (A.\\ Gould 2008, P.\\ Fouque 2008, private communication), implying that events with these features might be common.  From detailed investigation, \\citet{han08b} found that spike features commonly appear for planetary-lensing events affected by severe finite-source effect and thus they can be used for the diagnosis of the existence of a planet. However, they noted that such features can also be produced by a wide or a close binary companion and thus the existence of the feature does not necessarily confirm the planetary interpretation of the perturbation.   In this paper, we investigate systematic differences between the patterns of central perturbation induced by a planet and a binary companion.  From this,  we identify a diagnostic that can be used to immediately distinguish between the planetary and binary interpretations. We also investigate the origin of the difference. ",
        "conclusions": "We compared the patterns of central perturbations induced by a planet and a binary companion under severe finite-source effect.  We found that the most prominent difference shows up in the morphology of the edge features with negative excess, where the feature induced a binary companion forms a complete annulus, while the feature induced by a planet appears as several arc segments.  This difference provides a useful diagnostic for immediate discrimination of a planet-induced perturbation from that induced by a binary companion, where the absence of a well-developed dip in the residual from the single-lensing light curve at both or either of the moments of the caustic center's entrance into or exit from the source star surface indicates that the perturbation is produced by a planetary companion.  We found that that this difference is basically caused by the difference between  the shapes of the central caustics induced by the two different types of companions."
    },
    "0812/0812.0979_arXiv.txt": {
        "abstract": "GRB\\,060505 and GRB\\,060614 are nearby long-duration gamma-ray bursts (LGRBs) without accompanying supernovae (SNe) down to very strict limits. They thereby challenge the conventional LGRB-SN connection and naturally give rise to the question: are there other peculiar features in their afterglows which would help shed light on their progenitors? To answer this question, we combine new observational data with published data and investigate the multi-band temporal and spectral properties of the two afterglows. We find that both afterglows can be well interpreted within the framework of the jetted standard external shock wave model, and that the afterglow parameters for both bursts fall well within the range observed for other LGRBs. Hence, from the properties of the afterglows there is nothing to suggest that these bursts should have another progenitor than other LGRBs. Recently, {\\it Swift}-discovered GRB\\,080503 also has the spike $+$ tail structure during its prompt $\\gamma$-ray emission seemingly similar to GRB\\,060614. We analyse the prompt emission of this burst and find that this GRB is actually a hard-spike $+$ hard-tail burst with a spectral lag of 0.8$\\pm$0.4 s during its tail emission. Thus, the properties of the prompt emission of GRB\\,060614 and GRB\\,080503 are clearly different, motivating further thinking of GRB classification. Finally we note that, whereas the progenitor of the two SN-less bursts remains uncertain, \\emph{the core-collapse origin for the SN-less bursts would be quite certain if a wind-like environment can be observationally established, e.g, from an optical decay faster than the X-ray decay in the afterglow's slow cooling phase}. ",
        "introduction": "Gamma-ray bursts (GRBs) are the most luminous explosions in the Universe since the Big Bang. They fall into two (partially overlapping) populations according to their observed duration: $\\gamma$-ray durations (measured as the time in which 90\\% of the fluence is emitted) longer than 2~s are defined as long GRBs (LGRBs) while bursts with duration shorter than 2~s are defined as short GRBs (SGRBs; \\citealt{Kouveliotou93}). It is widely accepted that at least the majority of LGRBs are driven by the collapse of massive stars (e.g., \\citealt{WB06}), although some LGRBs may be generated by the merger of compact objects (e.g., \\citealt{KR98,RRD03}). The strongest evidence for the collapsar scenario is the detection of bright Ic SN component photometrically and spectroscopically associated with nearby LGRBs such as GRB\\,980425, GRB\\,030329, GRB\\,031203, and XRF\\,060218 (\\citealt{Galama98,Hjorth03,Malesani04,Sollerman06,Pian06,Campana06}). On the other hand, SGRBs may be powered by the merger of binary compact objects (e.g. \\citealt{Eich89,Narayan92}). This connection is observationally bolstered by the association of some SGRBs with old stellar populations and lack of accompanying bright SN components in cases such as GRB\\,050509B (\\citealt{Hjorth05a}), GRB\\,050709 (\\citealt{Fox05,Hjorth05b}) and GRB\\,050724 (\\citealt{Berger05,Malesani07}). However, challenging this simple picture some SGRBs displayed violent X-ray flares occurring at least $\\sim100$ s after the triggers, e.g., GRB\\,050709 (\\citealt{Fox05}) and GRB\\,050724 (\\citealt{Bart05,Campana06,Malesani07}). This suggests long-lasting activity of the central engine and hence the current understanding of the GRB progenitor mechanism may be too simple (e.g., \\citealt{Fan05,Dai06,Ross07}). There is also evidence for activity of the inner engine on much longer time-scales (several days) for GRB\\,050709 (\\citealt{Watson06}) and GRB\\,070707 (\\citealt{Piranomonte08}).  The whole picture became more complicated after the discovery of GRB\\,060505 and GRB\\,060614, because both bursts are nearby LGRBs according to the conventional taxonomy but they are observationally not associated with SNe down to very strict limits (\\citealt{Fynb06,DellaValle06,Gal-Yam06}). In this sense, they share the expected observational properties of both conventional LGRBs and SGRBs.  GRB\\,060505 had a fluence of $(6.2\\pm1.1)\\times10^{-7}$~erg~cm$^{-2}$ in the $15-150$~keV band, and a $T_{90}$ duration of $4\\pm1$~s (\\citealt{Hullinger06,McBreen08}). It was found to be associated with a bright, star-forming \\ion{H}{2} region within its host galaxy at $z=0.089$ (\\citealt{Ofek07,Thoene08}). In the compact-star merger scenario the diameter of the \\ion{H}{2} region and the location of the GRB within it suggest that the delay time from birth to explosion of GRB\\,060505 was $\\lesssim 10$ Myr. This is marginally matching the lower limit of the delay-time region for SGRBs (\\citealt{Ofek07}). On the other hand, the age of the \\ion{H}{2} region of $\\sim 6$ Myr (\\citealt{Thoene08}) is consistent with the expectation for core-collapse in a massive star. The prompt emission of GRB\\,060614 consisted of a hard-spectrum component lasting $\\sim5$~s followed by a soft-spectrum component lasting $\\sim100$~s. \\cite{Mangano07} reported a photon index $\\Gamma=1.63\\pm 0.07$ for the time interval [-2.83,-5.62] s since the BAT trigger ($\\chi^2/{\\rm dof}=48.2/56$) and $\\Gamma=2.21\\pm0.04$ for 5.62-97.0 s since the BAT trigger ($\\chi^2/{\\rm dof}=40.9/56$). The fluences in the two components are $(3.3\\pm0.1)\\times10^{-6}$~erg~cm$^{-2}$ and $(1.69\\pm0.02)\\times10^{-5}$~erg~cm$^{-2}$ in the 15-350 keV band, respectively (\\citealt{Gehrels06}). Its host galaxy has a redshift of $z=0.125$ (\\citealt{PBF06}). The host of GRB\\,060614 is very faint with an absolute magnitude of about $M_B=-15.3$ (\\citealt{Fynb06}). Its specific star-formation rate is quite low, but within the range covered by LGRB hosts (similar, e.g., to that of the GRB\\,050824 host galaxy, \\citealt{Sollerman07}).  The spectral lag has been invoked as a quantity that can be used to classify bursts such that SGRBs have zero lag and LGRBs fall on a well defined lag-luminosity relation \\citep{Norris00,Norris06}. For GRB\\,060614 \\cite{Gehrels06} found that the spectral lags for the hard and soft components in the prompt $\\gamma$-ray emission are both consistent with zero lag, falling entirely within the range for SGRBs. For GRB\\,060505, on the other hand, \\cite{McBreen08} found using the {\\it Suzaku}/WAM and {\\it Swift}/BAT data that the spectral lag for the prompt $\\gamma$-ray emission is $0.36\\pm0.05$~s, consistent with a LGRB identity. Furthermore, lags of LGRBs and SGRBs, regardless of their physical origins, appear to overlap quite significantly according to statistics of 265 {\\it Swift} bursts (see Fig. 1 in \\citealt{Bloom08}). Alternatively, the so-called Amati relation can be used to provide hints on which class GRB\\,060505 and GRB\\,060614 belong to. According to this relation derived from the observed GRBs with sufficient data, all SGRBs are outliers because of their relatively higher $\\nu F_\\nu$ peak energy, while all LGRBs, except the peculiar long GRB\\,980425, are consistent with this relation. Amati et al.\\ (2007) find that GRB\\,060614 follows the relation whereas GRB\\,060505 does not. Hence, based on properties of the prompt emission other than the duration it seems impossible to establish clear evidence about which class of bursts GRB\\,060505 and GRB\\,060614 belong to.  To gain further insight on this topic, in this work we add our own observational data to already published data and study the afterglows of GRB\\,060505 and GRB\\,060614. We aim to determine from the afterglow properties whether these two bursts differ from other LGRBs, besides being SN-less, to provide further clues to the nature of their progenitors. The observations of the two afterglows and data reduction are presented in \\S~2.1 and \\S~2.2. At the beginning of \\S~3, we briefly introduce the leading external shock wave model employed to explain GRB afterglows, and apply it to GRB\\,060505 in \\S~3.1. GRB\\,060614 has been studied by \\cite{Mangano07}; in \\S~3.2 and \\S~3.3 we re-analyze this burst and provide analytical and numerical constraints on the afterglow parameters, respectively. In \\S~4 we discuss possible progenitors of these two bursts in comparison with the recently discovered GRB\\,080503, which we define as the first hard-spike $+$ hard-tail {\\it Swift} GRB.  Throughout this paper we use the notation $F_{\\nu}(\\nu,t) \\propto t^{-\\alpha} \\nu^{-\\beta}$ for the afterglow monochromatic flux as a function of time, where $\\nu$ represents the observed frequency and $\\beta$ is the energy index which is related to the photon index $\\Gamma$ in the form of $\\beta = \\Gamma-1$. The convention $Q_x=Q/10^x$ has been adopted in cgs. In addition, we consider a standard cosmology model with $\\rm{H_0} = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M} = 0.3$, and $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "As shown in several early papers, for GRB\\,060505 (\\citealt{Fynb06}) and GRB\\,060614 (\\citealt{Fynb06,DellaValle06,Gal-Yam06}) there is no accompanying SN emission, down to limits hundreds of times fainter than the archetypical SN\\,1998bw that accompanied GRB\\,980425, and fainter than any Type Ic SN ever observed. Multi-wavelength observations of the early afterglow exclude the possibility of significant dust obscuration. For GRB\\,060505 the properties of the host galaxy (\\citealt{Ofek07,Thoene08}) as well as the spectral lag of the prompt emission (\\citealt{McBreen08}) is most consistent with the properties expected for the long-duration class of GRBs. For GRB\\,060614 the duration of the prompt emission places the burst firmly within the long-duration class of GRBs, but the negligible spectral lag (\\citealt{Gehrels06}) and the relatively modest star-formation activity of the host galaxy is more similar to the expected properties for the short-duration class of GRBs. In this paper we have investigated whether or not the properties of the two afterglows could provide some hints to the most likely progenitor types for these bursts.  For GRB\\,060505, the numerical fit of its afterglow shows that the standard jetted external shock wave model is consistent with the data, yielding a typical interstellar medium density of $n \\sim 1\\,{\\rm{cm}}^{-3}$, a wide jet angle $\\theta_j \\sim25^\\circ$, and a possible jet break at $t_b \\sim3$ days. For GRB\\,060614, the standard external shock wave model is again consistent with the data, but apparently needing to invoke energy injection. The afterglow shows the clearest achromatic jet break among all {\\swift} bursts studied so far, decaying in a broken power-law from $\\alpha_1\\sim-1.1$ to $\\alpha_2\\sim -2.5$ at $t_b\\sim 1.4$ days. An achromatic peak, especially in the {\\it UBVR} bands, occurs at $\\sim 0.3$ days, which we interpret as resulting from an episode of strong energy injection as \\cite{Mangano07} suggested. Numerical fit yields a circumburst density of $n \\sim 0.04 \\,{\\rm{cm}}^{-3}$, being consistent with the inferred value in \\cite{caito08},  and a jet angle of $\\theta_j\\sim 5^\\circ$. The inferred afterglow parameters for the two bursts fall within the range for other LGRBs.  If it had not been for the observed absence of associated SNe we would have no reason, from the afterglow properties, to question their classification as LGRBs.  After discovery of these two GRBs, to reconcile all SN-observed and SN-less GRBs within the conventional framework of short ($\\lesssim 2$ s) and long ($\\gtrsim 2$ s) GRBs, a new classification was proposed, in which GRBs featuring a short-hard spike and a (possible) long-soft tail would be ascribed to the conventional short class, or Type I in the new taxonomy, while all other GRBs, or Type II, would comprise the conventional long class (\\citealt{Zhang07}, see also \\citealt{Kann07}). The new classification expands the range of the conventional short class, and is applicable to GRB\\,060505 and GRB\\,060614, which then would be SN-less due to a merger-related progenitor rather than SN-less massive stellar death.  However, a recent burst, GRB\\,080503, seems to challenge the new classification. This burst also has a temporal spike $+$ tail structure in the prompt emission phase. The $T_{\\rm 90}$ values of the initial spike and the total emission in the 15-150 keV band are $0.32\\pm0.07$~s and $232$~s, respectively (\\citealt{Perley08}). The fluence of the non-spike emission measured from 5~s to 140~s after the BAT trigger in the 15-150 keV band is $(1.86\\pm 0.14)\\times 10^{-6}\\,{\\rm erg\\,cm^{-2}}$, being around 30 times that of the spike emission in the same band. This fluence ratio is much higher than the ratio of around 6 for GRB\\,060614, and higher than any previous similar {\\it Swift} GRB. For GRB\\,060614 and GRB\\,080503 we extracted the BAT lightcurves in different energy bins, shown in Fig.~\\ref{comparison}, for comparison study. For GRB\\,080503 we have analysed the spectral evolution during the prompt emission period by BAT and XRT and list our analysis and that of other groups in Table \\ref{080503}. From these we conclude that: (1) The results of different groups are fully consistent with each other. (2) The photon indices for the spike and non-spike emissions are consistent within their error regions ($90\\%$ confidence level). A strong spectral softening from the spike phase to the non-spike phase can be excluded. A cutoff power-law fit does not improve the fitting, yielding the error of $E_{\\rm peak}$ larger than 100\\%. (3) During the BAT-XRT overlap period, the XRT spectra are always harder than the BAT spectra, which implies that the BAT$+$XRT spectra (0.3-150 keV) would be harder than the BAT spectra (15-150 keV) alone. This adds more evidence that from spike to non-spike the spectra did not soften considerably. In addition, note that for all conventional LGRBs (e.g., from BATSE), there is a general trend that the spectra during the prompt emission would soften mildly from the beginning time to the ending time (\\citealt{Norris00}). Therefore, the prompt spectral evolution of GRB\\,080503 is different from that of GRB\\,060614.  In addition, we computed the spectral lags in different energy bands using the 64 ms binning lightcurves, following the method in \\cite{Norris00} and \\cite{Chen05}. The lag during the initial spike phase is consistent with zero. From 5 to 50 s since BAT trigger, the lags are $0.8^{+0.3}_{-0.4}$ s for the 25-50 keV vs. 15-25 keV band and $0.8^{+0.4}_{-0.5}$ s for the 50-100 keV vs. 15-25 keV band respectively, both well above the lag range for SGRBs. Again the lag of GRB\\,080503 is in contrast with that of GRB\\,060614. The redshift of GRB\\,080503 was not measured mainly because its optical and X-ray afterglows became very faint shortly after the BAT trigger (never exceeding 25 mag in deep observations starting at $\\sim 1$ hr since trigger), but the g-band photometric detection imposes a limit of approximately z$<$4 (\\citealt{Perley08}). In Fig.~\\ref{SL} we show the possible position of GRB\\,080503 in the spectral lag-peak luminosity plane relative to the positions of previous LGRBs and SGRBs. As can be seen, for the extended emission of GRB\\,080503, its position is outside the SGRB population at a very high confidence level regardless of its redshfit. For the spike emission of GRB\\,080503, its position is within the SGRB population because of a peak luminosity comparable to that of the extended emission, as well as a negligible spectral lag.  The very faint optical and X-ray afterglows of GRB\\,080503 may indicate the circumburst density is very low. This is consistent with the fact that the afterglow is located away from any host galaxy down to 28.5 mag in deep \\textit{Hubble Space Telescope} imaging (\\citealt{Perley08}). These observational signatures contribute to put GRB\\,080503 into the merger class. If we ignore the lag function in classifying GRBs and relax the restriction of ``soft-tail'' for a Type I GRB to ``either soft- or hard- tail'', then it would be (even more) ambiguous, also operationally difficult to define the type of spike$+$tail GRBs among all long GRBs, especially among those at $z\\gtrsim0.7$ for which a SN search in their afterglows is difficult or impossible. Bearing in mind that GRBs can be either luminous or under-luminous, and the redshift could be either high or low, then the problems would be: up to what duration should be classified as a spike, and how much should the flux ratio be for the spike component over the non-spike component? If GRB\\,080503 is interpreted as a merger burst, then it enhances the possibility that a merger could produce a long GRB, at least in the prompt emission phase, mimicking the one produced by a collapsar. GRB\\,080503 is a dark burst with the optical-to-Xray spectral slope $\\beta_{OX}$ well below 0.5, the critical value for defining a dark burst (\\citealt{Jakobsson04}), at 0.05 days after the burst. We speculate that some other dark bursts may have their progenitor same as GRB\\,080503.  The progenitors of GRB\\,060505 and GRB\\,060614 remain uncertain based on their observations and the current GRB and SN theory. Other than the lack of a SN component, their afterglows are actually not peculiar when compared with the afterglows of other LGRBs. According to the current theoretical study, in the core-collapse scenario the ``fallback''-formed black holes or progenitors with relatively low angular momentum could produce such SN-less GRBs (e.g., \\citealt{Frye06,Sumiyoshi06,Nakazato08,Kochanek2008,Valenti09} for observational existence of an extremely faint Ibc SN), or in the merger scenario the two compact objects also could produce such SN-less GRBs if the formed remnant, a differentially rotating neutron star or an uniformly rotating magnetar, has not collapsed into a black hole immediately \\citep{KR98,RRD03}. A difference existing in the afterglows between these two scenarios is that the collapsar model predicts a WIND-like circumburst medium created by the Wolf-Rayet progenitor star while the merger model does not. As a matter of fact the WIND signature is not clearly evident in most LGRBs, but this should not necessarily lead to the merger origin for these bursts because the definite WIND signature is an ideal case for a constant wind off a massive star. If lucky enough, the core-collapse origin for a SN-less GRB (no matter whether it has the hard-spike $+$ soft-tail structure) will be quite certain if in the afterglow, either the X-ray flux $F(t)$ decays as $\\propto t^{-\\alpha}$ with its spectrum as $\\propto \\nu^{-(2\\alpha-1)/3}$ for the $\\nu_m < \\nu_{X} < \\nu_c$ phase (normally in early afterglow), or the optical decay index is larger than the X-ray decay index by a factor of $\\sim1/4$ for the slow cooling phase of the WIND scenario."
    },
    "0812/0812.2467_arXiv.txt": {
        "abstract": "{We construct a new family of analytic models of black hole accretion disks in dynamical equilibria. Our construction is based on assuming distributions of angular momentum and entropy. For a particular choice of the distribution of angular momentum, we calculate the shapes of equipressure surfaces. The equipressure surfaces we find are similar to those in thick, slim and thin disks, and to those in ADAFs.} \\authorrunning{Qian Lei, M.A.\\,Abramowicz, P.C.\\,Fragile, J.\\,Hor{\\'a}k, M.\\,Machida \\& O.\\,Straub} \\titlerunning{The Polish doughnuts revisited} ",
        "introduction": "In accretion disk theory one is often interested in phenomena that occur on a ``dynamical'' timescale ${\\cal T}_{0}$ much shorter than the ``viscous'' timescale ${\\cal T}[{\\cal L}]$ needed for angular momentum redistribution and the ``thermal'' timescale ${\\cal T}[{\\cal S}]$ needed for entropy redistribution\\footnote{We use the spherical Boyer-Lindquist coordinates $t, \\phi, r, \\theta$, the geometrical units $c$ $=$ $1$ $=$ $G$ and the $+---$ signature. The Kerr metric is described by the ``geometrical'' mass $M$ and the ``geometrical'' spin parameter $0 < a < 1$, that relate to the ``physical'' mass and angular momentum by the rescaling, $M = GM_{\\rm phys}/c^2$, $a = J_{\\rm phys}/(M\\,c)$. Partial derivatives are denoted by $\\partial_i$ and covariant derivatives by $\\nabla_i$.}, \\begin{equation} \\label{timescales} {\\cal T}_{0} \\ll {\\rm min}\\left({\\cal T}[{\\cal L}], {\\cal T}[{\\cal S}] \\right). \\end{equation} The question whether it is physically legitimate to approximately describe the black hole accretion flows (at least in some ``averaged'' sense) in terms of stationary (independent on $t$) and axially symmetric (independent on $\\phi$) dynamical equilibria, is not yet resolved. While observations seem to suggest that many real astrophysical sources experience periods in which this assumption is quite reasonable, several authors point out that the results of recent numerical simulations seem to indicate that the MRI and other instabilities make the black hole accretion flows genuinely non-steady and non-symmetric, and that the very concept of the separate timescales (\\ref{timescales}) may be questionable in the sense that locally ${\\cal T}_{0} \\approx {\\cal T}[{\\cal L}] \\approx {\\cal T}[{\\cal S}]$. However, this assumption has been made in {\\it all} existing comparisons between theory and observations, be they by detailed spectral fitting \\citep[e.g.][and references there]{sha-2007, sha-2008}, line profile fitting \\citep[e.g.][]{fab-2003}, or studying small amplitude oscillations \\citep[see][for references]{abr-2005}. It seems that the present understanding of the black hole accretion phenomenon rests, in a major way, on studies of stationary and axially symmetric models.  From the point of view of mathematical self-consistency, in modeling of these stationary and axially symmetric dynamical equilibria, distributions of the {\\it conserved} angular momentum and entropy, \\begin{equation} \\label{distribution-lagrangian} \\ell = \\ell(\\xi, \\eta),~~~ s = s(\\xi, \\eta), \\end{equation} may be considered as being {\\it free functions} of the Lagrangian coordinates \\citep{ost-1970, abr-1970, bar-1970}. The Lagrangian coordinates $\\xi, \\eta$ are defined by demanding that a narrow ring of matter $(\\xi, \\xi + d\\xi)$, $(\\eta, \\eta + d\\eta)$ has the rest mass $dM_0 = \\rho_0(\\xi, \\eta)d\\xi d\\eta$ with $\\rho_0$ being the rest mass density. In the full physical description, the form of the functions in (\\ref{distribution-lagrangian}) is not arbitrary but given by the dissipative processes, like viscosity and radiative transfer. At present, several important aspects of these processes are still unknown, so there is still no practical way to calculate physically consistent models of accretion flows from first principles, without involving some ad hoc assumptions, or neglecting some important processes. Neither the hydrodynamical simulations (that e.g. use the ad hoc $\\alpha\\,=\\,$const viscosity prescription), nor the present day MHD simulations (that e.g. neglect radiative transfer) could be considered satisfactory. Furthermore, the simplifications made in these simulations are mathematically equivalent to guessing free functions (such as the entropy distribution). Bohdan Paczy{\\'n}ski pointed out that it could often be more pragmatic to make a physically motivated guess of the final result, e.g. to guess the form of the angular momentum and entropy distributions.  In practice, it is far easier to guess and use the coordinate distributions of the {\\it specific} angular momentum and entropy, \\begin{eqnarray} {\\cal L} &=& {\\cal L}(r, \\theta), \\label{momentum-distribution}\\\\ {\\cal S} &=& {\\cal S}(r, \\theta), \\label{entropy-distribution} \\end{eqnarray} than the Lagrangian distributions (\\ref{distribution-lagrangian}). However, one does not known a priori the relation between the conserved $\\ell$ and specific ${\\cal L}$ angular momenta (and entropy), or the functions,  $\\xi = \\xi(r,\\theta)$, $\\eta = \\eta(r, \\theta)$. Thus, assuming (\\ref{momentum-distribution}) and (\\ref{entropy-distribution}) is not equivalent to assuming (\\ref{distribution-lagrangian}), and usually it should be a subject to some consistency conditions. We shall return to this point in Section \\ref{discussion}.  In several ``astrophysical scenarios'' one indeed guesses a particular form of (\\ref{momentum-distribution}) and (\\ref{entropy-distribution}). For example, the celebrated \\citet{sha-sun-1974} {\\it thin disk} model assumes the Keplerian distribution of angular momentum, \\begin{equation} \\label{Keplerian} {\\cal L}(r, \\theta) = {\\cal L}_K(r) \\equiv \\frac{M^{1/2}\\,\\left(r^2 - 2aM^{1/2}r^{1/2} + a^2\\right)} {r^{3/2} - 2Mr^{1/2} + aM^{1/2}}, \\end{equation} and the popular {\\it cold-disk-plus-hot-corona} model assumes a low entropy flat disk surrounded by high entropy, more spherical corona. These models contributed considerably to the understanding of black-hole accretion physics.  The mathematically simplest assumption for the angular momentum and entropy distribution is, obviously, \\begin{eqnarray} {\\cal L}(r, \\theta) &=& {\\cal L}_0 = {\\rm const}, \\label{constant-momentum} \\\\ {\\cal S}(r, \\theta) &=& {\\cal S}_0 = {\\rm const}. \\label{constant-entropy} \\end{eqnarray} This was used by Paczy{\\'n}ski and his Warsaw team to introduce the {\\it thick disk} models \\citep{abr-1978, koz-1978, jar-1980, pac-wii-1980, abr-1980, abr-1981, pac-1982}. Thick disks have characteristic toroidal shapes, resembling a doughnut. Probably for this reason, Martin Rees coined the name of {\\it Polish doughnuts}\\footnote{However, real Polish doughnuts (called {\\it p{\\c a}czki} in Polish) have spherical shapes. They are definitely non-toroidal --- see e.g. http://en.wikipedia.org/wiki/Paczki ~.} for them.  Figure \\ref{analytic-numerical} shows a comparison of a state-of-art MHD simulation of black-hole accretion (time and azimuth averaged) with a Polish doughnut corresponding to a particular ${\\cal L}_0$. Both models show the same characteristic features of black hole accretion: (i) a funnel along the rotation axis, relevant for jet collimation and acceleration; (ii) a pressure maximum, possibly relevant for epicyclic oscillatory modes; and (iii) a cusp-like self-crossing of one particular equipressure surface, relevant for an inner boundary condition, and for stabilization of the Papaloizou-Pringle \\citep{bla-1987}, thermal, and viscous instabilities \\citep{abr-1971}. The cusp is located between the radii of marginally stable and marginally bound circular orbits, \\begin{equation} \\label{cusp} r_{mb} < r_{cusp} < r_{ms} \\equiv {\\rm ISCO}. \\end{equation} Polish doughnuts have been useful in semi-analytic studies of the astrophysical appearance of super-Eddington accretion \\citep[see e.g.][]{sik-1971, mad-1988, szu-1996} and in analytic calculations of small-amplitude oscillations of accretion structures in connection with QPOs \\citep[see e.g.][]{bla-2006}. In the same context, numerical studies of their oscillation properties for different angular momentum distributions were first carried out by \\citet{rez-2003a, rez-2003b}. Moreover, Polish doughnuts are routinely used as convenient starting initial configurations in numerical simulations \\citep[e.g.][]{haw-2001,dev-2003}. Recently, \\citet{kom-2006} has constructed analytic models of magnetized Polish doughnuts. \\begin{figure} \\centering \\includegraphics[width=9cm]{1518.F01.eps} \\caption{Equipressure surfaces in a very simple and analytic Polish doughnut (left, with linear spacing), and a sophisticated, state-of-art full 3D MHD numerical simulation (right, with logarithmic spacing). Although the shapes of equipressure surfaces are remarkably similar, in the numerical model the pressure gradient is seriously larger, and visibly enhanced along roughly conical surfaces, approximately $30^{\\circ}$ from the equatorial plane. \\citep[Figure taken from][]{abr-fra-2008}} \\label{analytic-numerical} \\end{figure} However, a closer inspection of Figure \\ref{analytic-numerical} reveals that the numerically constructed model of accretion has a (much) larger ``vertical'' pressure gradient than the analytic Polish doughnut, and that in the numerical model the gradient is visibly enhanced along roughly conical surfaces, approximately $30^{\\circ}$ from the equatorial plane. This (and several other) detailed features of the accretion structure cannot be modeled by either the Keplerian nor the constant angular momentum assumption alone. We suggest and discuss in this paper a simple but flexible ansatz, that is a combination of the two standard distributions, Keplerian (\\ref{Keplerian}) and constant (\\ref{constant-momentum}). The new ansatz preserves the virtues of assuming the standard distributions where this is appropriate, but leads to a far richer variety of possible accretion structures, as are indeed seen in numerical simulations. ",
        "conclusions": "The new ansatz (\\ref{ansatz-general}) captures two essential features of the angular momentum distribution in black hole accretion disks: \\begin{enumerate} \\item On the equatorial plane and far from the black hole, the angular momentum in the disk differs only little from the Keplerian one being slightly sub-Keplerian, but closer in it becomes (slightly) super-Keplerian and still closer, in the plunging region, sub-Keplerian again and nearly constant. \\item Angular momentum may significantly decrease off the equatorial plane, and become very low (even close to zero, in a non-rotating ``corona''). \\end{enumerate} Models of tori described here may be useful not only for accretion disks but also for tori that form in the latest stages of neutron star binary mergers. This is relevant for gamma ray bursts \\citep{wit-1994} and gravitational waves \\citep{bai-2008}."
    },
    "0812/0812.3461_arXiv.txt": {
        "abstract": "A four-dimensional statistical description of electromagnetic radiation is developed and applied to the analysis of radio pulsar polarization. The new formalism provides an elementary statistical explanation of the modal broadening phenomenon in single pulse observations. It is also used to argue that the degree of polarization of giant pulses has been poorly defined in past studies. Single and giant pulse polarimetry typically involves sources with large flux densities and observations with high time resolution, factors that necessitate consideration of source-intrinsic noise and small-number statistics. Self noise is shown to fully explain the excess polarization dispersion previously noted in single pulse observations of bright pulsars, obviating the need for additional randomly polarized radiation. Rather, these observations are more simply interpreted as an incoherent sum of covariant, orthogonal, partially polarized modes. Based on this premise, the four-dimensional covariance matrix of the Stokes parameters may be used to derive mode-separated pulse profiles without any assumptions about the intrinsic degrees of mode polarization. Finally, utilizing the small-number statistics of the Stokes parameters, it is established that the degree of polarization of an unresolved pulse is fundamentally undefined; therefore, previous claims of highly polarized giant pulses are unsubstantiated. ",
        "introduction": "\\label{sec:review}  This section reviews the relevant algebra of both Jones and Mueller representations of polarimetric transformations.  Unless otherwise noted, the notation and terminology synthesizes that of the similar approaches to polarization algebra presented by \\citet{clo86}, \\citet{bri00}, and \\cite{ham00}.  \\subsection{Jones Calculus} \\label{sec:jones}  The polarization of electromagnetic radiation is described by the second-order statistics of the transverse electric field vector, $\\mbf{e}$, as represented by the complex $2\\times2$ coherency matrix, $\\mbf{\\rho}\\equiv\\langle\\mbf{e\\, e}^\\dagger\\rangle$ \\citep{bw80}.  Here, the angular brackets denote an ensemble average, $\\mbf{e}^\\dagger$ is the Hermitian transpose of $\\mbf{e}$, and an outer product is implied by the treatment of $\\mbf{e}$ as a column vector.  A useful geometric relationship between the complex two-dimensional space of the coherency matrix and the real four-dimensional space of the Stokes parameters is expressed by the following pair of equations: \\begin{eqnarray} {\\mbf\\rho} & = & S_k\\,\\pauli{k} / 2     \\label{eqn:combination} \\\\ S_k & = & \\trace(\\pauli{k}\\mbf{\\rho}).  \\label{eqn:projection} \\end{eqnarray} Here, $S_k$ are the four Stokes parameters, Einstein notation is used to imply a sum over repeated indeces, $0\\le k \\le3$, $\\pauli{0}$ is the $2\\times2$ identity matrix, $\\pauli{1-3}$ are the Pauli matrices, and $\\trace$ is the matrix trace operator.  The Stokes four-vector is composed of the the total intensity $S_0$ and the polarization vector, $\\mbf{S}=(S_1,S_2,S_3)$.  \\Eqn{combination} expresses the coherency matrix as a linear combination of Hermitian basis matrices; \\eqn{projection} represents the Stokes parameters as the projections of the coherency matrix onto the basis matrices. These properties are used to interpret the well-studied statistics of random matrices through the familiar Stokes parameters.  Linear transformations of the electric field vector are represented using complex $2\\times2$ Jones matrices.  Substitution of $\\mbf{e}^\\prime={\\bf J}\\mbf{e}$ into the definition of the coherency matrix yields the congruence transformation, \\begin{equation} {\\mbf{\\rho}^\\prime}={\\bf{J}}\\mbf{\\rho}{\\bf{J}}^\\dagger, \\label{eqn:congruence} \\end{equation} which forms the basis of the various coordinate transformations that are exploited throughout this work. If ${\\bf{J}}$ is non-singular, it can be decomposed into the product of a Hermitian matrix and a unitary matrix known as its polar decomposition, \\begin{equation} \\label{eqn:polar} {\\bf J} = J \\, \\boost \\, \\rotat, \\end{equation} where $J=|{\\bf J}|^{1/2}$, \\boost\\ is positive-definite Hermitian, and \\rotat\\ is unitary; both \\boost\\ and \\rotat\\ are unimodular.  Under the congruence transformation of the coherency matrix, the Hermitian matrix \\begin{equation} \\label{eqn:Boost} \\boost = \\exp (\\beta\\,\\mbf{\\hat{m}\\cdot\\sigma}) = \\pauli{0}\\cosh\\beta + \\mbf{\\hat{m}\\cdot\\sigma}\\sinh\\beta \\end{equation} effects a Lorentz boost of the Stokes four-vector along the $\\mbf{\\hat m}$ axis by a hyperbolic angle $2\\beta$. As the Lorentz transformation of a spacetime event mixes temporal and spatial dimensions, the polarimetric boost mixes total and polarized intensities, thereby altering the degree of polarization. In contrast, the unitary matrix \\begin{equation} \\label{eqn:Rotation} \\rotat = \\exp (\\Ci\\phi\\,\\mbf{\\hat{n}\\cdot\\sigma}) = \\pauli{0}\\cos\\phi + \\Ci\\mbf{\\hat{n}\\cdot\\sigma}\\sin\\phi \\end{equation} rotates the Stokes polarization vector about the $\\mbf{\\hat n}$ axis by an angle $2\\phi$. As the orthogonal transformation of a vector in Euclidean space preserves its length, the polarimetric rotation leaves the degree of polarization unchanged.  These geometric interpretations promote a more intuitive treatment of the matrix equations that typically arise in polarimetry. Boost transformations can be utilized to convert unpolarized radiation into partially polarized radiation, and rotation transformations can be used to choose the orthonormal basis that maximizes symmetry. These properties are exploited in \\Sec{joint} to simplify the relevant mathematical expressions that describe the four-dimensional joint distribution of the Stokes parameters.   \\subsection{Eigen Decomposition} \\label{sec:eigen}  It proves useful in \\Sec{joint} to express the coherency matrix as a similarity transformation known as its eigen decomposition, \\begin{equation} \\mbf{\\rho}={\\bf{R}}\\left( \\begin{array}{cc} \\lambda_0 & 0 \\\\ 0 & \\lambda_1 \\end{array}\\right){\\bf{R}}^{-1}. \\label{eqn:eigen} \\end{equation} Here, ${\\bf{R}}=(\\mbf{e}_0\\,\\mbf{e}_1)$ is a $2\\times2$ matrix with columns equal to the eigenvectors of $\\mbf{\\rho}$, and $\\lambda_m$ are the corresponding eigenvalues, given by $\\lambda=(S_0\\pm|\\mbf{S}|)/2=(1 \\pm P)S_0/2$, where $P=S_0/|\\mbf{S}|$ is the degree of polarization.  If the signal is completely polarized, then $\\lambda_1=0$.  If the signal is unpolarized, then there is a single 2-fold degenerate eigenvalue, $\\lambda=S_0/2$ and ${\\bf R}$ is undefined.  If the eigenvectors are normalized such that ${\\mbf{e}_k}^\\dagger\\mbf{e}_k=1$, then \\eqn{eigen} is equivalent to a congruence transformation by the unitary matrix, ${\\bf{R}}$. In the natural basis defined by ${\\bf{R}}^\\dagger$, the eigenvalues $\\lambda_m$ are equal to the variances of two uncorrelated signals received by orthogonally polarized receptors described by the eigenvectors. The total intensity, $S_0=\\lambda_0+\\lambda_1$; the polarized intensity, $S_1=|\\mbf{S}|=\\lambda_0-\\lambda_1$; and $S_2=S_3=0$. That is, ${\\bf{R}}^\\dagger$ rotates the basis such that the mean polarization vector points along $S_1$, providing cylindrical symmetry about this axis.   \\subsection{Mueller Calculus} \\label{sec:mueller}  The congruence transformation of the coherency matrix by any Jones matrix {\\bf J} may be represented by an equivalent linear transformation of the Stokes parameters by a real-valued $4\\times4$ Mueller matrix \\begin{equation} \\label{eqn:Mueller} M_i^k = \\frac{1}{2}\\trace(\\pauli{i}\\,{\\bf J}\\,\\pauli{k}\\,{\\bf J}^\\dagger), \\end{equation} such that \\begin{equation} S^\\prime_i = M_i^k S_k. \\end{equation} Mueller matrices that have an equivalent Jones matrix are called pure, and such transformations are related to the Lorentz group \\citep[e.g.][]{bar63,bri00}. This motivates the definition of an inner product \\begin{equation} \\inner{A}{B} \\equiv A^k B_k = \\eta^{kk}A_k B_k = A_0B_0 - \\mbf{A \\cdot B}, \\end{equation} where $A$ and $B$ are Stokes four-vectors and $\\eta_{ij}$ is the Minkowski metric tensor with signature $(+,-,-,-)$. The Lorentz invariant of a Stokes four-vector is equal to four times the determinant of the coherency matrix; that is, \\begin{equation} \\inv{S} \\equiv \\inner{S}{S} = S_0^2 - |\\mbf{S}|^2 = 4 |\\mbf{\\rho}| \\end{equation} Similarly, the Euclidean norm \\norm{S} is twice the Frobenius norm of the coherency matrix \\norm{\\mbf{\\rho}}; i.e. \\begin{equation} \\norm{S}^2 \\equiv S_0^2 + |\\mbf{S}|^2 = 4 \\norm{\\mbf{\\rho}}^2. \\end{equation} The coherency matrix is a positive semi-definite Hermitian matrix; therefore, the Lorentz invariant of any physically realizable source of radiation is greater than or equal to zero.  It is equal to zero only for completely polarized radiation, when the degree of polarization is unity ($S_0=|\\mbf{S}|$).  As with the spacetime null interval, no linear transformation of the electric field can alter the degree of polarization of a completely polarized source. ",
        "conclusions": "\\label{sec:conclusion}  A four-dimensional statistical description of polarization is presented that exploits the homomorphism between the Lorentz group and the transformation properties of the Stokes parameters. Within this framework, a generalized expression for the covariance matrix of the Stokes parameters is developed and applied to the analysis of single-pulse polarization fluctuations. The consideration of source-intrinsic noise renders randomly polarized radiation unnecessary, explains the coincidence between increased amplitude modulation and orthogonal mode fluctuations, and indicates that orthogonally polarized modes are only partially polarized. Furthermore, the four-dimensional covariance matrix of the Stokes parameters enables estimation of the mode intensities, degrees of polarization, intensity correlation coefficient, and effective degrees of freedom of covariant, orthogonal, partially polarized modes. Measured as a function of pulse phase, these parameters may be used to produce mode-separated profiles without any assumptions about the intrinsic degree of mode polarization.  The formalism is also used to derive the first and second moments of the degree of polarization as a function of the intrinsic degree of polarization and the number of degrees of freedom of the stochastic process. These are used to demonstrate that giant pulse polarimetry is fundamentally limited by systematic bias due to insufficient statistical degrees of freedom. The discussions of amplitude modulation in \\Sec{modulation} and wave coherence in \\Sec{coherence} serve to illustrate the difficulties in defining an appropriate average when the signal is unresolved. In this regard, the approach employed in this paper is complimentary to previous analyses that make use of auto-correlation functions and fluctuation spectra to separately measure the statistics of the total and polarized intensities. Useful new results may be derived by extending the techniques employed in these works to include the cross-correlation terms that describe the statistical dependences between the Stokes parameters."
    },
    "0812/0812.3182_arXiv.txt": {
        "abstract": "We present an inflationary model that is geodesically complete and does not suffer from the transplanckian problem.  In most inflationary models, massless (conformal) scalar field fluctuations in a  deSitter background gives rise to a scale-invariant spectrum. In this work, we realize scale invariant perturbations from thermal fluctuations in (conformal) radiation during a  radiation dominated contraction era prior to inflation. As the modes exit the Hubble radius during the contraction phase, scale invariant fluctuations are indeed generated.  After many cycles, we enter into a power-law inflationary phase, that stretches the modes produced in the previous contraction phase to scales that we observe today. ",
        "introduction": "While the inflationary paradigm is successful in resolving the classic puzzles of the Standard Big Bang (SBB) cosmology and providing a mechanism for generating near scale-invariant fluctuations, it suffers from specific theoretical/technical issues. {\\bf (i)} Perhaps the most daunting is the fact that inflationary space-times are geodesically incomplete and, hence, do not address the issue of the BB singularity~\\cite{vilenkin}. ({\\bf ii}) In many inflationary scenarios  one encounters the transplanckian problem~\\cite{transplanck} where the scales one observes today becomes smaller than the Planck length at the beginning of inflation; it is not clear that one can trust the assumption of perturbation theory in the transplanckian regime  (see~\\cite{transplanck,trans} for a few different paradigms where attempts have been made to quantify the corrections). {\\bf (iii)} It is typically challenging to build a phenomenologically viable inflationary model where the potential remains flat for a long enough period. {\\bf (iv)} Finally, it has been realized for some time now that inflation initiates in a region which is  extremely low in entropy.  For inflation to work, although  one only requires a microscopic ``smooth'' patch where the scalar field is homogeneous and isotropic, a priori, in the context of a continuum field theory this is not guaranteed~\\cite{low-entropy};  consequently, it has been argued that the ``initial conditions'' necessary to begin inflation may not be ``sufficiently likely'', although there are counter-arguements~\\cite{linde-ic} and proposals~\\cite{way-around} for getting around this difficulty.  In this paper we explore an extension of inflation that addresses most of the aforementioned problems.  Cyclic models offer an interesting alternative to the inflationary paradigm; they evade the conceptual and technical problems with a ``beginning of time'' when the universe starts from a singularity by positing time to be ``eternal'' in either direction (past and future)\\footnote{It is rather surprising to note that in the context of four dimensional isotropic and homogeneous cosmology, apart from having $\\spadesuit$ an infinite number of cycles, there are only three other viable non-singular alternatives in the past: $\\heartsuit$ the bouncing universe where there is a single bounce and our current phase of expansion is preceded by a phase of contraction as for instance advocated in pre-big bang scenarios~\\cite{pre-bigbang} and non-local gravity models~\\cite{bms}. For a related interesting approach where at the bounce point the arrow of time reverses,  see~\\cite{moffat}. $\\diamondsuit$ the emergent universe scenario where the universe asymptotes to a constant space-time in the infinite past~\\cite{emerge}. $\\clubsuit$ The universe has  a finite number of cycles where the first one begins either as $\\heartsuit$ or $\\diamondsuit$.}. This alternative view has motivated physicists to revisit cyclic universes time and again~\\cite{tolman,narlikar,ekcyclic,barrow,dabrowski,phantom,kanekar,saridakis} (for a comprehensive review see~\\cite{review}). Further, most cyclic models automatically address some of the classic problems of Big Bang (BB) cosmology which were the motivations for inventing inflation in the first place.  Thus although cyclic and inflationary models are not mutually exclusive, it is natural to try to attempt to replace inflation altogether with ``cyclicity''.  In this paper, however, we take a slightly different approach, by exploring whether by embedding inflation in a cyclic universe setting, some of it's problems viz. {\\bf (i-iv)} can be alleviated.  Our main idea is to merge inflation with cyclic cosmology where the universe undergoes an infinite number of cycles before bouncing into a final power-law inflationary phase.  However, unlike conventional inflation where one exploits the fact that massless scalar field fluctuations in a conformally invariant deSitter background give rise to a scale-invariant spectrum, in this model a different ``conformal phase'' is employed to render scale-invariant fluctuations.  It is quite remarkable that these fluctuations generically produce a scale-invariant spectrum, unlike inflationary or ekpyrotic models which rely on special kind of  scalar field potentials. First alluded to by Peebles~\\cite{peebles}, this was elaborated on later in the context of a bouncing universe in~\\cite{pogosian}. The problem with this scenario  is that in a symmetric bounce the fluctuations corresponding to cosmic structures that we see today is generated at very low temperatures (during contraction) and hence  are extremely suppressed. This is where the power-law inflationary phase comes into play. As is well known, inflation can stretch  fluctuations generated at much smaller length scales (and therefore at a much higher temperature and amplitude) to exponentially larger lengths, so that they can become  the size of the Hubble radius today. In other words, the entropy production at the end of inflation  makes the ``adjacent'' contraction and expansion phase sufficiently asymmetric to ``amplify'' the thermal fluctuations.  ({\\bf iii}) Consequently, one only needs a power-law inflationary phase, $a\\sim t^p$, with $p> 1.3$, a requirement that can be satisfied with exponential potentials naturally arising in supergravity/string theory compactifications~\\cite{exp-pot,flux,bj,munoz,bbc}. ({\\bf ii}) A second rather unique feature of this model is that the number of e-foldings in the inflationary phase actually controls the amplitude of fluctuations, and the observed value in the sky automatically guarantees that this model does not suffer from the  transplanckian problem. We shall see that the fluctuations are produced when both classical thermodynamics and General Relativity are expected to be valid. ({\\bf iv}) Further, in this model the inflationary phase is preceded by a phase of thermal matter which in some sense prepares the initial condition for inflation to begin.  Although, whether this can really address the ``low entropy'' problem,  is not clear. Finally, by embedding inflation in the ``emergent cyclic universe'' scenario we will be able to address {\\bf (i)} concretely.  While cyclic models are employed to describe past-eternal non-singular cosmologies, it is remarkably difficult to achieve this ``in practice''.  Firstly, one needs to come up with a theoretically consistent (ghost and instability free) new physics to resolve the Big Bang singularity. In General Relativity there are strict ``no-go'' theorems prohibiting  ``bounces''~\\cite{paris-ec} where contraction gives way to expansion. Fortunately encouraging progress have been made in the recent past, mostly using non-local and/or non-perturbative physics, such as in Loop Quantum Cosmology\\footnote{In  Brane-world scenarios with extra time-like directions~\\cite{sahni}, one also has an ``effective'' negative energy contribution which goes like $\\rho^2$ similar to what is found in Loop Quantum cosmology which can therefore also lead to singularity free bouncing universes.}~\\cite{loop}, string inspired non-local modifications of gravity~\\cite{bms},  stringy toy models using AdS/CFT ideas~\\cite{turok}, tachyon dynamics~\\cite{tachyon}, and mechanisms involving ghost condensation~\\cite{ghost}, and fermion condensations (both classically~\\cite{greene} and via quantum  BCS-like gap formation~\\cite{ab}).  Secondly, even if we are able to avoid the BB singularity, in cyclic models we are confronted with the second law of thermodynamics according to which the total entropy in the universe can only increase monotonically. As first pointed out by Tolman~\\cite{tolman}, this immediately tells us that the evolution can at most be ``quasi-periodic'' where the length of the cycle typically  monotonically increases with the increase in entropy from cycle to cycle (see~\\cite{columbia} for a counter-example). Moreover,  the problem of the ``beginning of time'' comes back to haunt us, as  the cycles  either become vanishingly small at a finite proper time  in the past~\\cite{tolman} or is geodesically incomplete, as in the past-eternal inflationary scenarios~\\cite{vilenkin}.  In~\\cite{emergent} a new ``emergent'' cyclic universe model was presented where it was argued that String theory may have the ingredients to cure both the Big Bang and Tolman's Entropy problem. Repulsive Casimir energies coming from scalar massless string states were invoked to obtain the bounce. (A more detailed accounting of all other different contributions to the Casimir energy is  currently under study~\\cite{radu}.)  In~\\cite{emergent} it was also argued how the existence of a thermal Hagedorn phase in string theory at high enough energies/temperatures can  provide a nice solution to Tolman's entropy problem. The important implications of the Hagedorn phase in String-Gas-Cosmology~\\cite{vafa} (see~\\cite{string-gas-review} for recent reviews) has been discussed several times~\\cite{tseytlin,string-thermo,anupam,hindmarsh} in the past, but only recently in the context of cyclic models~\\cite{emergent,columbia}. In particular in~\\cite{emergent} it was shown how for a large class of closed universe models, as we go back in the past (cycles) the universe starts spending more and more time in the thermal Hagedorn phase where entropy remains constant. As a result, in the infinite past the universe assymptotes to an almost periodic  evolution with short time periods (string time scale and energy densities) and constant but non-zero entropy, see fig.~\\ref{figemergent}.  In each of these cycles some amount of entropy is always produced in the non-thermal phase (gas of radiation + non-relativistic matter) via exchange of energy from matter to radiation and as a result the cycles and the universe gradually becomes longer and larger respectively~\\cite{emergent}. Effectively, because of entropy production in every cycle the universe expands a little more than it contracts and this process can continue forever.  In this paper we explore the much richer dynamics that can emerge once we include scalar fields, and in particular with exponential potentials~\\cite{exp-pot,flux,bj,munoz,bbc} and couplings~\\cite{vafa,tseytlin,bbc,string-gas} to matter, which are quite common in string and supergravity theories. Depending upon the exponents, we will find that there can now be two different phases: In the initial ``emergent'' phase the scalar field is essentially a spectator, the cycles grow gradually and the evolution is qualitatively similar to what  was discussed in~\\cite{emergent}. There is no inflationary phase in these cycles. However, when the entropy of the universe reaches a critical value, the scalar field suddenly becomes important and  starts to dominate the energy density  of the universe. Further, if the scalar field potential is flat  enough, then one enters a phase of  inflation which prevents the spatial curvature driven turn-around, making it the last  cycle!  It is already well known that the exponential potentials that one obtains in stringy compactifications can indeed give rise to a ``power-law'' phase of inflation, especially in the context of string gas cosmology~\\cite{bbem,bbmm}, but typically with a very small power which is unable to reproduce the spectral tilt observed in CMB. Fortunately as mentioned before, in our scenario we do not require inflation to produce the scale invariant perturbations.  Finally,  our scenario can be completely consistent with the recent observation of dark energy by  including the usual tiny cosmological constant.  In the next section we will start with a review of the emergent cyclic universe model that was presented in~\\cite{emergent}, and discuss the evolution in the early short cycles. Next in section \\ref{spectrum}, we will discuss how in this cyclic universe model one can incorporate a ``new'' mechanism for producing the fluctuations that we see in our sky today. In the subsequent section \\ref{toy}, we will provide a specific toy model realization of the above mechanism via a phase of power-law inflation in the last cycle.   Finally in section \\ref{conclusions}, we conclude with an outlook towards future direction. ",
        "conclusions": "The aim of this paper was to present a qualitatively new phenomenologically viable model of inflation which has some advantages over the conventional slow-roll counterpart. In particular we were able to provide a non-singular completion to inflation in the infinite past, circumvent the trans-Planckian problem, and  most crucially be consistent with CMB with just a low power-law inflationary phase. This was achieved by embedding a string gas inspired power-law inflationary phase in the ``emergent cyclic universe'' scenario which was recently proposed to solve Tolman's entropy problem. The analysis was most certainly done at a toy model level, but was inspired by several ingredients in string theory, such as the existence of a thermal Hagedorn phase in the UV, the presence of  moduli fields  which can give rise to repulsive Casimir energies and the fact that often some of the scalar moduli fields have exponential potentials and couplings to matter. This gives one hope that one may be able to realize this  scenario, or slight variants of it, in a more concrete string theory framework.   In order to achieve this we need in future to examine different aspects of equilibrium and non-equilibrium string thermodynamics as well as include specific moduli dynamics/stabilization. In principle this should also explain how the observed three dimensions became large as compared to the extra dimensions. It would be interesting to revisit the BV mechanism~\\cite{vafa} proposed in this context in the light of the new cyclic cosmology discussed here. Another important ingredient in string theory is the existence of higher dimensional objects such as branes, which may in fact play an important role  in explaining the  hierarchy in the sizes between the large observed   and the smaller internal manifolds if they can wrap around the latter~\\cite{branes}. We were not able to address how we start from an initial homogeneous and isotropic universe. If however ``initially'' (at $t=-\\infty$) the universe has a small anisotropy then in the course of subsequent evolution this anisotropy is only going to get further diluted as the universe, on the whole,  keeps expanding due to entropy production. So, at least if the ``initial conditions'' are fine tuned one should be able to avoid a mixmaster like chaotic evolution. Intriguingly, there are suggestions that a brane gas phase can naturally lead to isotropization of the different directions~\\cite{brane-isotropy}, which could essentially predate the emergent cyclic behavior discussed here. There are certainly indications that one may  have new thermal phases involving brane states which can later on pave way to the more conventional thermal Hagedorn phase~\\cite{samir} and it would be interesting to study these possibilities in the future. With respect to inhomogeneities, near the bounce in the Hagedorn phase the sub-Hubble\\footnote{By sub-Hubble we really mean here modes whose physical wave lengths are smaller than the time scale of the bounce.} thermal fluctuations remain small as the temperature which controls their amplitude is always less than the Planck scale, while the super-Hubble fluctuations freeze out before and are even smaller in amplitude. Thus inhomogenieties should also remain under control around the bounce in congruence with the BKL conjecture~\\cite{bkl} according to which it is really the anisotropy rather than the inhomogeneity that one needs to worry about near the bounce. In any case, it is clear that the issue of anisotropy and inhomogenieties, both ``initial'' and their growth, demands a much more detailed analysis in the future.     From a technical point of view, we seem to require fine-tuning of the parameter $v_2$ which resulted from the fact that we did not want the modes to re-enter the Hubble radius before inflation could start. One may however be able to alleviate this problem significantly by considering  longer decay times, and just ensuring that in the Penultimate cycle enough radiation is produced so that it starts to dominate by the time the energy density increases to $\\sim T_{\\mth}^4$ in the contraction phase. In order to test these scenarios a more detailed numerical analysis will be required.  Finally, from the observational view point, it is clear that a very distinctive feature of the model is that the spectrum will make a transition from being nearly scale-invariant (when we see the modes that exited during contraction) to a very red spectrum at smaller length scales (which exit the Hubble radius during the power-law phase of inflation). This should manifest itself as a running in the spectral index, for which there is now some evidence~\\cite{casas}. We do however need to make these predictions precise. Also, although it is clear that there are various effects which can give small corrections to scale-invariance of the modes that we observe in CMB\\footnote{The biggest effect is expected due to deviations from a radiation dominated universe near the turnaround. If the turnaround temperature is close to the temperatures when the physically relevant modes that we observe today left the Hubble radius in the contracting phase, then we will see corrections to the scale-invariant spectrum. This seems a likely possibility for phenomenological reasons as discussed in section \\ref{fine-tuning}.}, we  haven't yet investigated them in any detail to make connections to observations. We leave these issues for future."
    },
    "0812/0812.1652_arXiv.txt": {
        "abstract": "We have observed with {\\it XMM--Newton} four radiatively efficient active type 1 galaxies with black hole masses $<10^6~M_\\odot$, selected optically from the Sloan Digital Sky Survey and previously detected in the Rosat All Sky Survey. Their X--ray spectra closely resemble those of more luminous Seyferts and quasars, powered by accretion onto much more massive black holes and none of the objects is intrinsically absorbed by cold matter totally covering the source. We show here that their soft X--ray spectrum exhibits a soft excess with the same characteristics as that observed ubiquitously in radio--quiet Seyfert 1 galaxies and type 1 quasars, both in terms of temperatures and strength. This is highly surprising because the small black hole mass of these objects should lead to a thermal disc contribution in the soft X--rays and not in the UV (as for more massive objects) with thus a much more prominent soft X--ray excess.Moreover, even when the soft X--ray excess is modelled with a pure thermal disc, its luminosity turns out to be much lower than that expected from accretion theory for the given temperature. While alternative scenarios for the nature of the soft excess (namely smeared ionized absorption and disc reflection) cannot be distinguished on a pure statistical basis, we point out that the absorption model produces a strong correlation between absorbing column density and ionization state, which may be difficult to interpret and is most likely spurious. Moreover, as pointed out before by others, absorption must occur in a fairly relativistic wind which is problematic, especially because of the enormous implied mass outflow rate. As for reflection, it does only invoke standard ingredients of any accretion model for radiatively efficient sources such as a hard X--rays source and a relatively cold (though partially ionized) accretion disc, and therefore seems the natural choice to explain the soft excess in most (if not all) cases. The reflection model is also consistent with the additional presence of a thermal disc component with the theoretically expected temperature (although, again, with smaller--than--expected total luminosity).  We also studied in some detail the X--ray variability properties of the four objects. The observed active galaxies are among the most variable in X--rays and their excess variance is among the largest. This is in line with their relatively small black hole mass and with expectations from simple power spectra models. ",
        "introduction": "Two families of black holes have been observed so far: stellar--mass black holes are seen in X--ray binaries, while black holes with masses $10^6$--$10^9~M_\\odot$ are ubiquitous in the centers of galaxies, sometimes revealing themselves as Active Galactic Nuclei (AGN). The large mass gap between the two families is still scarcely populated. Dynamical evidence for a black hole with mass $\\sim 1.7\\times 10^4~M_\\odot$ has been reported in the G1 globular cluster (Gebhardt, Rich \\& Ho 2005) and a black hole with a mass of $\\sim 4\\times 10^4~M_\\odot$ may be present in Omega Centauri as well (Noyola et al. 2008). As for AGN, a limited number of low luminosity Seyfert--like active galaxies has been shown to harbour intermediate--mass black holes (IMBH, arbitrarily defined here as black holes in galaxy centers with $M_{\\rm BH} < 10^6~M_\\odot$). NGC~4395 is a late--type spiral with no bulge component and small central stellar velocity dispersion (Filippenko \\& Ho 2003) which exhibits broad permitted optical/UV emission lines, a compact radio core, and a highly variable point--like X--ray source unambiguously pointing towards a Seyfert identification (Filippenko, Ho \\& Sargent 1993; Ho \\& Ulvestad 2001; Ho et al 2001; Vaughan et al 2005). Mass estimates from the H$\\beta$ line width--luminosity relationship (Filippenko \\& Ho 2003) and from X--ray variability (Vaughan et al 2005) suggest a BH of $10^4-10^5~M_\\odot$, consistent with the upper limit ($10^5~M_\\odot$) inferred from the $M$--$\\sigma$ relation, but slightly lower than the BH mass of $3.6\\times 10^5~M_\\odot$, obtained through reverberation mapping (Peterson et al. 2005). The dwarf elliptical POX~52 also exhibits Seyfert~1 like broad and narrow emission lines (Kunth, Sargent \\& Bothun 1987; Barth et al. 2004) and the estimate of the BH mass obtained via the H$\\beta$ line width--luminosity relationship and the $M$--$\\sigma$ relationship point toward a BH of the order of $10^5~M_\\odot$. The main difference in the AGN properties of the two IMBH is that NGC~4395 is a low Eddington ratio source ($<10^{-2}$), while POX~52 appears to be radiating at nearly its maximum rate. Finally, the dwarf Seyfert~1 SDSS~J160531.84+174826.1 has a relatively broad H$\\alpha$ line component which suggests (via luminosity and line width) a black hole mass of $\\sim 7\\times 10^4~M_\\odot$ (Dong et al. 2007b).  Greene \\& Ho (2004) have defined a sample of 19 IMBH candidates in AGN using the first data release of the Sloan Digital Sky Survey (SDSS). Given the adopted selection criteria (low BH mass and prominent AGN--like spectrum) the sample mostly comprises objects with high Eddington ratio. The Greene \\& Ho work was recently extended to the fourth data release by Dong et al (2007a) and Greene \\& Ho (2007b) who have increased the sample of IMBH by about one order of magnitude. Interestingly, about 15 per cent of the larger sample(s) are suspected to have Eddington ratios below 10 per cent. Due to selection criteria, and with the exception of NGC~4395, all AGN with IMBH are actually radiating a very significant fraction of their Eddington luminosity. The presence of relatively low Eddington ratio sources in the new sample(s) will make it possible to investigate NGC~4395--like nuclei in the future, opening a new important window on IMBH activity.  The original smaller Greene \\& Ho sample (hereafter the GH sample) has been studied in the X--rays by making use of 5~ks long {\\it Chandra} snapshot observations (Greene \\& Ho 2007a). The sample is characterized by X--ray luminosities in the range of $10^{41}$--$10^{43}$~erg~s$^{-1}$ in the soft 0.5--2~keV X--ray band and the soft X--ray spectral properties are consistent with those of more luminous sources harbouring more massive black holes, showing that black hole mass is certainly not one of the main drivers of soft X--rays spectral properties.  Here we report results from {\\it XMM--Newton} 40~ks--long observations of four of the IMBHs from the GH sample which were previously detected in the Rosat All Sky Survey. The larger effective area and longer exposure of these observations with respect to the {\\it Chandra} snapshots allow us to consider in some detail the X--ray spectral properties of these IMBH also above 2~keV, and to compare them with their more massive and more luminous counterparts (e.g. the PG quasars).  Moreover, we were able to also study their X--ray variability properties filling (although with only four objects) a relatively poorly studied range in black hole masses. Very recently, the four observations we consider here have been included in a submitted work by Dewangan et al (2008). Their analysis is mainly devoted to NGC~4395 and POX~52. As for the common sources, the main difference with our work is that we focus much more our study on the spectral analysis and soft excess, exploring different possible models and interpretations, and comparing the X--ray properties of the four objects with AGN powered by accretion onto much more massive black holes. ",
        "conclusions": "We have observed with {\\it XMM--Newton} 4 IMBH from the original sample of 19 objects selected by Greene \\& Ho (2004) from the SDSS. The X--ray properties of the IMBH considered here do not significantly differ from those of their more massive counterparts such as the PG quasars. The IMBHs could represent a population of not yet fully grown super--massive black holes and, if in a rapid growing phase, their accretion properties would be expected to significantly depart from those of standard and fully--grown accreting supermassive black holes. However, as already suggested by their Eddington ratios (high because of the sample selection strategy, but not that unusual), their accretion properties are in all respect similar to those of more massive AGN and quasars.  The X--ray spectra of our objects are characterized by the same X--ray slope as PG quasars in both the hard and soft X--ray bands. No (neutral) absorption in excess of the Galactic one is observed in any of the objects. Three out of four objects exhibit a soft excess with the same properties (uniform temperature and strength) as larger black hole mass samples (standard Seyfert~1 and quasars). This is surprising if one considers that the small black hole mass of the IMBHs should allow thermal emission from the accretion disc to dominate the soft X--ray band, producing soft X--ray spectra that should significantly depart from those of more massive objects in which the disc quasi--blackbody emission is expected to contribute in the UV only. This is however not observed, and if the soft excess is interpreted as thermal disc emission, the inferred temperature in AGN with black hole masses of $\\leq 10^6~M_\\odot$ is the same as that of quasars with black hole masses two (or more) orders of magnitude larger. So far, the so--called soft excess problem was mainly defined as the fact that we observe soft X--ray emission in accreting supermassive black hole which is too hot (and too uniform) to be reasonably associated with thermal disc emission. Here, we point out that the problem is even more puzzling: when we observe sources in which thermal disc emission should largely extend into the soft X--ray band (such as the IMBH), we see exactly the same spectral shape as in more massive AGN and, if these IMBH were rotating (Kerr), the inferred temperature would even be too cold to be associated with the accretion disc.  \\subsection{Smeared absorption}  Gierli\\'nski \\& Done (2004) proposed that the soft excess is an artifact due to absorption (of an intrinsic steep power law spectrum) by partially ionized gas in the form of a highly relativistic wind. Since the observed smoothness of the soft excess is not consistent with the expected sharp absorption features, the model requires the gas to be characterized by strong velocity gradients to smear them out via Doppler smearing. The large velocity gradients, most likely imply that the wind is launched very close to the black hole.  As mentioned by Schurch \\& Done (2006), the wind model described above is unlikely to be a standard UV--line driven disc wind, because such winds are preferentially produced at $i>75^\\circ$ of the axis, which is in contrast with the observational fact that most low--inclination Seyfert 1 galaxies exhibit a soft excess. Moreover, the mass outflow rate ($\\dot M_{\\rm out}$) inferred from the large final velocities required by the model is generally hundreds times larger than the mass accretion rate ($\\dot M_{\\rm acc}$) in among the most luminous X--ray sources in the Universe (but see the idea of a failed wind proposed by Schurch \\& Done 2006). More recently, Schurch \\& Done (2007) have shown that when a more detailed smearing profile is adopted, the shape of the model does not describe well the observed smoothness of the soft excess, casting severe doubts on the overall model.  In addition, here we have shown that, even if the smeared absorption model can reproduce the spectral shape, the uniform soft excess strength is obtained only if a tight correlation between column density and ionization parameter is enforced. The $N_H-\\xi$ correlation inferred from fitting the smeared absorption model to X--ray data is however not a natural product of photo--ionization and is most likely spurious, either due to a degeneracy in the parameter space (which implies the model does not provide any constraint on the absorber properties) or, as mentioned above, to the need of producing a uniform soft excess strength in sources which do require different column densities.   \\subsection{Disc reflection}   The smearing of the reflection features by high velocities is a natural consequence of the presence of an inner accretion disc. Undoubtedly, an inner accretion disc must exists in sources accreting above a few per cent of the Eddington ratio and thus, unlikely the smeared absorption model, the disc reflection model does not introduce any new (so far unobserved, if not for the soft excess interpretation itself) component to explain the soft excess. The reflection interpretation just makes use of the basic ingredients of accretion theory, a powerful primary source of hard X--rays, and an optically thick accretion disc which reprocesses the irradiating flux into a reflection spectrum. An additional external absorber may obviously mask this contribution and this may well be the case in some most extreme sources. However, since disc reflection only requires the presence of an inner accretion disc and of an illuminating X--ray source, it is an unavoidable component in any reasonable accretion model for radiatively efficient AGN (and X--ray binaries). In fact (see e.g.  Middleton, Done \\& Gierli\\'nski 2007) when the smeared absorption model is applied to PG quasars and NLS1s, it always does require also a disc reflection component. If the disc reflection component alone can account for the spectral shape and variability of the sources, it is difficult to understand the need for an additional smeared absorber with problematic launching mechanism, extreme mass outflow rate, and rather unnatural relationship between column density and ionization.  Since disc reflection invokes the presence of the accretion disc down to the last stable orbit, thermal emission from the disc should also be present and, given the small black hole masses and relatively high accretion rates in our IMBHs, such emission should peak in the soft X--rays. Indeed, when a multicolor disc blackbody is included in the spectral models, a solution in which the soft excess is partly due to reflection and partly to disc blackbody with the theoretically expected temperature is found. However, we point out that if thermal disc emission is present in the soft X--ray spectrum of these objects, its luminosity is far below the expected one for the given temperatures not only when reflection is invoked, but even when all the soft excess is modelled as a pure disc blackbody.  It is actually difficult to distinguish between the two competing models (smeared absorption and disc reflection) by using spectral and variability information in the limited band--pass of {\\it XMM--Newton} (0.3--10~keV band) even in bright, well observed sources. Moreover, our mini--sample of small mass BH is not well suited to investigate the problem due to the relatively low quality of the data. High--energy data are crucial to test whether the X--ray spectrum above 10~keV is just the high--energy unabsorbed tail of the intrinsic power law continuum, or whether it is instead characterized by the presence of a Compton hump around 20--30~keV, related to X--ray reflection. While signatures for the presence of X--ray reflection from the inner accretion disc are sometimes detected, the number of well observed sources above 10~keV is still too small to draw any clear--cut conclusion (e.g. Miniutti et al 2007; Reeves et al 2007). Observations with the {\\it Suzaku} X--ray mission and with future missions such as Simbol--X (Ferrando et al. 2006) will most likely play a crucial role in that direction in the near future.  \\subsection{X--ray variability}  Given their relatively small black hole mass (i.e. small size), the observed IMBH are not surprisingly among the most variable in X--rays. In particular, their excess variance $\\sigma^2_{\\rm NXS}$ is among the largest obtained from AGN X--ray light curves. Our observations begin to fill a relatively poorly explored range of black hole masses in the $\\sigma^2_{\\rm NXS}$--$M_{\\rm BH}$ relationship and show that the X--ray variability properties of IMBH smoothly join with those of more massive Seyfert galaxies. The $\\sigma^2_{\\rm NXS}$--$M_{\\rm BH}$ is consistent with a simple (well known) anti--correlation that can be explained with a universal PSD model in which the break frequencies (most importantly the high--frequency one) scale with $M^{-1}_{\\rm BH}$. However, such functional form for $\\nu_{\\rm H}$ is known to be only a zeroth--order approximation (see M$^{\\rm c}$Hardy et al 2006) and the universal PSD model suffers also for other uncertainties that do not allow to properly account for the scatter in the relationship. Future longer observations of these and other IMBH in X--rays would allow to explicitly search for the high frequency break in the PSD, enabling us to extend the $\\nu_{\\rm H}$--$M_{\\rm BH}$--$\\dot M_{\\rm{acc}}$ relationship pointed out by M$^{\\rm c}$Hardy et al (2006) to the most crucial black hole mass range of $M_{\\rm BH} \\leq 10^6~M_\\odot$."
    },
    "0812/0812.4288_arXiv.txt": {
        "abstract": "We show that a suitably defined marked correlation function can be used to break degeneracies in halo-occupation distribution modeling.  The statistic can be computed on both 3D and 2D data sets, and should be applicable to all upcoming galaxy surveys.  A proof of principle, using mock catalogs created from N-body simulations, is given. ",
        "introduction": "In recent years our ability to describe galaxy clustering has advanced dramatically. The halo model (\\citealt{Sel00,PeaSmi00}, see e.g.~\\citealt{CooShe} for a review) has provided us with a physically informative and flexible means of describing galaxy bias - the relation between galaxies and the underlying dark matter halos. The key insight is that an accurate prediction of galaxy clustering requires knowledge of how galaxies are apportioned between and distributed within halos - the halo occupation distribution or HOD. Combined with the theoretically predicted spatial distribution of halos from e.g.~N-body simulations, a specified HOD makes strong predictions about a wide array of galaxy clustering statistics.  The formalism is now widely used in the interpretation of galaxy clustering and to infer cosmological parameters from large-scale galaxy surveys.  Much of the recent work on fitting HODs has used the two-point galaxy correlation function as the observation of choice. While the galaxy correlation function provides very strong constraints on the HOD, there exist degeneracies between the inferred HOD and the underlying cosmology. Much of this degeneracy arises because a change in the cosmology, and hence the halo population, can be compensated to a large extent by a change in the galaxy halo occupation.  This modified halo occupation apportions galaxies differently amongst halos of different mass than the fiducial model.  Not surprisingly, combining the galaxy correlation function with a second observable with different sensitivities to the HOD can lift such degeneracies \\citep{ZheWei07} tightening the constraints and allowing one to simultaneously constrain the cosmological world model and HOD \\cite[see e.g.][]{Aba05}. A number of such observables have been considered in the literature. Galaxy-galaxy lensing has the potential to directly measure the mass of the halos hosting a galaxy population. Cross-correlating with another galaxy sample selected to live in high (or low) mass halos can help to break degeneracies, as can redshift space distortions (which are very sensitive to the satellite fraction) or peculiar velocities. Another observation is the abundance of rich clusters of galaxies, which can constrain the number density of massive halos. While all of these approaches are certainly valid, and will continue to be used in the future, a disadvantage is that they generally require additional observations or measurements, with the associated modeling and additional systematics that must be calibrated. A natural question, therefore, is whether one can break these degeneracies using only the data going into the clustering statistics themselves?  Higher order clustering measurements provide one such approach \\citep[e.g.][]{Kul07}. A particularly convenient choice for our purposes is a marked correlation function, where the mark is determined from the galaxy spatial distribution. While this represents new information, it is information which is readily available. We demonstrate here that appropriately chosen marked two-point correlation functions \\citep{BeiKer00,BeiKerMec02,Got02,SheTor04,SheConSki05,Har06,Wec06} can lift degeneracies in the HOD or between the HOD and the cosmology. Such marked correlation functions are straightforward to compute with the same set of observations, and require no additional data nor understanding of the survey or algorithms beyond those required for basic clustering statistics. This paper serves as a proof-of-principle, by demonstrating that the degeneracy between HOD and the amplitude of the primordial fluctuation spectrum is broken with a simple density mark in simulations. ",
        "conclusions": "\\label{sec:conclusions}  Although the galaxy two-point correlation function has proved to be extremely useful in modeling the relationship between galaxies and dark matter, it does not exhaust the information in the data. One degeneracy that cannot be broken by the correlation function alone is between the HOD and cosmology - within the context of the correlation function, one is free to re-apportion galaxies to compensate for differences in the halo mass function. The number of such degeneracies will only increase as we attempt more detailed mappings of galaxies to dark matter in the future. This note provides proof of principle that such degeneracies can be lifted by marked correlation functions.  An important advantages of marked correlations is that they do not involve multiple data sets.  Indeed, for the example presented here, the mark required {\\it no\\/} additional information beyond the spatial distribution of galaxies (and survey mask), which one needed to compute the correlation function in the first place. Using the same data set considerably simplifies any modeling step. Furthermore the marked correlation function can be computed with the same code and at the same time as the standard correlation function, for which optimized algorithms exist. This is a non-negligible advantage when one considers the need to repeatedly compute it for mock samples while modeling, estimating errors, etc.   Rapid advances in computational power and algorithm development have made it reasonably straightforward to simulate the distribution of dark matter halos in large volumes for almost any cosmological model. Combined with a halo occupation approach this makes ``forward modeling'' of almost any galaxy statistic possible. This implies that statistics which use all of the galaxy information, in the manner of standard large-scale structure $n$-point statistics, can be just as useful as those which try to identify special subsets of galaxies (`cluster' vs.~`field'). As our ability to model a wider array of observations (e.g.~weak lensing, Sunyaev-Zel'dovich decrement or X-ray flux) matures, similar methods can be applied to these observations, bypassing the need to relate particular features in a map with individual 3D structures."
    },
    "0812/0812.4307_arXiv.txt": {
        "abstract": "Neutrons from ($\\alpha$,n) reactions through thorium and uranium decays are important sources of  background for direct dark matter detection.  The neutron yields and energy spectra from a range of materials that are used to build dark matter detectors are calculated and tabulated.  In addition to thorium and uranium decays, we found that $\\alpha$ particles from samarium, often the dopant of the window materials of photomultiplier tubes (PMT), are also an important source of neutron yield. The results in this paper can be used as the input to Monte Carlo simulations for many materials that will be used for next generation experiments. ",
        "introduction": "Neutron induced elastic scattering processes represent an important background for direct dark matter detection experiments searching for Weakly Interacting Massive Particles (WIMPs), which may constitute the dark matter in the universe~\\cite{dns, wfr, mwg, gjm}. Direct searches for WIMPs have been carried out by many experiments including CDMS~\\cite{cdms1}, EDELWEISS~\\cite{edel1}, Xenon10~\\cite{xenon10}, ArDM~\\cite{argon}, DAMA~\\cite{dama1}, CRESST~\\cite{cresst1}, PICASSO~\\cite{pica1}, NAIAD~\\cite{naia}, and ZEPLIN~\\cite{zep1}. Among these experiments, DAMA/NaI~\\cite{bern1}  and DAMA/LIBRA~\\cite{bern2} have claimed that they have observed a model independent annual modulation signature. The DAMA collaboration interprets this annual modulation as the signature induced by dark matter (DM) particles~\\cite{bern1,bern2}.  However, this claim is at odds with other experimental results if one assumes standard WIMP interactions and halo models. With the best limits set by CDMS II~\\cite{cdms2} and Xenon10, WIMPs remain unobserved. Large scale next generation detectors utilizing noble liquids to continue the direct search for WIMPs are underway. The key to these experiments lies in the ability to reduce various background to unprecedented low levels. Among all possible sources of background, neutron induced nuclear recoil is identified as a major source of background for this type of experiments.  There are three sources of underground neutrons: 1) those produced by ($\\alpha,n)$ reactions through thorium, uranium, and other radioactive isotope decays in the materials that surround or constitute the detector; 2) those from spontaneous uranium fission; and 3) those from cosmic ray muon-induced processes. In general, the ($\\alpha,n)$ neutrons dominate the total neutron contributions to the measured background for an underground experiment. This is because the origins of ($\\alpha, n)$ neutrons range from surrounding rock, external shielding, inner shielding, detector components, and detector target. In particular, the ($\\alpha,n)$ neutrons produced in the detector components and target are hard to cope with. These neutrons need to be understood very well in terms of their origin, transport, and interaction with materials. Calculations of the neutron yield and neutron energy spectrum in different materials are critical to dark matter experiments. The total neutron yield indicates the number of neutrons that enter or are produced in the target. The neutron energy spectrum determines the total background events in the region of interest (ROI).  Therefore, a full simulation of neutron background must take into account both neutron yield and energy spectrum. This paper is aimed at providing the ($\\alpha,n)$ neutron yield and energy spectrum for a number of materials that are used for the construction of dark matter detectors. ",
        "conclusions": "We have calculated the ($\\alpha$,n) neutron yield and energy spectrum for many elements that are used to build low background experiments. Both neutron yields and energy spectra are important to the contribution in the energy region of interest. Therefore, it is critical to have information from both in the Monte Carlo simulation to predict the possible contributions in the energy region of interest from the ($\\alpha$,n) neutrons. The neutron yield from many elements are compared to Heaton {\\it et al.}~\\cite{heat1990}. Good overall agreement is obtained."
    },
    "0812/0812.3711_arXiv.txt": {
        "abstract": "A survey for the molecular clouds in the Galaxy with NANTEN mm telescope has discovered molecular loops in the Galactic center region. The loops show monotonic gradients of the line of sight velocity along the loops and the large velocity dispersions towards their foot points. It is suggested that these loops are explained in terms of the buoyant rise of magnetic loops due to the Parker instability. We have carried out global three-dimensional magneto-hydrodynamic simulations of the gas disk in the Galactic center. The gravitational potential is approximated by the axisymmetric potential proposed by \\citet{miy1975}. At the initial state, we assume a warm ($ \\sim 10^4{\\rm K}$) gas torus threaded by azimuthal magnetic fields. Self-gravity and radiative cooling of the gas are ignored. We found that buoyantly rising magnetic loops are formed above the differentially rotating, magnetically turbulent disk. By analyzing the results of global MHD simulations, we have identified individual loops, about 180 in the upper half of the disk, and studied their statistical properties such as their length, width, height, and velocity distributions along the loops. Typical length and height of a loop are 1kpc and 200pc, respectively. The line of sight velocity changes linearly along a loop and shows large dispersions around the foot-points. Numerical results indicate that loops emerge preferentially from the region where magnetic pressure is large. We argue that these properties are consistent with those of the molecular loops discovered by NANTEN. ",
        "introduction": "Spiral galaxies generally possess large-scale magnetic fields (e.g., \\cite{sof1986, bec1996}). These magnetic fields are usually considered to be amplified and maintained by large-scale dynamo action driven by collective inductive effects of turbulence and differential rotation \\citep{par1971}. In our Galaxy, the strength of the magnetic field is about a few $\\mu$ G. Inside 200pc from the Galactic center, however, radio observations indicate that magnetic fields are as strong as a few mG (e.g., \\cite{mor1996}). Recent observations of other spiral galaxies indicate that the typical averaged equipartition strength of the total magnetic fields is about $10 \\mu$G \\citep{bec2008}.   Based on a survey for the molecular gas in the Galaxy by NANTEN (e.g., \\cite{miz2004}). \\citet{fuk2006} found loop-like structures of dense molecular gas showing large line-of-sight velocity gradients ($\\sim 30 {\\rm km/s}$) along the loops and large velocity dispersions at their foot-points. The total mass of the molecular gas estimated from the observation is about $1.7 \\times 10^5 M_{\\odot}$ when LTE at 50K was assumed at a distance of 8.5 kpc. The kinetic energy of the loops estimated from their line-of-sight velocity gradient exceeds $10^{51}$ ergs. This energy exceeds the energy injected by a supernova and requires other energetic mechanisms driving the molecular gas.  The velocity gradient along the loops and large velocity dispersions at their foot-points are typical features of solar magnetic loops emerging from  below the photosphere. The emergence of solar magnetic loops from the photosphere to the corona is driven by the buoyancy created by sliding the gas along the loop. Even when the unperturbed atmosphere is convectively stable, the magnetic loop can rise when the buoyancy at the loop top exceeds the restoring magnetic tension. This instability is called the Parker instability. Parker instability was originally proposed as a mechanism of the formation of the interstellar clouds  \\citep{par1966}. \\citet{mat1988} carried out two-dimensional magneto-hydrodynamic (MHD) simulations of the Parker instability in astrophysical disks and showed that in the nonlinear stage of the Parker instability, shock waves are formed near the foot-points of the magnetic loops because gas slides down supersonically(e.g., \\cite{shi1989}). Such shocks can be an origin of large velocity dispersions observed near the foot-points of Galactic center molecular loops \\citep{fuk2006}.  Two-dimensional MHD simulations of the Parker instability in the Cartesian coordinate system of the gas disk can reproduce loop-like structures and velocity gradients similar to the observed loops \\citep{fuk2006}. In their simulations, however, the Galactic rotation was not taken into account. In rotating stratified magnetized layer such as galactic gas disk, Coriolis force twists the magnetic loops. Since the twisting of magnetic fields enhances magnetic tension, the growth rate of the Parker instability is suppressed. \\citet{cho1997} carried out local three-dimensional MHD simulations of the Parker instability by using a frame corotating with the disk to consider the effect of the Coriolis force. They confirmed that the Parker instability grows in rotating disks, although growth rate becomes smaller than that without rotation.  In the gas disk rotating differentially, another instability, magneto-rotational instability (MRI) grows \\citep{bal1991}. Three-dimensional local 3D MHD simulations adopting the shearing sheet approximation confirmed that MRI drives magnetic turbulence inside the disk (e.g., \\cite{haw1995, mat1995, bra1995, sto1996, san1999}). When  we take into account the vertical gravity, buoyantly rising magnetic loops can be formed by the Parker instability. In the Galactic center region, since the loop size must be comparable to the distance from the Galactic center, non-local effects such as curvature becomes important. Global MHD simulations of the buoyant rise of magnetic loops from differentially rotating gas disks were first carried out by \\citet{mac2000}. They showed that magnetic loops emerge from the turbulent disk. They assumed the gravitational force created by the central point mass. In Galactic gas disks, however, gravity is determined by stars and the dark matter.  \\citet{nis2006} presented the results of global three-dimensional MHD simulations of Galactic gaseous disks to study how the Galactic magnetic fields are amplified and maintained. They assumed a steady axisymmetric gravitational potential given by \\citet{miy1975}. Their simulations showed that magnetic fields are amplified up to $\\mu$ G and are maintained for 10 billion years. They also found that mean azimuthal magnetic fields reverse their direction in every 1Gyr. This field reversal is driven by the buoyant escape of magnetic flux from the disk to the Galactic halo. Their simulation region ($0.8 {\\rm kpc} < \\varpi < 30 {\\rm kpc}$), however, did not include the galactic center. Here, $\\varpi$ is the distance from the Galactic center.  \\citet{bae2008} proposed that collision of the high velocity cloud with the Galactic disk explains the formation of the molecular loop structure seen at the Galactic center, when the high velocity cloud has an oblique collision with the disk plane. They showed that the size was similar to that of the molecular loop, but the velocity gradient was smaller than that of the observation.  In this paper, we report the results of global 3D MHD simulations of Galactic central gas disks inside 1kpc from the Galactic center. In section \\ref{sec2}, we present basic equations and describe the numerical model. Numerical results will be presented in section \\ref{sec3}. Section \\ref{sec4} is devoted for discussion. Summary will be given in section \\ref{sec5}.   \\newpage ",
        "conclusions": "\\label{sec4} \\subsection{Loop Formation and the Origin of Their Concentration} As shown in figures \\ref{fig:mfl3d} and \\ref{fig:ftbz}, magnetic loops are formed in the halo region ($z>0.1$). We identified about 400 loops summing up the upper and lower coronal region. These magnetic loops are concentrated both in certain azimuthal range and radial range (figure \\ref{fig:hist}c, and \\ref{fig:hist}d). Magnetic fields inside the gaseous disk have turbulent component whose amplitude is comparable to the mean azimuthal magnetic fields. Therefore, the magnetic energy shows random patchy structure (see figure \\ref{fig:ftbz}). On the other hand, one-armed non-axisymmetric pattern develops for density. Magnetic pressure in the low density region becomes stronger than that in the higher density region, because the total pressure is balanced in the disk. Since the magnetic loops with $\\beta \\sim 1$ are the most unstable to the Parker instability, buoyantly rising magnetic loops are formed in such lower density region. Thus, non-axisymmetric density distribution results in the azimuthal concentration of magnetic loops.   \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){fig9.ps} \\end{center} \\caption{Schematic picture showing the formation of magnetic loops. \\label{fig:sch}} \\end{figure}  Figure \\ref{fig:sch} shows schematically how magnetic loops are formed. As MRI grows, channel flows (e.g., \\cite{haw1992, san2001, mac2003}) create large scale horizontal magnetic fields (figure \\ref{fig:sch}b). When the magnetic energy accumulated in the channel flow becomes comparable to the thermal energy, magnetic loops emerge from the disk by buoyancy (figure \\ref{fig:sch}c, d). This process creates chains of magnetic loops.   \\subsection{Comparison with Observations}\\label{sub:comp} From their millimeter-wave CO observations,  \\citet{fuk2006} found two molecular features that look like a loop with a length of several hundred parsecs and width of $\\sim 30$ pc within $\\sim 1$kpc from the center of our Galaxy. The kinetic energy of these CO loops is about $\\sim 10^{51}$ erg. As shown in figure \\ref{fig:ftbz} and  \\ref{fig:hist}, our numerical simulation produced magnetic loops whose height $\\sim 200$pc, width $\\sim 50-300$pc, and length $\\sim 0.1-2$kpc. To compare our simulation with observations, we assumed that the initial gas density is $\\rho_{\\rm b}= 1.6 \\times 10^{-24} {\\rm g/cm}^3$. In this case the total energy in a loop is about $10^{51}$ {\\rm erg}.    \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){fig10.ps} \\end{center} \\caption{Position-velocity diagram of the mass elements in the loops depicted in green in figure \\ref{fig:ftbz}. This viewgraph shows the P-V diagram observed from the direction  $\\theta = 3\\pi /8 $. \\label{fig:pvd} } \\end{figure}  In CO loops observed in the Galactic center, the line of sight velocity decreases linearly with the length along the loop (see fig.1 in \\cite{fuk2006}). Figure \\ref{fig:pvd} shows a  position-velocity diagram obtained by our simulation. We plotted magnetic loops colored in green in figure \\ref{fig:ftbz}. The horizontal axis shows the $x$-direction in the Cartesian coordinate system and the vertical axis shows $y$-component of the velocity. We assumed that we observed the loops from infinity. Figure \\ref{fig:pvd} shows the linear decrease of the line of sight velocity along the loop. Numerical results indicate that the rotation of the Galactic center gas disk creates this velocity gradient. From figure \\ref{fig:pvd}, we can clearly see that foot points show the larger velocity dispersion than the loop top. It can be explained by the large velocity gradient around the foot-points of the magnetic loops where supersonic flow along the magnetic loops collides with the dense disk gas.  Figure \\ref{fig:vren} shows that a large-scale poloidal magnetic field is created in the innermost region of the Galactic gas disk. \\citet{kat2004} showed that magnetic towers are formed by twisting of magnetic loops emerging from a magnetically turbulent accretion disk. Such magnetic towers can be formed even when the initial magnetic field is purely azimuthal \\citep{mac2008}. The large scale poloidal magnetic fields penetrate the gas disk.  \\subsection{Mass Supply to the Central 100pc}  In our simulations, the unit of mass accretion rate is given by \\begin{equation} \\dot{M}_0  = 33 \\left( \\frac{\\varpi}{1{\\rm kpc}} \\right)^2 \\left( \\frac{v}{207 {\\rm km/s}}  \\right) \\left( \\frac{\\rho} {1.66\\times 10^{-24} {\\rm g/cm^3}} \\right) ~ M_{\\odot} /{\\rm year} ~ . \\end{equation} According to figure \\ref{fig:taccr}b, mass accretion rate at $200 {\\rm pc}$ in a quasi-steady state is $\\dot{M} \\sim 10^{-5} \\dot{M}_0  \\sim 3.3 \\times 10^{-4} M_{\\odot} / {\\rm yr}$. This mass accretion rate is determined not by the mass of the central black hole but by the mass of the interstellar gas. Note that the initial density of the 1kpc gas ring we assumed as an initial condition is arbitrary in our simulation. Thus, when dense gas infalls in this region, the accretion rate can be much higher.  The mass accretion rate near the event horizon of the Galactic central black hole ${\\rm SgrA^*}$  is estimated to be less than $ \\leq 10^{-8} M_{\\odot} {\\rm yr}^{-1}$, assuming an equipartition magnetic field (e.g., \\cite{bag2003}). Our simulation indicates the mass accretion rate at $ \\varpi = 200$ pc becomes about $\\simeq 10^{-4} M_{\\odot} {\\rm yr}^{-1}$. This may show that the mass accretion rate decreases from $10^{-4} M_{\\odot} {\\rm /yr}$ to $10^{-8}M_{\\odot}/ {\\rm yr}$ in $0.01 {\\rm pc} < \\varpi < 200 {\\rm pc}$ by some mechanism such as star formation, supernova explosions, or by mass outflows.    \\subsection{Effects of Cooling} In this paper, we ignored the cooling of the gas. Therefore, the gas temperature of the gas disk exceeds $10^5 {\\rm K}$. Since this temperature is higher than that of the warm component observed in the Galactic gaseous disk, the disk scale height may be overestimated. Since the most unstable wavelength of the Parker instability is  equal to 10 times of the disk scale height, the loop length in our simulation can be longer than that of the real galaxy. When we consider the radiative cooling, the interstellar gas will decompose into cold phase and the warm phase. Dense molecular clouds will be formed in the foot-points of magnetic loops where interstellar gas accumulates.  In section \\ref{sub:comp}, we compared the position-velocity diagram expected from our simulation with the observation. The velocity dispersion at the foot-points is smaller in our simulations than observations. Since the gas temperature and sound speed are overestimated in our simulations, shock waves formed at the loop foot points are weaker than that expected for simulations including cooling. Thus, when we include the cooling, the position-velocity diagram may indicate a large jump similar to that observed in the Galactic center molecular loops."
    },
    "0812/0812.1714_arXiv.txt": {
        "abstract": "{} {We attempt to understand the presence of gas phase CO below its sublimation temperature in circumstellar disks. We study two promising mechanisms to explain this phenomenon: turbulent mixing and photodesorption.} {We compute the chemical evolution of circumstellar disks including grain surface reactions with and without turbulent mixing and CO photodesorption.} {We show that photodesorption significantly enhances the gas phase CO abundance, by extracting CO from the grains when the visual extinction remains below about 5 magnitudes. However, the resulting dependence of column density on radial distance is inconsistent with observations so far. We propose that this inconsistency could be the result of grain growth. On the other hand, the influence of turbulent mixing is not found to be straightforward. The efficiency of turbulent mixing depends upon a variety of parameters, including  the disk structure. For the set of parameters we chose, turbulent mixing is not found to have any significant influence on the CO column density.} {} ",
        "introduction": "In circumstellar disks, because of the high densities, carbon monoxide is expected to stick onto grains at temperatures below about 17 K. However, millimeter interferometric observations of several circumstellar disks around T Tauri stars (e.g. DM Tau, GM Aur, and LkCa15) have revealed that a large amount of CO remains gaseous at temperatures as low as 10 K \\citep{Dartois03,Pietu07}. These authors show that the bulk of CO is at temperatures below 17 K between 100 and 800 AU from the central star. Since these disks are much older than the sticking timescale of CO onto grains ($\\sim 10^2$ yr), some mechanism must prevent the complete depletion of CO.  \\cite{Aikawa06} proposed that the presence of cold CO in disks can be a consequence of vertical turbulent mixing \\citep[See also Fig. 4 of][]{Willacy06}. \\cite{Semenov06} studied in detail the influence of both vertical and radial mixing on the abundance of cold CO. They found that vertical mixing cannot account for the observed abundance of cold CO, while the situation improves slightly when considering the full 2D mixing (vertical and radial). However, \\cite{Aikawa07} computed analytically the influence of vertical mixing on the abundance of cold CO and concluded that vertical mixing was sufficient to achieve an abundance of CO consistent with that observed by \\cite{Dartois03}. The apparent discrepancy with \\cite{Semenov06} can be understood as a consequence of the different efficiencies of grain surface reactions in their studies \\citep{Aikawa07}.  On the other hand, \\cite{Oberg07} measured  the photodesorption rate of CO in the laboratory. They found a rate similar to that found for H$_2$O by \\cite{Westley95}, about 2 orders of magnitude higher than previously available theoretical estimates \\citep{Draine79}, implying that this non-thermal desorption process was an appealing candidate for explaining the abundance of cold CO in various astrophysical environments. This motivated us to investigate the influence of both vertical mixing and photodesorption on the CO column density in disks, using a parametric model consistent with interferometric observations of circumstellar disks.  In Sect. \\ref{model}, we describe the physical and chemical model used, while Sect. \\ref{results} is devoted to the results of our computations. In Sects. \\ref{discussion} and \\ref{conclusions}, we discuss our results and present our conclusions. ",
        "conclusions": "\\label{conclusions}  We have shown that CO photodesorption appears to be a good candidate to explain the abundance of cold gas phase CO in circumstellar disks. However, it tends to create a flatter radial profile for the CO column density, in contrast to that observed in disks. Therefore, more work is necessary to solve this problem. Consideration of the photodissociating UV flux and variations in its penetration inside the disk due to grain growth and geometry seem to provide an avenue to improve our modelling of circumstellar disks. Moreover, we have shown that the influence of vertical mixing on the abundance of CO is not straightforward.  Rather, the efficiency of vertical mixing depends upon many parameters, including the details of the disk structure (density and temperature) and the level of grain growth, all of which remain uncertain today."
    },
    "0812/0812.2927_arXiv.txt": {
        "abstract": "We present deep {\\it Spitzer} mid-infrared spectroscopy, along with 16, 24, 70, and 850\\,$\\micron$\\ photometry, for 22 galaxies located in the Great Observatories Origins Deep Survey-North (GOODS-N) field. The sample spans a redshift range of $0.6\\la z \\la 2.6$, 24~$\\mu$m flux densities between $\\sim$0.2$-$1.2 mJy, and consists of submillimeter galaxies (SMGs), X-ray or optically selected active galactic nuclei (AGN), and optically faint ($z_{AB}>25$\\,mag) sources. We find that infrared (IR; $8-1000~\\micron$) luminosities derived by fitting local spectral energy distributions (SEDs) with 24~$\\micron$ photometry alone are well matched to those when additional mid-infrared spectroscopic and longer wavelength photometric data is used for galaxies having $z\\la1.4$ and 24~$\\micron$-derived IR luminosities typically $\\la 3\\times 10^{12}~L_{\\sun}$. However, for galaxies in the redshift range between $1.4\\la z \\la 2.6$, typically having 24~$\\micron$-derived IR luminosities $\\ga 3\\times 10^{12}~L_{\\sun}$, IR luminosities are overestimated by an average factor of $\\sim$5 when SED fitting with 24~$\\micron$ photometry alone. This result arises partly due to the fact that high redshift galaxies exhibit aromatic feature equivalent widths that are large compared to local galaxies of similar luminosities. Using improved estimates for the IR luminosities of these sources, we investigate whether their infrared emission is found to be in excess relative to that expected based on extinction corrected UV star formation rates (SFRs), possibly suggesting the presence of an obscured AGN. Through a spectral decomposition of mid-infrared spectroscopic data, we are able to isolate the fraction of IR luminosity arising from an AGN as opposed to star formation activity. This fraction is only able to account for $\\sim$30\\% of the total IR luminosity among the entire sample and $\\sim$35\\% of the ``excess\" IR emission among these sources, on average, suggesting that AGN are not the dominant cause of the inferred ``mid-infrared excesses\" in these systems. Of the sources identified as having mid-infrared excesses, half are accounted for by using proper bolometric corrections while half show the presence of obscured AGN. This implies sky and space densities for Compton-thick AGN of $\\sim$1600~deg$^{-2}$ and $\\sim$$1.3\\times 10^{-4}$~Mpc$^{-3}$, respectively. We also note that IR luminosities derived from SED fitting the mid-infrared {\\it and} 70~$\\micron$ broadband photometry agree within $\\sim$50\\% to those values estimated using the additional mid-infrared spectroscopic and submillimeter data. An inspection of the FIR-radio correlation shows no evidence for evolution over this redshift range. However, we find that the SMGs have IR/radio ratios which are a factor of $\\sim$3 lower, on average, than what is measured for star-forming galaxies in the local Universe. ",
        "introduction": "A reliable estimation of the star formation rates (SFRs), stellar mass, and active galactic nuclei (AGN) activity at all redshifts is required in order to obtain a comprehensive understanding of galaxy evolution. Various studies have shown that the SFR density increases by a factor of $\\sim$5$-$10 between $z\\sim0$ and $z\\sim1$ and that $\\gtrsim$60\\% of the total SFR density at $z\\sim1$ is obscured by dust \\citep{de02,ce01,el05,bm09}. Until now, these results have relied on deep  surveys at 15\\,$\\micron$ and 24\\,$\\micron$ which sample the redshifted mid-infrared spectral energy distribution (SED) of galaxies.  The mid-infrared ($5-40~\\micron$) SED is a complex interplay of broad emission features thought to arise from polycyclic aromatic hydrocarbon (PAH) molecules, silicate absorption features at 9.7 and 18~$\\micron$, and a mid-infrared continuum from very small grains \\citep{lp84,atb85}. Mid-infrared luminosities measured near $\\sim$8~$\\micron$ have been found to correlate with the total infrared (IR; $8-1000$) luminosities of galaxies in the local Universe \\citep{ce01, de02}, which itself is a measure of a galaxy's SFR. While a correlation is found, there does exist a large amount of scatter \\citep{dd05,jd07,la07} and systematic departures for low metallicity systems \\citep{ce05,sm06}. It is therefore uncertain if these empirical relations between mid-infrared and total IR luminosities ($L_{\\rm IR}$) for galaxies in the local Universe are applicable for galaxies at higher redshifts.  Recently, data from deep far-infrared surveys such as the Far-infrared Deep Legacy Survey (FIDEL; PI: M. Dickinson) are confirming the same redshift evolution of the SFR density seen in the mid-infrared and do not show any evolution in the SED of infrared luminous galaxies \\citep[e.g.][]{bm09}. Furthermore, individual lensed Lyman break galaxies (LBGs) at $z\\sim2.5$ \\citep{bs08, gw08, ag08}, seem to show bolometric corrections similar to those of local galaxies. In contrast, \\citet{jr08} have provided evidence that the high redshift lensed mid-infrared selected galaxies might show a factor of $\\sim$2 stronger rest-frame 8~$\\micron$ emission compared to their total IR luminosities. Part of the contradiction might be due to the intrinsic scatter in the physical properties of the galaxies, particularly the range of dust color temperatures and emissivities seen for galaxies in the local Universe \\citep[e.g.][]{ld00}. Individual galaxies detected in the far-infrared or submillimeter at high redshift, for which the comparisons are made, might be sampling the scatter in the ratio rather than the median trend which the empirical galaxy templates follow.  In order to assess the accuracy of mid-infrared derived bolometric luminosities as a tracer of star-formation at $z\\sim2$, \\citet{ed07a,ed07b} compared mid-infrared (i.e. observed-frame 24~$\\micron$) based SFRs to those derived from extinction corrected UV emission and radio emission. They found a systematic offset, whereby the mid-infrared derived bolometric luminosities were systematically higher than alternate SFR estimates, especially among the most intrinsically luminous mid-infrared sources. Furthermore, they found that stacking the X-ray data on these ``mid-infrared excess\" galaxies revealed the presence of a hard X-ray source suggestive of obscured AGN activity. Since mid-infrared wavelengths are sensitive to warm dust continuum emission arising from AGN activity, \\citet{ed07b} concluded that these mid-infrared excesses arise from deeply obscured AGN and report a space density of Compton thick AGN which is twice that of all X-ray detected AGN at $z\\sim2$.  The detection of a source in hard X-rays does not necessarily imply that an AGN is dominating its energetics as is found for submillimeter galaxies (SMGs) which appear to be both star-formation driven and commonly detected in the X-rays \\citep{da05, ap08a}. Furthermore, evidence for ``mid-infrared excess\" relative to extinction corrected UV luminosities might simply be evidence for optically thick star-formation, a phenomenon which is ubiquitous among ultraluminous infrared galaxies (ULIRGs; $L_{\\rm IR} \\geq 10^{12}~L_{\\sun}$) in the local Universe \\citep{bv07}. Mid-infrared spectroscopy of galaxy samples selected by their mid-infrared flux densities provides a powerful technique to isolate these various selection effects among infrared luminous sources. By decomposing these spectra, into PAH and continuum components, one can get a handle on the relative fraction of power being emitted by AGN activity within such sources relative to the output related to star formation \\citep[e.g.][]{as07,ap08a}.  The Great Origins Observatories Deep Survey-North \\citep[GOODS-N;][]{md03} field, due to the wealth of multiwavelength data spanning the ultraviolet to radio wavelengths, is the optimal field in which to undertake such a relatively unbiased survey. Using mid-infrared spectroscopy for 22 sources selected at 24~$\\micron$ in the GOODS-N field, along with existing 850~$\\micron$ and additional 70~$\\micron$ imagery obtained as part of FIDEL, we aim to improve our understanding of the star formation and AGN activity within a diverse group of $0.6\\la z \\la2.6$ galaxies. This is done through a more proper estimate of IR luminosities and the ability to decompose these measurements into star-forming and AGN components.  The paper is organized as follows: In $\\S$2 we introduce the sample and describe the observations. The determination of IR luminosities from SED fitting, calculation of SFRs, and the decomposition of AGN and star-formation power using the mid-infrared spectroscopy are described in $\\S$3. The results are presented in $\\S$4 while their implications are discussed in $\\S$5. Finally, in $\\S$6, we summarize our conclusions. ",
        "conclusions": ""
    },
    "0812/0812.3260_arXiv.txt": {
        "abstract": "{ We apply the scale-length method to several three dimensional samples of the Two degree Field Galaxy Redshift Survey. This method allows us to map in a quantitative and powerful way large scale structures in the distribution of galaxies controlling systematic effects.  By determining the probability density function of conditional fluctuations we show that large scale structures are quite typical and correspond to large fluctuations in the galaxy density field. We do not find a convergence to homogeneity up to the samples sizes, i.e. $\\approx 75$ Mpc/h.  We then measure, at scales $r\\ltapprox 40$ Mpc/h, a well defined and statistically stable power-law behavior of the average number of galaxies in spheres, with fractal dimension $D=2.2 \\pm 0.2$.  We point out that standard models of structure formation are unable to explain the existence of the large fluctuations in the galaxy density field detected in these samples. This conclusion is reached in two ways: by considering the scale, determined by the linear perturbation analysis of a self-gravitating fluid, below which large fluctuations are expected in standard models and through the determination of statistical properties of mock galaxy catalogs generated from cosmological N-body simulations of the Millenium consortitum} \\begin{document} ",
        "introduction": "In the past twenty years observations have provided growing evidences that galaxy distribution is organized in a complex network of structures and voids \\cite{devac1970,cfa2,colless01,york}.  Despite the fact that large scale galaxy structures, of size of the order of several hundreds of Mpc/h\\footnote{We use $H_0=100 h$ km/sec/Mpc for the value of the Hubble constant.}, have been observed to be the typical feature of the distribution of visible matter in the local universe, the statistical analysis measuring their properties has identified a characteristic scale which has only slightly changed since its discovery fourthy years ago in angular catalogs. This scale, $r_0$, was measured to be the one at which fluctuations in the galaxy density field are about twice the value of the sample density and it was indeed determined to be $r_0 \\approx$ 5 Mpc/h in the Shane and Wirtanen angular catalog \\cite{tk69}. Subsequent measurements of this scale --- see e.g. \\cite{dp83,davis88,park,benoist,norbergxi01, norbergxi02,zehavietal02,zehavietal04} --- found a similar value, although in several samples larger values of $r_0$ have been found (i.e. $r_0 \\approx 6-12$ Mpc/h).  This variation was then ascribed to a luminosity dependent effect --- see e.g. \\cite{davis88,park,benoist,zehavietal02}.  However, recently in a CCD survey of bright galaxies within the Northern and Southern strips of the 2dF Galaxy Redshift Survey (2dFGRS) \\cite{colless01} conclusive evidences where found that there are fluctuations of the order $ \\sim 30\\%$ in galaxy counts as a function of apparent magnitude \\cite{busswell03} (see also \\cite{frith03,frith06} for similar observations in other galaxy samples).  Further since in the angular region toward the Southern galactic cap (SGC) a deficiency, with respect to the Northern galactic cap (NGC), in the counts below magnitude $\\sim 17$ (in the $B$ filter) was found, persisting over the full area of the APM and APMBGC catalogs, this would be an evidence that there is a large void of radius of about $150$ Mpc/h implying that there are spatial correlations extending to scales larger than the scale detected by the 2dFGRS correlation function \\cite{norbergxi01,norbergxi02}. Indeed, by considering the two-point correlation function, and thus by normalizing the amplitude of fluctuations to the estimation of the sample density, the length-scale $r_0 \\approx 6-8$ Mpc/h was derived \\cite{norbergxi01,norbergxi02}.  Structures and fluctuations at scales of the order of 100 Mpc/h or more are at odds with the prediction of the concordance model of galaxy formation \\cite{busswell03,frith03,frith06}, while the small value of the correlation length is indeed compatible. In what follows we try to clarify this puzzling situation, i.e. the coexistence of the small typical length scales measured by the two-point correlation function analysis with the large fluctuations in the galaxy density field on large scales as measured by the simple galaxy counts. Because of the difference in the counts amplitude, and thus in the sample density between the NGC and the SGC samples, the estimation of the sample density is not stable and thus one must critically consider the significance of the normalization of fluctuations amplitude to the estimation of the sample density as used in the correlation analysis employed to measure the length scale $r_0$.   More generally the problem of the statistical characterization these structures in a finite sample, of volume $V$ containing $M$ galaxies, can be rephrased as the problem of measuring volume averaged statistical quantities. The basic issue concerns whether these are meaningful descriptors, i.e. whether they give or not stable statistical estimations of ensemble averaged quantities \\cite{book}. In general it is assumed that galaxy distribution is an ergodic stationary stochastic process \\cite{book}, which means that it is statistically translationally and rotationally invariant, thus avoiding special points or directions.  Stationary stochastic distributions satisfy these conditions also when they have zero average density in the infinite volume limit \\cite{book}.  The assumption of ergodicity implies that in a single realization of the microscopic number density field $n(\\vec{r})$ the average density $n_0$ in the infinite volume is well defined and equal to the ensemble average density \\cite{book}.  The constant $n_0$ is strictly positive for homogeneous distributions and it is zero for infinite inhomogeneous ones \\cite{book}. The infinite volume limit must be considered in the definition of probabilistic properties, but in physical systems one is concerned only with finite volumes and statistical determinations.  For inhomogeneous distributions, in a finite sample, the estimation of the average mass density gives a large relative error with respect to the ensemble value and it is thus systematically biased \\cite{book}.  This situation occurs as long as the sample size is smaller than the scale $\\lambda_0$ at which the distribution turns to homogeneity, i.e.  beyond which density fluctuations are small \\cite{book}.  In the finite sample analysis it is then necessary to study the conditional scaling properties of statistical quantities, by an analysis of fluctuations and correlations which explicitly considers whether a distribution can be or not homogeneous. Before turning to the description of the methods employed to study galaxy distributions and mock galaxy samples we discuss the properties of the samples considered. ",
        "conclusions": "In summary, by applying the SL method to the 2dFGRS samples we detect large density fluctuations of considerable spatial extension.  At scales $r \\ltapprox 40$ Mpc/h we find statistically stable power-law correlations with fractal dimension $D=2.2 \\pm 0.2$ in agreement with previous determinations \\cite{slmp98,book,hogg,dr4_paper,2df_paper,paper1}.  For $r>40$ Mpc/h we find that the galaxy distribution is strongly inhomogeneous and fluctuations are large up to the samples sizes, in agreement with a similar analysis of the SDSS data \\cite{paper1}.  Persistent large scale density fluctuations are compatible \\cite{book} with fractal power-law correlations extending to scales $r>$ 40 Mpc/h but incompatible with homogeneity at $\\lambda_0 \\le 75$ Mpc/h. On the other hand, standard models of galaxy formation, normalized to CMBR anisotropies, predict $\\lambda_0^m \\approx 10$ Mpc/h \\cite{pee80,springel05}, i.e.  smaller than our lower limit $\\lambda_0 > 75$ Mpc/h.  This prediction is in agreement with the results we found in mock galaxy catalogs where we measured that, fluctuations are more smoother than in the 2dFGRS samples, and their PDF rapidly converges to a Gaussian function for $r>10$ Mpc/h.   Our results are in contrast with the standard determinations that the characteristic length scale of galaxy distribution, marking the transition to the regime of small fluctuations, is of the order of $10$ Mpc/h \\cite{zehavietal02,zehavietal04,norbergxi02}. This is because this length scale is derived by measuring the amplitude of two-point correlation function $\\xi(r)$.  When considering this quantity, which is normalized to the estimation of the sample density, it is implicitly assumed that the distribution is homogeneous (i.e. with small amplitude fluctuations) well inside the sample volume, i.e. $\\lambda_0 \\ll V^{1/3}$ \\cite{book}. When fluctuations are large, as in the case of the 2dFGRS samples, this descriptor is systematically biased by finite size effects \\cite{book,2df_paper,paper1} and so is the characteristic length scale derived from its amplitude. On the other hand our results fairly agree with studies of galaxy counts as a function of the apparent magnitude $N(m)$, which indirectly probe radial distance fluctuations. These show large fluctuations around the average behavior: particularly $N(m)$ in the SGC are down by $30\\%$ relative to the NGC counts \\cite{busswell03}.  These behaviors can be now directly related to large scale galaxy structures and particularly to the fact that in the NGC samples there are more structures, and thus an higher amplitude of $\\overline{N(r;R,\\Delta r)}$ (Fig.\\ref{fig_NR}), than in SGC samples.  Finally it is worth noticing that our results agree with the conclusion of \\cite{einasto} who found that large scale structures (e.g. super-clusters) are more frequent in observed samples than in the simulations."
    },
    "0812/0812.4489_arXiv.txt": {
        "abstract": "\\vspace{1cm} \\centerline{\\bf ABSTRACT}\\vspace{2mm} Unparticle physics has been an active field since the seminal work of Georgi. Recently, many constraints on unparticles from various observations have been considered in the literature. In particular, the cosmological constraints on the unparticle dark component put it in a serious situation. In this work, we try to find a way out of this serious situation, by including the possible interaction between dark energy and the unparticle dark component. ",
        "introduction": "\\label{sec1} Recently, the so-called unparticle physics has been an active field since the seminal work of Georgi~\\cite{r1,r2}. It was based on the hypothesis that there could be an exact scale invariant hidden sector resisted at a high energy scale. A prototype model of such a sector is given by the Banks-Zaks theory which flows to an infrared fixed point at a lower energy scale $\\Lambda_{\\rm u}$ through dimensional transmutation~\\cite{r3}. Recently, there has been a flood of papers on the unparticle phenomenology. We refer to e.g.~\\cite{r3,r4} for some brief reviews. In fact, soon after the seminal work of Georgi~\\cite{r1,r2}, some authors considered the constraints on unparticles from various observations, such as the collider experiments~\\cite{r5}, the new long range force experiments~\\cite{r6,r7}, the solar and reactor neutrinos data~\\cite{r8}, as well as the observations from astrophysics and cosmology~\\cite{r7,r9,r10,r11,r23}. These observations put fairly stringent constraints on the unparticle.  The unparticle does not have a definite mass and instead has a continuous spectral density as a consequence of scale invariance~\\cite{r1,r2} (see also e.g.~\\cite{r11,r12,r24}), \\be{eq1} \\rho(P^2)=A_{d_{\\rm u}}\\theta(P^0)\\theta(P^2)(P^2)^{d_{\\rm u}-2}, \\ee where $P$ is the 4-momentum; $A_{d_{\\rm u}}$ is the normalization factor; $d_{\\rm u}$ is the scaling dimension. When $d_{\\rm u}=1$, Eq.~(\\ref{eq1}) reduces to the familiar one of a massless particle. The scaling dimension $d_{\\rm u}$ is constrained by the unitarity of the conformal algebra and the requirement to avoid the singular behavior~\\cite{r1,r4}. The theoretical bounds are $1\\leq d_{\\rm u}\\leq 2$ (for bosonic unparticles) or $3/2\\leq d_{\\rm u}\\leq 5/2$ (for fermionic unparticles)~\\cite{r4}. Typically, values $1\\leq d_{\\rm u}\\leq 2$ have been extensively considered in the literature.  The pressure and energy density of thermal unparticles have been derived in~\\cite{r12}. They are \\bea &&p_{\\rm u}=g_s T^4\\left(\\frac{T}{\\Lambda_{\\rm u}} \\right)^{2\\left(d_{\\rm u}-1\\right)}\\frac{{\\cal C} \\left(d_{\\rm u}\\right)}{4\\pi^2}\\,,\\label{eq2}\\\\ &&\\rho_{\\rm u}=\\left(2d_{\\rm u}+1\\right)g_s T^4\\left(\\frac{T} {\\Lambda_{\\rm u}}\\right)^{2\\left(d_{\\rm u}-1\\right)} \\frac{{\\cal C}\\left(d_{\\rm u}\\right)}{4\\pi^2}\\,,\\label{eq3} \\eea where ${\\cal C}\\left(d_{\\rm u}\\right)=B\\left(3/2,d_{\\rm u}\\right) \\Gamma\\left(2d_{\\rm u}+2\\right)\\zeta\\left(2d_{\\rm u}+2\\right)$, while $B$, $\\Gamma$, $\\zeta$ are the Beta, Gamma and Zeta functions, respectively. These are the results for the bosonic unparticles. The pressure and energy density of fermionic unparticles can be obtained by replacing ${\\cal C}\\left(d_{\\rm u}\\right)$ by $\\left[1-2^{-\\left(2d_{\\rm u} +1\\right)}\\right]{\\cal C}\\left(d_{\\rm u}\\right)$. Therefore, the equation-of-state parameter (EoS) of both the bosonic and fermionic unparticles is given by~\\cite{r12} \\be{eq4} w_{\\rm u}\\equiv\\frac{p_{\\rm u}}{\\rho_{\\rm u}} =\\frac{1}{2d_{\\rm u}+1}. \\ee When $d_{\\rm u}=1$, we have $w_{\\rm u}=1/3$, which is the same as radiation. When $d_{\\rm u}\\to\\infty$, we find $w_{\\rm u}\\to 0$, which approaches the EoS of pressureless matter. In the intermediate case, the EoS of unparticles is different from the one of radiation or cold dark matter, and generically lies in between. Since the unparticle interacts weakly with standard model particles, it is ``dark'' in this sense. Some authors regard the unparticle as a candidate of dark matter~\\cite{r13} (see also~\\cite{r11,r12}). Noting that the EoS of unparticles ($w_{\\rm u}>0$) is different from the one of cold dark matter ($w_{\\rm m}=0$), we instead call it ``dark component'' to avoid confusion.  In~\\cite{r11}, the cosmological constraints on the unparticle dark component have been considered, by using the type Ia supernovae (SNIa), the shift parameter of the cosmic microwave background (CMB), and the baryon acoustic oscillation (BAO). The authors of~\\cite{r11} found that $d_{\\rm u}>60$ at $95\\%$ confidence level (C.L.) for the $\\Lambda$UDM model in which the unparticle is the sole dark matter. As mentioned above, however, the theoretical bounds on unparticles are $1\\leq d_{\\rm u}\\leq 2$ (for bosonic unparticles) or $3/2\\leq d_{\\rm u}\\leq 5/2$ (for fermionic unparticles)~\\cite{r4}. The situation is serious. Even for the $\\Lambda$UCDM model in which the unparticle dark component co-exists with cold dark matter, they found that the unparticle dark component can at most make up a few percent of the total cosmic energy density if $d_{\\rm u}<10$; so that it cannot be a major component~\\cite{r11}. In fact, it is easy to understand these results. Since the unparticle dark component~scales as $a^{-3(1+w_{\\rm u})}$ whereas cold dark matter scales as $a^{-3}$ (here $a$ is the scale factor of the universe), for $w_{\\rm u}>0$, the energy density of the unparticle dark component decreases faster than the one of cold dark matter. So, it is not surprising that the energy density of the unparticle dark component cannot be comparable with cold dark matter at the present epoch. To make a considerable contribution to the total cosmic energy density, $d_{\\rm u}$ should be very large to make $w_{\\rm u}\\simeq 0$ so that the unparticle dark component could mimic cold dark matter.  In this work, we try to find a way out of the serious situation mentioned above. Our physical motivation is very simple. If dark energy, the major component of the universe, can decay into unparticles, the energy density of the unparticle dark component should decrease {\\em more slowly}. Since both dark energy and unparticle are unseen, the interaction between them is not prevented. If the dilution rates of unparticle dark component and cold dark matter are comparable, it is possible to have a considerable energy density of the unparticle dark component at the present epoch, without requiring $d_{\\rm u}$ to be very large. In fact, this is similar to the key point of the interacting dark energy models which are extensively considered in the literature to alleviate the cosmological coincidence problem (see e.g.~\\cite{r14,r25} and references therein).  Here, we consider a flat universe which contains dark energy, the unparticle dark component and pressureless matter (including cold dark matter and baryons). We assume that dark energy and unparticle dark component exchange energy according to (see e.g.~\\cite{r14,r25} and references therein) \\bea &&\\dot{\\rho}_{\\rm de}+3H\\rho_{\\rm de}(1+w_{\\rm de})=-Q,\\label{eq5}\\\\ &&\\dot{\\rho}_{\\rm u}+3H\\rho_{\\rm u}(1+w_{\\rm u})=Q,\\label{eq6} \\eea whereas pressureless matter (including cold dark matter and baryons) evolves independently, i.e., \\be{eq7} \\dot{\\rho}_{\\rm m}+3H\\rho_{\\rm m}=0. \\ee So, the total energy conservation equation $\\dot{\\rho}_{\\rm tot}+3H\\rho_{\\rm tot}(1+w_{\\rm tot})=0$ is preserved. $H\\equiv\\dot{a}/a$ is the Hubble parameter; $a=(1+z)^{-1}$ is the scale factor of the universe (we set $a_0=1$); the subscript ``0'' indicates the present value of the corresponding quantity; $z$ is the redshift; a dot denotes the derivative with respect to cosmic time $t$. The interaction forms extensively considered in the literature (see for instance~\\cite{r14,r25} and references therein) are $Q\\propto H\\rho_{\\rm m}$, $H\\rho_{\\rm de}$, $H\\rho_{\\rm tot}$, $\\kappa\\rho_{\\rm m}\\dot{\\phi}$ (where $\\kappa^2\\equiv 8\\pi G$), and so on. In this work, for simplicity, we consider the interaction term \\be{eq8} Q=3\\alpha H\\rho_{\\rm u}, \\ee where $\\alpha$ is a constant. So, Eq.~(\\ref{eq6}) becomes $\\dot{\\rho}_{\\rm u}+3H\\rho_{\\rm u}\\left(1+w_{\\rm u}^{\\rm eff}\\right)=0$, where $w_{\\rm u}^{\\rm eff}\\equiv w_{\\rm u}-\\alpha$. It is easy to find that \\be{eq9} \\rho_{\\rm u}=\\rho_{\\rm u0}\\,a^{-3\\left(1+ w_{\\rm u}^{\\rm eff}\\right)}. \\ee On the other hand, from Eq.~(\\ref{eq7}), we have $\\rho_{\\rm m}=\\rho_{\\rm m0}\\,a^{-3}$. For convenience, we introduce the fractional energy density $\\Omega_i\\equiv\\left(8\\pi G\\rho_i\\right)/\\left(3H^2\\right)$, where $i=\\rm de$, u and m.  In the following sections, similar to~\\cite{r11}, we consider the cosmological constraints on the unparticle dark component which interacts with dark energy, by using the observations of SNIa, the shift parameter $R$ of CMB, and the distance parameter $A$ from BAO. In Sec.~\\ref{sec2}, we present the cosmological data and the methodology used to constrain the models. In Sec.~\\ref{sec3} and \\ref{sec4}, we consider the observational constraints on the interacting $\\Lambda$UCDM model and the interacting XUCDM model, respectively. As expected, we find that the cosmological constraints on the unparticle dark component can be significantly relaxed, thanks to the possible interaction between dark energy and unparticle dark component. The tension between theoretical bounds and cosmological constraints could be removed. Finally, some concluding remarks are given in Sec.~\\ref{sec5}.  \\vspace{-2mm}  % ",
        "conclusions": "\\label{sec5} Unparticle physics has been an active field since the seminal work of Georgi~\\cite{r1,r2}. Recently, many constraints on unparticles from various observations were considered in the literature. In particular, the cosmological constraints on the unparticle dark component put it in a serious situation~\\cite{r11}. In the present work, we try to find a way out of this serious situation, by including the possible interaction between dark energy and the unparticle dark component. In fact, as we have shown above, the cosmological constraints on the unparticle dark component can be relaxed in the interacting $\\Lambda$UCDM model with $w_{\\rm u}^{\\rm eff}=-1$, and both the interacting XUCDM models with $w_{\\rm u}^{\\rm eff}\\not=w_{\\rm x}$ and $w_{\\rm u}^{\\rm eff}=w_{\\rm x}$. In these interacting models, the unparticle dark component can have a considerable contribution to the total cosmic energy density (namely $\\Omega_{\\rm u0}>0.1$; even $\\Omega_{\\rm u0}$ can be larger than $0.2$ in the interacting XUCDM model with $w_{\\rm u}^{\\rm eff}\\not=w_{\\rm x}$) while $1\\leq d_{\\rm u}\\leq 2$ (for bosonic unparticles) or $3/2\\leq d_{\\rm u}\\leq 5/2$ (for fermionic unparticles), thanks to the possible interaction between dark energy and unparticle dark component. The tension between theoretical bounds and cosmological constraints could be removed.  It is easy to understand these results physically. When dark energy, the major component of the universe, can decay into unparticles, the energy density of the unparticle dark component should decrease {\\em more slowly}. Since both dark energy and unparticle are unseen, the interaction between them is not prevented. If the dilution rates of unparticle dark component and cold dark matter are comparable, it is possible to have a considerable energy density of the unparticle dark component at the present epoch, without requiring $d_{\\rm u}$ to be very large. In fact, this is similar to the key point of the interacting dark energy models which are extensively considered in the literature to alleviate the cosmological coincidence problem (see e.g.~\\cite{r14,r25} and references therein).  As shown above, the cosmological constraints on the unparticle dark component can be relaxed in the interacting $\\Lambda$UCDM model with $w_{\\rm u}^{\\rm eff}=-1$, and both the interacting XUCDM models with $w_{\\rm u}^{\\rm eff}\\not=w_{\\rm x}$ and $w_{\\rm u}^{\\rm eff}=w_{\\rm x}$. However, one might note that the unparticle dark component with cosmological constant is constrained by the observational data relatively more severe than in the case with dark energy. This is also easy to understand. Notice that the EoS of dark energy $w_{\\rm x}$ is a {\\em free} model parameter, whereas the EoS of the cosmological constant is a constant $-1$. Therefore, the number of free model parameters in the interacting XUCDM model is larger than the one in the interacting $\\Lambda$UCDM model. As is well known, statistically, the constraints should be looser in this case."
    },
    "0812/0812.1323_arXiv.txt": {
        "abstract": "% The Orion Nebula is one of the most frequently observed nearby ($<1$~kiloparsec) star forming regions and, consequently, the subject of a large bibliography of observations and interpretation.  The summary in this chapter is bounded spatially by the blister H\\textsc{II} region, with sources beyond the central nebula that are part of the same dynamical clustering covered in other chapters in this book.  Herein are discussed panchromatic observations of the massive OB stars, the general T~Tauri population, the sub-stellar sources and variable stars within the Orion Nebula.  First, a brief history of 400 years of observation of the Nebula is presented. As this history is marked clearly by revelations provided in each age of new technology, recent ultra-deep X-ray surveys and high resolution multi-epoch monitoring of massive binary systems and radio stars receive special attention in this review. Topics discussed include the kinematics, multiplicity, mass distribution, rotation, and circumstellar characteristics of the pre-main sequence population. Also treated in depth are historical and current constraints on the distance to the Orion Nebula Cluster; a long standing 10-20\\% uncertainty has only recently begun to converge on a value near $\\sim400$ parsecs.  Complementing the current review of the stellar population is a companion chapter reviewing the molecular cloud, ionized H\\textsc{II} region and the youngest protostellar sources. ",
        "introduction": "\\label{sec:history}  During the fifty years after the development of the telescope in 1608 the nebular nature and stellar content of {Orion Nebula} were independently discovered by a handful of observers.  The observing logs of Nicolas-Claude Fabri de Peiresc (1610) and of Johann Baptist Cysat (at the latest 1618) ~\\citep{1854AN.....38..109W,1882USNOM..18A...1H} represent the earliest written records that the {Sword} of Orion contained a ``fog'' or ``milky nebulosity,'' which here borrows the words used by W. Herschel (1802) to describe the region.  The first hand drawn charts of the region include those of Galileo (1617), who did not distinguish the Nebula, of Giovanni Battista Hodierna (drawn sometime before 1654), who did, and the more famous 1656 drawing by Huyghens, which is the most widely known. His accurate rendition of the central nebula surrounding the {Trapezium} provides an origin for the term ``Huyghenian region'' (Figure \\ref{fig:Figures_huyghens}).% \\begin{figure}[t] \\centering \\includegraphics[angle=135,totalheight=3in]{onc1_fig_03.pdf}  \\caption{The 1656 sketch of the inner {Orion Nebula} by Huyghens. The figure has been re-oriented from its publication form such that North is up and East is to the left. Image is reproduced from \\citet{1882USNOM..18A...1H}.} \\label{fig:Figures_huyghens} \\end{figure}  A nearly complete journal of 273 years of telescope aided visual observation of the Nebula is provided in Edward S. Holden's \\textit{The Monograph of the Central Parts of the Nebula of Orion} (1882). Holden's monograph includes the reproduction of dozens of hand drawn sketches as well as observing logs each in their original language.  Variation in the reproduction of the Nebula is remarkable. The first photographic plates of Orion made by Henry Draper between 1880 and 1882 were included as an addendum to Holden's work as well as a discussion of the processes of obtaining these images.  Figure \\ref{fig:onc_draper} is a reproduction of that image and is captioned with Holden's description. \\begin{figure} [ht*]  \\centering \\includegraphics[angle=180, totalheight=3.5in]{onc1_fig_04.pdf}  \\caption{Draper's 1882 photographic image of the {Orion Nebula.} North is up and East is to the left. An approximately 2 hour visible photograph obtained by H. Draper with the 11 inch Clark telescope. An excerpt from Edward S. Holden's ``Monograph of the Central Parts of the Nebula of Orion'' (Washington 1882) p.226: \\textit{''The first photograph of the nebula of Orion was made by Dr. Henry Draper in September, 1880, and the unavoidable delay which has occurred in printing the present memoir enables me to include an account of the astonishing results which he has attained. A wood-cut which I had prepared from his first photograph was found to be so unsatisfactory that Dr. Draper most generously offered to supply the necessary photolitographic reproductions of his last negative (taken March 14, 1882) to accompany the brief account I had prepared. The full page photolitograph is here given...''} \\label{fig:onc_draper}}  \\end{figure}  The next $\\sim100$ years of photographic observation of the Nebula included quantitative studies of its variable stars, the discovery of a cluster of faint stars in the Nebula's core, very broad censuses and an expansion in the role of the Nebula as a testbed for new observing techniques.  Numerous variable star studies were performed using the Harvard Plates by H. Leavitt (published by Pickering) or confirmed by M. Applegate (published by Shapley).  In the 1930s, deep red photographic plates revealed that the Nebula contained a substantial cluster of fainter stars in addition to the brightest members \\citep{1931PASP...43..255T,1937ApJ....86..119B}.  The broad surveys of \\citet{1935POLyo...1...12B} and \\citet{1954TrSht..25....1P} provided excellent photographic updates to the ~\\citet{1867AnHar...5....1B} visual census of stars in the Nebula.  Parenago, a Russian astronomer, published a major, lengthy analysis of the region in 1954. A translation of sections of this publication was undertaken for this review and it indicates that he relied on prior and concurrent work of female Russian astronomers (e.g., Barkahatova, Uranova, Kirillova) some of whose names do not reappear in the literature outside of his book.  His analysis extended to the topics of parallax, proper motions, astrometry and photometry, covering all of the {Sword} of Orion. Most important he found a clear evidence for a ``cloud'' of members lying above the main sequence. Perhaps because of the Cold War and the lack of translations from his work from Russian to other languages, his work is extremely poorly cited in the literature. The disparaging of his work by \\citet{1956ApJS....2..365W} was not unnoticed by the author as revealed in the posthumous publication, \\citet{1961AJ.....66..103P}. Nevertheless, this work was not considered seriously in most subsequent studies, not appearing, for example, in the otherwise meticulous study by Goudis (see below). Even today this significant work has garnered a mere 59 citations\\footnote{Derived via the NASA Astrophysical Data Service circa May 2008 by merging the results of citations to \\citet{1954TrSht..25....1P} and \\citet{1954TrSht..25....3P}.}. His valuable data tables were converted to machine readable formats by \\citet{1992BICDS..40...13M} and ingested into an electronic format in 1997 \\citep{1997yCat.2171....0P}.  In 1982, approximately one hundred years after the Holden monograph and the first photographic images by Draper, two useful summarizing publications appeared. C. Goudis's \\textit{The Orion Complex: A Case Study of Interstellar Matter,} focused on the structure and nature of the Nebula itself, details previous approaches to studying the Nebula and its content, including infrared, radio and spectroscopy, and provides useful tables of past observations and results. Second, a conference was held on the {Orion Nebula} and honoring Henry Draper \\citep{1982NYASA.395.....G}; the conference proceedings include 32 articles on all current aspects of study of the Nebula and records of the participants discussion about each contribution. In addition there are a number of articles that provide a history, more detailed than that of Holden, about Draper's photographic work obtaining these images of the Nebula as well as his subsequent scientific legacy after his death in late 1882.  Reviews of the literature over the subsequent $\\sim20$~years of CCD observations of the Nebula include \\citet{1989ARA&A..27...41G} and \\citet{2001ARA&A..39...99O}.  The Nebula was also included in a review that encompassed all of the Orion star forming region by~\\citet{1991lmsf.book....1B}. Its recent study has also been the focus of a book, \\citet{2003onws.book.....O}.  \\defcitealias{1918ApJ....47..104K}{a} \\defcitealias{1918ApJ....47..146K}{b} \\defcitealias{1918ApJ....47..255K}{c} \\begin{table}[p]  \\caption{Summary of published distances to the Orion Id association. Distances are segregated as corresponding to the Ic or Id associations with the notation Ic* used to indicate that part of Ic right around Id and to distinguish it from the much larger Ic as defined by WH78. Unfortunately the notation Ic* was also used by many authors before the Ia,b,c,d subgroups were defined to indicate the entire ``Sword'' region; thus some ambiguities may remain.  \\label{tab:onc_dist} } \\smallskip  \\begin{center} {\\tiny \\begin{tabular} {l@{\\hskip5pt}l@{\\hskip5pt}c@{\\hskip5pt}c@{\\hskip5pt}c@{\\hskip5pt}c@{\\hskip5pt}c@{\\hskip5pt}c@{\\hskip5pt}c}  \\tableline \\noalign{\\smallskip} Author Name & Year of & Region & Distance & Error & Stellar & Method & Data & Number \\\\ & Pub. & Desig. & Modulus & & Types &   &   & of Stars\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip}  \\citeauthor{1917HarCi.205....1P} & 1917 & Id+Ic* & 11.5 & & B3 & ZAMS & pg & ? \\\\ \\citeauthor{1918ApJ....47..104K} & 1918\\citetalias{1918ApJ....47..104K,1918ApJ....47..146K,1918ApJ....47..255K} & Ic & 6.34 & & B & PM & pg & ? \\\\ \\citeauthor{1919PASP...31...86P} & 1919 & Id+Ic* & 8.5 & & B3 & ZAMS & pg & ? \\\\ \\citeauthor{1929PUAms...2....1P} & 1929 & all & 7.6 & & B & ZAMS & pg & ? \\\\ \\citeauthor{1931PASP...43..255T} & 1931 & Id & 8.5 & & O9-A2 & ZAMS & pg & 17 \\\\ \\citeauthor{1946PASP...58..356M} & 1946 & Id & 7.38 & & O7-B1 & ZAMS & pg & 3 \\\\ \\citeauthor{1949AJ.....54..111M} & 1949 & Id & 8.58 & 0.35 & B1-B3 & ZAMS & pg & 17 \\\\ \\citeauthor{1952ApJ...116..251S} & 1952,4& Id+Ic & 8.5 & 0.30 & B Stars & ZAMS & pe & 190 \\\\ $\\vdots$ & 1952 & Id & 8.6 & 0.30 & B Stars & ZAMS & & ? \\\\ \\citeauthor{1954TrSht..25....1P} & 1954 & Id & 8.0 & & & ZAMS & pg & ? \\\\ \\citeauthor{1956ApJ...123..267J} & 1956 & Id & 8.0 & & B8-A0 & ZAMS & & ? \\\\ \\citeauthor{1958ApJ...128...14S} & 1958 & Id & 8.6 & & O6-K2? & PM / RV& plate & 20 \\\\ \\citeauthor{1962ApJ...136..767S} & 1962 & Id & 8.2 & & B Stars & & $UBV,H\\gamma$ & 180 \\\\ \\citeauthor{1964BAN....17..358B} & 1964 & Id & 8.33 & 0.11 & & ZAMS & 7-filter & 5 \\\\ \\citeauthor{1965ApJ...142..964J} & 1965 & Id & 7.9 & & & PM / RV & & 21 \\\\ \\citeauthor{1966VA......8...83M} & 1966 & Ic* & 8.1 & & & ZAMS & & ? \\\\ \\citeauthor{1968ApJ...152..905L} & 1968 & Ic* & 8.5 & 0.1 & & ZAMS & & 14 \\\\ \\citeauthor{1969ApJ...155..447W} & 1969 & Id+Ic* & 8.37 & 0.05 & B2-B9 & ZAMS & $UBV$ & 51 \\\\ \\citeauthor{1973ApJ...183..505P} & 1973 & Id & 7.8 & 0.15 & B stars & ZAMS & $BV$ & 15 \\\\ \\citeauthor{1975MNRAS.171..219P} & 1975 & Id & 8.1 & 0.13 & B stars & ZAMS & $BV$ & ? \\\\ $\\vdots$ & 1975 & Id & 7.71 & 0.21 & B stars & ZAMS & $VI$ & ? \\\\ $\\vdots$ & 1975 & Id & 7.98 & 0.12 & B stars & ZAMS & $V$;SpT & ? \\\\ \\citeauthor{1977ApJS...34..115W} & 1977 & Id & 8.42 & 0.53 & B stars & & $UBV;ubvy;H_{\\beta}$ & 6 \\\\ $\\vdots$ & 1977 & Ic* & 8.16 & 0.49 & B stars & & $UBV;ubvy;H_{\\beta}$ & 44 \\\\ \\protect{\\citeauthor{1981A&AS...44..467M}} & 1981 & Id+Ic* & 8.20 & 0.15 & & ZAMS? & $UBV?$ & ? \\\\ \\citeauthor{1981ApJ...244..884G} & 1981 & K-L region& 8.41 & 0.40 & H$_2$0 masers&stat. parallax & VLBI $pm+rv$ & $\\sim 30$ \\\\ \\citeauthor{1981ApJ...248..963B} & 1981 & Id+Ic & 8.0 & 0.5 & & ZAMS & $BV$ & ? \\\\ \\citeauthor{1982AJ.....87.1213A} & 1982 & Ic & 7.87 & 0.09 & B stars & ZAMS & $M_V; H_{\\beta}$ & 41 \\\\ $\\vdots$ & 1982 & Id+Ic* & 8.19 & 0.10 & B stars & ZAMS & $M_V; H_{\\beta}$ & 15 \\\\ \\citeauthor{1990AJ....100.1994W} & 1990 & Ic & 7.7 & 0.50 & B Stars & ZAMS & co-$H_{\\gamma}$ & ? \\\\ $\\vdots$ & 1990 & Id & 8.2 & 0.03 & B Stars & ZAMS & co-$H_{\\gamma}$ & 2 \\\\ \\protect{\\citeauthor{1994A&A...289..101B}} & 1994 & Ic & 8.0 & 0.49 & B Stars & ZAMS & $VBLUW$ & 34 \\\\ $\\vdots$ & 1994 & Id & 7.9 & 0.25 & B Stars & ZAMS & $VBLUW$ & 3 \\\\ $\\vdots$ & 1994 & cloud & 8.1 & 0.48 & B Stars & Red. & $VBLUW$ & \\\\ \\citeauthor{1999osps.conf..411B} & 1998 & Ic & 8.32 & 0.17 & B Stars & trig. parallax & Hipparcos & 34 \\\\ \\citeauthor{1999AJ....117..354D} & 1999 & Ic & 8.52 & 0.25 & B Stars & trig. parallax & Hipparcos & 34 \\\\ \\citeauthor{2004ApJS..151..357S} & 2004 & Ic* & 7.96 & 0.10 & ec.~binary & Radius & - & 1 \\\\ \\citeauthor{2005AJ....129..856H} & 2005 & Ib+Ic & 8.23 & 0.08 & B5-F Stars& trig. parallax & Hipparcos & 121 \\\\ $\\vdots$ & 2005 & Ib+Ic & 7.97 & 0.10 & B5-F Stars& ZAMS & $BV$/Hipparcos & 111 \\\\ \\protect{\\citeauthor{2005A&A...430..523W}} & 2005 & Id+Ic & 8.34 & 0.32 & Stars & Red. & CO/Hipparcos & ? \\\\ \\citeauthor{2006Natur.440..311S} & 2006a & Ic* & 8.19 & 0.30 & ec.~binary & Radius & - & 1 \\\\ \\citeauthor{2007MNRAS.376.1109J} & 2007a & Id & 8.22 & 0.16 & G6-M2 &$R\\,sin(i)$ & Various & 74 \\\\ $\\vdots$ & 2007a & Id & 7.97 & 0.17 & G6-M2 &$R\\,sin(i)$ & Various & 34 \\\\ \\protect{\\citeauthor{2007A&A...466..649K}} & 2007 & Id & 8.19 & 0.06 & O &dyn. parallax & binary orbit & 1 \\\\ $\\vdots$ & 2007 & Id & 7.94 & 0.06 & O &dyn. parallax & binary orbit & 1 \\\\ \\citeauthor{2007PASJ...59..897H}  & 2007 & K-L region & 8.20 & 0.09 & H$_2$O masers &trig. parallax & VERA & 1 \\\\ \\citeauthor{2007ApJ...667.1161S}  & 2007 & Id & 7.95 & 0.13 & radio star &trig. parallax & VLBA & 1 \\\\ \\protect{\\citeauthor{2007A&A...474..515M}}  & 2007 & Id & 8.08 & 0.03 & radio stars &trig. parallax & VLBA & 4 \\\\ \\citeauthor{2008arXiv0801.4085M}  & 2008  & Id &  7.96 & 0.06 & B1-A0 & MS Fitting & $VI$ & $\\sim$20  \\\\ \\tableline \\end{tabular}  } \\end{center}  \\end{table} ",
        "conclusions": ""
    },
    "0812/0812.0505_arXiv.txt": {
        "abstract": "In this paper we review the modeling of the Local Bubble (LB) with special emphasis on the progress we have made since the last major conference ``The Local Bubble and Beyond (I)'' held in Garching in 1997. Since then new insight was gained into the possible origin of the LB, with a moving group crossing its volume during the last 10 - 15 Myr being most likely responsible for creating a local cavity filled with hot recombining gas. Numerical high resolution 3D simulations of a supernova driven inhomogeneous interstellar medium show that we can reproduce both the extension of the LB and the O{\\sc vi} column density in absorption measured with FUSE for a LB age of 13.5 - 14.5 Myr. We further demonstrate that the LB evolves like an ordinary superbubble expanding into a density stratified medium by comparing analytical 2D Kompaneets solutions to Na{\\sc i} contours, representing the extension of the local cavity. These results suggest that LB blow-out into the Milky Way halo has occurred roughly 5 Myr ago. ",
        "introduction": "At the time of the last Local Bubble (LB) conference, ``The Local Bubble and Beyond (I)'', as it should be called now, held in Garching in 1997, modeling of the LB was no more than guess work, because the crucial boundary conditions, like e.g. how many supernovae exploded and when, were unknown (cf. \\cite{cs01}). Therefore the only constraint was to match EUV and soft X-ray data. In the meantime significant observational progress has been made, warranting detailed 3D high resolution numerical simulations of the origin and evolution of the LB, which will be discussed below.  In an attempt to outline the morphology of the local cavity deficient of H{\\sc i}, Na{\\sc i} was used as a sensitive tracer in absorption against background stars \\citep{sl99,lw03}. It has been found that the cavity is elongated and inclined by $\\sim 20^\\circ$ with respect to the Galactic midplane (similar to the Gould's Belt), and open towards the poles, suggesting it to be a so-called Local Chimney. These chimneys have been proposed long ago to be the result of spatially and temporally correlated supernova (SN) explosions \\citep{ni89}, generating high pressure bubbles, which blow out of the Galactic disk. Meridional sections through the LB also show a variation of the contours with longitude, arguing for an inhomogeneous ambient medium. Since the extension of the soft X-ray emitting LB is much less certain known than that of the local cavity, we will for convenience not distinguish between the two in the following.  To establish the temperature and density of the LB from soft X-ray and EUV observations has been extremely difficult, because there was a noticeable deficit of lines at low energies, putting the standard assumption of collisional ionization equilibrium (CIE) into question. Subsequently, models based on non-equilibrium ionization (NEI) have been suggested by Breitschwerdt \\& Schmutzler \\citep{bs94}, which could explain the lack of lines by an overionized plasma due to delayed recombination of a rapidly adiabatically cooling superbubble, but were also proceeding from unknown initial conditions. While CIE models do not depend on the evolution and hence the thermodynamic path the plasma has taken, NEI models are crucially determined by the underlying astrophysical model. Hence, if a CIE model for known abundances, temperature and absorbing column density does not fit, CIE can be ruled out. Not so for NEI models. The failure of a particular NEI model just rules out that particular astrophysical model which was responsible for it, but not NEI altogether. However, it would be an exhausting task to devise series of NEI models trying to fit the data, without even knowing the initial conditions of the problem. The situation has been exacerbated by an as yet uncertain amount of soft X-rays produced locally by solar wind charge exchange reactions (SWCX) onto heliospheric neutrals (\\cite{cr00}; see also Koutroumpa, Shelton, these proceedings). Although some protagonists of the SWCX contribution to soft X-rays seem to believe that nearly all of the emission in the Galactic plane could be accounted for, there are still photons from higher Galactic latitudes, which need to be explained, as well as the state of the then not X-ray emitting plasma, which has to fill the local cavity. In summary, it seems reasonable to assume that multiple SN explosions are responsible for a significant amount of the soft X-ray and UV emitting plasma, and to search for the origin of the ``smoking gun'', because there is no young stellar cluster found inside the LB. Investigating moving groups of the local ISM (Bergh\\\"ofer \\& Breitschwerdt \\cite{bb02}), which was later extended to a complete and unbiased analysis of Hipparcos data of a larger volume of 400 pc in diameter, and combining these data by ARIVEL catalogue data (Fuchs et al. \\cite{fu06}), it turned out that indeed a moving group crossed the present day LB region about 10 - 15 Myr ago, with its surviving members being now part of the Sco Cen association (see also \\cite{ma01}). During this passage about 14 to 19 SNe exploded and energized the LB. Although the trajectory of the moving group passes close to the LB boundary in the Galactic Center direction, it is not unrealistic that the expansion of the LB is predominantly in the anticenter direction, because the presence of the Loop I bubble, which has been generated almost coevally with the LB, strongly inhibits the LB expansion in the center direction. This effect has been observed by an annular X-ray shadow in the soft ROSAT bands due to a compressed interaction shell between the two bubbles with a column density of $\\sim 7 \\times 10^{20} \\, {\\rm cm}^{-2}$ \\citep{ea95}.  Bearing these findings in mind, we were able to start high resolution 3D numerical simulations with well constrained initial conditions \\citep{ba06}, and compare our results to recent observations, e.g. in the UV with FUSE. This review is organized as follows. In Section~2 we outline the physical and numerical setup of our LB model and discuss the results in Section 3. In Section 4, we show that basic features of the LB, like its extension, shell velocity etc. can also be derived from analytical 2D Kompaneets \\citep{ko60} solutions. In Section 5 our summary and conclusions on LB modeling will be presented. ",
        "conclusions": "\\label{conc} We have developed a \\textit{physical model} for the LB, which makes definite predictions for the state of the local plasma, and can thus be confronted with observations. The basic features of the model are: (i) the number of SNe and their explosion times are given, (ii) the ambient medium is \\textit{not} assumed to be homogeneous, but is calculated from a larger general SN driven ISM simulation, (iii) the evolution of the neighboring Loop I superbubble is calculated simultaneously, with SN explosions derived from its present stellar content, (iv) diffuse stellar background heating and radiative cooling are taken into account. We have shown that 3D high resolution numerical simulations can reproduce UV observations and the overall LB morphology for an inferred age of $13.5 - 14.5$ Myr. Moreover, the LB pressure turns out to be fairly low, i.e. $P/k_B \\sim 2000 - 3000 \\, {\\rm K} \\, {\\rm cm}^{-3}$, where $k_B$ is Boltzmann's constant, in agreement with observations (for a recent discussion see \\cite{je08}), resolving the earlier problem of a huge pressure imbalance between a hot LB and the Local Cloud, in which our solar system is embedded, and which has $P/k_B \\sim 1300 - 3000 \\, {\\rm K} \\, {\\rm cm}^{-3}$ (e.g. \\cite{fr95}). Since the last SN occurred about 0.5 Myr ago, the LB is becoming extinct and is in a state of recombination. We find at present a much smaller amount of gas at (or in excess of) $10^6$ K than any previous model, allowing for a sizeable contribution of SWCX to the soft X-ray emission. However, in a next steep we need to quantify the NEI structure and the fraction of X-rays due to delayed recombination, which is the subject of a forthcoming paper.  Comparing our LB runs to the general ISM simulations, we find it reassuring that the LB behaves quite similar to other superbubbles in the computational box, and therefore raises no questions about it being a special object, as has been hypothesized occasionally. This picture is completed by the fact that even 2D analytical solutions provide a reasonable fit to the LB morphology and correctly predict the formation of a Local Chimney.  What lies ahead is to calculate O{\\sc vi} column densities \\textit{in emission} and to produce soft X-ray maps in a full NEI simulation, and hope for quantitative results for the SWCX contribution from our colleagues, until the ``Local Bubble and Beyond (III)'' meeting.  \\begin{theacknowledgments} DB thanks the organizers for financial support to attend the meeting. VB acknowledges financial support from the research grant ``Computational Astrophysics'' funded by Vienna University under contract number FS 538001. \\end{theacknowledgments}"
    },
    "0812/0812.0991_arXiv.txt": {
        "abstract": "We present the results of a Spitzer IRAC and MIPS 24$\\mu$m study of extended Lyman-$\\alpha$ clouds (or Lyman-$\\alpha$ Blobs, LABs) within the SSA22 filamentary structure at $z = $ 3.09. We detect 6/26 LABs in all  IRAC filters, four of which are also detected at 24$\\mu$m, and find good correspondence with the 850$\\mu$m measurements of \\citet{gea05}.  An analysis of the rest-frame ultraviolet, optical, near- and mid-infrared colors reveals that these six systems exhibit signs of nuclear activity (AGN) and/or extreme star formation.  Notably, they have properties that bridge galaxies dominated by star formation (Lyman-break galaxies; LBGs) and those with AGNs (LBGs classified as QSOs). The LAB systems not detected in all four IRAC bands, on the other hand, are, as a group, consistent with pure star forming systems, similar to the majority of the  LBGs within the filament.      These results  indicate that the galaxies within LABs do not comprise a homogeneous population, though they are also  consistent with scenarios in which the gas halos are ionized through a common mechanism such as galaxy-scale winds driven by the galaxies within them,   or gravitational heating of the collapsing cloud itself. ",
        "introduction": "Understanding the complex physics of galaxy formation is one of the key goals of astronomy today.  Observations at  high-redshift show that the young universe was strikingly different than it is locally and  while many high-redshift galaxy populations have low-redshift analogues, significant evolution is seen in their global properties such as their clustering strength \\citep{coi04,lefev05}, the shape of their luminosity function \\citep[e.g.][]{cir08} and their overall contribution to the integrated background light \\citep[e.g.][]{web03} and the stellar mass density of the universe \\citep{rud06}.   Individually, the galaxies themselves are generally  more violent and dynamic objects than typical galaxies today.  They are more gas-rich, show higher levels of star formation and nucleic activity \\citep{cha05,fin06}, experience increased rates of major mergers \\citep{pat02}, and exhibit signs of energetic outflows \\citep{shap03}. The role that these processes play in building and shaping galaxies is still poorly understood, as is the link between different evolutionary phases and the importance of  properties such as mass and the local environment in driving evolution.  Hierarchical structure formation models predict, and observations confirm,  that the first objects to form in the early universe will be highly clustered \\citep[e.g.][]{kai84,gia98}. For this reason, many studies of galaxy evolution focus on single large structures at high-redshift which, in addition to isolating the most active  sites of structure formation at early times,   provides large samples of equidistant galaxies over a range in mass, luminosity and evolutionary stages.    Such studies of high density environments have recently detected a new population of extended Lyman-$\\alpha$ (Ly$\\alpha$) emitting nebulae at high-redshift \\citep{fyn99,keel99,ste00,fran01,pal04,mat04, dey05}.  The properties of these so-called Ly$\\alpha$ Blobs (LABs) indicate they are energetic phases in the formation of galaxies: they have projected physical extents of tens to several hundreds of  kpc,  Ly$\\alpha$ luminosities of $\\sim$ 10$^{44}$ erg s$^{-1}$,  and the Ly$\\alpha$ emission is often morphologically irregular and dynamically complex.     They are prevalent in high-density regions which exhibit other signs of intense star formation and active galactic nuclei (AGN) activity, though the extent to which this is a selection effect has not been adequately explored through unbiased studies \\citep[c.f.][]{pres08}.  The largest published sample of extended LABs to date is presented in \\citet[][hereafter M04]{mat04}, within the SSA22 structure at $z\\sim$ 3.09.  Originally discovered as an over-density of Lyman-break Galaxies (LBGs) \\citep{ste98}, subsequent narrow-band imaging of the SSA22 field \\citep{ste00} revealed two giant LABs within the concentration with sizes $>$ 100 kpc. Narrow-band imaging over a larger area by M04 assembled a sample  35 LABs  with sizes above a few tens of kpc and luminosities of 10$^{42-44}$ erg s$^{-1}$.  They are imbedded within a large ($>$ 60 Mpc) structure,  traced by the fainter and smaller Ly$\\alpha$ Emitter (LAE) population \\citep{haya04,mat05} consisting of at least three dynamically distinct filaments. The two brightest two LABs are located at the filamentary intersection.  Clearly, the SSA22 structure is a concentration of evolutionary activity and the large number of LABs within it provides us with an opportunity to conduct a more systematic study these enigmatic systems than has been done before. The primary outstanding question concerns the mechanism responsible for their intense Ly$\\alpha$ luminosities and super-galactic sizes. While an obvious explanation  is photo-ionization by an embedded source, many LABs do not contain UV-bright galaxies of sufficient luminosity (M04).  Neither are they powered by interactions between radio-jets and the ambient inter-galactic medium as they are radio-quiet,  unlike the similar Ly$\\alpha$ halos surrounding high-redshift radio galaxies \\citep{reu03}.    Alternative scenarios hold that these systems are powered by shock heating from galaxy-wide superwinds \\citep{gea05,col06}, driven by intense starbursts or AGN \\citep{keel99}, or through gravitational heating during the collapse of a large primordial cloud \\citep[e.g.][]{smi07}.   Although exceptional UV and radio sources are not in general found within LABs, recent observations have shown that many (but not necessarily all)  LABs are associated with infrared luminous galaxies \\citep{cha01,gea05,col06,bel08}.  The work of \\citet{cha01} and \\citet{gea05} detected 5 of the M04 LABs at 850$\\mu$m ($>$ 3$\\sigma$ significance).   It is notable that the Submillimeter Array has failed to detect the most luminous of these, LAB01 \\citep[$S_{850{\\mu}m}$ = 17 mJy;][]{ste00} indicating either extended submillimeter emission ($>$ 4{\\arcsec}) or multiple submillimeter emitters confused within the beam \\citep{mat07}, a point we will return to in \\S 4.2.   Work with the MIPS camera on board Spitzer has also revealed infrared-luminous counterparts to a number of LABs including the largest currently known \\citep{dey05,col06}.  These results all imply  classifications of ultra-/hyper-luminous infrared galaxies (ULIRGs/HyLIRGs), with infrared luminosities of $\\sim$10$^{12-13}$ L$_\\odot$. The  relationship between the infrared luminous systems and their respective LABs is unclear, but a weak correlation between the counterpart bolometric luminosity and the Ly$\\alpha$ luminosity suggests that the size and luminosity of the LAB is controlled by global, galactic sized processes such as super-winds.    In this paper we present the first mid-infrared study of a large sample of LABs within the SSA22 filament  with IRAC \\citep{faz04} and MIPS \\citep{riek04} on board the Spitzer Space Telescope \\citep{wer04}.   We use these observations, along with supplementary data from the literature, archives, and our own optical observations,  to identify the galaxies which are associated with the LABs and to investigate their nature and relation to the clouds. At $z = $ 3.09 these data sample the rest-frame near- and mid-infrared (NIR/MIR), and are sensitive to stellar emission, thermally radiating hot dust, and polycyclic aromatic hydrocarbon features (PAHs). The broad shape of the spectral energy distribution (SED) in this region provides a    useful diagnostic of the energy production mechanism within each galaxy that is complementary to the optical/UV and far-infrared observations.  The paper is organized as follows:  \\S 2 describes the observations and data analysis techniques of the mid-infrared, optical and archival X-ray data;  in Section \\S 3 we present the infrared counterparts to the LABs;  and \\S 4 contains a discussion of the constraints on the nature of the counterparts,  comparisons to previous studies, and implications for our understanding of the LABs and their relation to galaxy formation. ",
        "conclusions": "\\subsection{Observed Mid-Infrared Colors: Constraints on the Nature of the   LAB Counterparts}  At $z\\sim$ 3 the IRAC filters sample the rest-frame NIR. In this region pure stellar populations are characterized by the stellar bump at  1.6$\\mu$m, with decreasing emission toward longer wavelengths.  If cold dust is present within a galaxy, band emission from PAHs becomes important over the range of $\\sim$ 6 - 11 $\\mu$m while  warm dust  produces emission that continues to rise as a power-law longward of 1.6$\\mu$m \\citep[e.g.][]{alonso06}.   These distinctly different SED shapes can, in principle, be exploited to diagnose the dominant power source of a galaxy \\citep[e.g.][]{lacy04,stern05}, but do suffer from important ambiguities and limitations, in particular beyond $z\\sim$ 2. In particular, red rest-frame NIR colors only indicate the presence of warm dust, and cannot discriminate between an AGN, with a  nuclear warm dust component and  an extreme infrared-luminous starburst whose warm dust is associated with HII regions \\citep[e.g.][]{yun08}. Moreover, the separation of infrared-luminous galaxies into AGN and star-bursting systems is regardless an over-simplification, since the most luminous  ULIRGs at low and high redshifts  show evidence for both  AGN activity and intense star formation and  the optical/NIR emission will contain contributions from both processes. Nevertheless, an investigation of the rest-frame NIR-MIR SED shapes can provide basic clues to the nature of the galaxies residing within LABs.   In Figure \\ref{cc} we show an IRAC color-color diagram for  LAB counterparts detected in all four IRAC filters \\citep[e.g.][]{lacy04}. As discussed above, infrared power-law galaxies at low and high redshift inhabit a specific region of the diagram that is generally distinct from galaxies  dominated by stellar emission in the optical/infrared, but which overlaps with the location of  high-redshift star-forming ULIRGs. To illustrate this we plot the observed infrared colors of two representative SEDs with redshift, from $z=$ 0 to $z=$ 3.09: the ULIRG Arp220 and a typical Irregular starburst galaxy from \\citet{cww80}.  Also shown are the observed infrared colors of 8$\\mu$m-bright LBGs within the SSA22 structure \\citep{ste03} and the general 8$\\mu$m-selected population in the SSA22 field.  The LBGs exhibit  rest-frame optical/NIR colors spanning almost an order of magnitude; while the bulk of the objects have colors consistent with a $z\\sim$ 3 blue starburst galaxy, a number extend into the region of the diagram inhabited by ULIRGs and AGN. The LAB counterparts on the other hand are found without exception in the region of the redder LBGs  and AGN/ULIRG templates and are not well described by a low-extinction starburst.   In addition to the ULIRG/AGN ambiguity,  it is also difficult to constrain the redshifts using the observed NIR colors:  pure power-law systems in particular have a flat color-redshift relation.  Nevertheless, one can see from Figure \\ref{cc} that among the 8$\\mu$m-selected population  power-law (and power-law-like) galaxies  are rare ($\\sim$10\\%)  and thus the likelihood that all five LABs have a randomly aligned unassociated power-law galaxy (and two in the case of LAB01)  within the extent of the Lyman-$\\alpha$ emission is remote ($\\lesssim$ 1\\% assuming a generous alignment distance of 5$\\arcsec$), and we therefore conclude that all seven systems have $z = $ 3 and  are associated with their parent LAB.  Although the majority of the LABs were not detected in all four IRAC filters, nor MIPS 24$\\mu$m, we can still place constraints on their average properties through a stacking analysis of their infrared emission.  For this exercise we take as the source position the coordinates of the 3.6$\\mu$m counterpart when available (the majority of the LABs) or if no counterpart exists at 3.6$\\mu$m we simply take the peak of the Ly$\\alpha$ emission;  only 1 LAB is too extended to identify a clear peak/center and has no clear counterpart  (LAB27), resulting in a sample of 20 systems, including LAB02-a. We detect stacked emission at $>$ 5$\\sigma$ in all four IRAC filters and  show the average color of the infrared-faint LAB systems in Figure \\ref{cc}. As we see, they are not simply lower luminosity versions of the 8$\\mu$m detected LABs, but rather are similar in color to the blue 8$\\mu$m-detected LBG population, implying they are more moderate starburst systems.  Returning to the infrared-bright LABs, the observed NIR colors require the presence of hot dust, heated by either  extreme  dust-enshrouded star formation or feedback from a central AGN.  As we now discuss however, the X-ray and MIR data indicate that AGN are present in some of these systems. First, turning to the X-ray results, two of the ten LBGs in this region are previously known optically confirmed QSOs \\citep{ste03} which are strong X-ray emitters, and two of the seven LAB counterparts have weak X-ray emission (see \\S 3).    Here we use the \\citet{bau04}  AGN/starburst classification scheme where a  galaxy is considered an AGN if it satisfies any one of five X-ray or spectral conditions.   Two of these conditions are L$_X>$ 3$\\times$10$^{42}$ erg/s and $f_{0.5-8.0 keV}/f_R > 0.1$  which are satisfied by the two X-ray detected LABs ( L$_X\\sim$ 4$\\times$10$^{43}$ erg/s and $f_{0.5-8.0 keV}/f_R \\sim$ 1.5).    The individual and stacked X-ray limits for the remaining LABs are generous (L$_X<$ 8$\\times$10$^{43}$ erg/s and L$_X<$ 6$\\times$10$^{42}$ erg/s respectively) and are consistent with AGN according to these criteria.  Deeper X-ray studies (Geach et al. 2008, in preparation) will clarify this issue.   In Figure \\ref{iracmips} we look at the constraints placed by the 24$\\mu$m data through an investigation of the IRAC-MIPS colors.  Here we are limited to the LAB and LBG comparison systems which are also detected at 24$\\mu$m. AGN-dominated systems generally inhabit the lower right of this diagram: hot dust enhances their observed  8$\\mu$m flux (2$\\mu$m in the rest-frame) over the shorter wavelength stellar emission, but the lack of strong PAH features results in blue 24$\\mu$m-8$\\mu$m colors, with some variation due to the steepness of the power-law.  Star forming systems, on the other hand have less enhanced observed 8$\\mu$m (rest-frame 2$\\mu$) emission, and thus blue 8.0$\\mu$m-4.5$\\mu$m colors, but exhibit  a wide variety of PAH strengths and therefore a range in 24$\\mu$m-8$\\mu$m colors.   Such a diagram has been used by other authors to attempt to identify the power sources within the very luminous SMG population \\citep[e.g][]{ega04,pope06}  which contain AGN and star-formation contributions, sometimes within the same objects, and, as seen in Figure \\ref{iracmips}, tend to populate the entire color space.  We see a clear separation of AGN and star-forming systems within the 24$\\mu$m-bright LBG population.  It is interesting to see that the LAB counterparts occupy a small region of the color-space, and transition  between these two populations.  Thus while smoking-gun evidence of AGN in these systems is lacking, they appear distinct from the presumably pure star-forming LBG galaxies, but are also exhibit bluer NIR colors than AGN. We again perform a stacking analysis on the infrared faint LABs, this time restricting the sample to those which fall on the MIPS image. No MIPS flux is detected from these systems as a whole and we place an upper-limit on their observed 24$\\mu$m/8$\\mu$m color on Figure \\ref{iracmips}. Here the separation between the infrared-bright and -faint LABs is less clear, though the infrared-faint LABs again lie blue-ward of the infrared-bright systems, toward the star-forming region of the diagram.    \\begin{figure} \\epsscale{1.1} \\plotone{f3.eps} \\caption{The IRAC color-color diagram.   The grey points correspond to the entire 8$\\mu$m-selected catalog in the extended SSA22 field ($>$10$\\sigma$ in all 4 IRAC filters) and the filled black circles are those associated with the LAB counterparts ($>$5$\\sigma$ in all 4 IRAC filters). Overlaid are two template SED tracks with redshift:  red denotes Arp220 and  blue shows an irregular  (Im) galaxy \\citep{cww80}. Tracks run from $z$ = 0 to $z$ = 3.09 with the higher redshift end denoted by an open square. Also shown (black line) is the location of pure power-law galaxies ($S ~ {\\propto} ~ \\nu^\\alpha$) with $\\alpha$ = [-0.5,-2.0].  LBGs in the SSA22 field with 8$\\mu$m detections are shown by the blue triangles; the two red stars are those LBGs confirmed as QSOs \\citep{ste03}. The stacked detection of the LABs which are not individually detected at 8$\\mu$m is shown by the open black circle.  % \\label{cc}} \\end{figure}     \\begin{figure} \\epsscale{1.1} \\plotone{f4.eps} \\caption{An IRAC/MIPS color-color diagram, restricted to those objects detected in all four IRAC filters, as well as  at 24$\\mu$m.   Filled black circles correspond to the LAB counterparts;  filled blue triangles denote the LBGs;  filled red stars show the LBGs which are confirmed QSO; and the solid green squares show SMGs with IRAC/MIPS detections which lie at $z\\sim$ 3 from \\citet{ega04} and \\citet{pope06}. Also shown is the stacked detection of the infrared-faint LAB (open black circle).  Overlaid, to show the approximate distinction between the colors of AGN and ULIRGs are the template SEDs M82, Arp 220 (star-forming galaxies) and NGC 5506, NGC 1068 (dusty AGN).  \\label{iracmips}} \\end{figure}  \\subsection{Agreement with Submillimeter Results} \\citet{gea05}  observed 25 of the 35 LABs from \\citet{mat04} at 850$\\mu$m and detected 5 at $>$ 3$\\sigma$ significance. One of these, LAB10, lies outside of our imaged region. We find an almost one to one correspondence between the 8$\\mu$m and 24$\\mu$m detections presented here and the submillimeter detections.  We recover all 4  of the 850$\\mu$m sources (for which we have data) at 8$\\mu$m. At 24$\\mu$m we fail to detect only LAB05 (limits are given in Table 1).  The 24$\\mu$m-850$\\mu$m flux ratios or limits are in good agreement with those measured  for $z\\sim$ 3 SMGs ($<$ 0.02) by \\citet{pope06}.   As extensively discussed by \\citet{pope06} such low values are consistent with galaxies such as Arp220 and Mrk273, which are not typical ULIRGs,  and indicate higher levels of extinction and/or cooler dust temperatures than generally seen for ULIRGs.   Cooler dust temperatures may be reached if the dust is spatially extended  compared to local compact ULIRGs.  This explanation is particularly interesting for LAB01: recent observations with the SMA \\citep{mat07} have failed to detect the $S_{850{\\mu}m} \\sim$ 17$\\mu$Jy source in the blob and this places a lower limit on its spatial extent of $>$ 4{\\arcsec} in the submillimeter (assuming it is  signal dust enshrouded system).   Such an explanation however cannot account for the failure of recent ASTE-AzTec observations to detect this source with a 5$\\sigma$ limit of $\\sim$ 10mJy \\citep{koh08}.  As we showed in the previous section the 8$\\mu$m flux of the 8$\\mu$m-detected systems cannot be due only to stellar emission, but must be enhanced by the presence of hot dust,  associated with prodigious star formation or due to a central dust enshrouded AGN.  We have not detected any 8$\\mu$m source within an LAB which is not also detected at 850$\\mu$m, with the exception of LAB16, whose 850$\\mu$m limit is 4$\\times$ higher than the 4 other 850$\\mu$m sources.  Adopting the average 850$\\mu$m-8$\\mu$m color for the infrared-bright LABs, a non-detection at 850$\\mu$m of $<$ 3mJy (the approximate depth of Geach et al.) would lie below our 8$\\mu$m detection limit and therefore, based on this color alone, the infrared-faint LABs could simply be lower luminosity versions of the same kinds of systems.  A non-detection at 850$\\mu$m \\citep{gea05} and 24$\\mu$m (this work) does not place very stringent limits on the bolometric luminosities of these systems however, and they remain compatible with ULIRG-level luminosities. Alternatively, the infrared-faint systems could be starburst galaxies with low extinction levels and correspondingly lower 850$\\mu$m-8$\\mu$m ratios. This scenario is in agreement with the stacked measurement on Figure \\ref{cc} and with the additional color results discussed in the following sections.   We can rule out with certainty however,  dust enshrouded QSOs  of similar rest-frame 2$\\mu$m luminosities, but which lack an extended cold dust component.    \\subsection{The rest-frame UV-Optical-NIR properties}   The bulk of the LAB sample is not detected longward of 4.5$\\mu$m, and in addition to the stacking analysis of the previous sections we can also look to the combination of 3.6$\\mu$m and/or 4.5$\\mu$m data with the NIR and optical imaging.   In Figure \\ref{ccr} we show a rest-frame UV/optical/NIR color-color diagram for the LBGs and LABs.  Here again, we see a clear separation between 8$\\mu$m-detected and non-detected systems, be they LABs or LBGs. The infrared-bright galaxies continue to show red colors into the rest-frame UV, indicating either older stellar populations and/or high levels of extinction.  Notably, as seen in Fig \\ref{cc} the infrared-faint LABs are well-separated in color from the infrared-bright systems, and are very similar to the general infrared-faint LBGs.  The exceptions are the two infrared-bright LBG-QSOs, which show extremely blue UV/optical colors.  In Figure \\ref{cm} we show the observed 8.0$\\mu$m/3.6$\\mu$m flux ratio as a function of observed 8.0$\\mu$m and, alternately, 3.6$\\mu$m flux.  Again we compare to the LBG population, and add to the analysis the SMGs and QSOs as in Figure \\ref{iracmips}.  In Figure \\ref{cm}(a) we see a correlation between the rest-frame 2.0$\\mu$m and the rest-frame NIR/optical colors, in the sense that galaxies which are more luminous at 2.0$\\mu$m exhibit redder colors, and the stacking analysis confirms that this relation continues to fainter levels. Visually, the slope of the relation appears to lie between a flat color-luminosity relation and that of the line representing constant 3.6$\\mu$m flux; that is, the relation results from an overall increase in luminosity and a enhancement of the 2.0$\\mu$m flux over the rest-frame optical.    The rest-frame $\\sim$ 2.0$\\mu$m is often taken as a proxy for stellar mass, however we have already argued that this wavelength is contaminated by AGN or intense star-formation in many of these systems.  Indeed, the most luminous systems at observed 8.0$\\mu$m are the confirmed QSOs, followed by the LABs and then the infrared-bright LBG population. Though clearly not a complete comparison sample this does again hint that the LABs are transition objects between pure AGN and star-formation dominated systems.  In Figure \\ref{cm}(b) we plot the same color, but now against observed 3.6$\\mu$m flux.   In addition to illustrating the spread in 3.6$\\mu$m luminosities, this allows us to probe the observed 8$\\mu$m flux of the comparison LBG population  as a function of 3.6$\\mu$m flux.  The statistical detection of the 8$\\mu$m-faint LABs fits within  the general behavior of the LBGs, which become markedly bluer in the optical/NIR with decreasing 3.6$\\mu$m luminosity.  Thus,  the fainter, presumably less massive systems contain lower levels of dust, star formation and AGN activity.  \\begin{figure} \\epsscale{1.1} \\plotone{f5.eps} \\caption{ A rest-frame optical/UV color-color diagram. The small grey points correspond to the 3.6$\\mu$m-selected population in the SSA22 field. Points are as defined in \\ref{cc} and \\ref{cm} with the exception of the open triangles (individual 8$\\mu$m-faint LBGs) and the open circles (individual 8$\\mu$m-faint LABs). Here we denote 8$\\micron$-detected LABs and LBGs as infrared-luminous (IL) and 8$\\micron$ non-detections as infrared-faint (IF).  Also shown are the SED tracks (from $z =$ 0 - 3.09)  for Arp 220 (red) and an irregular galaxy \\citep{cww80} with the open square denoting $z =$ 3.09. \\label{ccr}} \\end{figure}   \\begin{figure} \\epsscale{1.1} \\plotone{f6.eps} \\caption{(a) The optical/NIR color as a function of rest-frame $\\sim$ 2$\\mu$m.  The vertical dotted line denotes the approximate 5$\\sigma$ limit of the 8$\\mu$m data. The diagonal dotted line corresponds to the 8$\\mu$m/3.6$\\mu$m flux ratio, at the 3.6$\\mu$m detection limit. Points are the same as in Figure 4. The filled black circles are the 8$\\mu$m-detected LAB counterparts; the filled blue triangles are the 8$\\mu$m-detected LBGs, with the two solid red stars corresponding to the confirmed QSOs, and the solid green squares are  SMGs \\citep{ega04,pope06} for comparison. The open black circle and blue triangle  below the 8$\\mu$m detection limit correspond to stacked measurements for the 8$\\mu$m-faint LABs and LBGs respectively.  (b)    The optical/NIR color as a function of rest-frame $\\sim$0.9$\\mu$m. The dotted line shows the color limit at the 8$\\mu$m depth. The points are as in (a), however the LBG stacked detections include both 8$\\mu$m-faint  and  8$\\mu$m-bright objects. \\label{cm}} \\end{figure}         \\subsection{Ly$\\alpha$ excitation mechanisms and implications for galaxy evolution within the SSA22 structure }  The two central questions concerning LABs are:  what drives the extended Ly$\\alpha$ emission and, more generally, how are these structures related to galaxy formation?  There are two basic scenarios for the relationship between the LABs and their counterpart galaxies.  In the first, the embedded object is the excitation power source through processes such as photoionization by a strong UV emitter or by shock-heating due to super-galactic scale winds from a very  strong starburst or a central AGN.  In the second scenario, the Ly$\\alpha$ emision could be due to heating through gravitational collapse; in such a picture, the LAB and counterpart have no causal relationship  (that is, the galaxy is not driving the Ly$\\alpha$ emission), but rather both mark sites of galaxy formation activity. In either case we might expect a correlation between the properties of the counterpart and the Ly$\\alpha$ halo; in the first case the correlation is obvious and causal while in the second case we might, for example, expect the most massive and luminous LABs to trace regions of correspondingly massive galaxy formation activity such as major-mergers, AGN, and extreme infrared starbursts.   Much work has been done to attempt to separate these different scenarios, but no common picture which works for all LABS has yet emerged.    M04 extensively discussed the possible excitation mechanisms of the 35 LABs in SSA22.  Using simple assumptions, such as a Salpeter IMF and an unhindered line of sight,  they showed that $\\sim$ 1/3 of the LABs do not contain galaxies which are sufficiently UV-luminous to power the LABs. Of   the six 8$\\mu$m-detected LABs discussed here,   four (LAB01, LAB02, LAB05, LAB14) are not sufficiently bright in the rest-frame UV  unless the bulk of the ionizing radiation is emitted away from our line of sight. This is not an unreasonable explanation since all six objects are dusty and IR-luminous, as evidenced by their 8$\\mu$m, 24$\\mu$m, and/or 850$\\mu$m emission.  The remaining  LABs however, have UV/optical/IR colors which imply lower levels of dust extinction and are therefore unlikely to be highly obscured, but still lack obvious UV emitters.  By probing both hot and cold dust emission a number of authors \\citep{gea05, col06,dey05} have shown that a small fraction of LABs host infrared-bright systems of ULIRG and HyLIRG  luminosity levels and importantly, that the inferred bolometric luminosity of these systems is weakly correlated with the Ly$\\alpha$ luminosity, though roughly three orders of magnitude larger.  To the infrared-bright sample of \\citet{gea05} we add LAB16;  although not detected at 850$\\mu$m we may extrapolate from the 24$\\mu$m flux (assuming a similar SED as the other infrared-bright LABs) and find that LAB16 follows the roughly same L$_{Ly\\alpha}$-L$_{bol}$ relation. As argued by \\citet{gea05}, this suggests that the process which heats the dust, i.e.~a powerful AGN or intense star formation, is also responsible for exciting the gas halo, and a natural mechanism for this is galaxy-scale winds; however, this relation could simply indicate a natural non-causal correlation between the size of the collapsing gas cloud and the galaxy formation activity occurring within it.  We can use the data presented here to add to these studies by directly probing the nature of the LAB counterparts, though as with the earlier work  this cannot prove a causal connection between the embedded galaxy and the gas emission. Throughout the paper however we have included similar analyses on the LBG population within the SSA22 structure, with the motivation that a comparative study of those systems which do not exhibit extended Ly$\\alpha$ emission might further elucidate the processes responsible for the LABs.  We can divide the sample into two distinct groups:  those systems, be they LBGs or LABs, which are detected at 8$\\mu$m, hereafter referred to as  infrared-luminous (IL), and galaxies which are not detected at 8$\\mu$m,  or infrared-faint (IF).  Let us first discuss the IL sample. The MIR/NIR colors of the LBGs and LABs within this group are similar and are best described by either a dusty AGN power-law or an intense starburst ULIRG.  Some of these have been detected in the X-ray with emission levels consistent with though not uniquely indicative of  AGN. Differences begin to arise when we look at the UV/optical colors or include the rest-frame 6$\\mu$m in the analysis (Figures \\ref{ccr} and \\ref{iracmips}).   The bulk of the LBG population is systematically bluer in UV/optical color than the LABs, while in MIR color there is a clear continuum of colors between the bulk of the LBGs, the LABs and the LBG-QSOs.    Finally, in Figure \\ref{cm} we see hints of a similar gradient in rest-frame 2$\\mu$m luminosity, such that the LABs lie between the confirmed LBG-QSO and the remainder of the LBG population. This hints that, like the SMG population, IL-LABs may harbor deeply enshrouded AGN, in addition to intense star formation.  The following toy-model might therefore be proposed for these systems. The most luminous LABs are powered through galaxy-scale superwinds driven by intense star formation and AGN activity. As the AGN grows it eventually becomes powerful enough to completely blow-out the gas cloud and enshrouding dust and emerges as a naked AGN, still luminous at the shorter IR wavelengths but no longer an LAB. It is not clear from these data how the IL-LBGs fit into this picture.  As suggested by \\citet{dey05}, one might speculate that the brightest LBGs are older systems  whose AGN has turned off, hence their lower 8$\\mu$m fluxes, but whether a LAB phase is common to all LBGs is not yet clear.  This scenario could however account for the variety of morphologies of the LABs and the location of the galaxies within them. Take for example the three largest LABs, LAB01, LAB02 and the LAB of \\citet{dey05}:  all have likely AGN counterparts near the edge of the Ly$\\alpha$ emission, and these may be systems caught in the process of emerging from their parent cloud.  One might posit that an AGN is required to power the intense Ly$\\alpha$ emission,  and this is the root of the difference between the LBG without extended Ly$\\alpha$ emission and LAB populations.  This seems at odds however, with the IF systems.   Though the data are sparse for these groups, we see no difference between the properties of the IF LAB counterparts and the IF LBGs: both appear to be pure star-forming systems with roughly similar masses (now using the 8$\\mu$m luminosity as a proxy for mass). In these systems AGN are not driving the Ly$\\alpha$ emission, nor can the presence of a Ly$\\alpha$ halo for one group but not the other be explained by obvious differences in  their counterpart properties.  It seems more likely therefore that the difference is in evolutionary phase over a cross-section of masses;  those galaxies with extended Ly$\\alpha$ emission are in  earlier phases of formation than the older LBGs.  In this case the Ly$\\alpha$ emission could be powered by superwinds, as discussed above, or may be unrelated to (but still correlated with) the embedded galaxies and driven instead by gravitational heating, as suggested, for example, by \\citet{dey05} and \\citet{smi07}.  Again, however, such a scenario would seem to imply that an LAB phase is common to all systems, unless secondary factors such as morphology play a role.      \\subsection{Conclusions} We have used mid-infrared imaging from Spitzer IRAC and MIPS 24$\\mu$m to investigate the nature of a large sample of  LAB systems within the SSA22 filament at $z=$ 3.09.   We have detected 6/26 of the objects in all four IRAC filters, with four also detected with MIPS 24$\\mu$m, and find excellent correspondence with the 850$\\mu$m detections of \\citet{gea05}.  By analyzing the rest-frame UV, optical, and near/mid/far-infrared colors, and archival X-ray imaging, we show that these six systems exhibit signs of AGN activity and/or intense and dusty star formation;  notably, they exhibit colors which bridge the star-forming LBG galaxies within the SSA22 structure and objects optically identified as QSOs. Through a stacking analysis see that the infrared faint LABs are markedly different than the infrared-luminous systems, and are consistent with blue star-forming galaxies, with no evidence of an AGN contribution. These results are in line with a model in which the LAB Ly$\\alpha$ luminosity is powered by large-scale superwinds, driven by star-formation in the smaller  systems and a combination of AGN and extreme starbursts in the most massive. The presence of such a halo may be dictated by the evolutionary phase of the galaxy, such that the halo-less LBGs are older systems than the LABs. A scenario in which the LAB is the result of gravitational heating and the galactic activity is not driving the LAB luminosity cannot be ruled out, however."
    },
    "0812/0812.3395_arXiv.txt": {
        "abstract": "We identify a new secular instability of eccentric stellar disks around supermassive black holes. We show that retrograde precession of the stellar orbits, due to the presence of a stellar cusp, induces coherent torques that amplify deviations of individual orbital eccentricities from the average, and thus drive all eccentricities away from their initial value. We investigate the instability using $N$-body simulations, and show that it can propel individual orbital eccentricities to significantly higher or lower values on the order of a precession time-scale.  This physics is relevant for the Galactic center, where massive stars are likely to form in eccentric disks around the SgrA* black hole. We show that the observed bimodal eccentricity distribution of disk stars in the Galactic center is in good agreement with the distribution resulting from the eccentricity instability and demonstrate how the dynamical evolution of such a disk results in several of its stars acquiring high ($1-e\\ll 0.1$) orbital eccentricity. Binary stars on such highly eccentric orbits would get tidally disrupted by the SgrA* black hole, possibly producing both S-stars near the black hole and high-velocity stars in the Galactic halo. ",
        "introduction": "Most well-observed galaxies with stellar bulges show evidence for supermassive black holes (SMBHs) in their centers \\citep*{Kormendy:2004p2865}. An important feature of stellar dynamics in galactic nuclei is that for orbits contained well within the SMBH's radius of influence, the precession time $t_{\\rm prec}$ is many orders of magnitude larger than their period $P$.\\footnote{This is in contrast to the rapidly precessing stellar orbits typical for globular clusters or galaxies.} Therefore, for a study of long-term (secular) gravitational dynamics, it is useful to think of the stars not as particles but as massive elliptical wires stretched along the stars' eccentric orbits  \\citep{Rauch:1996p270}. Gravitational torques between the wires can become the dominant mode of orbital eccentricity evolution, as is seen in the case of resonant relaxation \\citep{Rauch:1996p270,  Gurkan:2007p42, Touma:2009p4317,Eilon:2008p2923}. Secular effects have a significant impact  on the stability of stellar cusps \\citep{Tremaine:2005p2846, Polyachenko:1981p3004, Polyachenko:2008p2804},  on gravitational wave emission from in-spiraling compact objects \\citep*{Hopman:2006p1864, Hopman:2006p1668}, on tidal disruption of stars \\citep{Rauch:1998p2872} and, as we show here, on the origin of the S-stars in the Galactic center. \\\\  In this Letter we describe a mechanism for generating high-eccentricity stars via a fast secular instability of an eccentric stellar disk (\\S\\ref{sec:ins}). We use $N$-body simulations to investigate the non-linear development of the instability (\\S\\ref{sec:sims} and \\S\\ref{sec:results}) and to explore its dependence on the initial eccentricity and mass of the disk. We then show that the instability may be relevant for young stars in the Galactic center, and may provide an interesting channel for formation of the S-stars and of high-velocity stars in the Galactic halo (\\S\\ref{sec:sstars}). We critically discuss our results in \\S\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  So far, we have described the eccentricity instability for an idealized initial disk configuration. We now demonstrate the effectiveness of the instability for a broader scope of parameters and extend our disk model to include (1) a significant range in both magnitude and direction of the stellar eccentricity vectors, (2) heavily inclined stellar orbits, and (3) different values of $\\gamma$, the power-law index of the cusp, which determines the precession rate within the disk. We discuss each of these in turn. \\\\  (1) We simulate disks with initial eccentricity vectors scattered between a range of opening angles $\\theta \\in [0, 2 \\pi ]$. As anticipated, significant spreading of the vector decreases the effectiveness of the instability. However, we find sufficiently high eccentricity orbits to generate S-stars up to $\\theta \\sim \\pi \\rads$, which suggests that the instability remains important for substantially less favorable circumstances. Next we vary the magnitudes of the eccentricities throughout the disk, basing our initial conditions on Run 1 and 2 of \\citet{Bonnell:2008p3000}. Run 1 (0.6 $<$ $e$ $<$ 0.7) is in excellent agreement with our previous results. The instability in Run 2 (0 $<$ $e$ $<$ 0.6) however, is not efficient at generating the highest $e$ orbits and hence the S-stars. These conclusions are transparent; the stars which experience the greatest torque are those at small $a$ (see \\S\\ref{sec:sims}) and if initially circular, need a greater torque to produce the highest $e$ orbits. We find that the most important condition, for $M_{\\rm disk} = 10^4$, is for innermost orbits to have $e > 0.4$. Future work on gas dynamics of eccentric disks will tell whether this presents a serious problem for our scenario.\\\\  (2) We study disks with stellar inclinations $h/r$ in the range $[-0.3, 0.3]$ and find that the instability is robust in all cases. However, above the limit $[-0.2, 0.2]$, the instability appears ineffective in generating the highest $e$ ($> 1 - 10^{-3}$) orbits. \\\\  (3) We vary the power-law slope of the cusp. For $\\gamma < 3/2$, due to the shallowness of the cusp, the innermost orbits lag behind in precession and are more susceptible to being pushed to high $e$ orbits. While this improves the ability of the instability to produce S-stars, the instability is not as effective at generating low $e$ orbits. Conversely, an increase in $\\gamma$ $(> 3/2)$ results in the innermost orbits precessing faster than the bulk of the stars and hence suppresses the generation of the highest $e$ orbits and lengthens the time-scale on which it occurs. We conclude that the production of S-stars is unlikely for a cusp with $\\gamma > 7/4$. In all cases examined ($1 < \\gamma < 2$), the instability proves robust and, as emphasized in \\S\\ref{sec:ins}, there is both observational and theoretical evidence that the slope of the cusp in the Galactic center is sufficiently close to $3/2$ for the instability to be highly effective. \\\\  The eccentricity instability described in this work is relevant for an eccentric disk embedded in a stellar cusp. We note that the eccentric disk in M31 \\citep*{Tremaine:1995p3101} is not subject to our instability, since it is not embedded in any visible cusp. Indeed we find that such a disk remains coherent over several precession times, although the innermost stars can undergo significant angular momentum changes as described by \\citet{Cuadra:2008p3486}. The situation is different in the Galactic center, where massive stars were likely born from an eccentric gaseous disk, and where there is strong observational evidence for a heavy spherical stellar cusp \\citep{Genzel:2003p870, Schodel:2007p1797}. We have shown that the instability may have strongly influenced the development of the eccentricities of the massive stars observed, and in particular, can account for the bimodality in their eccentricity distribution. Furthermore, combined with Hills' (\\citeyear{Hills:1988p3075}) mechanism, the eccentricity instability may have played a role in the formation of the S-stars and of the high-velocity stars in the Galactic halo."
    },
    "0812/0812.4104_arXiv.txt": {
        "abstract": "We study the Wouthuysen-Field coupling at early universe with numerical solutions of the integrodifferential equation describing the kinetics of photons undergoing resonant scattering. The numerical solver is developed based on the weighted essentially non-oscillatory (WENO) scheme for the Boltzmann-like integrodifferential equation. This method has perfectly passed the tests of analytic solution and conservation property of the resonant scattering equation. We focus on the time evolution of the Wouthuysen-Field (W-F) coupling in relation to the 21 cm emission and absorption at the epoch of reionization. We especially pay attention to the formation of the local Boltzmann distribution, $e^{-(\\nu-\\nu_0)/kT}$, of photon frequency spectrum around resonant frequency $\\nu_0$ within width $\\nu_l$, i.e. $|\\nu-\\nu_0| \\leq \\nu_l$. We show that a local Boltzmann distribution will be formed if photons with frequency $\\sim \\nu_0$ have undergone a ten thousand or more times of scattering, which corresponds to the order of 10$^3$ yrs for neutral hydrogen density of the concordance $\\Lambda$CDM model. The time evolution of the shape and width of the local Boltzmann distribution actually doesn't dependent on the details of atomic recoil, photon sources, or initial conditions very much. However, the intensity of photon flux at the local Boltzmann distribution is substantially time-dependent. The time scale of approaching the saturated intensity can be as long as 10$^{5}$-$10^{6}$ yrs for typical parameters of the $\\Lambda$CDM model. The intensity of the local Boltzmann distribution at time less than 10$^{5}$ yrs is significantly lower than that of the saturation state. Therefore, it may not be always reasonable to assume that the deviation of the spin temperature of 21 cm energy states from cosmic background temperature is mainly due to the W-F coupling {\\it if} first stars or their emission/absorption regions evolved with a time scale equal to or less than Myrs. ",
        "introduction": "It is generally believed that the physical state of the universe at the epoch of reionization can be probed by detecting the redshifted 21 cm signals from the ionized and heated regions around the first generation of stars (e.g. Furlanetto, Oh, \\& Briggs 2006).  The emission and absorption of 21 cm are caused by the deviation of the spin temperature $T_s$ of neutral hydrogen from the temperature of cosmic microwave background (CMB) $T_{\\rm CMB}$ at the considered redshift $z$. Many calculations have been done on  the 21 cm emission and absorption from ionized halos of the first stars (Chuzhoy et al. 2006; Cen 2006; Liu et al. 2007). A common assumption of these calculations is that the deviations of $T_s$ from $T_{\\rm CMB}$ are mainly due to the Wouthuysen-Field (W-F) coupling (Wouthuysen, 1952; Field, 1958, 1959). That is, the resonant scattering of Ly$\\alpha$ photons with neutral hydrogen atoms locks the color temperature $T_c$ of the photon spectrum around the Ly$\\alpha$ frequency to be equal to the kinetic temperature of hydrogen gas $T$. Consequently, the spin degree of freedom is determined by the kinetic temperature of hydrogen gas $T$.  The W-F coupling is from the kinetics of photons undergoing resonant scattering, which is described by a Boltzmann integrodifferential equation. All the above-mentioned calculations are based on time-independent solution of the resonant scattering kinetic equation with Fokker-Planck approximation (Chen \\& Miralda-Escude, 2004; Hirata, 2006; Furlanetto \\& Pritchard 2006; Chuzhoy \\& Shapiro 2006). This is equal to assume that the time scale of the onset of the W-F coupling is less than all time scales related to the 21 cm emission/absorption. However, even in the first paper of the W-F coupling, the problem of time scale has been addressed as follows: ``One can infer from this fact that the photons (in a box), after an infinite number of scattering processes on gas atoms with kinetic temperature $T$, will obtain a statistical distribution over the spectrum proportional to the Planck-radiation spectrum of temperature $T$. After a finite but large number of scattering processes, the Planck shape will be produced in a region around the initial frequency'' (Wouthuysen, 1952). That is, the W-F coupling is onset only ``after a finite but large number of scattering''. One cannot assume that the time-independent solution is available for the W-F coupling if the time-scales of the evolution of the first stars and their emission/absorption regions are short. A study on the time evolution of radiation spectrum due to resonant scattering is necessary.  This problem is especially important for the 21 cm signal from the first stars, as the life times of the first stars are short. The ionized and heated regions around the first stars are strongly time-dependent (Cen 2006; Liu et al. 2007). The 21 cm emission/absortion regions are located in a narrow shell just outside the ionized region. On the other hand, the speed of the ionization-front (I-front) is rather high, even comparable to the speed of light. The time scales of the formation and the evolution of the 21 cm regions can be estimated by $\\sim d/c$, $d$ being the thickness of the shells of 21 cm emission and absorption.  The time-independent solution would be proper only if Ly$\\alpha$ photons approach the time-independent state in a time shorter than that of the 21 cm region evolution.  Very few works have been done on the time dependent behavior of the W-F coupling. There is the lack of a time dependent solution even for the Fokker-Planck approximation. The existed time-dependent solvers (e.g. Meiksin, 2006) cannot pass the tests of analytical solutions (Field, 1959). On the other hand, the WENO algorithm is found to be effective to solve the Boltzmann equation (Carrillo et al. 2003) and radiative transfer (Qiu et al. 2006, 2007, 2008). In this paper, we will study the time-dependent behavior of the W-F coupling with the WENO method by numerically solving the integrodifferential equation. In this context, we also develop a numerical solver in accordance with the term of resonant scattering.  The paper is organized as follows. Section 2 presents the basic equations of the resonant scattering of photons. Section 3 very briefly mentions the numerical solver of the WENO scheme and its test with Field's analytic solutions, leaving the details of algorithmic issues to the Appendix. Section 4 presents the time-dependent W-F coupling with a static background. Section 5 shows the numerical results of the W-F coupling in an expanding background. Finally, conclusions are given in Section 6, in which the application to the 21 cm problem is also addressed. ",
        "conclusions": "\\subsection{Summary}  The onset of the W-F coupling, or the formation of local Boltzmann distribution is similar to the process of approaching a statistically thermal equilibrium state via collisions or scattering. The particle distribution in the statistical equilibrium is independent of time, initial distribution, and the details of collision. The equilibrium distribution is maintained only by the enough collisions among particles. A local Boltzmann distribution is formed once the number of resonant scattering is large enough. Like other statistically thermal equilibrium, the features of the local Boltzmann distribution are independent of time, photon source, initial photon distribution, and etc.  In an expanding universe, photons are moving in the frequency space with ``speed\" given by the redshift. The formation of the local Boltzmann distribution depends on the competition between the resonant scattering and the redshift. If photons have undergone enough scattering during their path through the frequency space around the resonant frequency, a local Boltzmann distribution will be formed. Otherwise local statistical equilibrium cannot be approached.  In our work, we use the Gaussian profile eq.(\\ref{eq3}), but not the Voigt profile. Our major results on the formation of local Boltzmann distribution will not be affected by using the Voigt profile if the width $x_l$ is not larger than the Doppler thermal broadening. We found that the solution of eq.(9) with rest background is not affected by Voigt profile if the ratio $a$ between the natural and Doppler broadening is equal to 10$^{-3}$, which corresponds to $T\\simeq 650$ K. For large $a$, or small $T$, the Doppler thermal broadening is small. In this case, the width $x_l$, in which the color temperature $T_c$ of Ly$\\alpha$ photons is locked to the kinetic temperature $T$ of hydrogen atoms, would also be small.  \\subsection{Applications to the 21 cm problem}  A basic problem for the 21 cm signal from the ionized and heated region around first stars is the conditions on which one can estimate the 21 cm emission and absorption with the W-F coupling, which forces the internal (spin-)degree of freedom to be determined by the thermal motion of the atoms. In this case, the relative occupation of the two hyperfine-structure components of the ground state depends only upon the shape of the spectrum near the Ly$\\alpha$ frequency. Therefore, we need a local Boltzmann distribution of Ly$\\alpha$ photon with frequency width equal to or larger than $|\\nu_f -\\nu_0|\\geq \\nu_{21}=1420$ MHz, or \\begin{equation} x_f \\geq 0.014 \\left (\\frac{10^4}{T}\\right)^{1/2}. \\end{equation} Thus, from Figure 7 and eq.(10) one can conclude that in regions with $f_{\\rm HI}\\leq 10^{-4}$, where $\\gamma > 10^{-1}$, the W-F coupling will not work. Although electron-hydrogen and proton-hydrogen collisions can be important; the 21 cm signal will be incredibly small. On the other hand, in the primarily neutral IGM, W-F coupling is very efficient.  From eqs.(7) and (10), we have \\begin{equation} t=0.26\\times 10^5 h^{-1} \\left(\\frac{T}{10^4}\\right)^{1/2}\\left (\\frac{10}{1+z}\\right )^{3/2}\\left(\\frac{0.25}{\\Omega_M}\\right)^{1/2}\\gamma \\tau \\hspace{3mm} yrs. \\end{equation} We know that the saturation, or time independent solution around resonant frequency, can be used only when the time $\\tau$ is larger than the time of photon moving over a frequency space from $-x_l$ to $x_l$. From Figures 6 and 9, we have $\\tau \\gamma$ equal to about 10 at the saturation. Therefore, eq.(19) yields that the time independent solution is available only if the time scale of the evolution of the ionized sphere of first stars is larger than about $3\\times 10^5$ yrs. This gives a constraint on the 21 cm emission region, as such regions are very narrow, the time scale of the evolution being comparable to $10^6$ yr, or even less (Liu et al. 2007).  Actually we may not need a saturation state. What we used for estimating the 21 cm signals is the frequency distribution $J(x,\\tau)$ to show a Boltzmann-like shape in the central part $|x|\\simeq 0$, i.e. the onset of the W-F coupling. The time-scale of the W-F coupling onset is equal to about 10$^3$ yrs for neutral hydrogen density of the concordance $\\Lambda$CDM model [eq.(7)]. This time scale is much less than that of the evolution of 21 cm region of first stars. It seems to indicate that we can safely use the W-F coupling in the 21 cm estimation of first stars.  However, we should mention the effect of the photon intensity. The coupling coefficient between Ly$\\alpha$ photons and spin temperature is proportional to the intensity $J$ (e.g. Furlanetto, Oh, \\& Briggs 2006). The W-F coupling would generally be suppressed due to the fact that the flux at the resonant frequency $J(x=0)$ is always less than the flux at other frequencies. In a saturated state this suppression is small (Figure 6). However, before approaching the saturated state, the intensity $J$ generally is significantly less than its saturated value. That is, although the local Boltzmann distribution is formed at the time of the order of $\\tau\\simeq 10^4$, the intensity at that time would still be low, and the W-F coupling is not enough to produce the deviation of $T_s$ from $T_{\\rm CMB}$. Therefore, it may not be always reasonable to assume that the deviation of $T_s$ from $T_{\\rm CMB}$ is mainly due to the W-F coupling if first stars or their emission/absorption regions evolved with time scale equal to or less than Myrs.  The WENO algorithm revealed the time evolution of photons undergoing resonant scattering, whose information is generally lost in the asymptotic solutions, or the time-independent solution. Although time-independent solutions provide useful guidance, they do not show the conditions for the efficiency of the W-F coupling at different times. The asymptotic solution probably is never reached for short life-time objects. It would be impossible to correctly estimate observable 21 cm signal from ionized and heated halos of first stars without a correct understanding of the time evolution of the W-F coupling."
    },
    "0812/0812.4873_arXiv.txt": {
        "abstract": "Magnetospheres of neutron stars are anchored in the rigid crust and can be twisted by sudden crustal motions (``starquakes''). The twisted magnetosphere does not remain static and gradually untwists, dissipating magnetic energy and producing radiation. The equation describing this evolution is derived, and its solutions are presented. Two distinct regions coexist in untwisting magnetospheres: a potential region where $\\nabla\\times\\bB=0$ (``cavity'') and a current-carrying bundle of field lines with $\\nabla\\times\\bB\\neq 0$ (``j-bundle''). The cavity has a sharp boundary, which expands with time and eventually erases all of the twist. In this process, the electric current of the j-bundle is sucked into the star. Observational appearance of the untwisting process is discussed. A hot spot forms at the footprints of the j-bundle. The spot shrinks with time toward the magnetic dipole axis, and its luminosity and temperature gradually decrease. As the j-bundle shrinks, the amplitude of its twist $\\ta$ can grow to the maximum possible value $\\tamax\\sim 1$. The strong twist near the dipole axis increases the spindown rate of the star and can generate a broad beam of radio emission. The model explains the puzzling behavior of magnetar \\XTE~ --- a canonical example of magnetospheric evolution following a starquake. We also discuss implications for other magnetars. The untwisting theory suggests that the nonthermal radiation of magnetars is preferentially generated on a bundle of extended closed field lines near the dipole axis. ",
        "introduction": "Neutron stars are highly conducting and strongly magnetized. Their extended magnetospheres are anchored deep in the rigid crust and corotate with the star. The magnetosphere is usually assumed to be static in the co-rotating frame or evolving very slowly as the star ages. Electric currents are confined to a narrow bundle of open field lines that connect the star to its light cylinder (Goldreich \\& Julian 1969). The main, closed, part of the magnetosphere is usually assumed to be current-free and potential, $\\nabla\\times\\bB=0$.  The standard picture of a static potential magnetosphere is reasonable for ordinary pulsars, yet apparently it does not describe all neutron stars. In particular, neutron stars with ultrastrong fields $B\\simgt 10^{14}$~G (magnetars) are inferred to have {\\it dynamic} magnetospheres (see e.g. reviews by Woods \\& Thompson 2006; Kaspi 2007; Mereghetti 2008). Their activity is believed to be caused by crustal motions relieving internal stresses.\\footnote{Well-studied ordinary pulsars (Crab, Vela) also show glitches in their spindown rates, which are associated with sudden crustal deformations.} In contrast to ordinary pulsars, these objects show huge temporal variations in luminosity, spectrum, and spindown rate. Theoretically, stresses are expected to build up in the deep crust as the star ages (e.g. Ruderman 1991; Thompson \\& Duncan 1995). In particular, the large Ampere forces $\\bj\\times\\bB/c$ inside magnetars can break the crust and shear it in a catastrophic way.\\footnote{In some cases, the crust may move plastically.} Such a starquake twists the magnetic field anchored in the crust, creating $\\nabla\\times\\bB\\neq 0$ and inducing electric currents in the closed magnetosphere (Thompson et al. 2000). The currents are approximately force-free and flow along the magnetic field lines, $\\bj\\times\\bB=0$. They emerge from the deep crust sheared in the starquake.  Twisted force-free magnetospheric configurations were studied extensively in the context of the solar corona, and these models can be applied to neutron stars. A simple example is the self-similarly twisted dipole. It was constructed by Wolfson (1995) and applied to neutron stars by Thompson, Lyutikov, \\& Kulkarni (2002). This and similar force-free configurations are {\\it magnetostatic} solutions. A sequence of such configurations may be constructed by changing their boundary conditions, i.e. displacing the footpoints of the magnetic field lines. If the footpoints freeze, the configuration freezes as well. At a first glance, this seems to suggest that the implanted twist must freeze when the starquake ends, and wait for another starquake.  In fact, the magnetosphere must evolve after the starquake, even though it remains anchored in the motionless deep crust. Indeed, energy is continually dissipated in the twisted magnetosphere, because the twist current $\\bj=(c/4\\pi)\\nabla\\times\\bB$ is maintained by a voltage $\\Phi_e\\neq 0$ established along the magnetic field lines. Thompson et al. (2000) estimated voltage $\\Phi_e$ assuming that the currents are carried by electrons and ions lifted from the star's surface against gravity. Beloborodov \\& Thompson (2007; hereafter BT07) found that $\\Phi_e$ is regulated by an $e^\\pm$ discharge.\\footnote{ The discharge on closed field lines differs from that on open field lines (see Arons 2008 and Beloborodov 2008 for a recent discussion of of the polar-cap discharge in ordinary pulsars). In both cases, however, the discharge is ultimately driven by the magnetic twist $\\nabla\\times\\bB\\neq 0$ that imposes an electric current. } The voltage is significant --- comparable to 1~GeV --- and implies a modest lifetime of the twist, comparable to one year.  The untwisting dynamics of the magnetosphere remained, however, unknown. Usually, resistivity in a plasma leads to diffusion of currents across the magnetic field. Voltage $\\Phi_e\\neq 0$ implies an effective resistivity, and one could expect the decaying twist to spread diffusively across the magnetosphere. This expectation is incorrect as will be shown below.  The goal of this paper is to develop an electrodynamic theory of twisted magnetospheres that describes their evolution. We focus on axially symmetric configurations. In this case, the twist is created through a latitude-dependent azimuthal rotation of the crust. An introductory description of twisted magnetic configurations is given in \\S~2, and their untwisting dynamics is qualitatively discussed in \\S~3. In \\S~4 we derive the electrodynamic equation for axisymmetric magnetospheres.  \\S~5 presents solutions to the evolution equation and explores the mechanism of untwisting. Observational effects of this process are described in \\S~6. \\S~7 compares the theory with the recent observations of a starquake in the anomalous X-ray pulsar \\XTE. \\S~8 summarizes the results of the paper and discusses implications for magnetars. ",
        "conclusions": "Electrodynamics of untwisting may be summarized as follows:  1. --- The twist evolution following a starquake is not diffusive spreading that was pictured previously. Instead, the twist current is sucked into the star, a cavity immediately forms in the inner magnetosphere and grows until it erases all of the twist (Fig.~6). A sharp, step-like drop in current density $j$ is maintained at the boundary of the cavity. It is caused by the threshold nature of the discharge that conducts magnetospheric currents.  2. --- As the cavity expands and the j-bundle shrinks toward the magnetic axis, the twist amplitude $\\ta$ in the j-bundle can grow. The growth occurs if $d\\V/du>0$, i.e. the discharge voltage is smaller on field lines extending farther from the star. This is plausible if the current is conducted through $e^\\pm$ discharge.\\footnote{In contrast, if the current-carrying charges were lifted from the star by voltage (\\ref{eq:Vei}), the voltage would be larger for field lines that extend to larger altitudes, i.e. $dV/du<0$. Then the twist on the j-bundle would diminish with time rather than grow.} The j-bundle with the growing twist is shrinking faster with time, so that the total twist energy $\\Etor$ decreases consistently with the Ohmic dissipation rate.  3. --- The growing twist that has reached the threshold for instability $\\tamax={\\cal O}(1)$ is expected to ``boil over'' and drive an intermittent outflow of magnetic energy from the star. The twist in the j-bundle then remains near $\\tamax$ for the rest of its lifetime. It may be regulated by the limit-cycle instability --- the repeated growth of $\\ta$ to $\\tamax$ followed by a sudden reduction of $\\ta$ below $\\tamax$. The value of $\\tamax$ and the nonlinear behavior of the outer magnetosphere at $\\ta\\sim\\tamax$ needs to be studied further using numerical simulations.  4. --- In addition to rare large starquakes, the magnetosphere can be gradually twisted by the continued motion of its footpoints, either plastic or through a sequence of small starquakes. The continued footpoint motion leads to either a very strong twist or a tiny (negligible) twist that has no observational effects. A quasi-steady state with a non-negligible twist amplitude is possible only with $\\ta\\sim\\tamax$.  Recent observations of \\XTE~ are particularly useful for testing the untwisting theory, for a few reasons: (1) \\XTE~ displayed a clean post-starquake evolution, which was observed for years uninterrupted by new starquakes. (2) The low level of quiescent luminosity from this object allows one to see clearly the shrinking hot spot on the star, and its evolving luminosity was tracked from $10^{35}$~erg/s down to $10^{33}$~erg/s. (3) The detection of radio pulsations and detailed measurements of spindown torque make this object yet more interesting for testing theoretical models. For these reasons, this paper focused mainly on \\XTE (\\S~7). Recently, an outburst was detected in the similar radio magnetar \\AXP~ (Camilo et al. 2008; Halpern et al. 2008). Its behavior appears to be more complicated than that of \\XTE; apparently, episodes of repeated (and overlapping in time) activity occurred a few months after the outburst. The analysis of this object is deferred to a future work.  Other, more active, AXPs and SGRs display a diverse and complicated behavior of X-ray luminosity, pulse profile, and spindown rate (see Woods \\& Thompson 2006; Kaspi 2007; Mereghetti 2008 for reviews). Repeated starquakes of various amplitudes and possible plastic deformations of the crust make these objects more difficult to analyze. The continuing injection of a magnetospheric twist can slow down the decay of its luminosity. It can also affect the star's spindown in a more complicated way than described in \\S~\\ref{sc:torque} for an isolated starquake. In general, twist injection should lead to higher X-ray activity and faster spindown (Thompson et al. 2002). Such a general correlation exists in the magnetar population (Marsden \\& White 2001) but not always observed in individual objects. Note that spindown is controlled by the X-ray-dim bundle of open field lines. The impact of a starquake on this narrow bundle may be immediate, occur with a delay, or never occur, depending on the starquake geometry and amplitude (cf. the case of a ``ring'' starquake in \\S~\\ref{sc:ring}). Repeated starquakes may lead to a non-trivial relation between spindown and X-ray emission.  Despite the diverse behavior of active magnetars, some general features may be inferred. It is clear that the observed sources do not have strong global twists, for two reasons. First, the evolution timescale of such twists would be too long, $\\tev\\sim(10-10^2)\\V_9^{-1}$~yr, where $\\V_9$ is the discharge voltage in units of $10^9$~V (see eq.~\\ref{eq:tev}). In contrast, the magnetospheres of observed magnetars usually evolve on timescales $\\sim 1$~yr or even shorter.\\footnote{The theoretical $\\tev$ would be reduced for a larger discharge voltage $\\V\\gg 10^9$~V, however it appears impossible to sustain such a voltage as it leads to runaway $e^\\pm$ creation (BT07). Even if $\\V\\gg 10^9$~V were theoretically possible, strong global twists would still be ruled out by observations because they are overluminous.} Second, the luminosity produced by a strong global twist would be too high, $L\\sim 10^{37}\\V_9$~erg~s$^{-1}$ (see eq.~\\ref{eq:lum}). It is 2 orders of magnitude higher than the typical observed luminosities of magnetars, $L\\sim 10^{35}$~erg/s (e.g. Durant \\& van Kerkwijk 2006). This leaves two possibilities: the twist is weak ($\\ta\\ll 1$) or localized to a narrow bundle of field lines. The changing spindown rate suggests a strong twist, at least near the magnetic dipole axis. Thus, observations appear to support the picture of a strongly twisted j-bundle near the dipole axis. It is certainly supported by the observations of \\XTE~ (\\S~\\ref{sc:XTE}).  The localized strong twist may be created by starquakes localized to the polar region. It also tends to form dynamically from a weak global twist as the j-bundle shrinks to the axis. The luminosity and evolution timescale of an untwisting magnetosphere with $\\uf=\\sin^2\\theta_\\star\\sim 0.1$ and $\\ta\\sim 1$ is consistent with typical $L\\sim 10^{35}$~erg/s and $\\tev\\sim 1$~yr of active magnetars (see eqs.~\\ref{eq:lum} and \\ref{eq:tev}). Both nonthermal X-ray components in the magnetar spectra, 1-20~keV and 20-300~keV, can be produced by the j-bundle. They are likely emitted at different radii. If the hard component is produced near the star, where the j-bundle is narrow, a relatively narrow 20-300~keV pulse is expected.  The resonant-scattering model for 1-20~keV radiation is consistent with the picture of a narrow j-bundle near the dipole axis. The cyclotron resonance for $e^\\pm$ with keV photons takes place at radii $r\\sim 10R$. This means that the resonant scattering is confined to the field-line bundle with $u=R/\\Rmax\\sim 0.1$ (\\S~\\ref{sc:nonth}). The luminosity of upscattered radiation $L_{\\rm sc}$ is supplied by Ohmic dissipation of electric currents in this bundle.  The footprints of a bundle with $u\\sim 0.1$ form $\\sim 3$~km spot on the star. A significant fraction of energy released in the j-bundle may be transported to its footprints and emitted there quasi-thermally. If the spot radiates a significant part of the bundle luminosity, $L\\sim 10^{35}$~erg/s, then it must have a temperature $kT\\approx 1$~keV. Such hot spots were observed in \\XTE~ and \\AXP, but were not reported for other, brighter magnetars. The typical reported temperatures of blackbody components are $0.3-0.6$~keV (e.g. Perna et al. 2001). The disappearance of hot spots in the presence of large $L_{\\rm sc}$ may be caused by the outward drag created by resonant scattering, which suppresses the transport of released energy to the footprints of the j-bundle."
    },
    "0812/0812.1667_arXiv.txt": {
        "abstract": "{}{As an extension of a previous work, we present a comparison of four methods of filtering solar-like variability to increase the efficiency of detection of Earth-like planetary transits by means of box-shaped transit finder algorithms. Two of these filtering methods are the harmonic fitting method and the iterative non-linear filter that, coupled respectively with the Box Least-Square (BLS) and Box Maximum-Likelihood algorithms, demonstrated the best performance during the first detection blind test organized inside the CoRoT consortium. The third method, the 3-spot model, is a simplified physical model of Sun-like variability and the fourth is a simple sliding boxcar filter.} { We apply a Monte Carlo approach by simulating a large number of 150-day light curves (as for CoRoT long runs) for different planetary radii, orbital periods, epochs of the first transit and standard deviations of the photon shot noise. Stellar variability is given by the Total Solar Irradiance variations as observed close to the maximum of solar cycle 23. { After filtering solar variability, transits are searched for by means of the BLS algorithm.}} {We find that the iterative non-linear filter is the best method to filter light curves of solar-like stars when a suitable window can be chosen. As the performance of this filter depends critically on the length of its window, we point out that the window must be as long as possible, according to the magnetic activity level of the star. We show an automatic method to choose the extension of the filter window from the power spectrum of the light curves.} {The iterative non-linear filter, when used with a suitable choice of its window, has a better performance than more complicated and computationally intensive methods of fitting solar-like variability, like the 200-harmonic fitting or the 3-spot model.} ",
        "introduction": "To date more than 50 transiting planets over about 300 known extrasolar planets have been discovered\\footnote{See http://exoplanet.eu/}. They are the most interesting to study since transits remove the degeneracy \\-bet\\-we\\-en orbital inclination and amplitude of the radial velocity curve providing us with information on planetary masses and radii. Only for transiting planets it is possible to develop accurate models of their internal structures, study their atmospheres through transmission spectroscopy or infrared emission and detect possible spin-orbit misalignments by measuring the Rossiter-McLaughlin effect.  The new frontier is the search for transiting terrestrial pla\\-nets. Very recently four super-Earths have been discovered, three of them by means of the radial velocity technique \\citep{Mayoretal08} and one by microlensing \\citep{Bennettetal08}. The latter is a $3.3$ Earth-mass planet and is \\-pro\\-ba\\-bly the lowest mass exoplanet found to date, apart from planets around pulsars \\citep{WolszczanFrail92, Wolszczan94}. To detect the transits of terrestrial \\-pla\\-nets we need to go to space since the photometric precision from the ground is limited to the millimagnitude level because of the scintillation effects produced by the Earth's atmosphere. The CoRoT space mission, currently operating, can reach the \\-pho\\-to\\-me\\-tric precision to detect transits of Earth-size planets in short period orbits around solar-like stars \\citep{Baglin03, Bordeetal03}. The Kepler space mission, whose launch is planned in spring 2009, has been specifically designed to find also Earthly twins, i.e., Earth-size planets in Earth-like orbits, transiting solar-type stars \\citep{Boruckietal04}.  One common approach to searching for planetary transits in light curves of main-sequence late-type stars  requires two main steps: first, the filtering of stellar va\\-ria\\-bi\\-li\\-ty to remove distortions produced by the presence of photospheric cool spots and bright faculae, whose visibility is modulated by stellar rotation; secondly, the search for transits in the filtered light curve by means of suitable detection algorithms. The intrinsic stellar variability represents the main source of astrophysical noise in the detection of transits of terrestrial planets even in relatively inactive stars \\citep[cf., e.g., ][]{Defayetal01, Jenkins02, AigrainIrwin04, Aigrainetal04}.  Several methods to reduce the impact of solar-like \\-va\\-ria\\-bi\\-li\\-ty on the detection of planetary transits have been developed \\citep [see e.g.] [] {Defayetal01, Jenkinsetal02, Carpanoetal03, AigrainIrwin04, Moutouetal05, Reguloetal07}. Some of them were tested on simulated light curves during the first CoRoT blind test carried out by \\citet{Moutouetal05}. The best performance in terms of reduction of missed detections and false alarms was achieved by team 3 who made use of a linear combination of 200 harmonic functions to filter stellar variability and the Box-fitting Least Square (BLS) algorithm by \\citet{Kovacsetal02} to search for transits in the filtered light curves. Team 5 got the second best performance using the iterative non-linear filter by \\citet{AigrainIrwin04} in combination with the Box Maximum Likelihood transit finder algorithm (see \\citealt{AigrainIrwin04}, Sect. 2).  In a recent paper, \\citet{BonomoLanza08} proposed a different method to treat stellar variability, called the 3-spot model, based on the rotational modulation of the flux produced by three point-like active regions \\citep{Lanzaetal03, Lanzaetal07}. By analysing a large number of simulated light curves with photon noise, Sun-like variability and planetary transits, they compared its performance with that of the 200-harmonic fitting \\-u\\-sing the same transit detection algorithm (BLS), to search for transits after the filtering process. They found that the 3-spot model has a better performance than the 200-harmonic fitting when the standard deviation of the noise is 2--4 times larger than the central depth of the transits. On the other hand, the 200-harmonic fitting reduces more efficiently the impact of stellar variability when the standard deviation of the noise is comparable to the transit depth. \\citet{BonomoLanza08} showed that the poor performance of the 200-harmonic fitting in the former case is due to the use of orthogonal functions to fit stellar variability, which makes it significantly affected by the Gibbs phenomenon \\citep{MorseFeshbach54}. This latter reduces the depth of the transits in the filtered light curves thus lowering the efficiency of detection in the presence of noise (see \\citealt{BonomoLanza08}, Sect.5).  The Gibbs phenomenon also affects other filtering methods proposed to reduce the impact of stellar microvariability, such as, e.g., those of \\citet{Carpanoetal03}, or the Wiener-like discrete filters by \\citet{AigrainIrwin04}. Wavelet-based methods \\citep[e.g., ][]{Jenkins02, Reguloetal07} may be useful to overcome the problems related to the Gibbs phenomenon, given the non-orthogonal nature of their basis functions, but they are negatively affected by gaps in the time series. On the other hand, the iterative non-linear filter proposed by \\citet{AigrainIrwin04} is practically insensitive to gaps or irregular sampling and does not make use of orthogonal functions. Therefore, in the present paper, we extend our comparison to it and to another simple sliding boxcar filter which has been recently applied by, e.g., \\citet{Bordeetal03} and \\citet{CarpanoFridlund08}. ",
        "conclusions": "We have performed extensive numerical experiments to compare the performance of four different variability filters for the detection of Earth-like planetary transits by means of a box-shaped transit finder algorithm. The INL filter has proved to be the best method to filter light curves of quiet solar-like stars when a sufficiently long window can be chosen. We have shown that the choice of the window of the filter is  critical since its performance depends significantly on it. We point out that the window must be as long as possible, according to the magnetic activity level of the star. A method to choose the extension of the window, similar to that proposed by \\citet{Reguloetal07}, is shown in Sect.~\\ref{disc}. The INL filter, when used with a sui\\-ta\\-ble choice of its window, has a better performance than more complicated and \\-com\\-pu\\-ta\\-tio\\-nal\\-ly intensive methods of fitting solar-like variability, like the 200-harmonic fitting or the 3-spot model."
    },
    "0812/0812.3662_arXiv.txt": {
        "abstract": "We present a homogeneous X-ray analysis of all 318 Gamma Ray Bursts detected by the X-ray Telescope on the \\swift\\ satellite up to 2008 July 23; this represents the largest sample of X-ray GRB data published to date. In Sections 2--3 we detail the methods which the \\swift-XRT team has developed to produce the enhanced positions, light curves, hardness ratios and spectra presented in this paper. Software using these methods continues to create such products for all new GRBs observed by the \\swift-XRT. We also detail web-based tools allowing users to create these products for any object observed by the XRT, not just GRBs. In Sections 4--6 we present the results of our analysis of GRBs, including probability distribution functions of the temporal and spectral properties of the sample. We demonstrate evidence for a consistent underlying behaviour which can produce a range of light curve morphologies, and attempt to interpret this behaviour in the framework of external forward shock emission. We find several difficulties, in particular that reconciliation of our data with the forward shock model requires energy injection to continue for days to weeks. ",
        "introduction": "\\label{sec:intro}  The \\swift\\ satellite (Gehrels \\etal2004) has revolutionised our understanding of Gamma Ray Bursts (GRBs), and the X-ray telescope (XRT, Burrows \\etal2005) has played a major role in this process. For example, XRT data have shown that the early X-ray light curve is much more complex than originally thought, often containing a shallow-decay `plateau' phase interpreted variously as energy injection in a forward shock (Zhang \\etal2006), reverse shock emission (Uhm \\&\\ Beloborodov 2007) and dust-scattering of the prompt GRB emission (Shao \\&\\ Dai 2007; although this has recently been ruled out by Shen \\etal2008). The high impact of the XRT has been made possible by \\swift's unique rapid slewing capability: the large field-of-view Burst Alert Telescope (BAT, Barthelmy \\etal2005a) detects GRBs, and \\swift\\ then automatically repoints itself so that the narrow-field XRT and UV/Optical Telescope (UVOT; Roming \\etal2005) can observe the new burst within \\til60--90 seconds of the trigger. The XRT has detected \\til95\\% of \\swift-BAT triggered GRBs, most of those within minutes of the trigger, usually providing the most accurate rapidly available localisation of the GRB and enabling early follow up with ground telescopes. \\swift's X-ray afterglow light curves and spectra, starting at typically $<$100 s have been crucial to the discoveries made by \\swift\\ (see Zhang \\etal2007; O'Brien \\etal2006; Willingale \\etal2007, for reviews of the impact of \\swift\\ on GRB science).  Because GRBs fade rapidly, follow-up observers need to make quick decisions about whether or not to invest observing time on a given burst. This is especially true of potentially unusual bursts, such as high-redshift candidates or under-luminous GRBs. Thus it is desirable for GRB data to be rapidly available and analysed quickly, reliably and ideally in a uniform manner. \\swift\\ data are downlinked many times per day and are immediately processed by the Swift Data Center (SDC) at Goddard Space Flight Center and made available to the public via the SDC and data centres in the UK and Italy, minutes to hours after the downlink. We (the \\swift-XRT team) have developed software which, when data arrive at the UK Swift Science Data Centre (UKSSDC), automatically determines the best possible XRT GRB position and builds X-ray light curves and spectra; the results are then published on the internet. This provides a homogeneously generated catalogue of data, which we detail and discuss in this paper (Section~\\ref{sec:res}), following a description of the software.  There are two types of GRB follow-up data telemetered from \\swift: TDRSS and Malindi data (described below). Initially, our software only worked with Malindi data, however in 2008 February we modified our ground-based software to work with TDRSS data as well (Evans \\etal2008a); this provides better positions, positions of fainter GRBs, and more reliable light curves and spectra than those produced on board (although these should not be used for scientific analysis).    \\subsection{TDRSS data}  When \\swift\\ first detects and observes a GRB, some data are immediately telemetered to the ground via NASA's Tracking and Data Relay Satellite System (TDRSS). Among these `TDRSS data' are the XRT Single Pixel Event Report (`SPER') data. \\swift's XRT selects between Windowed Timing (WT, high time resolution but only 1-D spatial information) and Photon Counting (PC, lower time resolution but full spatial information) modes automatically, based on the count rate in the central portion of the CCD (Hill \\etal2005). When the XRT enters PC mode during the first look at a new GRB, event lists containing every single-pixel (grade 0) event above 0.55 keV detected within the central 200$\\times$200 pixel region, the SPER data, are delivered to the XRT team via the GCN system every 2 minutes until \\swift\\ slews away from the burst (up to \\til2 ks after the trigger). It must be noted that while we have made our products as reliable as possible, SPER data are not fully calibrated and the light curves and spectra are intended as quick-look products; they should not be used for scientific analysis.  \\subsection{Malindi data}  `Malindi data' are available several hours after a GRB trigger and comprise the full observation dataset. Nine or ten times per day, \\swift\\ passes over the Malindi ground station in Kenya and downlinks the data buffered on board. These are then processed at the SDC and delivered to archives in the US, UK and Italy, typically 90--120 minutes after the Malindi downlink. \\swift\\ also observes GRBs which did not trigger the BAT, but are uploaded as Targets of Opportunity (ToOs). For these bursts only Malindi data are telemetered.  Malindi data are grouped into \\emph{observations}, usually one per day, each with a unique ObsID. A single observation may contain many \\emph{snapshots\\/}; that is, pointings toward the source, since \\swift\\ is in a low Earth orbit and thus its targets get occulted by the Earth once per orbit. The terminology \\emph{observations\\/} and \\emph{snapshots} is standard \\swift\\ parlance\\footnote{http://heasarc.gsfc.nasa.gov/docs/swift/archive/archiveguide1/node3.html}, and will be used throughout this paper.  \\subsection{GRB afterglow models} \\label{sec:fireball}  In this paper we introduce tools to produce high-precision positions, light curves and spectra from XRT data, and describe the automatic application of these tools to GRBs. This gives us a catalogue of results for all 318 GRBs detected by the \\swift-XRT up to GRB 080723B, and we discuss the implications for afterglow science from this analysis. While our dataset can be used to test any models for GRB emission, we do this in the conext of the fireball model (e.g.\\ Rees \\&\\ M\\'esz\\'aros 1994, Sari, Piran \\&\\ Narayan 1998), which is the current consensus model.  In this model the GRB progenitor launches highly relativistic jets of material in a series of shells, of differing bulk Lorentz factors. Internal collisions (i.e.\\ within the jet) between shells cause shocks which radiate the GRB `prompt emission'. The X-ray data presented in this paper may contain the tail of this prompt emission, however it is thought to arise predominantly from the afterglow. This is emission from an external shock which forms where the jet is decelerated by the circumburst medium, and which propogates into that medium, cooling by synchrotron radiation as it does so. See Piran (2005) for a comprehensive review of the fireball model.  \\subsection{Layout of the paper} \\label{sec:roadmap}  This paper is laid out thus: \\begin{itemize} \\item{\\S2 (pp \\pageref{sec:prods}--\\pageref{sec:ngrb}): XRT automatic analysis tools: \\begin{itemize} \\item{\\S2.1 (pp \\pageref{sec:spec}--\\pageref{sec:astrom}): Spectra.} \\item{\\S2.2 (pp \\pageref{sec:astrom}--\\pageref{sec:lc}): Positions.} \\item{\\S2.3 (pp \\pageref{sec:lc}--\\pageref{sec:ngrb}): Light curves \\&\\ hardness ratios.} \\end{itemize} } \\item{\\S3 (pp \\pageref{sec:ngrb}--\\pageref{sec:res}): Web tools to analyse any XRT source} \\item{\\S4 (pp \\pageref{sec:res}--\\pageref{sec:discuss}): Results of XRT GRB analysis} \\item{\\S5 (pp \\pageref{sec:discuss}--\\pageref{sec:interpret}): A canonical X-ray light curve?} \\item{\\S6 (pp \\pageref{sec:interpret}--\\pageref{sec:usage}): Understanding the X-ray afterglow} \\end{itemize}  Throughout this paper, errors quoted in all tables are at the 90\\%\\ confidence level. Error bars on data in all figures are 1-$\\sigma$ uncertainties. ",
        "conclusions": "We have developed software to automatically produce light curves and hardness ratios, spectra and high-precision enhanced XRT positions of GRBs. Preliminary versions of these are available within minutes of a trigger, and the full versions available within a few hours. Users can interact with and customise these products as desired. We also provide general purpose versions of these tools to run for any object observed by XRT, available via a web interface.  These products are available online:  \\begin{itemize} \\item{Index: http://www.swift.ac.uk/xrt\\_products} \\item{Positions: http://www.swift.ac.uk/xrt\\_positions} \\item{Light curves: http://www.swift.ac.uk/xrt\\_curves} \\item{Spectra: http://www.swift.ac.uk/xrt\\_spectra} \\end{itemize}  Using this software we have performed a homogeneous analysis of all GRBs observed by the XRT to date, and presented positions and temporal and spectral indices, in various formats. Analysis of these data show that a variety of light curve morphologies exist, and the so-called `canonical' curve, while the most common case, accounts for less than half of the light curves seen by \\swift. Defining a subsample of 162 GRBs with sufficient coverage to detect the canonical shape, if it existed, we found:  \\begin{itemize} \\item{8 (5 \\%) have no breaks.} \\item{49 (30 \\%) have one break (25 shallow, 24 steepen).} \\item{67 (41 \\%) are canonical.} \\item{38 (24 \\%) are oddballs.} \\end{itemize}  We have, however, demonstrated that this range of morphologies can be explained by a single underlying behaviour; the two-component model suggested by Willingale \\etal(2007), which involves a `prompt' component and an `afterglow' component. To achieve this we require a range of prompt-to-afterglow emission ratios, and a range of energy injection rates, both of which are easy to accept given the variations seen from burst to burst.  If the afterglow emission is due to the external forward shock model then in many cases this scenario  can only explain the data if energy injection continues beyond the plateau phase, and lasts for days to weeks after the GRB. The data also require a mechanism which can cause a light curve break (i.e.\\ the end of the plateau) without terminating energy injection, and without causing a change in the X-ray spectrum.  \\subsection{Usage policy} \\label{sec:usage}  Anybody is welcome to use the products and tools details in this paper for their work. Although we have verified these tools as far as possible, we still strongly advise users to `sanity check' their results, particularly with regards to light curve binning (Section~\\ref{sec:ngrblc}).  If these products or tools are used in any publication, we ask that this paper be cited, and that users include the following statement in the acknowledgements:  ``This work made use of data supplied by the UK Swift Science Data Centre at the University of Leicester.''    \\clearpage  \\begin{table} \\begin{center} \\begin{tabular}{ll} \\hline GRB  &  Times ignored from fit \\\\ & (s since trigger) \\\\ \\hline GRB 050406  &  115--430 \\\\ GRB 050410  &  45000--100000 \\\\ GRB 050502B  &  298--1680, 24800--120000 \\\\ GRB 050607  &  223--428 \\\\ GRB 050712  &  179--960 \\\\ GRB 050713A  &  96--152, 157--203 \\\\ GRB 050714B  &  237--486 \\\\ GRB 050724  &  220--350, 15831--103776 \\\\ GRB 050726  &  0--195, 200--320 \\\\ GRB 050730  &  191--791, 72464--119752 \\\\ GRB 050801  &  160--700, 3000--6000 \\\\ GRB 050814  &  1161--1627, 1813--2718 \\\\ GRB 050820A  &  171--260 \\\\ GRB 050822  &  115--193, 218--760 \\\\ GRB 050904  &  308--600, 1190--1378, 2900--55000 \\\\ GRB 050908  &  298--477 \\\\ GRB 050915A  &  131--180 \\\\ GRB 050916  &  6501--29640 \\\\ GRB 050922B  &  612--1517 \\\\ GRB 051117A  &  150--200, 300--700, 800--2000 \\\\ GRB 051227  &  101--121 \\\\ GRB 060111A  &  79--515 \\\\ GRB 060115  &  292--564 \\\\ GRB 060124  &  250--1200,  \\\\ GRB 060202  &  154--163, 35137--52020, 110145--127545 \\\\ GRB 060204B  &  104--138, 246--413 \\\\ GRB 060206  &  1776--11829 \\\\ GRB 060210  &  103--612 \\\\ GRB 060218  &  159--3000 \\\\ GRB 060223A  &  800--2000 \\\\ GRB 060312  &  66--190, 475--817 \\\\ GRB 060319  &  198--474, 740--1861 \\\\ GRB 060413  &  490--950 \\\\ GRB 060418  &  107--215 \\\\ GRB 060510B  &  732--1007 \\\\ GRB 060512  &  147--380 \\\\ GRB 060526  &  208--647 \\\\ GRB 060602B  &  120--300 \\\\ GRB 060604  &  118--237 \\\\ GRB 060607A  &  90--400 \\\\ GRB 060707  &  180--400 \\\\ GRB 060712  &  250--400 \\\\ GRB 060714  &  107--190 \\\\ GRB 060719  &  170--400 \\\\ GRB 060729  &  165--188 \\\\ GRB 060814  &  118--294 \\\\ GRB 060904A  &  160--397, 572--1120 \\\\ GRB 060904B  &  91--726 \\\\ GRB 060926  &  290--1000 \\\\ GRB 060929  &  280--1100 \\\\ GRB 061121  &  66--80 \\\\ GRB 061202  &  126--200 \\\\ GRB 070107  &  250--547, 1245--1661, 64099--116566 \\\\ GRB 070110  &  30000--80000 \\\\ GRB 070129  &  218--1139 \\\\ GRB 070318  &  182--429 \\\\ GRB 070330  &  142--374 \\\\ GRB 070419B  &  87--118, 180--1000 \\\\ GRB 070517  &  700--2000 \\\\ GRB 070518  &  89--227 \\\\ GRB 070520B  &  120--380 \\\\ \\end{tabular} \\caption{Times identified as deviating from power-law decays, and excluded from the light curve fits, see Section~\\ref{sec:flares} for details.} \\label{tab:flares} \\end{center} \\end{table}  \\begin{table} \\begin{center} \\begin{tabular}{ll} \\hline GRB  &  Times ignored from fit \\\\ & (s since trigger) \\\\ \\hline GRB 070616  &  480--550, 714--995 \\\\ GRB 070621  &  135--220 \\\\ GRB 070704  &  257--658 \\\\ GRB 070721B  &  124--141, 187--858 \\\\ GRB 070724A  &  90--150 \\\\ GRB 071021  &  148--300, 1000--9000 \\\\ GRB 071031  &  109--700 \\\\ GRB 071112C  &  523--726 \\\\ GRB 071118  &  189--225, 321--900 \\\\ GRB 071122  &  398--498 \\\\ GRB 080210  &  165--379 \\\\ GRB 080212  &  169--491 \\\\ GRB 080229A  &  100--156 \\\\ GRB 080310  &  130--980 \\\\ GRB 080319D  &  264--640 \\\\ GRB 080320  &  245--500, 650--1000 \\\\ GRB 080325  &  158--400 \\\\ GRB 080506  &  364--1200 \\\\ GRB 080604  &  18000--100000 \\\\ GRB 080607  &  116--210, 3000--7000 \\\\ \\end{tabular} \\contcaption{} \\end{center} \\end{table}   \\pagebreak  \\begin{table} \\begin{center} \\begin{tabular}{lccc} \\hline GRB  &  RA (J2000.0)  & Dec (J2000.0) & Error$^1$ \\\\ \\hline GRB050124  &  12 51 30.55  &  +13 02 39.6  &  1.6 \\\\ GRB050128  &  14 38 17.73  &  $-$34 45 55.2  &  1.7 \\\\ GRB050219A  &  11 05 39.01  &  $-$40 41 03.1  &  2.1 \\\\ GRB050219B  &  05 25 16.05  &  $-$57 45 29.3  &  1.5 \\\\ GRB050223  &  18 05 32.38  &  $-$62 28 21.9  &  4.3 \\\\ GRB050315  &  20 25 54.20  &  $-$42 36 01.5  &  1.5 \\\\ GRB050318  &  03 18 50.99  &  $-$46 23 45.0  &  1.4 \\\\ GRB050319  &  10 16 47.88  &  +43 32 55.3  &  1.4 \\\\ GRB050326  &  00 27 48.50  &  $-$71 22 18.5  &  2.2 \\\\ GRB050406  &  02 17 52.12  &  $-$50 11 16.4  &  1.9 \\\\ GRB050408  &  12 02 17.27  &  +10 51 09.7  &  1.4 \\\\ GRB050410  &  05 59 13.39  &  +79 36 11.8  &  2.0 \\\\ GRB050416A  &  12 33 54.60  &  +21 03 27.2  &  1.4 \\\\ GRB050421  &  20 29 02.79  &  +73 39 17.9  &  1.7 \\\\ GRB050502B  &  09 30 10.08  &  +16 59 48.0  &  1.4 \\\\ GRB050504  &  13 24 01.31  &  +40 42 15.0  &  2.8 \\\\ GRB050505  &  09 27 03.31  &  +30 16 24.3  &  1.4 \\\\ GRB050509A  &  20 42 20.00  &  +54 04 17.2  &  2.1 \\\\ GRB050509B  &  12 36 13.75  &  +28 59 03.4  &  3.3 \\\\ GRB050525A  &  18 32 32.68  &  +26 20 22.5  &  1.5 \\\\ GRB050603  &  02 39 56.93  &  $-$25 10 54.7  &  1.4 \\\\ GRB050701  &  15 09 01.55  &  $-$59 24 52.5  &  3.0 \\\\ GRB050712  &  05 10 48.19  &  +64 54 48.2  &  1.5 \\\\ GRB050713A  &  21 22 09.36  &  +77 04 28.9  &  1.4 \\\\ GRB050713B  &  20 31 15.54  &  +60 56 43.8  &  1.4 \\\\ GRB050714B  &  11 18 47.63  &  $-$15 32 48.9  &  1.5 \\\\ GRB050716  &  22 34 20.79  &  +38 41 04.2  &  1.4 \\\\ GRB050717  &  14 17 24.45  &  $-$50 32 00.4  &  1.5 \\\\ GRB050721  &  16 53 44.54  &  $-$28 22 52.0  &  1.7 \\\\ GRB050724  &  16 24 44.29  &  $-$27 32 26.9  &  1.5 \\\\ GRB050726  &  13 20 11.91  &  $-$32 03 51.3  &  1.4 \\\\ GRB050730  &  14 08 17.14  &  $-$03 46 18.4  &  1.4 \\\\ GRB050802  &  14 37 05.81  &  +27 47 11.5  &  1.4 \\\\ GRB050803  &  23 22 37.91  &  +05 47 09.1  &  1.4 \\\\ GRB050813  &  16 07 56.97  &  +11 14 57.8  &  2.9 \\\\ GRB050814  &  17 36 45.29  &  +46 20 22.2  &  1.4 \\\\ GRB050815  &  19 34 22.94  &  +09 08 48.5  &  1.5 \\\\ GRB050819  &  23 55 01.65  &  +24 51 38.8  &  1.5 \\\\ GRB050820A  &  22 29 38.13  &  +19 33 37.5  &  1.4 \\\\ GRB050822  &  03 24 27.24  &  $-$46 01 59.6  &  1.4 \\\\ GRB050824  &  00 48 56.25  &  +22 36 33.3  &  1.5 \\\\ GRB050826  &  05 51 01.62  &  $-$02 38 36.6  &  1.5 \\\\ GRB050827  &  04 17 09.62  &  +18 12 01.1  &  1.5 \\\\ GRB050904  &  00 54 50.88  &  +14 05 10.4  &  1.4 \\\\ GRB050908  &  01 21 50.79  &  $-$12 57 18.5  &  1.5 \\\\ GRB050915A  &  05 26 44.87  &  $-$28 00 58.8  &  1.4 \\\\ GRB050915B  &  14 36 25.90  &  $-$67 24 33.4  &  1.6 \\\\ GRB050916  &  09 03 57.26  &  $-$51 25 42.7  &  1.5 \\\\ GRB050922B  &  00 23 13.40  &  $-$05 36 17.9  &  1.7 \\\\ GRB050922C  &  21 09 33.02  &  $-$08 45 30.6  &  1.4 \\\\ GRB051001  &  23 23 48.79  &  $-$31 31 21.9  &  1.7 \\\\ GRB051006  &  07 23 14.12  &  +09 30 20.0  &  1.5 \\\\ GRB051008  &  13 31 29.59  &  +42 05 53.0  &  1.4 \\\\ GRB051016A  &  08 11 16.68  &  $-$18 17 53.1  &  1.5 \\\\ GRB051016B  &  08 48 27.87  &  +13 39 19.7  &  1.4 \\\\ GRB051021A  &  01 56 36.28  &  +09 04 01.6  &  1.5 \\\\ GRB051021B  &  08 24 12.24  &  $-$45 32 28.5  &  1.6 \\\\ GRB051022  &  23 56 04.06  &  +19 36 23.9  &  1.4 \\\\ GRB051028  &  01 48 15.02  &  +47 45 11.4  &  1.7 \\\\ GRB051109A  &  22 01 15.25  &  +40 49 22.6  &  1.4 \\\\ GRB051109B  &  23 01 50.33  &  +38 40 46.5  &  1.5 \\\\ GRB051111  &  23 12 33.08  &  +18 22 27.9  &  1.5 \\\\ GRB051117A  &  15 13 34.09  &  +30 52 11.8  &  1.6 \\\\ \\hline \\end{tabular} \\caption{All available enhanced XRT positions of GRBs observed by \\swift. The positions are also available online at http://www.swift.ac.uk/xrt\\_positions, which is updated automatically with new GRBs. \\newline $^1$arc seconds, 90\\%\\ confidence radius. } \\label{tab:enhpos} \\end{center} \\end{table}  \\begin{table} \\begin{center} \\begin{tabular}{lccc} \\hline GRB  &  RA (J2000.0)  & Dec (J2000.0) & Error$^1$ \\\\ \\hline GRB051117B  &  05 40 43.39  &  $-$19 16 26.9  &  1.7 \\\\ GRB051210  &  22 00 41.33  &  $-$57 36 49.4  &  1.7 \\\\ GRB051211B  &  23 02 41.53  &  +55 04 50.7  &  1.6 \\\\ GRB051221A  &  21 54 48.58  &  +16 53 26.0  &  1.4 \\\\ GRB051221B  &  20 49 35.25  &  +53 02 10.7  &  2.1 \\\\ GRB051227  &  08 20 58.31  &  +31 55 31.0  &  1.6 \\\\ GRB060105  &  19 50 00.68  &  +46 20 55.4  &  1.4 \\\\ GRB060108  &  09 48 02.13  &  +31 55 07.2  &  1.5 \\\\ GRB060109  &  18 50 43.65  &  +31 59 26.5  &  1.4 \\\\ GRB060111A  &  18 24 49.12  &  +37 36 14.0  &  1.5 \\\\ GRB060111B  &  19 05 42.64  &  +70 22 33.0  &  1.4 \\\\ GRB060115  &  03 36 08.29  &  +17 20 42.7  &  1.4 \\\\ GRB060116  &  05 38 46.24  &  $-$05 26 13.9  &  1.5 \\\\ GRB060121  &  09 09 52.09  &  +45 39 47.2  &  1.6 \\\\ GRB060124  &  05 08 25.86  &  +69 44 26.5  &  1.4 \\\\ GRB060202  &  02 23 22.97  &  +38 23 03.1  &  1.4 \\\\ GRB060203  &  06 54 03.72  &  +71 48 38.9  &  1.4 \\\\ GRB060204B  &  14 07 15.04  &  +27 40 37.0  &  1.4 \\\\ GRB060206  &  13 31 43.44  &  +35 03 02.9  &  1.6 \\\\ GRB060210  &  03 50 57.33  &  +27 01 33.7  &  1.4 \\\\ GRB060211A  &  03 53 32.59  &  +21 29 19.1  &  1.5 \\\\ GRB060211B  &  05 00 17.11  &  +14 56 57.3  &  1.6 \\\\ GRB060218  &  03 21 39.65  &  +16 52 01.3  &  1.4 \\\\ GRB060219  &  16 07 21.49  &  +32 18 57.1  &  1.5 \\\\ GRB060223A  &  03 40 49.57  &  $-$17 07 49.9  &  1.5 \\\\ GRB060306  &  02 44 22.83  &  $-$02 08 55.1  &  1.4 \\\\ GRB060313  &  04 26 28.43  &  $-$10 50 41.5  &  1.4 \\\\ GRB060319  &  11 45 33.05  &  +60 00 39.6  &  1.4 \\\\ GRB060323  &  11 37 45.13  &  +49 59 08.3  &  1.6 \\\\ GRB060403  &  18 49 21.60  &  +08 19 46.7  &  2.3 \\\\ GRB060413  &  19 25 07.85  &  +13 45 29.9  &  1.4 \\\\ GRB060418  &  15 45 42.66  &  $-$03 38 20.4  &  1.4 \\\\ GRB060421  &  22 54 33.07  &  +62 43 51.4  &  1.7 \\\\ GRB060427  &  08 17 04.03  &  +62 40 16.8  &  1.4 \\\\ GRB060428A  &  08 14 10.81  &  $-$37 10 11.7  &  1.4 \\\\ GRB060428B  &  15 41 25.70  &  +62 01 30.1  &  1.4 \\\\ GRB060502A  &  16 03 42.50  &  +66 36 02.6  &  1.5 \\\\ GRB060502B  &  18 35 45.25  &  +52 37 53.9  &  5.2 \\\\ GRB060505  &  22 07 03.38  &  $-$27 48 52.9  &  1.9 \\\\ GRB060507  &  05 59 50.67  &  +75 14 54.8  &  1.5 \\\\ GRB060510B  &  15 56 29.25  &  +78 34 12.1  &  1.5 \\\\ GRB060512  &  13 03 05.73  &  +41 11 26.6  &  1.5 \\\\ GRB060515  &  08 29 09.61  &  +73 34 04.8  &  3.8 \\\\ GRB060522  &  21 31 44.89  &  +02 53 10.3  &  1.5 \\\\ GRB060526  &  15 31 18.34  &  +00 17 05.9  &  1.4 \\\\ GRB060602B  &  17 49 31.83  &  $-$28 08 04.5  &  1.6 \\\\ GRB060604  &  22 28 55.10  &  $-$10 54 55.7  &  1.4 \\\\ GRB060605  &  21 28 37.35  &  $-$06 03 30.3  &  1.4 \\\\ GRB060607A  &  21 58 50.41  &  $-$22 29 47.1  &  1.4 \\\\ GRB060614  &  21 23 32.19  &  $-$53 01 36.5  &  1.4 \\\\ GRB060707  &  23 48 19.11  &  $-$17 54 17.6  &  1.7 \\\\ GRB060708  &  00 31 13.83  &  $-$33 45 32.2  &  1.4 \\\\ GRB060712  &  12 16 16.16  &  +35 32 17.6  &  1.6 \\\\ GRB060714  &  15 11 26.39  &  $-$06 33 59.0  &  1.4 \\\\ GRB060717  &  11 23 21.61  &  +28 57 03.0  &  3.2 \\\\ GRB060719  &  01 13 43.76  &  $-$48 22 50.1  &  1.4 \\\\ GRB060729  &  06 21 31.84  &  $-$62 22 12.1  &  1.4 \\\\ GRB060801  &  14 12 01.30  &  +16 58 54.0  &  1.5 \\\\ GRB060804  &  07 28 49.32  &  $-$27 12 55.0  &  1.5 \\\\ GRB060805A  &  14 43 43.45  &  +12 35 11.6  &  1.6 \\\\ GRB060805B  &  18 10 05.41  &  +58 09 19.0  &  1.7 \\\\ GRB060807  &  16 50 02.63  &  +31 35 30.2  &  1.4 \\\\ GRB060813  &  07 27 35.38  &  $-$29 50 49.3  &  1.4 \\\\ \\hline \\end{tabular} \\contcaption{$^1$arc seconds, 90\\%\\ confidence radius. Also available online at http://www.swift.ac.uk/xrt\\_positions, updated automatically with new GRBs.} \\end{center} \\end{table}   \\begin{table} \\begin{center} \\begin{tabular}{lccc} \\hline GRB  &  RA (J2000.0)  & Dec (J2000.0) & Error$^1$ \\\\ \\hline GRB060814  &  14 45 21.32  &  +20 35 09.2  &  1.4 \\\\ GRB060825  &  01 12 29.43  &  +55 47 51.8  &  2.0 \\\\ GRB060901  &  19 08 37.81  &  $-$06 38 05.7  &  2.9 \\\\ GRB060904B  &  03 52 50.56  &  $-$00 43 30.1  &  1.4 \\\\ GRB060906  &  02 43 00.90  &  +30 21 43.0  &  1.4 \\\\ GRB060908  &  02 07 18.42  &  +00 20 32.4  &  1.4 \\\\ GRB060912A  &  00 21 08.11  &  +20 58 19.2  &  1.4 \\\\ GRB060919  &  18 27 41.74  &  $-$51 00 52.5  &  1.7 \\\\ GRB060923A  &  16 58 28.10  &  +12 21 37.9  &  1.5 \\\\ GRB060926  &  17 35 43.70  &  +13 02 16.6  &  1.6 \\\\ GRB060927  &  21 58 12.02  &  +05 21 48.6  &  1.6 \\\\ GRB060929  &  17 32 28.92  &  +29 50 06.2  &  1.5 \\\\ GRB061002  &  14 41 23.51  &  +48 44 28.6  &  2.2 \\\\ GRB061004  &  06 31 10.91  &  $-$45 54 24.3  &  1.4 \\\\ GRB061006  &  07 24 07.83  &  $-$79 11 55.3  &  1.9 \\\\ GRB061007  &  03 05 19.70  &  $-$50 30 02.8  &  1.4 \\\\ GRB061019  &  06 06 31.01  &  +29 34 12.3  &  1.5 \\\\ GRB061021  &  09 40 36.17  &  $-$21 57 05.1  &  1.4 \\\\ GRB061025  &  20 03 36.87  &  $-$48 14 38.3  &  2.3 \\\\ GRB061028  &  06 28 54.69  &  +46 17 57.5  &  1.6 \\\\ GRB061110A  &  22 25 09.85  &  $-$02 15 31.1  &  1.6 \\\\ GRB061110B  &  21 35 40.43  &  +06 52 31.9  &  1.7 \\\\ GRB061121  &  09 48 54.61  &  $-$13 11 42.0  &  1.4 \\\\ GRB061122  &  20 15 19.83  &  +15 31 00.8  &  1.4 \\\\ GRB061126  &  05 46 24.55  &  +64 12 39.8  &  1.4 \\\\ GRB061201  &  22 08 32.22  &  $-$74 34 48.2  &  1.4 \\\\ GRB061202  &  07 02 06.20  &  $-$74 41 54.9  &  1.4 \\\\ GRB061217  &  10 41 37.10  &  $-$21 07 33.8  &  3.4 \\\\ GRB061222A  &  23 53 03.37  &  +46 31 57.5  &  1.4 \\\\ GRB061222B  &  07 01 24.70  &  $-$25 51 35.3  &  1.7 \\\\ GRB070103  &  23 30 13.75  &  +26 52 33.8  &  1.4 \\\\ GRB070107  &  10 37 36.56  &  $-$53 12 47.6  &  1.4 \\\\ GRB070110  &  00 03 39.27  &  $-$52 58 28.5  &  1.4 \\\\ GRB070125  &  07 51 17.80  &  +31 09 04.8  &  1.5 \\\\ GRB070129  &  02 28 00.93  &  +11 41 03.3  &  1.4 \\\\ GRB070208  &  13 11 32.77  &  +61 57 54.9  &  1.5 \\\\ GRB070219  &  17 20 46.19  &  +69 22 12.6  &  1.7 \\\\ GRB070220  &  02 19 06.82  &  +68 48 16.5  &  1.4 \\\\ GRB070223  &  10 13 48.37  &  +43 08 01.9  &  1.8 \\\\ GRB070224  &  11 56 06.61  &  $-$13 19 49.8  &  1.6 \\\\ GRB070227  &  08 02 19.33  &  $-$46 18 51.3  &  1.9 \\\\ GRB070306  &  09 52 23.22  &  +10 28 55.3  &  1.4 \\\\ GRB070309  &  17 34 39.86  &  $-$37 55 50.3  &  4.4 \\\\ GRB070311  &  05 50 08.22  &  +03 22 29.4  &  1.5 \\\\ GRB070318  &  03 13 56.72  &  $-$42 56 47.0  &  1.4 \\\\ GRB070328  &  04 20 27.68  &  $-$34 04 00.6  &  1.4 \\\\ GRB070330  &  17 58 10.24  &  $-$63 47 34.7  &  1.5 \\\\ GRB070411  &  07 09 19.90  &  +01 03 52.4  &  1.5 \\\\ GRB070412  &  12 06 10.03  &  +40 08 35.8  &  1.5 \\\\ GRB070419A  &  12 10 58.83  &  +39 55 31.1  &  1.7 \\\\ GRB070419B  &  21 02 49.84  &  $-$31 15 48.1  &  1.5 \\\\ GRB070420  &  08 04 55.12  &  $-$45 33 21.2  &  1.4 \\\\ GRB070429A  &  19 50 48.94  &  $-$32 24 17.4  &  2.1 \\\\ GRB070429B  &  21 52 03.85  &  $-$38 49 40.8  &  2.3 \\\\ GRB070506  &  23 08 52.30  &  +10 43 21.2  &  1.8 \\\\ GRB070508  &  20 51 12.00  &  $-$78 23 04.7  &  1.4 \\\\ GRB070509  &  15 51 50.40  &  $-$78 39 07.5  &  1.9 \\\\ GRB070518  &  16 56 47.68  &  +55 17 42.6  &  1.6 \\\\ GRB070520A  &  12 53 26.05  &  +74 59 23.3  &  3.4 \\\\ GRB070520B  &  08 07 31.08  &  +57 36 30.0  &  1.8 \\\\ GRB070521  &  16 10 38.60  &  +30 15 22.9  &  1.4 \\\\ GRB070529  &  18 54 58.30  &  +20 39 34.3  &  1.4 \\\\ GRB070531  &  00 26 58.58  &  +74 18 46.6  &  1.8 \\\\ \\hline \\end{tabular} \\contcaption{$^1$arc seconds, 90\\%\\ confidence radius. Also available online at http://www.swift.ac.uk/xrt\\_positions, updated automatically with new GRBs.} \\end{center} \\end{table}   \\begin{table} \\begin{center} \\begin{tabular}{lccc} \\hline GRB  &  RA (J2000.0)  & Dec (J2000.0) & Error$^1$ \\\\ \\hline GRB070611  &  00 07 58.15  &  $-$29 45 20.4  &  1.8 \\\\ GRB070612B  &  17 26 54.46  &  $-$08 45 04.7  &  1.8 \\\\ GRB070616  &  02 08 36.40  &  +56 56 45.0  &  1.4 \\\\ GRB070621  &  21 35 10.17  &  $-$24 49 02.6  &  1.4 \\\\ GRB070628  &  07 41 06.06  &  $-$20 16 45.3  &  1.4 \\\\ GRB070704  &  23 38 47.31  &  +66 15 12.7  &  1.7 \\\\ GRB070714A  &  02 51 43.36  &  +30 14 35.1  &  1.8 \\\\ GRB070714B  &  03 51 22.30  &  +28 17 51.3  &  1.4 \\\\ GRB070721A  &  00 12 39.20  &  $-$28 33 00.0  &  1.6 \\\\ GRB070721B  &  02 12 32.98  &  $-$02 11 39.8  &  1.5 \\\\ GRB070724A  &  01 51 14.08  &  $-$18 35 38.8  &  1.8 \\\\ GRB070724B  &  01 09 56.68  &  +57 40 37.6  &  2.0 \\\\ GRB070729  &  03 45 16.04  &  $-$39 19 19.9  &  2.5 \\\\ GRB070802  &  02 27 35.91  &  $-$55 31 39.1  &  1.5 \\\\ GRB070808  &  00 27 03.40  &  +01 10 35.4  &  1.5 \\\\ GRB070810A  &  12 39 51.18  &  +10 45 02.8  &  1.4 \\\\ GRB071025  &  23 40 17.02  &  +31 46 42.4  &  1.4 \\\\ GRB071028A  &  07 59 16.66  &  +21 29 05.0  &  1.5 \\\\ GRB071031  &  00 25 37.40  &  $-$58 03 33.9  &  1.5 \\\\ GRB071104  &  19 42 26.78  &  +14 36 35.8  &  1.5 \\\\ GRB071112C  &  02 36 50.89  &  +28 22 16.9  &  1.4 \\\\ GRB071117  &  22 20 10.40  &  $-$63 26 36.6  &  1.5 \\\\ GRB071118  &  19 58 51.77  &  +70 07 27.9  &  1.4 \\\\ GRB071122  &  18 26 25.26  &  +47 04 29.5  &  1.8 \\\\ GRB071227  &  03 52 31.09  &  $-$55 59 03.3  &  1.6 \\\\ GRB080120  &  15 01 03.29  &  $-$10 52 30.5  &  1.8 \\\\ GRB080123  &  22 35 46.16  &  $-$64 54 02.7  &  1.7 \\\\ GRB080129  &  07 01 08.13  &  $-$07 50 47.8  &  1.4 \\\\ GRB080130  &  17 26 27.27  &  $-$53 11 14.4  &  4.3 \\\\ GRB080205  &  06 33 00.43  &  +62 47 31.2  &  1.5 \\\\ GRB080207  &  13 50 03.02  &  +07 30 08.0  &  1.4 \\\\ GRB080210  &  16 45 04.03  &  +13 49 36.3  &  1.5 \\\\ GRB080212  &  15 24 35.44  &  $-$22 44 30.5  &  1.4 \\\\ GRB080218B  &  11 51 49.66  &  $-$53 05 49.0  &  1.8 \\\\ GRB080229A  &  15 12 52.31  &  $-$14 42 16.2  &  1.4 \\\\ GRB080303  &  07 28 14.18  &  $-$70 14 03.2  &  1.5 \\\\ GRB080307  &  09 06 30.73  &  +35 08 19.8  &  1.4 \\\\ GRB080310  &  14 40 13.89  &  $-$00 10 30.4  &  1.4 \\\\ GRB080319A  &  13 45 19.97  &  +44 04 49.1  &  1.7 \\\\ GRB080319B  &  14 31 41.04  &  +36 18 09.2  &  1.4 \\\\ GRB080319C  &  17 15 55.48  &  +55 23 31.7  &  1.7 \\\\ GRB080319D  &  06 37 53.47  &  +23 56 33.7  &  1.5 \\\\ GRB080320  &  11 50 56.37  &  +57 09 25.3  &  1.4 \\\\ GRB080325  &  18 31 34.40  &  +36 31 26.1  &  1.7 \\\\ GRB080328  &  05 21 55.81  &  +47 30 38.0  &  1.4 \\\\ GRB080409  &  05 37 19.20  &  +05 05 05.1  &  1.6 \\\\ GRB080411  &  02 31 54.82  &  $-$71 18 07.5  &  1.4 \\\\ GRB080413A  &  19 09 11.73  &  $-$27 40 39.6  &  1.5 \\\\ GRB080413B  &  21 44 34.62  &  $-$19 58 51.2  &  1.4 \\\\ GRB080426  &  01 45 59.79  &  +69 28 06.0  &  1.5 \\\\ GRB080430  &  11 01 14.63  &  +51 41 07.8  &  1.4 \\\\ GRB080503  &  19 06 28.74  &  +68 47 35.8  &  1.6 \\\\ GRB080506  &  21 57 41.79  &  +38 59 06.7  &  1.4 \\\\ GRB080507  &  15 34 43.41  &  +56 26 08.4  &  2.3 \\\\ GRB080514B  &  21 31 22.61  &  +00 42 30.3  &  1.7 \\\\ GRB080516  &  08 02 34.16  &  $-$26 09 35.3  &  1.5 \\\\ GRB080517  &  06 48 58.00  &  +50 44 06.3  &  1.6 \\\\ GRB080520  &  18 40 46.50  &  $-$54 59 30.4  &  1.6 \\\\ GRB080523  &  01 23 11.53  &  $-$64 01 50.9  &  1.5 \\\\ GRB080602  &  01 16 42.18  &  $-$09 13 55.4  &  1.7 \\\\ GRB080603A  &  18 37 37.95  &  +62 44 39.6  &  1.6 \\\\ GRB080603B  &  11 46 07.46  &  +68 03 38.5  &  1.7 \\\\ GRB080604  &  15 47 51.64  &  +20 33 28.5  &  1.5 \\\\ \\hline \\end{tabular} \\contcaption{$^1$arc seconds, 90\\%\\ confidence radius. Also available online at http://www.swift.ac.uk/xrt\\_positions, updated automatically with new GRBs.} \\end{center} \\end{table}   \\begin{table} \\begin{center} \\begin{tabular}{lccc} \\hline GRB  &  RA (J2000.0)  & Dec (J2000.0) & Error$^1$ \\\\ \\hline GRB080605  &  17 28 30.04  &  +04 00 56.9  &  1.5 \\\\ GRB080607  &  12 59 47.14  &  +15 55 09.6  &  1.4 \\\\ GRB080613A  &  14 13 04.98  &  +05 10 21.8  &  2.0 \\\\ GRB080613B  &  11 35 11.51  &  $-$07 06 17.5  &  2.5 \\\\ GRB080623  &  15 50 38.65  &  $-$62 02 57.3  &  1.5 \\\\ GRB080625  &  19 53 37.09  &  +56 16 41.7  &  1.7 \\\\ GRB080701  &  03 03 21.31  &  +75 28 29.3  &  1.7 \\\\ GRB080702A  &  20 52 12.21  &  +72 18 46.7  &  1.8 \\\\ GRB080702B  &  23 42 32.07  &  $-$05 31 51.7  &  2.2 \\\\ GRB080703  &  06 47 12.43  &  $-$63 13 08.3  &  1.4 \\\\ GRB080707  &  02 10 28.48  &  +33 06 35.1  &  1.5 \\\\ GRB080710  &  00 33 05.68  &  +19 30 05.8  &  1.4 \\\\ GRB080714  &  12 32 25.29  &  $-$60 16 40.2  &  1.5 \\\\ GRB080721  &  14 57 55.76  &  $-$11 43 24.9  &  1.4 \\\\ GRB080723A  &  18 17 57.70  &  $-$06 53 58.3  &  1.4 \\\\ GRB080723B  &  11 47 19.77  &  $-$60 14 27.9  &  1.5 \\\\ \\hline \\end{tabular} \\contcaption{$^1$arc seconds, 90\\%\\ confidence radius. Also available online at http://www.swift.ac.uk/xrt\\_positions, updated automatically with new GRBs.} \\end{center} \\end{table}   \\clearpage     \\begin{table*} \\begin{center} \\begin{tabular}{lccccccccc} \\hline GRB  & Target ID & $T_{90}$  &  F$_{11}$  & $\\alpha_1$ & $\\alpha_2$ & $\\alpha_3$ & $\\alpha_4$ & $\\alpha_5$ & $\\alpha_6$ \\\\ \\hline GRB 041223    &    00100585   &    109.1    &    &   $1.90^{+0.40}_{-0.39}$ \\\\ GRB 050124    &    00103647   &    4.0    &  1.12e-12  &   $1.50^{+0.15}_{-0.15}$ \\\\ GRB 050126    &    00103780   &    24.8    &  1.22e-13  &   $2.57^{+0.99}_{-0.46}$    &    $0.94^{+0.24}_{-0.31}$ \\\\ GRB 050128    &    00103906   &    19.2    &  2.25e-12  &   $0.955^{+0.038}_{-0.108}$    &    $1.361^{+0.075}_{-0.059}$ \\\\ GRB 050215B    &    00106107   &    8.1    &  3.10e-13  &   $0.92^{+0.15}_{-0.13}$ \\\\ GRB 050219A    &    00106415   &    23.7    &  8.73e-13  &   $2.91^{+0.25}_{-0.25}$    &    $0.754^{+0.089}_{-0.083}$ \\\\ GRB 050219B    &    00106442   &    30.7    &  3.53e-12  &   $1.254^{+0.038}_{-0.037}$ \\\\ GRB 050223    &    00106709   &    22.5    &  1.17e-13  &   $0.89^{+0.27}_{-0.22}$ \\\\ GRB 050315    &    00111063   &    95.6    &  4.24e-12  &   $3.51^{+0.24}_{-0.21}$    &    $0.257^{+0.034}_{-0.038}$    &    $1.216^{+0.064}_{-0.063}$ \\\\ GRB 050318    &    00111529   &    32    &  8.25e-13  &   $1.21^{+0.13}_{-0.17}$    &    $1.86^{+0.15}_{-0.13}$ \\\\ GRB 050319    &    00111622   &    152.5    &  3.26e-12  &   $4.77^{+0.24}_{-1.05}$    &    $0.493^{+0.048}_{-0.049}$    &    $0.96^{+0.20}_{-0.12}$    &    $3^{+26}_{-1}$ \\\\ GRB 050326    &    00112453   &    29.3    &  8.48e-13  &   $1.657^{+0.051}_{-0.049}$ \\\\ GRB 050401    &    00113120   &    33.3    &  3.52e-12  &   $0.569^{+0.022}_{-0.022}$    &    $1.428^{+0.095}_{-0.077}$ \\\\ GRB 050406    &    00113872   &    5.4    &  8.90e-14  &   $0.386^{+0.062}_{-0.266}$    &    $1^{+9}_{-2}$ \\\\ GRB 050408    &    00020004   &    N/A    &  1.40e-12  &   $0.960^{+0.026}_{-0.024}$ \\\\ GRB 050410    &    00114299   &    42.5    &  6.78e-13  &   $0.68^{+0.14}_{-0.10}$    &    $1.28^{+1.11}_{-0.28}$ \\\\ GRB 050412    &    00114485   &    26.5    &  4.44e-14  &   $1.39^{+0.25}_{-0.23}$ \\\\ GRB 050416A    &    00114753   &    2.5    &  8.27e-13  &   $1.64^{+1.18}_{-0.65}$    &    $0.29^{+0.26}_{-0.70}$    &    $0.859^{+0.023}_{-0.023}$ \\\\ GRB 050421    &    00115135   &    15.0    &  4.28e-18  &   $0.69^{+0.15}_{-1.10}$    &    $3.33^{+0.24}_{-0.21}$ \\\\ GRB 050422    &    00115214   &    59.3    &  5.48e-14  &   $4.84^{+1.35}_{-0.97}$    &    $0.91^{+0.13}_{-0.23}$ \\\\ GRB 050502B    &    00116116   &    17.7    &  3.99e-13  &   $0.743^{+0.049}_{-0.056}$    &    $1.74^{+2.77}_{-0.58}$ \\\\ GRB 050505    &    00117504   &    58.9    &  2.70e-12  &   $0.32^{+0.22}_{-0.10}$    &    $1.189^{+0.074}_{-0.069}$    &    $1.734^{+0.154}_{-0.096}$ \\\\ GRB 050509A    &    00118707   &    11.4    &  3.38e-13  &   $1.143^{+0.102}_{-0.092}$ \\\\ GRB 050520    &    00020010   &    N/A    &  4.35e-14  &   $1.73^{+0.60}_{-0.47}$ \\\\ GRB 050525A    &    00130088   &    8.8    &  1.18e-12  &   $0.646^{-0.724}_{-0.030}$    &    $1.586^{-0.040}_{-0.042}$ \\\\ GRB 050603    &    00131560   &    12.4    &  2.19e-12  &   $1.659^{+0.092}_{-0.087}$ \\\\ GRB 050607    &    00132247   &    26.4    &  1.98e-13  &   $-0.37^{+0.61}_{-0.91}$    &    $5^{+2}_{-1}$    &    $0.59^{+0.13}_{-0.17}$    &    $1.17^{+0.21}_{-0.13}$ \\\\ GRB 050701    &    00143708   &    21.8    &  3.10e-13  &   $1.22^{+0.19}_{-0.18}$ \\\\ GRB 050712    &    00145581   &    51.6    &  6.72e-13  &   $0.833^{+0.070}_{-0.082}$    &    $1.25^{+0.21}_{-0.17}$ \\\\ GRB 050713A    &    00145675   &    124.7    &  1.77e-12  &   $8.43^{+0.51}_{-2.78}$    &    $2.37^{+0.11}_{-0.19}$    &    $-0.22^{+0.12}_{-0.26}$    &    $1.157^{+0.029}_{-0.028}$ \\\\ GRB 050713B    &    00145754   &    54.2    &  5.47e-12  &   $3.08^{+0.13}_{-0.13}$    &    $-0.11^{+0.11}_{-0.13}$    &    $1.038^{+0.038}_{-0.038}$ \\\\ GRB 050714B    &    00145994   &    46.7    &  1.62e-13  &   $6.91^{+0.89}_{-0.86}$    &    $0.573^{+0.080}_{-0.079}$ \\\\ GRB 050716    &    00146227   &    69.1    &  5.34e-13  &   $0.21^{+0.22}_{-0.22}$    &    $2.116^{+0.071}_{-0.069}$    &    $0.974^{+0.060}_{-0.055}$ \\\\ GRB 050717    &    00146372   &    85.0    &  6.56e-14  &   $1.95^{+0.19}_{-0.10}$    &    $1.316^{+0.064}_{-0.054}$    &    $3^{+16}_{-1}$ \\\\ GRB 050721    &    00146970   &    98.4    &  5.18e-13  &   $1.592^{+0.094}_{-0.088}$    &    $0.91^{+0.16}_{-0.15}$ \\\\ GRB 050724    &    00147478   &    96.0    &  1.56e-13  &   $1.702^{+0.065}_{-0.069}$    &    $6.24^{+0.40}_{-0.32}$    &    $0.82^{+0.16}_{-0.12}$ \\\\ GRB 050726    &    00147788   &    49.9    &  3.52e-13  &   $0.982^{+0.041}_{-0.043}$    &    $1.83^{+0.19}_{-0.15}$ \\\\ GRB 050730    &    00148225   &    156.5    &  3.60e-12  &   $0.512^{+0.016}_{-0.016}$    &    $2.570^{+0.051}_{-0.050}$ \\\\ GRB 050801    &    00148522   &    19.4    &  1.57e-13  &   $1.022^{+0.086}_{-0.113}$    &    $1.27^{+1.36}_{-0.18}$ \\\\ GRB 050802    &    00148646   &    19.0    &  1.41e-12  &   $0.700^{+0.063}_{-0.061}$    &    $1.593^{+0.088}_{-0.049}$ \\\\ GRB 050803    &    00148833   &    87.9    &  2.35e-12  &   $5.35^{+0.65}_{-0.58}$    &    $0.515^{+0.044}_{-0.045}$    &    $1.677^{+0.080}_{-0.075}$ \\\\ GRB 050814    &    00150314   &    150.9    &  9.31e-13  &   $3.343^{+0.080}_{-0.069}$    &    $0.527^{+0.070}_{-0.069}$    &    $1.80^{+0.57}_{-0.26}$ \\\\ GRB 050815    &    00150532   &    2.9    &  7.34e-14  &   $0.15^{+0.12}_{-0.48}$    &    $1.92^{+0.34}_{-0.28}$ \\\\ GRB 050819    &    00151131   &    37.7    &  1.49e-13  &   $3.97^{+0.53}_{-0.40}$    &    $0.752^{+0.114}_{-0.100}$ \\\\ GRB 050820A    &    00151207   &    26    &  9.40e-12  &   $2.00^{+0.36}_{-0.32}$    &    $0.215^{+0.051}_{-0.243}$    &    $1.265^{+0.018}_{-0.017}$ \\\\ GRB 050822    &    00151486   &    103.4    &  1.65e-12  &   $3.00^{+0.34}_{-0.23}$    &    $0.206^{+0.076}_{-0.111}$    &    $1.052^{+0.042}_{-0.040}$ \\\\ GRB 050824    &    00151905   &    22.6    &  8.85e-13  &   $0.342^{+0.106}_{-0.079}$    &    $0.879^{+0.097}_{-0.070}$ \\\\ GRB 050826    &    00152113   &    35.5    &  3.06e-13  &   $1.30^{+0.19}_{-0.22}$    &    $0.75^{+0.17}_{-0.20}$    &    $1.60^{+2.18}_{-0.45}$ \\\\ GRB 050827    &    00152325   &    50.0    &    &   $1.56^{+0.15}_{-0.14}$ \\\\ GRB 050904    &    00153514   &    174.2    &  1.96e-13  &   $2.012^{+0.050}_{-0.048}$    &    $1.49^{+0.39}_{-1.06}$ \\\\ GRB 050908    &    00154112   &    19.4    &  6.28e-14  &   $1.293^{+0.069}_{-0.064}$ \\\\ GRB 050915A    &    00155242   &    52.0    &  2.94e-13  &   $0.911^{+0.032}_{-0.032}$ \\\\ GRB 050915B    &    00155284   &    40.9    &  2.61e-13  &   $5.32^{+0.32}_{-0.32}$    &    $0.690^{+0.041}_{-0.040}$ \\\\ GRB 050916    &    00155408   &    49.5    &  1.21e-13  &   $4.88^{+3.42}_{-0.98}$    &    $0.998^{+0.093}_{-0.081}$ \\\\ GRB 050922B    &    00156434   &    150.9    &  9.51e-13  &   $3.513^{+0.096}_{-0.091}$    &    $0.349^{+0.084}_{-0.075}$    &    $1.71^{+0.37}_{-0.32}$ \\\\ GRB 050922C    &    00156467   &    4.5    &  6.19e-13  &   $0.985^{+0.063}_{-0.063}$    &    $1.372^{+0.052}_{-0.049}$ \\\\ GRB 051001    &    00157870   &    189.1    &  2.33e-13  &   $0.84^{+0.90}_{-0.60}$    &    $3.235^{+0.111}_{-0.088}$    &    $0.752^{+0.075}_{-0.078}$ \\\\ \\hline \\end{tabular} \\caption{Power-law decay indices from light curve fits. A positive value of $\\alpha$ indicates a decay. The break times between indices are given in Table~\\ref{tab:breaks}. For bursts where the light curve straddles $T0+11$ hrs, the model flux at 11 hours is given (in erg \\cms\\ s$^{-1})$). The \\swift\\ target ID and BAT $T_{90}$ (taken from the Swift Data Table) are given for reference.} \\label{tab:alphas} \\end{center} \\end{table*} \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lccccccccc} \\hline GRB  & Target ID & $T_{90}$  &  F$_{11}$  & $\\alpha_1$ & $\\alpha_2$ & $\\alpha_3$ & $\\alpha_4$ & $\\alpha_5$ & $\\alpha_6$ \\\\ \\hline GRB 051006    &    00158593   &    34.8    &  2.91e-14  &   $1.608^{+0.073}_{-0.063}$ \\\\ GRB 051008    &    00158855   &    $>$32    &  9.91e-13  &   $0.966^{+0.081}_{-0.085}$    &    $2.064^{+0.110}_{-0.097}$ \\\\ GRB 051016A    &    00159913   &    23.0    &  9.66e-14  &   $3.86^{+0.67}_{-0.43}$    &    $0.894^{+0.049}_{-0.047}$ \\\\ GRB 051016B    &    00159994   &    4.0    &  7.22e-13  &   $7.10^{+0.50}_{-1.23}$    &    $0.143^{+0.076}_{-0.086}$    &    $1.009^{+0.049}_{-0.048}$ \\\\ GRB 051021A    &    00020018   &    N/A    &  8.62e-13  &   $1.058^{+0.072}_{-0.063}$ \\\\ GRB 051021B    &    00160672   &    46.5    &  8.94e-14  &   $1.933^{+0.104}_{-0.098}$    &    $0.73^{+0.45}_{-0.56}$ \\\\ GRB 051022    &    00020019   &    N/A    &  8.65e-12  &   $1.547^{+0.075}_{-0.072}$    &    $2.27^{+0.45}_{-0.34}$ \\\\ GRB 051028    &    00020020   &    66.9    &  1.04e-12  &   $1.15^{+0.16}_{-0.12}$ \\\\ GRB 051109A    &    00163136   &    37.2    &  2.82e-12  &   $3.23^{+0.53}_{-0.42}$    &    $0.513^{+0.060}_{-0.149}$    &    $1.221^{+0.025}_{-0.023}$ \\\\ GRB 051109B    &    00163170   &    14.3    &  1.61e-13  &   $3.66^{+6.34}_{-0.69}$    &    $-0.09^{+0.26}_{-0.15}$    &    $1.033^{+0.077}_{-0.067}$ \\\\ GRB 051111    &    00163438   &    46.1    &  3.12e-13  &   $1.60^{+0.11}_{-0.10}$ \\\\ GRB 051117A    &    00164268   &    136.3    &  5.62e-13  &   $1.055^{+0.016}_{-0.016}$    &    $19^{+107}_{-14}$    &    $0.917^{+0.042}_{-0.041}$ \\\\ GRB 051117B    &    00164279   &    9.0    &  2.84e-15  &   $1.71^{+0.33}_{-0.31}$ \\\\ GRB 051210    &    00171931   &    1.3    &  9.42e-17  &   $2.53^{+0.13}_{-0.12}$ \\\\ GRB 051211B    &    00020025   &    N/A    &  5.63e-13  &   $0.982^{+0.064}_{-0.062}$ \\\\ GRB 051221A    &    00173780   &    1.4    &  7.50e-13  &   $1.442^{+0.061}_{-0.059}$    &    $0.07^{+0.20}_{-0.21}$    &    $1.390^{+0.079}_{-0.070}$ \\\\ GRB 051221B    &    00173904   &    39.9    &  1.65e-28  &   $8^{+3}_{-2}$ \\\\ GRB 051227    &    00174738   &    114.6    &  9.50e-14  &   $2.54^{+0.28}_{-0.13}$    &    $1.125^{+0.067}_{-0.082}$ \\\\ GRB 060105    &    00175942   &    54.4    &  1.87e-12  &   $1.175^{+0.162}_{-0.099}$    &    $0.798^{+0.021}_{-0.021}$    &    $2.282^{+0.073}_{-0.073}$    &    $0.26^{+0.43}_{-0.45}$    &    $2.01^{+0.22}_{-0.15}$ \\\\ GRB 060108    &    00176453   &    14.3    &  5.35e-13  &   $1.53^{+0.55}_{-0.34}$    &    $0.32^{+0.13}_{-0.19}$    &    $1.17^{+0.17}_{-0.11}$ \\\\ GRB 060109    &    00176620   &    115.4    &  4.77e-13  &   $4.56^{+0.35}_{-0.29}$    &    $0.108^{+0.088}_{-0.088}$    &    $1.56^{+0.11}_{-0.10}$ \\\\ GRB 060110    &    00176702   &    26.0    &    &   $1.86^{+1.28}_{-0.24}$ \\\\ GRB 060111A    &    00176818   &    13.2    &  4.89e-13  &   $1.530^{+0.043}_{-0.039}$    &    $0.830^{+0.054}_{-0.052}$ \\\\ GRB 060111B    &    00176918   &    58.8    &  2.45e-13  &   $3.42^{+0.48}_{-0.44}$    &    $0.967^{+0.087}_{-0.086}$    &    $1.33^{+0.18}_{-0.13}$ \\\\ GRB 060115    &    00177408   &    139.6    &  3.87e-13  &   $3.09^{+0.20}_{-0.22}$    &    $0.838^{+0.036}_{-0.035}$ \\\\ GRB 060116    &    00177533   &    105.9    &  2.28e-13  &   $1.049^{+0.049}_{-0.044}$ \\\\ GRB 060121    &    00020027   &    N/A    &  8.05e-13  &   $1.204^{+0.060}_{-0.056}$ \\\\ GRB 060123    &    00020028   &    N/A    &    &   $1.23^{+0.13}_{-0.13}$ \\\\ GRB 060124    &    00178750   &    ~75    &  1.04e-11  &   $-0.31^{+0.37}_{-0.18}$    &    $1.097^{+0.066}_{-0.074}$    &    $1.465^{+0.047}_{-0.042}$ \\\\ GRB 060202    &    00179968   &    198.9    &  1.16e-12  &   $2.733^{+0.102}_{-0.098}$    &    $0.137^{+0.062}_{-0.064}$    &    $5.68^{+0.25}_{-0.23}$    &    $0.863^{+0.031}_{-0.029}$ \\\\ GRB 060203    &    00180151   &    68.2    &  4.21e-13  &   $0.972^{+0.060}_{-0.057}$ \\\\ GRB 060204B    &    00180241   &    139.4    &  6.79e-13  &   $3.86^{+0.30}_{-0.32}$    &    $0.670^{+0.085}_{-0.097}$    &    $1.492^{+0.084}_{-0.073}$ \\\\ GRB 060206    &    00180455   &    7.6    &  1.10e-12  &   $-0.22^{+0.53}_{-0.58}$    &    $0.944^{+0.024}_{-0.025}$ \\\\ GRB 060210    &    00180977   &    255.0    &  5.00e-12  &   $0.942^{+0.046}_{-0.049}$    &    $1.333^{+0.103}_{-0.051}$ \\\\ GRB 060211A    &    00181126   &    126.3    &  2.72e-13  &   $4.02^{+0.12}_{-0.12}$    &    $2.53^{+0.30}_{-0.47}$    &    $0.509^{+0.076}_{-0.079}$ \\\\ GRB 060211B    &    00181156   &    27.7    &  3.80e-13  &   $2.13^{+0.27}_{-0.18}$    &    $-0.07^{+0.53}_{-0.87}$ \\\\ GRB 060218    &    00191157   &    ~21    &  3.86e-13  &   $5.87^{+0.41}_{-0.56}$    &    $1.129^{+0.061}_{-0.058}$ \\\\ GRB 060219    &    00191512   &    62.1    &  2.08e-13  &   $10^{+2}_{-3}$    &    $0.527^{+0.085}_{-0.100}$    &    $1.31^{+0.33}_{-0.19}$ \\\\ GRB 060223A    &    00192059   &    11.3    &  2.97e-14  &   $0.28^{+0.11}_{-0.73}$    &    $1.34^{+0.12}_{-0.10}$ \\\\ GRB 060306    &    00200638   &    61.2    &  1.07e-12  &   $3.81^{+0.38}_{-0.42}$    &    $0.628^{+0.050}_{-0.063}$    &    $1.069^{+0.058}_{-0.060}$ \\\\ GRB 060312    &    00201391   &    50.6    &  4.28e-13  &   $2.515^{+0.242}_{-0.075}$    &    $0.744^{+0.045}_{-0.041}$ \\\\ GRB 060313    &    00201487   &    0.74    &  3.54e-13  &   $-0.36^{+0.16}_{-0.37}$    &    $3^{+4}_{-1}$    &    $0.743^{+0.079}_{-0.124}$    &    $1.59^{+0.13}_{-0.14}$ \\\\ GRB 060319    &    00202035   &    10.6    &  8.23e-13  &   $0.749^{+0.055}_{-0.077}$    &    $1.168^{+0.075}_{-0.069}$ \\\\ GRB 060323    &    00202505   &    25.4    &  2.07e-13  &   $1.027^{+0.047}_{-0.045}$ \\\\ GRB 060403    &    00203755   &    30.1    &  8.38e-14  &   $1.501^{+0.131}_{-0.081}$    &    $1.09^{+0.31}_{-1.56}$ \\\\ GRB 060413    &    00205096   &    147.7    &  5.26e-12  &   $2.836^{+0.073}_{-0.072}$    &    $0.209^{+0.058}_{-0.061}$    &    $2.83^{+0.17}_{-0.14}$ \\\\ GRB 060418    &    00205851   &    103.1    &  5.27e-13  &   $2.482^{+0.066}_{-0.070}$    &    $1.03^{+0.15}_{-0.19}$    &    $1.555^{+0.091}_{-0.075}$ \\\\ GRB 060421    &    00206257   &    12.2    &  2.42e-13  &   $1.036^{+0.054}_{-0.053}$ \\\\ GRB 060427    &    00207281   &    64.0    &  1.15e-13  &   $4.02^{+0.33}_{-0.31}$    &    $1.290^{+0.077}_{-0.075}$ \\\\ GRB 060428A    &    00207364   &    39.5    &  5.00e-12  &   $6.50^{+0.71}_{-1.30}$    &    $0.519^{+0.034}_{-0.036}$    &    $1.129^{+0.074}_{-0.068}$    &    $2.23^{+0.82}_{-0.34}$ \\\\ GRB 060428B    &    00207399   &    57.9    &  2.69e-13  &   $4.49^{+0.14}_{-0.14}$    &    $0.971^{+0.043}_{-0.044}$ \\\\ GRB 060502A    &    00208169   &    28.4    &  2.19e-12  &   $3.00^{+0.24}_{-0.23}$    &    $0.575^{+0.058}_{-0.070}$    &    $1.120^{+0.064}_{-0.063}$ \\\\ GRB 060505    &    00208654   &    ~4    &    &   $1.04^{+0.47}_{-0.24}$ \\\\ GRB 060507    &    00208870   &    183.3    &  6.00e-13  &   $0.08^{+0.48}_{-0.70}$    &    $1.098^{+0.075}_{-0.070}$ \\\\ GRB 060510A    &    00209351   &    20.4    &  9.53e-12  &   $2^{+1}_{-1}$    &    $-0.002^{+0.085}_{-0.091}$    &    $1.489^{+0.044}_{-0.042}$ \\\\ GRB 060510B    &    00209352   &    275.2    &  1.71e-13  &   $0.225^{+0.042}_{-0.044}$    &    $10.97^{+0.53}_{-0.60}$    &    $0.88^{+0.16}_{-0.13}$ \\\\ GRB 060512    &    00209755   &    8.5    &  1.45e-13  &   $3^{+1}_{-1}$    &    $1.139^{+0.058}_{-0.060}$ \\\\ \\hline \\end{tabular} \\contcaption{} \\end{center} \\end{table*} \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lccccccccc} \\hline GRB  & Target ID & $T_{90}$  &  F$_{11}$  & $\\alpha_1$ & $\\alpha_2$ & $\\alpha_3$ & $\\alpha_4$ & $\\alpha_5$ & $\\alpha_6$ \\\\ \\hline GRB 060515    &    00210084   &    52.0    &  2.70e-14  &   $1.79^{+0.52}_{-0.38}$ \\\\ GRB 060522    &    00211117   &    71.1    &  1.81e-13  &   $3.80^{+0.63}_{-0.49}$    &    $1.78^{+0.21}_{-0.21}$    &    $0.904^{+0.069}_{-0.054}$ \\\\ GRB 060526    &    00211957   &    298.2    &  8.15e-13  &   $1.222^{+0.073}_{-0.081}$    &    $0.36^{+0.15}_{-0.39}$    &    $0.947^{+0.161}_{-0.085}$    &    $2.91^{+2.19}_{-0.78}$ \\\\ GRB 060602A    &    00213180   &    75.0    &    &   $1.14^{+0.99}_{-0.32}$ \\\\ GRB 060602B    &    00213190   &    N/A    &  2.12e-13  &   $1.115^{+0.102}_{-0.086}$ \\\\ GRB 060604    &    00213486   &    95.0    &  9.56e-13  &   $2.99^{+0.41}_{-0.45}$    &    $0.53^{+0.12}_{-0.47}$    &    $1.238^{+0.080}_{-0.073}$ \\\\ GRB 060605    &    00213630   &    79.1    &  3.81e-13  &   $1.125^{+0.045}_{-0.199}$    &    $0.262^{+0.045}_{-0.279}$    &    $1.96^{+0.11}_{-0.10}$ \\\\ GRB 060607A    &    00213823   &    102.2    &  1.15e-12  &   $1.116^{+0.043}_{-0.047}$    &    $0.460^{+0.028}_{-0.039}$    &    $3.45^{+0.12}_{-0.12}$ \\\\ GRB 060614    &    00214805   &    108.7    &  6.28e-12  &   $1.91^{+0.14}_{-0.13}$    &    $3.573^{+0.075}_{-0.067}$    &    $0.077^{+0.082}_{-0.069}$    &    $1.863^{+0.070}_{-0.051}$ \\\\ GRB 060707    &    00217704   &    66.2    &  1.79e-12  &   $1.482^{+0.105}_{-0.099}$    &    $0.33^{+0.13}_{-0.60}$    &    $1.179^{+0.086}_{-0.089}$ \\\\ GRB 060708    &    00217805   &    10.2    &  6.62e-13  &   $4.21^{+0.22}_{-0.21}$    &    $0.675^{+0.046}_{-0.054}$    &    $1.308^{+0.076}_{-0.109}$ \\\\ GRB 060712    &    00218582   &    26.3    &  2.47e-13  &   $2.75^{+0.58}_{-0.66}$    &    $-0.199960^{+0.226000}_{-0.000036}$    &    $1.170^{+0.098}_{-0.089}$ \\\\ GRB 060714    &    00219101   &    115.0    &  8.33e-13  &   $9^{+1}_{-1}$    &    $0.47^{+0.18}_{-0.18}$    &    $1.236^{+0.052}_{-0.048}$ \\\\ GRB 060717    &    00219646   &    3.0    &  3.72e-14  &   $0.92^{+0.17}_{-0.20}$ \\\\ GRB 060719    &    00220020   &    66.9    &  4.94e-13  &   $6.46^{+0.45}_{-0.61}$    &    $0.21^{+0.11}_{-0.13}$    &    $1.225^{+0.079}_{-0.072}$ \\\\ GRB 060729    &    00221755   &    115.3    &  1.12e-11  &   $4.68^{+0.14}_{-0.20}$    &    $7.03^{+0.25}_{-0.24}$    &    $0.296^{+0.023}_{-0.023}$    &    $1.259^{+0.028}_{-0.029}$    &    $1.637^{+0.100}_{-0.078}$ \\\\ GRB 060801    &    00222154   &    0.49    &  1.63e-29  &   $-3.10^{+0.74}_{-6.89}$    &    $1.62^{+0.40}_{-0.34}$    &    $9^{+21}_{-4}$ \\\\ GRB 060804    &    00222546   &    17.8    &  8.34e-13  &   $-0.13^{+0.22}_{-0.10}$    &    $1.220^{+0.111}_{-0.073}$ \\\\ GRB 060805A    &    00222683   &    5.3    &  6.07e-14  &   $0.424^{+0.100}_{-0.105}$    &    $1.68^{+0.41}_{-0.44}$ \\\\ GRB 060805B    &    00020037   &    N/A    &    &   $0.86^{+0.37}_{-0.32}$ \\\\ GRB 060807    &    00223217   &    54.0    &  1.29e-12  &   $2.47^{+4.75}_{-0.49}$    &    $-0.035^{+0.080}_{-0.156}$    &    $1.105^{+0.087}_{-0.085}$    &    $2.02^{+0.30}_{-0.15}$ \\\\ GRB 060813    &    00224364   &    16.1    &  2.48e-12  &   $0.573^{+0.043}_{-0.043}$    &    $1.236^{+0.050}_{-0.053}$    &    $2.45^{+0.53}_{-0.38}$ \\\\ GRB 060814    &    00224552   &    145.3    &  2.30e-12  &   $3.42^{+0.21}_{-0.55}$    &    $2.162^{+0.088}_{-0.281}$    &    $3.27^{+0.56}_{-0.29}$    &    $0.624^{+0.053}_{-0.058}$    &    $1.388^{+0.042}_{-0.040}$ \\\\ GRB 060825    &    00226382   &    8.0    &  2.06e-13  &   $0.945^{+0.046}_{-0.048}$ \\\\ GRB 060901    &    00020039   &    N/A    &  8.05e-13  &   $1.37^{+0.41}_{-0.29}$ \\\\ GRB 060904A    &    00227996   &    80.1    &  1.83e-12  &   $3.98^{+0.12}_{-0.12}$    &    $0.58^{+0.13}_{-0.13}$    &    $1.155^{+0.144}_{-0.075}$ \\\\ GRB 060904B    &    00228006   &    171.5    &  4.54e-13  &   $0.762^{+0.057}_{-0.610}$    &    $1.354^{+0.063}_{-0.056}$ \\\\ GRB 060906    &    00228316   &    43.5    &  3.31e-13  &   $3.86^{+0.58}_{-1.30}$    &    $0.272^{+0.088}_{-0.094}$    &    $1.59^{+0.20}_{-0.16}$ \\\\ GRB 060908    &    00228581   &    19.3    &  1.87e-13  &   $0.578^{+0.052}_{-0.139}$    &    $1.58^{+0.12}_{-0.13}$ \\\\ GRB 060912A    &    00229185   &    5.0    &  2.64e-13  &   $0.75^{+0.17}_{-0.26}$    &    $1.126^{+0.056}_{-0.050}$ \\\\ GRB 060919    &    00230115   &    9.1    &  8.72e-14  &   $0.975^{+0.091}_{-0.083}$ \\\\ GRB 060923A    &    00230662   &    51.7    &  6.51e-13  &   $2.50^{+0.31}_{-0.26}$    &    $-0.18^{+0.29}_{-0.89}$    &    $1.213^{+0.053}_{-0.049}$ \\\\ GRB 060923B    &    00230702   &    8.6    &  1.17e-12  &   $0.622^{+0.077}_{-0.081}$ \\\\ GRB 060923C    &    00230711   &    75.8    &  5.86e-13  &   $3.16^{+0.23}_{-0.18}$    &    $0.591^{+0.060}_{-0.190}$    &    $2.00^{+0.90}_{-0.65}$ \\\\ GRB 060926    &    00231231   &    8.0    &  1.46e-13  &   $1.84^{+0.70}_{-0.89}$    &    $0.23^{+0.24}_{-0.63}$    &    $1.44^{+0.24}_{-0.19}$ \\\\ GRB 060927    &    00231362   &    22.5    &  7.65e-14  &   $0.730^{+0.070}_{-0.364}$    &    $1.76^{+0.82}_{-0.42}$ \\\\ GRB 060928    &    00020042   &    N/A    &    &   $0.72^{+1.70}_{-0.32}$ \\\\ GRB 060929    &    00231702   &    554.0    &  8.73e-14  &   $-0.04^{+0.71}_{-1.05}$    &    $1.01^{+0.13}_{-0.10}$ \\\\ GRB 061002    &    00231974   &    17.6    &  8.19e-14  &   $2.87^{+0.48}_{-0.37}$    &    $0.97^{+0.13}_{-0.16}$ \\\\ GRB 061004    &    00232339   &    6.2    &  2.68e-13  &   $1.72^{+0.19}_{-0.17}$    &    $0.39^{+0.15}_{-0.33}$    &    $1.18^{+0.12}_{-0.11}$ \\\\ GRB 061006    &    00232585   &    129.9    &  1.56e-13  &   $0.787^{+0.076}_{-0.071}$ \\\\ GRB 061007    &    00232683   &    75.3    &  1.07e-12  &   $1.945^{+0.053}_{-0.068}$    &    $1.643^{+0.022}_{-0.022}$    &    $2.52^{+0.93}_{-0.58}$    &    $1.497^{+0.051}_{-0.047}$ \\\\ GRB 061019    &    00234516   &    183.8    &  9.57e-13  &   $0.863^{+0.052}_{-0.078}$    &    $2.31^{+0.86}_{-0.69}$ \\\\ GRB 061021    &    00234905   &    46.2    &  3.03e-12  &   $2.001^{+0.111}_{-0.096}$    &    $0.541^{+0.074}_{-0.076}$    &    $1.075^{+0.017}_{-0.016}$ \\\\ GRB 061025    &    00020043   &    N/A    &  1.31e-13  &   $1.32^{+0.31}_{-0.31}$ \\\\ GRB 061028    &    00235715   &    106.2    &  3.67e-14  &   $4.68^{+0.21}_{-0.18}$    &    $0.68^{+0.26}_{-0.29}$ \\\\ GRB 061102    &    00236430   &    45.6    &  9.05e-14  &   $3.66^{+0.54}_{-0.51}$    &    $0.36^{+0.24}_{-0.23}$ \\\\ GRB 061110A    &    00238108   &    40.7    &  9.54e-14  &   $0.48^{+0.16}_{-0.13}$    &    $4.26^{+0.22}_{-0.14}$    &    $2.27^{+0.24}_{-0.29}$    &    $0.631^{+0.081}_{-0.084}$ \\\\ GRB 061110B    &    00238174   &    134.0    &  6.24e-14  &   $1.47^{+0.32}_{-0.28}$ \\\\ GRB 061121    &    00239899   &    81.3    &  7.36e-12  &   $2.80^{+0.29}_{-0.33}$    &    $8.07^{+0.59}_{-0.54}$    &    $3.90^{+0.56}_{-0.60}$    &    $0.265^{+0.044}_{-0.043}$    &    $1.010^{+0.071}_{-0.069}$    &    $1.553^{+0.046}_{-0.070}$ \\\\ GRB 061122    &    00020044   &    N/A    &  1.65e-12  &   $1.276^{+0.070}_{-0.065}$ \\\\ GRB 061126    &    00240766   &    70.8    &  3.36e-12  &   $1.350^{+0.027}_{-0.023}$    &    $1.276^{+0.044}_{-0.052}$ \\\\ GRB 061201    &    00241840   &    0.76    &  1.76e-13  &   $0.54^{+0.13}_{-0.14}$    &    $1.86^{+0.18}_{-0.15}$ \\\\ GRB 061202    &    00241963   &    91.2    &  3.20e-12  &   $2.55^{+0.13}_{-0.13}$    &    $-0.27^{+0.13}_{-0.13}$    &    $1.741^{+0.059}_{-0.055}$ \\\\ GRB 061210    &    00243690   &    85.3    &    &   $1.58^{+1.44}_{-0.33}$ \\\\ GRB 061222A    &    00252588   &    71.4    &  8.33e-12  &   $4.18^{+0.12}_{-0.11}$    &    $0.296^{+0.076}_{-0.076}$    &    $0.976^{+0.103}_{-0.059}$    &    $1.708^{+0.054}_{-0.052}$ \\\\ GRB 061222B    &    00252593   &    40.0    &  6.48e-17  &   $2.97^{+0.17}_{-0.16}$ \\\\ \\hline \\end{tabular} \\contcaption{} \\end{center} \\end{table*} \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lccccccccc} \\hline GRB  & Target ID & $T_{90}$  &  F$_{11}$  & $\\alpha_1$ & $\\alpha_2$ & $\\alpha_3$ & $\\alpha_4$ & $\\alpha_5$ & $\\alpha_6$ \\\\ \\hline GRB 070103    &    00254532   &    18.6    &  1.62e-13  &   $-0.64^{+0.22}_{-0.37}$    &    $0.86^{+0.18}_{-0.67}$    &    $1.43^{+0.12}_{-0.11}$ \\\\ GRB 070107    &    00255029   &    347.3    &  2.29e-12  &   $0.864^{+0.020}_{-0.021}$    &    $1.308^{+0.046}_{-0.040}$ \\\\ GRB 070110    &    00255445   &    88.4    &  6.83e-13  &   $2.558^{+0.051}_{-0.128}$    &    $0.612^{+0.586}_{-0.095}$    &    $0.132^{+0.037}_{-0.041}$    &    $9.65^{+0.35}_{-0.39}$    &    $0.892^{+0.045}_{-0.040}$ \\\\ GRB 070125    &    00020047   &    N/A    &    &   $0.65^{+0.36}_{-0.27}$    &    $2.01^{+0.18}_{-0.15}$ \\\\ GRB 070129    &    00258408   &    460.6    &  1.14e-12  &   $2.532^{+0.079}_{-0.070}$    &    $-0.093^{+0.230}_{-0.091}$    &    $1.034^{+0.051}_{-0.049}$    &    $2.41^{+1.38}_{-0.97}$ \\\\ GRB 070208    &    00259714   &    47.7    &  3.42e-13  &   $-0.98^{+0.20}_{-0.23}$    &    $0.99^{+0.13}_{-0.38}$    &    $1.58^{+0.59}_{-0.28}$ \\\\ GRB 070219    &    00261132   &    16.6    &  1.02e-13  &   $1.67^{+0.94}_{-0.55}$    &    $-0.27^{+0.75}_{-0.80}$    &    $0.99^{+0.20}_{-0.17}$ \\\\ GRB 070220    &    00261299   &    129.0    &  6.20e-13  &   $1.28^{+0.15}_{-0.17}$    &    $2.25^{+2.48}_{-0.27}$    &    $1.058^{+0.052}_{-0.040}$    &    $2.10^{+0.28}_{-0.14}$ \\\\ GRB 070223    &    00261664   &    88.5    &  2.58e-13  &   $0.666^{+0.053}_{-0.430}$    &    $2.48^{+0.44}_{-0.42}$    &    $0.916^{+0.087}_{-0.086}$ \\\\ GRB 070224    &    00261880   &    34.5    &  1.60e-13  &   $3.22^{+0.24}_{-0.21}$    &    $0.560^{+0.069}_{-0.089}$ \\\\ GRB 070227    &    00262347   &    7    &    &   $1.11^{+0.14}_{-0.13}$ \\\\ GRB 070306    &    00263361   &    209.5    &  8.85e-12  &   $-0.13^{+0.35}_{-0.58}$    &    $6.34^{+0.22}_{-0.20}$    &    $0.139^{+0.059}_{-0.057}$    &    $0.76^{+0.38}_{-0.25}$    &    $1.887^{+0.091}_{-0.085}$ \\\\ GRB 070311    &    00020052   &    N/A    &  6.71e-13  &   $1.26^{+0.20}_{-0.14}$    &    $-0.02^{+0.21}_{-10.00}$    &    $2.96^{+0.58}_{-0.47}$ \\\\ GRB 070318    &    00271019   &    74.6    &  1.33e-12  &   $1.083^{+0.011}_{-0.011}$ \\\\ GRB 070328    &    00272773   &    75.3    &  2.38e-12  &   $1.34^{+0.21}_{-0.13}$    &    $0.277^{+0.052}_{-0.052}$    &    $1.211^{+0.037}_{-0.111}$    &    $1.587^{+0.040}_{-0.034}$ \\\\ GRB 070330    &    00273180   &    9.0    &  3.18e-13  &   $0.881^{+0.046}_{-0.050}$ \\\\ GRB 070411    &    00275087   &    121.5    &  5.70e-13  &   $1.134^{+0.038}_{-0.037}$ \\\\ GRB 070412    &    00275119   &    33.8    &  3.13e-13  &   $3.25^{+0.49}_{-0.85}$    &    $0.881^{+0.053}_{-0.077}$    &    $1.48^{+6.21}_{-0.25}$ \\\\ GRB 070419A    &    00276205   &    115.6    &  3.56e-14  &   $1.92^{+0.16}_{-0.15}$    &    $3.71^{+0.18}_{-0.15}$    &    $0.53^{+0.47}_{-0.37}$ \\\\ GRB 070419B    &    00276212   &    236.4    &  1.07e-11  &   $1.1591^{+0.0093}_{-0.0089}$    &    $2.41^{+0.20}_{-0.16}$ \\\\ GRB 070420    &    00276321   &    76.5    &  3.01e-12  &   $4.50^{+0.19}_{-0.18}$    &    $0.351^{+0.085}_{-0.094}$    &    $2.78^{+0.48}_{-0.49}$    &    $0.87^{+0.16}_{-0.19}$    &    $1.87^{+0.21}_{-0.15}$ \\\\ GRB 070429A    &    00277571   &    163.3    &  5.61e-13  &   $7.01^{+0.28}_{-0.56}$    &    $4.13^{+0.35}_{-0.49}$    &    $0.466^{+0.057}_{-0.073}$    &    $3^{+7}_{-2}$ \\\\ GRB 070506    &    00278693   &    4.3    &  1.20e-12  &   $0.55^{+0.11}_{-0.11}$ \\\\ GRB 070508    &    00278854   &    20.9    &  3.00e-12  &   $0.476^{+0.035}_{-0.034}$    &    $1.149^{+0.045}_{-0.038}$    &    $1.654^{+0.087}_{-0.122}$ \\\\ GRB 070509    &    00278903   &    7.7    &  7.03e-14  &   $0.910^{+0.120}_{-0.099}$ \\\\ GRB 070517    &    00279494   &    7.6    &  1.85e-13  &   $0.408^{+0.061}_{-0.069}$    &    $2.33^{+7.67}_{-0.69}$ \\\\ GRB 070518    &    00279592   &    5.5    &  1.51e-13  &   $2.22^{+0.12}_{-0.11}$    &    $0.649^{+0.070}_{-0.079}$ \\\\ GRB 070520A    &    00279817   &    18.1    &  1.34e-13  &   $4.33^{+0.42}_{-0.42}$    &    $0.641^{+0.089}_{-0.090}$ \\\\ GRB 070520B    &    00279898   &    65.8    &  9.44e-14  &   $2.14^{+0.18}_{-0.16}$    &    $0.99^{+0.66}_{-0.57}$ \\\\ GRB 070521    &    00279935   &    37.9    &  9.45e-13  &   $-0.05^{+0.16}_{-0.21}$    &    $0.783^{+0.056}_{-0.060}$    &    $1.927^{+0.108}_{-0.098}$ \\\\ GRB 070529    &    00280706   &    109.2    &  3.45e-13  &   $9^{+2}_{-1}$    &    $0.799^{+0.152}_{-0.096}$    &    $1.306^{+0.086}_{-0.075}$ \\\\ GRB 070531    &    00280958   &    44.5    &    &   $1.46^{+0.14}_{-0.14}$ \\\\ GRB 070611    &    00282003   &    12.2    &  1.79e-13  &   $1.105^{+0.111}_{-0.097}$ \\\\ GRB 070612B    &    00282073   &    13.5    &  1.23e-13  &   $1.28^{+0.52}_{-0.47}$ \\\\ GRB 070616    &    00282445   &    402.4    &  4.63e-13  &   $-0.015^{+0.027}_{-0.027}$    &    $6.46^{+0.24}_{-0.21}$    &    $1.359^{+0.037}_{-0.037}$ \\\\ GRB 070621    &    00282808   &    33.3    &  2.90e-13  &   $3.645^{+0.036}_{-0.055}$    &    $1.37^{+0.19}_{-0.23}$    &    $0.958^{+0.043}_{-0.043}$ \\\\ GRB 070628    &    00283320   &    39.1    &  3.85e-12  &   $0.356^{+0.082}_{-0.084}$    &    $1.170^{+0.047}_{-0.047}$ \\\\ GRB 070704    &    00283791   &    38    &  6.61e-13  &   $2.75^{+0.62}_{-0.42}$    &    $0.839^{+0.080}_{-0.084}$ \\\\ GRB 070714A    &    00284850   &    2.0    &  3.83e-13  &   $0.635^{+0.055}_{-0.048}$ \\\\ GRB 070714B    &    00284856   &    64    &  5.97e-14  &   $1.28^{+0.19}_{-0.22}$    &    $2.29^{+0.22}_{-0.20}$    &    $0.36^{+0.46}_{-1.56}$    &    $1.565^{+0.088}_{-0.076}$ \\\\ GRB 070721A    &    00285653   &    3.4    &  1.38e-13  &   $2.82^{+0.57}_{-0.38}$    &    $0.798^{+0.057}_{-0.056}$ \\\\ GRB 070721B    &    00285654   &    34    &  3.63e-13  &   $5.43^{+0.84}_{-0.54}$    &    $0.780^{+0.041}_{-0.041}$    &    $2.34^{+0.27}_{-0.19}$ \\\\ GRB 070724A    &    00285948   &    0.4    &  1.37e-13  &   $2.50^{+0.20}_{-0.16}$    &    $0.70^{+0.21}_{-0.27}$ \\\\ GRB 070724B    &    00020055   &    N/A    &    &   $1.14^{+0.20}_{-0.17}$ \\\\ GRB 070802    &    00286809   &    16.4    &  3.41e-13  &   $2.36^{+0.58}_{-0.43}$    &    $0.157^{+0.086}_{-0.259}$    &    $1.18^{+0.15}_{-0.17}$ \\\\ GRB 070808    &    00287260   &    32    &  2.03e-13  &   $3.49^{+0.62}_{-0.47}$    &    $0.923^{+0.092}_{-0.079}$ \\\\ GRB 070809    &    00287344   &    1.3    &  5.06e-13  &   $0.511^{+0.055}_{-0.055}$ \\\\ GRB 070810A    &    00287364   &    11.0    &  2.86e-13  &   $0.49^{+0.31}_{-0.28}$    &    $1.29^{+0.12}_{-0.11}$ \\\\ GRB 070911    &    00290624   &    162    &    &   $1.516^{+0.075}_{-0.070}$ \\\\ GRB 070917    &    00291292   &    7.3    &    &   $0.87^{+0.58}_{-0.32}$ \\\\ GRB 070925    &    00020056   &    N/A    &  2.13e-13  &   $1.03^{+0.28}_{-0.21}$ \\\\ GRB 071001    &    00292826   &    58.5    &    &   $2.56^{+0.23}_{-0.19}$    &    $-0.07^{+0.17}_{-0.22}$ \\\\ GRB 071003    &    00292934   &    15    &  2.53e-12  &   $1.620^{+0.060}_{-0.056}$ \\\\ GRB 071010A    &    00293707   &    6    &  8.27e-13  &   $-0.34^{+0.55}_{-0.82}$    &    $1.79^{+0.30}_{-0.26}$ \\\\ GRB 071010B    &    00293795   &    $>$35.7    &  3.38e-12  &   $0.661^{+0.064}_{-0.064}$ \\\\ GRB 071011    &    00293924   &    61    &  2.98e-12  &   $0.927^{+0.046}_{-0.045}$    &    $2.45^{+2.09}_{-0.96}$ \\\\ GRB 071020    &    00294835   &    4.2    &  1.50e-12  &   $0.47^{+0.17}_{-0.22}$    &    $1.145^{+0.019}_{-0.017}$ \\\\ \\hline \\end{tabular} \\contcaption{} \\end{center} \\end{table*} \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lccccccccc} \\hline GRB  & Target ID & $T_{90}$  &  F$_{11}$  & $\\alpha_1$ & $\\alpha_2$ & $\\alpha_3$ & $\\alpha_4$ & $\\alpha_5$ & $\\alpha_6$ \\\\ \\hline GRB 071021    &    00294974   &    225    &  3.10e-13  &   $4.40^{+0.43}_{-0.26}$    &    $0.813^{+0.042}_{-0.044}$ \\\\ GRB 071025    &    00295301   &    109    &  5.72e-13  &   $3.59^{+0.18}_{-0.26}$    &    $2.25^{+0.14}_{-0.20}$    &    $1.322^{+0.054}_{-0.040}$    &    $1.668^{+0.064}_{-0.063}$ \\\\ GRB 071028A    &    00295527   &    27.0    &  1.83e-13  &   $3.18^{+0.19}_{-0.17}$    &    $-0.06^{+0.21}_{-0.59}$    &    $1.95^{+1.20}_{-0.85}$ \\\\ GRB 071031    &    00295670   &    18    &  2.51e-13  &   $4.84^{+0.48}_{-0.31}$    &    $0.95^{+0.14}_{-0.39}$    &    $1.30^{+8.70}_{-0.35}$ \\\\ GRB 071101    &    00295779   &    9.0    &  6.64e-14  &   $0.73^{+0.13}_{-0.15}$ \\\\ GRB 071104    &    00020058   &    N/A    &  8.88e-13  &   $1.34^{+0.26}_{-0.26}$ \\\\ GRB 071112C    &    00296504   &    15    &  3.21e-13  &   $1.424^{+0.096}_{-0.117}$    &    $0.81^{+0.13}_{-0.18}$    &    $1.535^{+0.048}_{-0.045}$ \\\\ GRB 071117    &    00296805   &    6.6    &  6.46e-13  &   $0.979^{+0.082}_{-0.082}$ \\\\ GRB 071118    &    00296856   &    71    &  9.01e-13  &   $2.24^{+0.29}_{-0.29}$    &    $0.672^{+0.069}_{-0.071}$    &    $2.05^{+6.52}_{-0.31}$ \\\\ GRB 071122    &    00297114   &    68.7    &    &   $2.537^{+0.108}_{-0.097}$ \\\\ GRB 071227    &    00299787   &    1.8    &  2.32e-14  &   $1.08^{+0.22}_{-0.24}$    &    $5.48^{+0.70}_{-0.62}$    &    $1.19^{+0.16}_{-0.13}$ \\\\ GRB 080120    &    00020060   &    N/A    &  9.14e-14  &   $2.96^{+0.66}_{-0.55}$ \\\\ GRB 080123    &    00301578   &    115    &  2.93e-14  &   $1.82^{+0.18}_{-0.20}$    &    $8.12^{+1.00}_{-0.80}$    &    $0.80^{+0.30}_{-0.25}$ \\\\ GRB 080129    &    00301981   &    48    &  7.99e-13  &   $0.61^{+0.16}_{-0.12}$    &    $1.30^{+0.13}_{-0.10}$ \\\\ GRB 080205    &    00302506   &    106.5    &  1.22e-13  &   $3.29^{+0.43}_{-0.66}$    &    $7.13^{+0.86}_{-1.41}$    &    $0.52^{+0.10}_{-0.13}$    &    $1.518^{+0.059}_{-0.056}$ \\\\ GRB 080207    &    00302728   &    34    &  6.49e-13  &   $-0.27^{+0.24}_{-0.24}$    &    $1.703^{+0.063}_{-0.060}$ \\\\ GRB 080210    &    00302888   &    45    &  4.81e-13  &   $1.172^{+0.046}_{-0.048}$ \\\\ GRB 080212    &    00303105   &    123    &  8.79e-13  &   $6.23^{+0.32}_{-0.31}$    &    $0.43^{+0.19}_{-0.45}$    &    $1.371^{+0.074}_{-0.070}$ \\\\ GRB 080218B    &    00303631   &    6.2    &  1.96e-13  &   $0.988^{+0.050}_{-0.048}$ \\\\ GRB 080229A    &    00304379   &    64    &  1.66e-11  &   $3.27^{+0.17}_{-0.18}$    &    $-0.036^{+0.067}_{-0.085}$    &    $0.897^{+0.037}_{-0.040}$    &    $1.50^{+0.65}_{-0.20}$ \\\\ GRB 080229B    &    00020064   &    N/A    &    &   $1.18^{+0.18}_{-0.17}$ \\\\ GRB 080303    &    00304549   &    67    &  2.25e-13  &   $1.606^{+0.075}_{-0.071}$    &    $0.788^{+0.079}_{-0.075}$ \\\\ GRB 080307    &    00305011   &    125.9    &  9.98e-14  &   $-0.48^{+0.14}_{-0.13}$    &    $1.945^{+0.046}_{-0.041}$    &    $0.43^{+0.14}_{-0.17}$ \\\\ GRB 080310    &    00305288   &    365    &  9.77e-13  &   $2.66^{+0.69}_{-0.67}$    &    $0.32^{+0.11}_{-0.12}$    &    $1.539^{+0.095}_{-0.084}$ \\\\ GRB 080319A    &    00306754   &    64    &  5.72e-13  &   $0.907^{+0.089}_{-0.070}$ \\\\ GRB 080319B    &    00306757   &    $>$5    &  3.29e-12  &   $1.352^{+0.035}_{-0.029}$    &    $1.588^{+0.016}_{-0.015}$    &    $2.89^{+0.21}_{-0.22}$    &    $1.027^{+0.061}_{-0.066}$    &    $2.35^{+0.61}_{-0.35}$ \\\\ GRB 080319C    &    00306778   &    34    &  2.18e-12  &   $-2.17^{+0.96}_{-0.56}$    &    $0.958^{+0.078}_{-0.078}$    &    $1.66^{+0.15}_{-0.13}$ \\\\ GRB 080319D    &    00306793   &    24    &  1.92e-13  &   $2.60^{+0.23}_{-0.19}$    &    $0.82^{+0.14}_{-0.18}$ \\\\ GRB 080320    &    00306858   &    14    &  1.75e-12  &   $1.88^{+1.26}_{-0.38}$    &    $0.618^{+0.057}_{-0.052}$    &    $1.190^{+0.133}_{-0.069}$ \\\\ GRB 080325    &    00307604   &    128.4    &  5.38e-13  &   $4.57^{+0.96}_{-0.87}$    &    $0.974^{+0.073}_{-0.074}$ \\\\ GRB 080328    &    00307931   &    90.6    &  2.73e-12  &   $5.77^{+0.48}_{-0.43}$    &    $0.706^{+0.045}_{-0.046}$    &    $1.227^{+0.059}_{-0.056}$ \\\\ GRB 080330    &    00308041   &    61    &  1.07e-12  &   $5.27^{+0.42}_{-0.41}$    &    $0.18^{+0.11}_{-0.19}$    &    $0.82^{+8.12}_{-0.27}$ \\\\ GRB 080405    &    00020065   &    N/A    &    &   $0.15^{+1.09}_{-0.33}$ \\\\ GRB 080409    &    00308812   &    20.2    &  2.38e-13  &   $-0.54^{+0.32}_{-0.84}$    &    $0.900^{+0.084}_{-0.079}$ \\\\ GRB 080411    &    00309010   &    56    &  1.93e-11  &   $0.925^{+0.083}_{-0.111}$    &    $1.323^{+0.023}_{-0.023}$    &    $1.73^{+1.25}_{-0.23}$ \\\\ GRB 080413A    &    00309096   &    46    &  2.25e-13  &   $2.79^{+0.20}_{-0.19}$    &    $1.175^{+0.057}_{-0.058}$ \\\\ GRB 080413B    &    00309111   &    8.0    &  2.34e-12  &   $0.556^{+0.067}_{-0.099}$    &    $0.976^{+0.024}_{-0.025}$    &    $1.428^{+0.115}_{-0.099}$ \\\\ GRB 080426    &    00310219   &    1.7    &  1.15e-13  &   $-0^{+1}_{-1}$    &    $1.258^{+0.067}_{-0.060}$ \\\\ GRB 080430    &    00310613   &    16.2    &  2.28e-12  &   $2.22^{+0.17}_{-0.15}$    &    $0.491^{+0.033}_{-0.034}$    &    $1.132^{+0.043}_{-0.042}$ \\\\ GRB 080503    &    00310785   &    17    &  2.69e-20  &   $1.46^{+0.16}_{-0.13}$    &    $2.51^{+0.20}_{-0.32}$    &    $4.61^{+0.28}_{-0.32}$ \\\\ GRB 080506    &    00311159   &    15    &  5.24e-13  &   $-0.49^{+0.75}_{-0.67}$    &    $8.08^{+0.73}_{-0.82}$    &    $1.81^{+0.12}_{-0.20}$    &    $0.61^{+0.20}_{-0.21}$ \\\\ GRB 080507    &    00020068   &    N/A    &    &   $1.89^{+1.67}_{-0.38}$ \\\\ GRB 080514B    &    00020069   &    N/A    &  2.62e-12  &   $1.53^{+0.27}_{-0.20}$ \\\\ GRB 080515    &    00311658   &    21    &    &   $0.77^{+0.64}_{-0.36}$ \\\\ GRB 080516    &    00311762   &    5.8    &  4.49e-13  &   $0.34^{+0.16}_{-0.17}$    &    $0.99^{+0.13}_{-0.13}$ \\\\ GRB 080517    &    00311874   &    64.6    &  9.60e-14  &   $2.56^{+0.23}_{-0.21}$    &    $0.65^{+0.22}_{-0.26}$ \\\\ GRB 080520    &    00312047   &    2.8    &  8.11e-14  &   $0.99^{+0.15}_{-0.12}$ \\\\ GRB 080523    &    00312242   &    102    &  1.81e-13  &   $1.58^{+0.14}_{-0.14}$    &    $2.56^{+0.17}_{-0.14}$    &    $0.75^{+0.18}_{-0.17}$ \\\\ GRB 080602    &    00312958   &    74    &    &   $4.50^{+0.81}_{-0.86}$    &    $-0.07^{+0.16}_{-0.23}$    &    $0.70^{+0.17}_{-0.16}$ \\\\ GRB 080603A    &    00020074   &    N/A    &  1.11e-12  &   $0.952^{+0.094}_{-0.075}$ \\\\ GRB 080603B    &    00313087   &    6    &    &   $6^{+2}_{-2}$    &    $3.14^{+0.16}_{-0.16}$    &    $0.832^{+0.040}_{-0.035}$ \\\\ GRB 080604    &    00313116   &    82    &  5.34e-14  &   $2.068^{+0.053}_{-0.049}$    &    $0.18^{+0.45}_{-1.68}$ \\\\ GRB 080605    &    00313299   &    2    &  1.69e-12  &   $0.520^{+0.042}_{-0.085}$    &    $0.958^{+0.142}_{-0.051}$    &    $1.708^{+0.089}_{-0.083}$ \\\\ GRB 080607    &    00313417   &    79    &  6.56e-13  &   $8.00^{+0.64}_{-0.72}$    &    $2.215^{+0.083}_{-0.087}$    &    $1.418^{+0.023}_{-0.023}$    &    $2.04^{+7.96}_{-0.48}$ \\\\ GRB 080613A    &    00020075   &    N/A    &  1.10e-12  &   $1.29^{+1.50}_{-0.46}$ \\\\ GRB 080613B    &    00313954   &    105    &  4.32e-14  &   $2.18^{+0.11}_{-0.11}$    &    $50^{+36}_{-23}$    &    $0.80^{+0.14}_{-0.11}$ \\\\ GRB 080623    &    00315080   &    15.2    &  5.97e-13  &   $1.142^{+0.051}_{-0.046}$ \\\\ \\hline \\end{tabular} \\contcaption{} \\end{center} \\end{table*} \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lccccccccc} \\hline GRB  & Target ID & $T_{90}$  &  F$_{11}$  & $\\alpha_1$ & $\\alpha_2$ & $\\alpha_3$ & $\\alpha_4$ & $\\alpha_5$ & $\\alpha_6$ \\\\ \\hline GRB 080625    &    00020076   &    N/A    &  1.88e-12  &   $1.409^{+0.085}_{-0.169}$    &    $2.39^{+1.29}_{-0.61}$ \\\\ GRB 080701    &    00315615   &    18    &  3.52e-13  &   $1.07^{+0.22}_{-0.15}$ \\\\ GRB 080702A    &    00315710   &    0.5    &  3.31e-14  &   $1.054^{+0.116}_{-0.095}$ \\\\ GRB 080703    &    00315819   &    3.4    &  8.08e-13  &   $0.559^{+0.036}_{-0.037}$    &    $2.30^{+0.29}_{-0.24}$ \\\\ GRB 080707    &    00316204   &    27.1    &  4.68e-13  &   $4.39^{+1.10}_{-0.57}$    &    $0.21^{+0.11}_{-0.12}$    &    $1.048^{+0.098}_{-0.092}$ \\\\ GRB 080710    &    00316534   &    12    &  8.31e-13  &   $1.050^{+0.081}_{-0.063}$    &    $1.80^{+0.23}_{-0.17}$ \\\\ GRB 080714    &    00316910   &    33    &  6.44e-13  &   $0.10^{+0.64}_{-0.92}$    &    $1.87^{+2.62}_{-0.49}$    &    $1.025^{+0.049}_{-0.047}$ \\\\ GRB 080721    &    00317508   &    16.2    &  7.26e-12  &   $0.710^{+0.046}_{-0.021}$    &    $-1.10^{+1.34}_{-0.40}$    &    $0.938^{+0.020}_{-0.020}$    &    $1.639^{+0.022}_{-0.022}$ \\\\ GRB 080723A    &    00317662   &    17.3    &  9.78e-13  &   $1.82^{+0.36}_{-0.31}$    &    $0.02^{+0.12}_{-0.25}$    &    $1.060^{+0.059}_{-0.054}$ \\\\ GRB 080723B    &    00020079   &    N/A    &  6.66e-12  &   $1.40^{+0.19}_{-0.20}$ \\\\ \\hline \\end{tabular} \\contcaption{} \\end{center} \\end{table*} \\clearpage       \\begin{table*} \\begin{center} \\begin{tabular}{lcccccc} \\hline GRB  & Obs times$^*$ & $t{_{\\rm break},1}$ & $t{_{\\rm break},2}$ & $t{_{\\rm break},3}$ & $t{_{\\rm break},4}$ & $t{_{\\rm break},5}$\\\\ \\hline GRB 041223 & 16.7--28.6 ks  \\\\ GRB 050124 & 11.1--4967.5 ks  \\\\ GRB 050126 & 128 s--93.2 ks   &  $495^{+540}_{-260}$ \\\\ GRB 050128 & 105 s--99.6 ks   &  $6492^{+2247}_{-3355}$ \\\\ GRB 050215B & 5.8--3011.3 ks  \\\\ GRB 050219A & 112 s--3154.9 ks   &  $339^{+53}_{-37}$ \\\\ GRB 050219B & 3.2--3205.2 ks  \\\\ GRB 050223 & 2.9--1047.2 ks  \\\\ GRB 050315 & 83 s--948.3 ks   &  $464^{+51}_{-47}$  &  $(5.87^{+0.62}_{+0.63}) \\tim{4}$ \\\\ GRB 050318 & 3.3--832.7 ks   &  $(1.21^{+1.09}_{+0.41}) \\tim{4}$ \\\\ GRB 050319 & 234 s--2436.1 ks   &  $396^{+38}_{-33}$  &  $(1.51^{+1.10}_{+0.49}) \\tim{4}$  &  $(4^{+3}_{+1}) \\tim{5}$ \\\\ GRB 050326 & 3.3--531.4 ks  \\\\ GRB 050401 & 133 s--1066.8 ks   &  $4554^{+621}_{-603}$ \\\\ GRB 050406 & 92 s--1407.0 ks   &  $(6^{+8}_{+6}) \\tim{5}$ \\\\ GRB 050408 & 2.6--3310.6 ks  \\\\ GRB 050410 & 1.9--916.2 ks   &  $(4^{+12}_{+3}) \\tim{4}$ \\\\ GRB 050412 & 104 s--545.7 ks  \\\\ GRB 050416A & 84 s--6435.5 ks   &  $214^{+191}_{-79}$  &  $992^{+783}_{-308}$ \\\\ GRB 050421 & 116 s--486.2 ks   &  $157^{+43}_{-8}$ \\\\ GRB 050422 & 115 s--1522.3 ks   &  $280^{+186}_{-63}$ \\\\ GRB 050502B & 68 s--916.3 ks   &  $(1.34^{+0.91}_{+0.39}) \\tim{5}$ \\\\ GRB 050505 & 2.8--1207.9 ks   &  $7182^{+2571}_{-1345}$  &  $(4.37^{+1.56}_{+0.66}) \\tim{4}$ \\\\ GRB 050509A & 3.7--1285.2 ks  \\\\ GRB 050520 & 7.7--226.9 ks  \\\\ GRB 050525A & 74 s--3021.5 ks   &  $6435^{+2293}_{-6170}$ \\\\ GRB 050603 & 33.8--1800.4 ks  \\\\ GRB 050607 & 96 s--2554.5 ks   &  $356^{+83}_{-74}$  &  $807^{+132}_{-122}$  &  $(1.16^{+1.12}_{+0.61}) \\tim{4}$ \\\\ GRB 050701 & 5.8--886.2 ks  \\\\ GRB 050712 & 171 s--1833.9 ks   &  $(1.04^{+0.87}_{+0.49}) \\tim{5}$ \\\\ GRB 050713A & 77 s--1711.4 ks   &  $86^{+8}_{-2}$  &  $343^{+40}_{-17}$  &  $1100^{+1056}_{-489}$ \\\\ GRB 050713B & 141 s--1424.1 ks   &  $729^{+76}_{-61}$  &  $(1.18^{+0.19}_{+0.15}) \\tim{4}$ \\\\ GRB 050714B & 157 s--955.1 ks   &  $362^{+51}_{-37}$ \\\\ GRB 050716 & 104 s--1765.2 ks   &  $180^{+8}_{-6}$  &  $1445^{+267}_{-228}$ \\\\ GRB 050717 & 91 s--601.9 ks   &  $311^{+81}_{-95}$  &  $(2^{+2}_{+2}) \\tim{4}$ \\\\ GRB 050721 & 194 s--3385.4 ks   &  $4708^{+3710}_{-2387}$ \\\\ GRB 050724 & 79 s--731.3 ks   &  $201^{+3}_{-5}$  &  $505^{+38}_{-37}$ \\\\ GRB 050726 & 130 s--669.8 ks   &  $7101^{+1427}_{-1319}$ \\\\ GRB 050730 & 132 s--496.7 ks   &  $9090^{+267}_{-266}$ \\\\ GRB 050801 & 69 s--624.4 ks   &  $(1^{+47}_{+1}) \\tim{4}$ \\\\ GRB 050802 & 311 s--1256.3 ks   &  $6520^{+1891}_{-1013}$ \\\\ GRB 050803 & 159 s--1371.3 ks   &  $264^{+17}_{-14}$  &  $(1.60^{+0.18}_{+0.17}) \\tim{4}$ \\\\ GRB 050814 & 165 s--1077.3 ks   &  $974^{+62}_{-67}$  &  $(6^{+3}_{+2}) \\tim{4}$ \\\\ GRB 050815 & 74 s--178.3 ks   &  $3844^{+1002}_{-3222}$ \\\\ GRB 050819 & 146 s--630.1 ks   &  $429^{+88}_{-71}$ \\\\ GRB 050820A & 88 s--5029.3 ks   &  $163^{+93}_{-14}$  &  $6877^{+575}_{-506}$ \\\\ GRB 050822 & 111 s--4565.4 ks   &  $814^{+159}_{-152}$  &  $(1.68^{+0.28}_{+0.26}) \\tim{4}$ \\\\ GRB 050824 & 6.1--2713.0 ks   &  $(6^{+10}_{+2}) \\tim{4}$ \\\\ GRB 050826 & 112 s--284.2 ks   &  $2828^{+3765}_{-1592}$  &  $(4^{+5}_{+2}) \\tim{4}$ \\\\ GRB 050827 & 64.1--881.0 ks  \\\\ GRB 050904 & 168 s--857.3 ks   &  $2107^{+309700}_{-659}$ \\\\ GRB 050908 & 112 s--578.5 ks  \\\\ GRB 050915A & 92 s--475.0 ks  \\\\ GRB 050915B & 149 s--955.6 ks   &  $419^{+33}_{-27}$ \\\\ GRB 050916 & 215 s--371.0 ks   &  $274^{+633}_{-54}$ \\\\ GRB 050922B & 349 s--2361.8 ks   &  $2169^{+188}_{-165}$  &  $(2.26^{+0.85}_{+1.87}) \\tim{5}$ \\\\ GRB 050922C & 115 s--590.9 ks   &  $2001^{+1221}_{-676}$ \\\\ GRB 051001 & 193 s--550.4 ks   &  $249^{+35}_{-9}$  &  $1610^{+222}_{-183}$ \\\\ \\hline \\end{tabular} \\caption{Times of the breaks in the light curve fits. The power-law indices are given in Table~\\ref{tab:alphas}. \\newline $^*$ Time range covered by the cleaned event lists. Zero is the BAT trigger time.} \\label{tab:breaks} \\end{center} \\end{table*} \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lcccccc} \\hline GRB  & Obs times$^*$ & $t{_{\\rm break},1}$ & $t{_{\\rm break},2}$ & $t{_{\\rm break},3}$ & $t{_{\\rm break},4}$ & $t{_{\\rm break},5}$\\\\ \\hline GRB 051006 & 113 s--358.1 ks  \\\\ GRB 051008 & 3.0--542.4 ks   &  $(1.74^{+0.23}_{+0.20}) \\tim{4}$ \\\\ GRB 051016A & 120 s--931.0 ks   &  $291^{+39}_{-41}$ \\\\ GRB 051016B & 83 s--1452.0 ks   &  $152^{+12}_{-9}$  &  $7103^{+1488}_{-1523}$ \\\\ GRB 051021A & 11.2--1430.5 ks  \\\\ GRB 051021B & 85 s--516.4 ks   &  $2584^{+2972}_{-1497}$ \\\\ GRB 051022 & 12.5--1263.3 ks   &  $(2.50^{+0.61}_{+0.55}) \\tim{5}$ \\\\ GRB 051028 & 25.7--900.2 ks  \\\\ GRB 051109A & 127 s--1549.5 ks   &  $200^{+38}_{-14}$  &  $5633^{+851}_{-1247}$ \\\\ GRB 051109B & 91 s--24611.2 ks   &  $200^{+34}_{-76}$  &  $1408^{+572}_{-297}$ \\\\ GRB 051111 & 5.6--405.3 ks  \\\\ GRB 051117A & 113 s--2054.3 ks   &  $(1.23^{+0.11}_{+0.51}) \\tim{4}$  &  $(1.48^{+0.27}_{+0.50}) \\tim{4}$ \\\\ GRB 051117B & 140 s--383.6 ks  \\\\ GRB 051210 & 87 s--117.0 ks  \\\\ GRB 051211B & 10.8--862.0 ks  \\\\ GRB 051221A & 93 s--1199.9 ks   &  $3212^{+813}_{-735}$  &  $(2.42^{+0.45}_{+0.30}) \\tim{4}$ \\\\ GRB 051221B & 281 s--4640.1 ks  \\\\ GRB 051227 & 101 s--625.5 ks   &  $424^{+118}_{-69}$ \\\\ GRB 060105 & 95 s--579.7 ks   &  $191^{+32}_{-28}$  &  $2948^{+191}_{-194}$  &  $(2.67^{+0.24}_{+0.28}) \\tim{4}$  &  $(5.21^{+0.77}_{+0.43}) \\tim{4}$ \\\\ GRB 060108 & 96 s--463.1 ks   &  $541^{+421}_{-199}$  &  $(1.50^{+1.04}_{+0.40}) \\tim{4}$ \\\\ GRB 060109 & 108 s--457.9 ks   &  $347^{+31}_{-30}$  &  $7616^{+942}_{-856}$ \\\\ GRB 060110 & 242.0--492.9 ks  \\\\ GRB 060111A & 74 s--761.8 ks   &  $6227^{+1762}_{-1636}$ \\\\ GRB 060111B & 87 s--444.8 ks   &  $172^{+15}_{-30}$  &  $(1.08^{+2.28}_{+0.52}) \\tim{4}$ \\\\ GRB 060115 & 120 s--468.4 ks   &  $579^{+102}_{-66}$ \\\\ GRB 060116 & 159 s--745.6 ks  \\\\ GRB 060121 & 10.6--1042.1 ks  \\\\ GRB 060123 & 75.1--1042.6 ks  \\\\ GRB 060124 & 111 s--2619.5 ks   &  $2221^{+1044}_{-578}$  &  $(6.22^{+1.93}_{+0.99}) \\tim{4}$ \\\\ GRB 060202 & 148 s--2735.6 ks   &  $308^{+7}_{-7}$  &  $744^{+6}_{-6}$  &  $1521^{+57}_{-53}$ \\\\ GRB 060203 & 3.0--605.1 ks  \\\\ GRB 060204B & 102 s--810.5 ks   &  $261^{+29}_{-23}$  &  $7868^{+1377}_{-1176}$ \\\\ GRB 060206 & 63 s--3697.3 ks   &  $358^{+494}_{-100}$ \\\\ GRB 060210 & 102 s--1879.5 ks   &  $(3.16^{+2.18}_{+0.66}) \\tim{4}$ \\\\ GRB 060211A & 184 s--794.0 ks   &  $493^{+108}_{-58}$  &  $2571^{+986}_{-485}$ \\\\ GRB 060211B & 88 s--186.9 ks   &  $1789^{+1980}_{-844}$ \\\\ GRB 060218 & 158 s--13673.0 ks   &  $9805^{+534}_{-436}$ \\\\ GRB 060219 & 125 s--694.6 ks   &  $178^{+27}_{-19}$  &  $(2.10^{+1.68}_{+0.94}) \\tim{4}$ \\\\ GRB 060223A & 91 s--128.5 ks   &  $386^{+160}_{-125}$ \\\\ GRB 060306 & 95 s--382.4 ks   &  $179^{+25}_{-13}$  &  $8422^{+7800}_{-2542}$ \\\\ GRB 060312 & 58 s--509.3 ks   &  $303^{+33}_{-53}$ \\\\ GRB 060313 & 84 s--428.8 ks   &  $190^{+27}_{-26}$  &  $324^{+200}_{-78}$  &  $7127^{+1602}_{-2444}$ \\\\ GRB 060319 & 139 s--3796.7 ks   &  $(3^{+2}_{+1}) \\tim{4}$ \\\\ GRB 060323 & 286 s--375.1 ks  \\\\ GRB 060403 & 52 s--208.4 ks   &  $6353^{+101400}_{-5969}$ \\\\ GRB 060413 & 120 s--522.4 ks   &  $776^{+60}_{-55}$  &  $(2.415^{+0.089}_{+0.074}) \\tim{4}$ \\\\ GRB 060418 & 83 s--920.4 ks   &  $343^{+27}_{-17}$  &  $1921^{+1460}_{-697}$ \\\\ GRB 060421 & 93 s--341.7 ks  \\\\ GRB 060427 & 138 s--98.0 ks   &  $302^{+37}_{-27}$ \\\\ GRB 060428A & 76 s--3356.5 ks   &  $126^{+14}_{-11}$  &  $(6^{+2}_{+1}) \\tim{4}$  &  $(8^{+4}_{+2}) \\tim{5}$ \\\\ GRB 060428B & 211 s--1487.1 ks   &  $663^{+43}_{-35}$ \\\\ GRB 060502A & 84 s--1601.1 ks   &  $245^{+32}_{-17}$  &  $(2.59^{+1.14}_{+0.81}) \\tim{4}$ \\\\ GRB 060505 & 51.7--1448.0 ks  \\\\ GRB 060507 & 2.8--893.5 ks   &  $6868^{+6428}_{-1570}$ \\\\ GRB 060510A & 97 s--572.8 ks   &  $149^{+35}_{-18}$  &  $4650^{+724}_{-592}$ \\\\ GRB 060510B & 126 s--1091.6 ks   &  $348^{+2}_{-3}$  &  $613^{+40}_{-36}$ \\\\ GRB 060512 & 109 s--346.2 ks   &  $195^{+288}_{-43}$ \\\\ \\hline \\end{tabular} \\contcaption{$^*$ Time range covered by the cleaned event lists. Zero is the BAT trigger time.} \\end{center} \\end{table*} \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lcccccc} \\hline GRB  & Obs times$^*$ & $t{_{\\rm break},1}$ & $t{_{\\rm break},2}$ & $t{_{\\rm break},3}$ & $t{_{\\rm break},4}$ & $t{_{\\rm break},5}$\\\\ \\hline GRB 060515 & 3.4--40.6 ks  \\\\ GRB 060522 & 147 s--509.1 ks   &  $226^{+37}_{-19}$  &  $927^{+236}_{-130}$ \\\\ GRB 060526 & 81 s--545.5 ks   &  $1433^{+1060}_{-295}$  &  $6592^{+9496}_{-4265}$  &  $(1.23^{+0.52}_{+0.33}) \\tim{5}$ \\\\ GRB 060602A & 152.6--870.0 ks  \\\\ GRB 060602B & 88 s--653.7 ks  \\\\ GRB 060604 & 116 s--1920.6 ks   &  $635^{+172}_{-134}$  &  $(2.50^{+0.87}_{+1.26}) \\tim{4}$ \\\\ GRB 060605 & 98 s--190.9 ks   &  $535^{+434}_{-161}$  &  $7465^{+597}_{-632}$ \\\\ GRB 060607A & 73 s--325.5 ks   &  $548^{+85}_{-48}$  &  $(1.239^{+0.035}_{+0.041}) \\tim{4}$ \\\\ GRB 060614 & 97 s--64643.1 ks   &  $179^{+8}_{-6}$  &  $1507^{+87}_{-95}$  &  $(4.45^{+0.44}_{+0.27}) \\tim{4}$ \\\\ GRB 060707 & 126 s--2765.5 ks   &  $1263^{+949}_{-265}$  &  $(2.71^{+1.22}_{+0.95}) \\tim{4}$ \\\\ GRB 060708 & 70 s--1336.7 ks   &  $184^{+14}_{-11}$  &  $(1.28^{+0.30}_{+0.58}) \\tim{4}$ \\\\ GRB 060712 & 189 s--1047.1 ks   &  $787^{+213}_{-480}$  &  $7956^{+6926}_{-1350}$ \\\\ GRB 060714 & 106 s--1498.6 ks   &  $273^{+15}_{-11}$  &  $3593^{+2272}_{-1022}$ \\\\ GRB 060717 & 221 s--39.6 ks  \\\\ GRB 060719 & 84 s--1094.3 ks   &  $191^{+20}_{-16}$  &  $8130^{+1720}_{-1332}$ \\\\ GRB 060729 & 130 s--13063.0 ks   &  $206^{+5}_{-6}$  &  $405^{+10}_{-10}$  &  $(7.60^{+0.42}_{+0.40}) \\tim{4}$  &  $(1.00^{+0.35}_{+0.19}) \\tim{6}$ \\\\ GRB 060801 & 70 s--312.9 ks   &  $109^{+51}_{-20}$  &  $437^{+59}_{-130}$ \\\\ GRB 060804 & 127 s--312.6 ks   &  $914^{+326}_{-240}$ \\\\ GRB 060805A & 98 s--328.3 ks   &  $8730^{+3041}_{-6633}$ \\\\ GRB 060805B & 126.5--377.4 ks  \\\\ GRB 060807 & 112 s--532.3 ks   &  $195^{+41}_{-51}$  &  $4613^{+422}_{-923}$  &  $(2.93^{+1.28}_{+0.62}) \\tim{4}$ \\\\ GRB 060813 & 84 s--260.3 ks   &  $1133^{+255}_{-233}$  &  $(5.16^{+1.23}_{+0.91}) \\tim{4}$ \\\\ GRB 060814 & 77 s--1383.7 ks   &  $89^{+4}_{-6}$  &  $325^{+60}_{-112}$  &  $750^{+81}_{-45}$  &  $(1.58^{+0.61}_{+0.17}) \\tim{4}$ \\\\ GRB 060825 & 71 s--589.4 ks  \\\\ GRB 060901 & 13.8--115.9 ks  \\\\ GRB 060904A & 71 s--1292.2 ks   &  $286^{+20}_{-18}$  &  $2820^{+4728}_{-1179}$ \\\\ GRB 060904B & 76 s--675.0 ks   &  $4804^{+1185}_{-4625}$ \\\\ GRB 060906 & 153 s--365.5 ks   &  $285^{+79}_{-52}$  &  $(1.12^{+0.18}_{+0.18}) \\tim{4}$ \\\\ GRB 060908 & 79 s--1087.5 ks   &  $848^{+210}_{-301}$ \\\\ GRB 060912A & 114 s--868.4 ks   &  $741^{+1112}_{-316}$ \\\\ GRB 060919 & 92 s--313.4 ks  \\\\ GRB 060923A & 86 s--1190.5 ks   &  $350^{+74}_{-73}$  &  $2523^{+1236}_{-1019}$ \\\\ GRB 060923B & 124 s--732.9 ks  \\\\ GRB 060923C & 211 s--1382.1 ks   &  $836^{+98}_{-95}$  &  $(2.43^{+0.69}_{+1.41}) \\tim{5}$ \\\\ GRB 060926 & 65 s--284.0 ks   &  $154^{+795}_{-81}$  &  $2367^{+4739}_{-1854}$ \\\\ GRB 060927 & 70 s--208.3 ks   &  $4344^{+1406}_{-2874}$ \\\\ GRB 060928 & 477.0--866.9 ks  \\\\ GRB 060929 & 98 s--619.5 ks   &  $438^{+2994}_{-188}$ \\\\ GRB 061002 & 137 s--46.8 ks   &  $427^{+164}_{-105}$ \\\\ GRB 061004 & 68 s--533.4 ks   &  $416^{+145}_{-87}$  &  $3072^{+1843}_{-1071}$ \\\\ GRB 061006 & 148 s--370.9 ks  \\\\ GRB 061007 & 86 s--1603.8 ks   &  $357^{+140}_{-52}$  &  $4840^{+830}_{-339}$  &  $6881^{+2959}_{-964}$ \\\\ GRB 061019 & 2.8--762.0 ks   &  $(1.56^{+0.68}_{+0.83}) \\tim{5}$ \\\\ GRB 061021 & 78 s--4384.1 ks   &  $330^{+56}_{-51}$  &  $7357^{+1469}_{-1014}$ \\\\ GRB 061025 & 5.6--190.3 ks  \\\\ GRB 061028 & 204 s--157.2 ks   &  $1061^{+368}_{-202}$ \\\\ GRB 061102 & 106 s--130.9 ks   &  $1032^{+560}_{-304}$ \\\\ GRB 061110A & 75 s--907.9 ks   &  $143^{+5}_{-3}$  &  $359^{+59}_{-37}$  &  $3005^{+1133}_{-699}$ \\\\ GRB 061110B & 3.0--75.1 ks  \\\\ GRB 061121 & 61 s--2199.4 ks   &  $116^{+3}_{-3}$  &  $150^{+4}_{-4}$  &  $212^{+11}_{-8}$  &  $2425^{+437}_{-397}$  &  $(2.22^{+0.47}_{+0.34}) \\tim{4}$ \\\\ GRB 061122 & 24.5--1267.4 ks  \\\\ GRB 061126 & 1.6--2819.5 ks   &  $(4^{+7}_{+3}) \\tim{4}$ \\\\ GRB 061201 & 84 s--160.5 ks   &  $2233^{+808}_{-594}$ \\\\ GRB 061202 & 125 s--635.6 ks   &  $1109^{+133}_{-121}$  &  $(1.537^{+0.094}_{+0.083}) \\tim{4}$ \\\\ GRB 061210 & 209.2--989.9 ks  \\\\ GRB 061222A & 104 s--1538.8 ks   &  $219^{+6}_{-6}$  &  $2084^{+188}_{-518}$  &  $(5.80^{+1.78}_{+0.68}) \\tim{4}$ \\\\ GRB 061222B & 151 s--416.6 ks  \\\\ \\hline \\end{tabular} \\contcaption{$^*$ Time range covered by the cleaned event lists. Zero is the BAT trigger time.} \\end{center} \\end{table*} \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lcccccc} \\hline GRB  & Obs times$^*$ & $t{_{\\rm break},1}$ & $t{_{\\rm break},2}$ & $t{_{\\rm break},3}$ & $t{_{\\rm break},4}$ & $t{_{\\rm break},5}$\\\\ \\hline GRB 070103 & 72 s--179.9 ks   &  $616^{+55}_{-180}$  &  $4628^{+1985}_{-2918}$ \\\\ GRB 070107 & 182 s--817.1 ks   &  $(1.76^{+0.39}_{+0.28}) \\tim{4}$ \\\\ GRB 070110 & 99 s--2306.2 ks   &  $317^{+45}_{-42}$  &  $4043^{+440}_{-2940}$  &  $(2.068^{+0.021}_{+0.021}) \\tim{4}$  &  $(2.716^{+0.082}_{+0.041}) \\tim{4}$ \\\\ GRB 070125 & 46.7--1614.9 ks   &  $(10^{+4}_{+2}) \\tim{4}$ \\\\ GRB 070129 & 139 s--1637.1 ks   &  $1786^{+284}_{-179}$  &  $(1.21^{+0.17}_{+0.22}) \\tim{4}$  &  $(8^{+5}_{+4}) \\tim{5}$ \\\\ GRB 070208 & 119 s--480.4 ks   &  $923^{+351}_{-277}$  &  $(1.43^{+1.32}_{+0.59}) \\tim{4}$ \\\\ GRB 070219 & 86 s--255.0 ks   &  $627^{+1167}_{-232}$  &  $3315^{+53930}_{-1827}$ \\\\ GRB 070220 & 85 s--237.9 ks   &  $167^{+10}_{-21}$  &  $291^{+34}_{-32}$  &  $(1.147^{+0.312}_{+0.092}) \\tim{4}$ \\\\ GRB 070223 & 116 s--1081.6 ks   &  $166^{+8}_{-14}$  &  $1458^{+408}_{-496}$ \\\\ GRB 070224 & 149 s--1388.5 ks   &  $826^{+218}_{-131}$ \\\\ GRB 070227 & 55.8--1301.0 ks  \\\\ GRB 070306 & 157 s--1146.0 ks   &  $185^{+3}_{-3}$  &  $424^{+15}_{-16}$  &  $(1.89^{+0.41}_{+0.25}) \\tim{4}$  &  $(3.99^{+1.10}_{+0.59}) \\tim{4}$ \\\\ GRB 070311 & 7.0--765.9 ks   &  $(3.43^{+1.85}_{+0.68}) \\tim{4}$  &  $(1.75^{+0.17}_{+0.66}) \\tim{5}$ \\\\ GRB 070318 & 69 s--1006.2 ks  \\\\ GRB 070328 & 94 s--1088.9 ks   &  $184^{+11}_{-17}$  &  $602^{+29}_{-24}$  &  $4990^{+1117}_{-1839}$ \\\\ GRB 070330 & 71 s--416.5 ks  \\\\ GRB 070411 & 463 s--727.6 ks  \\\\ GRB 070412 & 67 s--682.8 ks   &  $105^{+15}_{-18}$  &  $(1.45^{+1.89}_{+0.91}) \\tim{4}$ \\\\ GRB 070419A & 119 s--828.0 ks   &  $216^{+11}_{-8}$  &  $2336^{+1614}_{-1014}$ \\\\ GRB 070419B & 87 s--561.0 ks   &  $(5.02^{+0.60}_{+0.55}) \\tim{4}$ \\\\ GRB 070420 & 105 s--752.6 ks   &  $215^{+10}_{-8}$  &  $4669^{+423}_{-497}$  &  $8125^{+984}_{-599}$  &  $(3.78^{+1.04}_{+0.74}) \\tim{4}$ \\\\ GRB 070429A & 159 s--1347.7 ks   &  $215^{+26}_{-10}$  &  $808^{+121}_{-103}$  &  $(5^{+4}_{+2}) \\tim{5}$ \\\\ GRB 070506 & 413 s--436.1 ks  \\\\ GRB 070508 & 82 s--758.9 ks   &  $574^{+55}_{-42}$  &  $(1.02^{+0.38}_{+0.62}) \\tim{4}$ \\\\ GRB 070509 & 71 s--161.5 ks  \\\\ GRB 070517 & 108 s--389.2 ks   &  $(1.81^{+1.02}_{+0.29}) \\tim{4}$ \\\\ GRB 070518 & 76 s--1672.8 ks   &  $1029^{+303}_{-198}$ \\\\ GRB 070520A & 164 s--1939.5 ks   &  $735^{+162}_{-113}$ \\\\ GRB 070520B & 104 s--362.9 ks   &  $1884^{+1522}_{-1475}$ \\\\ GRB 070521 & 80 s--487.4 ks   &  $333^{+111}_{-81}$  &  $9002^{+1015}_{-940}$ \\\\ GRB 070529 & 137 s--729.9 ks   &  $162^{+16}_{-14}$  &  $2533^{+3265}_{-1053}$ \\\\ GRB 070531 & 138 s--1.1 ks  \\\\ GRB 070611 & 3.3--456.8 ks  \\\\ GRB 070612B & 3.3--409.1 ks  \\\\ GRB 070616 & 137 s--371.1 ks   &  $547^{+3}_{-3}$  &  $1107^{+33}_{-36}$ \\\\ GRB 070621 & 116 s--856.2 ks   &  $349^{+31}_{-42}$  &  $1211^{+369}_{-242}$ \\\\ GRB 070628 & 114 s--83.3 ks   &  $6613^{+824}_{-638}$ \\\\ GRB 070704 & 159 s--417.0 ks   &  $590^{+239}_{-156}$ \\\\ GRB 070714A & 66 s--175.2 ks  \\\\ GRB 070714B & 67 s--134.5 ks   &  $140^{+17}_{-19}$  &  $434^{+124}_{-30}$  &  $731^{+326}_{-125}$ \\\\ GRB 070721A & 87 s--654.7 ks   &  $295^{+85}_{-66}$ \\\\ GRB 070721B & 98 s--259.5 ks   &  $147^{+7}_{-8}$  &  $8981^{+803}_{-841}$ \\\\ GRB 070724A & 73 s--565.6 ks   &  $797^{+502}_{-309}$ \\\\ GRB 070724B & 68.9--718.7 ks  \\\\ GRB 070802 & 141 s--404.5 ks   &  $507^{+127}_{-103}$  &  $7282^{+3621}_{-1821}$ \\\\ GRB 070808 & 118 s--180.0 ks   &  $289^{+55}_{-45}$ \\\\ GRB 070809 & 80 s--71.4 ks  \\\\ GRB 070810A & 91 s--39.7 ks   &  $1552^{+6800}_{-718}$ \\\\ GRB 070911 & 52.2--1960.6 ks  \\\\ GRB 070917 & 237.4--1355.0 ks  \\\\ GRB 070925 & 12.3--371.8 ks  \\\\ GRB 071001 & 88 s--4.8 ks   &  $557^{+132}_{-87}$ \\\\ GRB 071003 & 22.3--909.7 ks  \\\\ GRB 071010A & 34.0--555.6 ks   &  $(8^{+2}_{+2}) \\tim{4}$ \\\\ GRB 071010B & 6.2--156.1 ks  \\\\ GRB 071011 & 3.0--1596.0 ks   &  $(5^{+1}_{+1}) \\tim{5}$ \\\\ GRB 071020 & 67 s--1480.6 ks   &  $168^{+46}_{-27}$ \\\\ \\hline \\end{tabular} \\contcaption{$^*$ Time range covered by the cleaned event lists. Zero is the BAT trigger time.} \\end{center} \\end{table*} \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lcccccc} \\hline GRB  & Obs times$^*$ & $t{_{\\rm break},1}$ & $t{_{\\rm break},2}$ & $t{_{\\rm break},3}$ & $t{_{\\rm break},4}$ & $t{_{\\rm break},5}$\\\\ \\hline GRB 071021 & 136 s--1000.7 ks   &  $447^{+51}_{-60}$ \\\\ GRB 071025 & 149 s--497.6 ks   &  $183^{+15}_{-10}$  &  $318^{+305}_{-18}$  &  $3557^{+1494}_{-3061}$ \\\\ GRB 071028A & 124 s--87.5 ks   &  $1228^{+529}_{-277}$  &  $(2.52^{+0.80}_{+1.03}) \\tim{4}$ \\\\ GRB 071031 & 109 s--596.7 ks   &  $1608^{+230}_{-450}$  &  $(6^{+53}_{+6}) \\tim{4}$ \\\\ GRB 071101 & 72 s--92.7 ks  \\\\ GRB 071104 & 20.6--67.1 ks  \\\\ GRB 071112C & 89 s--624.4 ks   &  $243^{+71}_{-35}$  &  $881^{+176}_{-159}$ \\\\ GRB 071117 & 2.9--62.8 ks  \\\\ GRB 071118 & 129 s--121.8 ks   &  $401^{+68}_{-71}$  &  $(1.28^{+0.41}_{+0.10}) \\tim{4}$ \\\\ GRB 071122 & 146 s--1.5 ks  \\\\ GRB 071227 & 85 s--358.1 ks   &  $185^{+21}_{-7}$  &  $373^{+62}_{-39}$ \\\\ GRB 080120 & 8.5--84.5 ks  \\\\ GRB 080123 & 108 s--58.9 ks   &  $216^{+6}_{-6}$  &  $433^{+54}_{-40}$ \\\\ GRB 080129 & 3.2--837.6 ks   &  $(1.19^{+0.49}_{+0.19}) \\tim{4}$ \\\\ GRB 080205 & 107 s--221.4 ks   &  $152^{+7}_{-10}$  &  $189^{+10}_{-7}$  &  $411^{+82}_{-74}$ \\\\ GRB 080207 & 130 s--324.6 ks   &  $279^{+42}_{-34}$ \\\\ GRB 080210 & 160 s--403.8 ks  \\\\ GRB 080212 & 112 s--195.6 ks   &  $247^{+62}_{-40}$  &  $7682^{+8097}_{-1244}$ \\\\ GRB 080218B & 930 s--574.3 ks  \\\\ GRB 080229A & 96 s--540.2 ks   &  $209^{+7}_{-7}$  &  $1233^{+74}_{-130}$  &  $(3^{+8}_{+2}) \\tim{4}$ \\\\ GRB 080229B & 74.9--919.8 ks  \\\\ GRB 080303 & 78 s--983.5 ks   &  $2424^{+1094}_{-739}$ \\\\ GRB 080307 & 105 s--6349.2 ks   &  $283^{+11}_{-9}$  &  $(2.59^{+1.22}_{+0.70}) \\tim{4}$ \\\\ GRB 080310 & 95 s--832.5 ks   &  $443^{+340}_{-127}$  &  $(1.07^{+0.16}_{+0.13}) \\tim{4}$ \\\\ GRB 080319A & 565 s--117.0 ks  \\\\ GRB 080319B & 64 s--2823.7 ks   &  $474^{+135}_{-71}$  &  $(1.58^{+0.16}_{+0.14}) \\tim{4}$  &  $(4.11^{+0.44}_{+0.35}) \\tim{4}$  &  $(8^{+3}_{+1}) \\tim{5}$ \\\\ GRB 080319C & 230 s--713.8 ks   &  $353^{+38}_{-31}$  &  $5768^{+3361}_{-2584}$ \\\\ GRB 080319D & 159 s--456.9 ks   &  $988^{+434}_{-212}$ \\\\ GRB 080320 & 174 s--2290.9 ks   &  $559^{+218}_{-227}$  &  $(6^{+3}_{+1}) \\tim{4}$ \\\\ GRB 080325 & 158 s--474.4 ks   &  $775^{+320}_{-123}$ \\\\ GRB 080328 & 105 s--42.6 ks   &  $158^{+5}_{-5}$  &  $2308^{+647}_{-508}$ \\\\ GRB 080330 & 76 s--585.9 ks   &  $166^{+9}_{-11}$  &  $1011^{+65440}_{-253}$ \\\\ GRB 080405 & 155.2--283.1 ks  \\\\ GRB 080409 & 87 s--126.6 ks   &  $501^{+196}_{-129}$ \\\\ GRB 080411 & 4.3--5539.0 ks   &  $(1.38^{+0.69}_{+0.29}) \\tim{4}$  &  $(1.25^{+1.10}_{+0.40}) \\tim{6}$ \\\\ GRB 080413A & 67 s--148.7 ks   &  $164^{+22}_{-17}$ \\\\ GRB 080413B & 134 s--662.2 ks   &  $624^{+201}_{-130}$  &  $(6^{+13}_{+1}) \\tim{4}$ \\\\ GRB 080426 & 226 s--94.7 ks   &  $412^{+641}_{-92}$ \\\\ GRB 080430 & 55 s--3469.8 ks   &  $328^{+49}_{-45}$  &  $(3.32^{+0.55}_{+0.58}) \\tim{4}$ \\\\ GRB 080503 & 81 s--278.4 ks   &  $140^{+10}_{-8}$  &  $249^{+53}_{-28}$ \\\\ GRB 080506 & 146 s--174.1 ks   &  $183^{+4}_{-3}$  &  $244^{+11}_{-9}$  &  $3947^{+1697}_{-1193}$ \\\\ GRB 080507 & 109.3--214.7 ks  \\\\ GRB 080514B & 37.0--228.3 ks  \\\\ GRB 080515 & 132.8--579.8 ks  \\\\ GRB 080516 & 86 s--47.7 ks   &  $2555^{+2345}_{-1106}$ \\\\ GRB 080517 & 129 s--69.9 ks   &  $1002^{+555}_{-317}$ \\\\ GRB 080520 & 106 s--167.5 ks  \\\\ GRB 080523 & 92 s--202.9 ks   &  $192^{+20}_{-12}$  &  $2076^{+658}_{-505}$ \\\\ GRB 080602 & 117 s--1.4 ks   &  $146^{+11}_{-9}$  &  $453^{+132}_{-110}$ \\\\ GRB 080603A & 10.4--606.4 ks  \\\\ GRB 080603B & 68 s--12.3 ks   &  $76^{+10}_{-3}$  &  $160^{+7}_{-10}$ \\\\ GRB 080604 & 125 s--441.2 ks   &  $(1.35^{+4.44}_{+0.65}) \\tim{4}$ \\\\ GRB 080605 & 96 s--949.1 ks   &  $274^{+112}_{-37}$  &  $4180^{+1267}_{-1015}$ \\\\ GRB 080607 & 85 s--376.5 ks   &  $102^{+2}_{-2}$  &  $334^{+27}_{-19}$  &  $(10^{+16}_{+5}) \\tim{4}$ \\\\ GRB 080613A & 24.2--146.8 ks  \\\\ GRB 080613B & 75 s--214.1 ks   &  $217^{+4}_{-3}$  &  $237^{+14}_{-8}$ \\\\ GRB 080623 & 1.6--610.6 ks  \\\\ \\hline \\end{tabular} \\contcaption{$^*$ Time range covered by the cleaned event lists. Zero is the BAT trigger time.} \\end{center} \\end{table*} \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lcccccc} \\hline GRB  & Obs times$^*$ & $t{_{\\rm break},1}$ & $t{_{\\rm break},2}$ & $t{_{\\rm break},3}$ & $t{_{\\rm break},4}$ & $t{_{\\rm break},5}$\\\\ \\hline GRB 080625 & 17.9--818.1 ks   &  $(1.67^{+0.88}_{+1.09}) \\tim{5}$ \\\\ GRB 080701 & 2.9--131.2 ks  \\\\ GRB 080702A & 74 s--121.6 ks  \\\\ GRB 080703 & 103 s--180.9 ks   &  $(2.14^{+0.23}_{+0.23}) \\tim{4}$ \\\\ GRB 080707 & 74 s--416.6 ks   &  $184^{+18}_{-26}$  &  $7703^{+2997}_{-6696}$ \\\\ GRB 080710 & 3.1--646.9 ks   &  $(2.21^{+0.55}_{+0.52}) \\tim{4}$ \\\\ GRB 080714 & 86 s--450.2 ks   &  $126^{+71}_{-27}$  &  $287^{+158}_{-72}$ \\\\ GRB 080721 & 113 s--1405.7 ks   &  $290^{+9}_{-27}$  &  $306^{+12}_{-9}$  &  $3990^{+346}_{-323}$ \\\\ GRB 080723A & 96 s--573.3 ks   &  $380^{+92}_{-71}$  &  $2880^{+1151}_{-613}$ \\\\ GRB 080723B & 13.4--40.6 ks  \\\\ \\hline \\end{tabular} \\contcaption{$^*$ Time range covered by the cleaned event lists. Zero is the BAT trigger time.} \\end{center} \\end{table*} \\clearpage       \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lcccc|cc} \\hline &                  &            & \\multicolumn{2}{c}{WT Mode} & \\multicolumn{2}{c}{PC Mode} \\\\ GRB  &  Galactic $N_{\\rm H}$  &  redshift  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$ \\\\ \\hline GRB 041218  &  27.4  &    &  $21.6^{+133.6}_{-21.6}$  &  $2.3^{+5.4}_{-2.0}$  &  $<50.27$  &  $0.55^{+1.23}_{-0.69}$  \\\\ GRB 041223  &  10.1  &    &    &  &  $3.1^{+4.5}_{-3.1}$  &  $1.13^{+0.19}_{-0.16}$  \\\\ GRB 050124  &  2.66  &    &  $0.96^{+8.00}_{-0.96}$  &  $0.54^{+0.35}_{-0.24}$  &  $2.1^{+2.6}_{-2.0}$  &  $0.94^{+0.21}_{-0.20}$  \\\\ GRB 050126  &  4.62  &  1.29  &    &  &  $<11.50$  &  $1.14^{+0.22}_{-0.15}$  \\\\ GRB 050128  &  5.19  &    &  $<125.82$  &  $0.41^{+1.80}_{-0.45}$  &  $4.9^{+2.0}_{-1.9}$  &  $0.968^{+0.087}_{-0.044}$  \\\\ GRB 050215B  &  1.85  &    &    &  &  $<8.82$  &  $0.65^{+0.59}_{-0.28}$  \\\\ GRB 050219A  &  9.50  &    &  $17.4^{+4.7}_{-4.4}$  &  $1.16^{+0.16}_{-0.15}$  &  $4.8^{+11.8}_{-4.8}$  &  $0.66^{+0.37}_{-0.32}$  \\\\ GRB 050219B  &  2.98  &    &  $21.9^{+5.0}_{-4.6}$  &  $1.09^{+0.18}_{-0.17}$  &  $12.3^{+4.2}_{-4.9}$  &  $1.07^{+0.19}_{-0.20}$  \\\\ GRB 050223  &  7.08  &  0.5915  &    &  &  $<16.31$  &  $0.87^{+0.50}_{-0.39}$  \\\\ GRB 050315  &  3.69  &  1.949  &  $117.8^{+87.2}_{-64.0}$  &  $1.11^{+0.35}_{-0.29}$  &  $76.2^{+18.0}_{-17.0}$  &  $1.021^{+0.042}_{-0.079}$  \\\\ GRB 050318  &  1.86  &  1.44  &  $15.3^{+137.9}_{-15.3}$  &  $1.48^{+3.13}_{-0.71}$  &  $11.0^{+8.3}_{-6.1}$  &  $0.992^{+0.080}_{-0.087}$  \\\\ GRB 050319  &  1.26  &  3.24  &    &  &  $10.2^{+22.9}_{-10.2}$  &  $0.998^{+0.066}_{-0.048}$  \\\\ GRB 050326  &  4.16  &    &  $6.6^{+6.9}_{-6.0}$  &  $1.05^{+0.33}_{-0.30}$  &  $13.5^{+7.2}_{-2.4}$  &  $1.01^{+0.14}_{-0.14}$  \\\\ GRB 050401  &  4.40  &  2.9  &  $155.6^{+19.2}_{-18.3}$  &  $0.895^{+0.041}_{-0.040}$  &  $139.3^{+64.9}_{-59.3}$  &  $0.79^{+0.11}_{-0.11}$  \\\\ GRB 050406  &  1.86  &    &  $<3.43$  &  $1.27^{+0.30}_{-0.16}$  &  $7.2^{+30.7}_{-7.2}$  &  $0.96^{+0.69}_{-0.81}$  \\\\ GRB 050408  &  1.73  &  1.236  &    &  &  $125.2^{+24.0}_{-19.9}$  &  $1.19^{+0.13}_{-0.12}$  \\\\ GRB 050410  &  6.95  &    &  $8.8^{+17.4}_{-8.8}$  &  $1.21^{+0.69}_{-0.55}$  &  $3.3^{+35.7}_{-3.3}$  &  $0.71^{+0.41}_{-0.31}$  \\\\ GRB 050412  &  1.92  &    &  $26.8^{+47.1}_{-23.3}$  &  $0.72^{+0.87}_{-0.18}$  &  $1.5^{+9.8}_{-1.5}$  &  $0.62^{+0.36}_{-0.29}$  \\\\ GRB 050416A  &  2.45  &  0.6535  &  $53.8^{+56.4}_{-43.6}$  &  $2.4^{+1.6}_{-1.4}$  &  $54.1^{+5.8}_{-7.8}$  &  $1.053^{+0.071}_{-0.052}$  \\\\ GRB 050421  &  15.3  &    &  $70.8^{+29.3}_{-23.4}$  &  $0.40^{+0.29}_{-0.27}$  &  $53.5^{+28.0}_{-20.7}$  &  $1.39^{+0.52}_{-0.40}$  \\\\ GRB 050422  &  91.1  &    &  $10.0^{+52.3}_{-10.0}$  &  $2.07^{+1.01}_{-0.49}$  &  $<22.79$  &  $1.39^{+0.46}_{-0.41}$  \\\\ GRB 050502B  &  3.59  &    &  $7.05^{+0.67}_{-0.66}$  &  $1.245^{+0.034}_{-0.034}$  &  $0.82^{+2.06}_{-0.82}$  &  $0.945^{+0.077}_{-0.100}$  \\\\ GRB 050504  &  1.11  &    &    &  &  $23.4^{+35.0}_{-23.4}$  &  $1.5^{+1.6}_{-1.3}$  \\\\ GRB 050505  &  1.71  &  4.27  &    &  &  $175.1^{+41.0}_{-31.4}$  &  $1.020^{+0.034}_{-0.053}$  \\\\ GRB 050509A  &  30.3  &    &    &  &  $23.0^{+45.4}_{-19.8}$  &  $1.06^{+0.79}_{-0.65}$  \\\\ GRB 050509B  &  1.55  &  0.225  &    &  &  $8.0^{+8.1}_{-8.0}$  &  $0.92^{+1.09}_{-0.60}$  \\\\ GRB 050520  &  1.26  &    &    &  &  $150.4^{+222.2}_{-75.2}$  &  $-0.4^{+1.0}_{-1.0}$  \\\\ GRB 050525A  &  9.07  &  0.606  &    &  &  $14.6^{+8.5}_{-6.8}$  &  $1.04^{+0.14}_{-0.12}$  \\\\ GRB 050603  &  2.23  &  2.821  &    &  &  $44.9^{+22.3}_{-29.0}$  &  $1.06^{+0.10}_{-0.11}$  \\\\ GRB 050607  &  12.6  &    &    &  &  $9.3^{+2.1}_{-4.6}$  &  $1.302^{+0.085}_{-0.175}$  \\\\ GRB 050701  &  150.0  &    &    &  &  $276.2^{+161.7}_{-132.9}$  &  $1.45^{+0.77}_{-0.35}$  \\\\ GRB 050712  &  12.3  &    &  $8.4^{+3.2}_{-3.0}$  &  $1.28^{+0.14}_{-0.13}$  &  $6.6^{+4.0}_{-1.9}$  &  $1.15^{+0.14}_{-0.16}$  \\\\ GRB 050713A  &  10.7  &    &  $23.7^{+2.2}_{-2.1}$  &  $1.173^{+0.066}_{-0.064}$  &  $26.0^{+5.2}_{-4.9}$  &  $1.14^{+0.14}_{-0.13}$  \\\\ GRB 050713B  &  19.5  &    &  $12.8^{+4.1}_{-3.9}$  &  $0.652^{+0.099}_{-0.097}$  &  $24.6^{+5.9}_{-5.5}$  &  $1.08^{+0.14}_{-0.13}$  \\\\ GRB 050714  &  32.9  &    &    &  &  $170.3^{+173.5}_{-145.2}$  &  $1.5^{+1.5}_{-1.3}$  \\\\ GRB 050714B  &  4.63  &    &  $77.9^{+12.6}_{-11.1}$  &  $5.21^{+0.69}_{-0.61}$  &  $28.9^{+6.5}_{-7.2}$  &  $2.64^{+0.38}_{-0.37}$  \\\\ GRB 050716  &  8.42  &    &  $2.2^{+1.8}_{-1.7}$  &  $0.306^{+0.057}_{-0.057}$  &  $8.1^{+1.8}_{-4.0}$  &  $1.08^{+0.12}_{-0.16}$  \\\\ GRB 050717  &  20.2  &    &  $19.8^{+5.4}_{-5.0}$  &  $0.480^{+0.097}_{-0.095}$  &  $16.7^{+9.6}_{-8.9}$  &  $0.76^{+0.16}_{-0.20}$  \\\\ GRB 050721  &  14.5  &    &  $11.2^{+5.2}_{-4.8}$  &  $0.61^{+0.14}_{-0.13}$  &  $14.8^{+7.2}_{-3.4}$  &  $0.86^{+0.18}_{-0.16}$  \\\\ GRB 050724  &  14.0  &  0.257  &  $43.4^{+3.7}_{-3.6}$  &  $0.663^{+0.042}_{-0.042}$  &  $17.0^{+8.2}_{-8.2}$  &  $0.81^{+0.14}_{-0.17}$  \\\\ GRB 050726  &  4.95  &    &  $4.8^{+2.6}_{-2.5}$  &  $1.20^{+0.14}_{-0.14}$  &  $1.4^{+2.2}_{-1.4}$  &  $1.024^{+0.112}_{-0.061}$  \\\\ GRB 050730  &  3.01  &  3.97  &  $52.2^{+22.6}_{-21.8}$  &  $0.598^{+0.033}_{-0.032}$  &  $20.5^{+25.1}_{-20.5}$  &  $0.615^{+0.035}_{-0.035}$  \\\\ GRB 050801  &  6.34  &    &    &  &  $<1.96$  &  $0.89^{+0.18}_{-0.12}$  \\\\ GRB 050802  &  1.85  &    &  $3.9^{+3.4}_{-3.1}$  &  $0.99^{+0.19}_{-0.18}$  &  $3.57^{+1.07}_{-0.62}$  &  $0.918^{+0.036}_{-0.062}$  \\\\ GRB 050803  &  5.04  &    &  $13.5^{+4.5}_{-4.2}$  &  $0.81^{+0.17}_{-0.15}$  &  $14.6^{+2.4}_{-2.4}$  &  $1.058^{+0.091}_{-0.089}$  \\\\ GRB 050813  &  4.45  &    &    &  &  $6.3^{+41.0}_{-6.3}$  &  $1.5^{+0.0}_{-1.6}$  \\\\ GRB 050814  &  2.28  &  5.3  &  $78.8^{+62.2}_{-59.3}$  &  $1.103^{+0.069}_{-0.067}$  &  $<55.21$  &  $0.922^{+0.074}_{-0.062}$  \\\\ GRB 050815  &  23.3  &    &    &  &  $4.1^{+10.8}_{-4.1}$  &  $0.61^{+0.14}_{-0.22}$  \\\\ GRB 050819  &  4.61  &    &  $<5.08$  &  $1.35^{+0.34}_{-0.17}$  &  $1.8^{+3.0}_{-1.8}$  &  $1.11^{+0.20}_{-0.22}$  \\\\ GRB 050820A  &  4.41  &  2.612  &  $<5.96$  &  $0.033^{+0.042}_{-0.027}$  &  $35.9^{+15.2}_{-14.4}$  &  $0.966^{+0.051}_{-0.050}$  \\\\ GRB 050820B  &  8.52  &  2.612  &    &  &  $50.5^{+371.6}_{-50.5}$  &  $0.65^{+0.51}_{-0.38}$  \\\\ GRB 050822  &  1.41  &    &  $13.0^{+1.4}_{-1.3}$  &  $1.583^{+0.074}_{-0.070}$  &  $6.6^{+2.1}_{-1.1}$  &  $1.139^{+0.076}_{-0.104}$  \\\\ GRB 050824  &  3.70  &  0.83  &    &  &  $3.7^{+8.2}_{-3.7}$  &  $0.93^{+0.14}_{-0.14}$  \\\\ GRB 050826  &  19.0  &  0.3  &    &  &  $86.4^{+17.1}_{-28.0}$  &  $1.26^{+0.18}_{-0.29}$  \\\\ GRB 050827  &  16.2  &    &    &  &  $22.3^{+5.1}_{-9.5}$  &  $0.873^{+0.095}_{-0.234}$  \\\\ GRB 050904  &  4.53  &  6.29  &  $375.5^{+95.2}_{-91.7}$  &  $0.351^{+0.034}_{-0.034}$  &  $261.9^{+96.9}_{-94.1}$  &  $0.908^{+0.049}_{-0.047}$  \\\\ \\hline \\end{tabular} \\caption{Best-fit results for time-averaged spectra. Column densities are in units of $10^{20}$ \\cms. WT and PC mode were fitted independently.} \\label{tab:avespec} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lcccc|cc} \\hline &                  &            & \\multicolumn{2}{c}{WT Mode} & \\multicolumn{2}{c}{PC Mode} \\\\ GRB  &  Galactic $N_{\\rm H}$  &  redshift  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$ \\\\ \\hline GRB 050908  &  2.38  &  3.35  &    &  &  $0.65^{+44.49}_{-0.65}$  &  $1.91^{+0.24}_{-0.13}$  \\\\ GRB 050915A  &  1.97  &    &  $11.5^{+5.7}_{-5.2}$  &  $1.18^{+0.27}_{-0.25}$  &  $9.2^{+3.5}_{-4.5}$  &  $1.29^{+0.16}_{-0.22}$  \\\\ GRB 050915B  &  22.2  &    &  $9.7^{+4.9}_{-4.5}$  &  $1.68^{+0.20}_{-0.19}$  &  $18.4^{+7.8}_{-10.3}$  &  $1.36^{+0.19}_{-0.27}$  \\\\ GRB 050916  &  81.0  &    &    &  &  $13.5^{+19.4}_{-12.2}$  &  $0.64^{+0.16}_{-0.23}$  \\\\ GRB 050922B  &  3.14  &    &  $13.3^{+1.6}_{-1.5}$  &  $1.700^{+0.086}_{-0.083}$  &  $13.4^{+2.7}_{-2.5}$  &  $1.58^{+0.13}_{-0.12}$  \\\\ GRB 050922C  &  5.37  &  2.198  &  $17.3^{+14.4}_{-13.7}$  &  $1.028^{+0.073}_{-0.072}$  &  $43.7^{+18.3}_{-10.6}$  &  $1.287^{+0.084}_{-0.077}$  \\\\ GRB 051001  &  1.15  &    &  $12.4^{+1.9}_{-1.8}$  &  $0.443^{+0.061}_{-0.060}$  &  $15.4^{+2.8}_{-6.3}$  &  $1.26^{+0.25}_{-0.27}$  \\\\ GRB 051006  &  8.10  &    &  $46.4^{+21.6}_{-17.5}$  &  $0.47^{+0.31}_{-0.29}$  &  $73.2^{+26.6}_{-21.8}$  &  $1.64^{+0.40}_{-0.40}$  \\\\ GRB 051008  &  0.959  &    &    &  &  $31.5^{+3.1}_{-3.5}$  &  $0.930^{+0.047}_{-0.088}$  \\\\ GRB 051012  &  10.0  &    &    &  &  $1.6^{+38.5}_{-1.6}$  &  $0.58^{+2.11}_{-0.74}$  \\\\ GRB 051016A  &  12.2  &    &  $27.4^{+21.7}_{-17.9}$  &  $2.00^{+0.85}_{-0.72}$  &  $36.6^{+7.6}_{-12.7}$  &  $2.35^{+0.27}_{-0.46}$  \\\\ GRB 051016B  &  3.20  &  0.9364  &  $85.2^{+33.4}_{-28.3}$  &  $4.01^{+0.98}_{-0.80}$  &  $55.4^{+6.1}_{-14.0}$  &  $0.937^{+0.055}_{-0.133}$  \\\\ GRB 051021A  &  4.61  &    &    &  &  $5.3^{+3.6}_{-5.3}$  &  $1.09^{+0.24}_{-0.26}$  \\\\ GRB 051021B  &  71.0  &    &  $<31.14$  &  $0.30^{+0.41}_{-0.27}$  &  $12.0^{+33.0}_{-12.0}$  &  $0.68^{+0.33}_{-0.31}$  \\\\ GRB 051022  &  3.76  &  0.8  &    &  &  $315.1^{+35.9}_{-33.8}$  &  $1.04^{+0.12}_{-0.11}$  \\\\ GRB 051028  &  11.5  &    &    &  &  $13.8^{+4.4}_{-8.6}$  &  $1.22^{+0.17}_{-0.31}$  \\\\ GRB 051109A  &  16.1  &  2.346  &  $21.5^{+48.7}_{-21.5}$  &  $1.13^{+0.16}_{-0.15}$  &  $49.6^{+29.5}_{-27.2}$  &  $1.062^{+0.079}_{-0.077}$  \\\\ GRB 051109B  &  11.5  &    &    &  &  $13.3^{+3.6}_{-6.8}$  &  $1.17^{+0.23}_{-0.23}$  \\\\ GRB 051111  &  5.26  &  1.55  &    &  &  $60.0^{+19.1}_{-23.6}$  &  $1.25^{+0.10}_{-0.15}$  \\\\ GRB 051117A  &  1.78  &    &  $11.52^{+0.42}_{-0.41}$  &  $1.075^{+0.018}_{-0.018}$  &  $9.2^{+2.8}_{-2.6}$  &  $1.277^{+0.136}_{-0.072}$  \\\\ GRB 051117B  &  4.67  &    &    &  &  $21.7^{+16.6}_{-15.6}$  &  $0.68^{+0.46}_{-0.34}$  \\\\ GRB 051210  &  1.91  &    &  $2.5^{+5.4}_{-2.5}$  &  $0.28^{+0.21}_{-0.18}$  &  $44.6^{+16.0}_{-29.4}$  &  $2.26^{+0.63}_{-0.92}$  \\\\ GRB 051211B  &  32.0  &    &    &  &  $6.8^{+6.3}_{-6.8}$  &  $1.04^{+0.11}_{-0.25}$  \\\\ GRB 051221A  &  5.66  &  0.5465  &    &  &  $15.9^{+3.8}_{-3.1}$  &  $1.059^{+0.087}_{-0.065}$  \\\\ GRB 051221B  &  46.1  &    &  $33.7^{+904.9}_{-33.7}$  &  $-0.02^{+4.05}_{-0.73}$  &  $87.0^{+135.0}_{-49.1}$  &  $1.06^{+1.17}_{-0.80}$  \\\\ GRB 051227  &  3.35  &    &  $18.0^{+8.8}_{-7.7}$  &  $0.40^{+0.22}_{-0.21}$  &  $4.6^{+6.6}_{-4.6}$  &  $0.99^{+0.33}_{-0.31}$  \\\\ GRB 060105  &  14.8  &    &  $18.34^{+0.99}_{-0.98}$  &  $0.946^{+0.027}_{-0.027}$  &  $20.4^{+2.5}_{-2.3}$  &  $1.085^{+0.079}_{-0.076}$  \\\\ GRB 060108  &  1.79  &  2.03  &  $<143.21$  &  $0.53^{+1.42}_{-0.96}$  &  $39.4^{+25.5}_{-27.6}$  &  $0.96^{+0.10}_{-0.13}$  \\\\ GRB 060109  &  10.4  &    &  $17.2^{+4.8}_{-4.5}$  &  $1.21^{+0.17}_{-0.16}$  &  $22.8^{+2.2}_{-5.0}$  &  $1.33^{+0.11}_{-0.15}$  \\\\ GRB 060110  &  21.8  &    &    &  &  $4.1^{+30.9}_{-4.1}$  &  $0.80^{+0.84}_{-0.58}$  \\\\ GRB 060111A  &  2.93  &    &  $18.63^{+0.97}_{-0.95}$  &  $1.267^{+0.039}_{-0.038}$  &  $13.2^{+4.7}_{-4.8}$  &  $1.21^{+0.20}_{-0.19}$  \\\\ GRB 060111B  &  7.66  &    &  $11.9^{+4.9}_{-4.5}$  &  $1.20^{+0.21}_{-0.19}$  &  $21.9^{+4.8}_{-5.7}$  &  $1.30^{+0.16}_{-0.18}$  \\\\ GRB 060115  &  9.48  &  3.53  &  $135.7^{+59.1}_{-55.6}$  &  $0.852^{+0.088}_{-0.084}$  &  $<42.08$  &  $1.124^{+0.099}_{-0.083}$  \\\\ GRB 060116  &  17.6  &    &  $49.7^{+273.6}_{-49.7}$  &  $0.7^{+2.5}_{-2.1}$  &  $51.2^{+7.8}_{-16.2}$  &  $0.89^{+0.16}_{-0.26}$  \\\\ GRB 060121  &  1.38  &    &    &  &  $8.3^{+4.5}_{-4.8}$  &  $1.28^{+0.24}_{-0.24}$  \\\\ GRB 060123  &  1.27  &    &    &  &  $10.5^{+8.4}_{-6.9}$  &  $0.75^{+0.31}_{-0.25}$  \\\\ GRB 060124  &  8.98  &  2.297  &  $37.2^{+7.3}_{-7.1}$  &  $0.423^{+0.019}_{-0.019}$  &  $64.3^{+16.4}_{-15.5}$  &  $1.021^{+0.054}_{-0.052}$  \\\\ GRB 060202  &  4.40  &    &  $39.7^{+1.2}_{-1.2}$  &  $1.159^{+0.031}_{-0.030}$  &  $41.6^{+2.3}_{-4.4}$  &  $1.747^{+0.061}_{-0.139}$  \\\\ GRB 060203  &  6.35  &    &    &  &  $7.1^{+2.3}_{-4.1}$  &  $1.20^{+0.18}_{-0.18}$  \\\\ GRB 060204B  &  1.66  &    &  $19.5^{+1.9}_{-1.8}$  &  $0.993^{+0.067}_{-0.066}$  &  $12.2^{+1.8}_{-1.6}$  &  $1.281^{+0.065}_{-0.099}$  \\\\ GRB 060206  &  9.33e-01  &  4.05  &  $<271.78$  &  $1.46^{+1.47}_{-0.84}$  &  $117.3^{+145.0}_{-117.3}$  &  $1.21^{+0.27}_{-0.24}$  \\\\ GRB 060210  &  6.08  &  3.91  &  $227.3^{+24.6}_{-24.0}$  &  $0.934^{+0.031}_{-0.031}$  &  $179.8^{+36.4}_{-34.5}$  &  $1.079^{+0.049}_{-0.048}$  \\\\ GRB 060211A  &  11.3  &    &  $7.4^{+1.9}_{-1.8}$  &  $0.942^{+0.072}_{-0.069}$  &  $6.7^{+5.7}_{-5.8}$  &  $0.96^{+0.12}_{-0.22}$  \\\\ GRB 060211B  &  17.0  &    &  $2.6^{+16.8}_{-2.6}$  &  $1.25^{+0.80}_{-0.39}$  &  $5.5^{+5.6}_{-5.5}$  &  $0.75^{+0.15}_{-0.28}$  \\\\ GRB 060218  &  9.37  &  0.0331  &  $26.35^{+0.45}_{-0.43}$  &  $0.7478^{+0.0102}_{-0.0099}$  &  $34.6^{+3.1}_{-4.9}$  &  $2.56^{+0.16}_{-0.18}$  \\\\ GRB 060219  &  2.34  &    &    &  &  $35.1^{+9.5}_{-8.8}$  &  $2.19^{+0.38}_{-0.33}$  \\\\ GRB 060223A  &  6.91  &  4.41  &    &  &  $273.9^{+146.4}_{-190.5}$  &  $1.06^{+0.17}_{-0.19}$  \\\\ GRB 060306  &  3.44  &    &  $37.9^{+10.9}_{-9.8}$  &  $1.15^{+0.29}_{-0.27}$  &  $36.8^{+2.2}_{-4.4}$  &  $1.275^{+0.053}_{-0.120}$  \\\\ GRB 060312  &  11.6  &    &  $10.2^{+2.5}_{-2.4}$  &  $1.172^{+0.097}_{-0.094}$  &  $<2.17$  &  $0.848^{+0.119}_{-0.086}$  \\\\ GRB 060313  &  4.93  &    &  $<7.85$  &  $0.74^{+0.40}_{-0.19}$  &  $<0.92$  &  $0.949^{+0.075}_{-0.074}$  \\\\ GRB 060319  &  1.46  &    &    &  &  $43.6^{+5.1}_{-5.2}$  &  $1.254^{+0.099}_{-0.131}$  \\\\ GRB 060322  &  10.7  &    &    &  &  $17.6^{+15.1}_{-17.6}$  &  $2.35^{+2.47}_{-0.96}$  \\\\ GRB 060323  &  1.48  &    &  $3.8^{+9.8}_{-3.8}$  &  $1.02^{+0.45}_{-0.31}$  &  $2.0^{+2.9}_{-2.0}$  &  $0.99^{+0.30}_{-0.24}$  \\\\ GRB 060403  &  44.0  &    &  $18.6^{+15.5}_{-13.8}$  &  $0.37^{+0.19}_{-0.18}$  &  $15.6^{+30.1}_{-15.6}$  &  $1.32^{+0.56}_{-0.68}$  \\\\ GRB 060413  &  108.0  &    &  $127.0^{+16.0}_{-15.6}$  &  $0.89^{+0.10}_{-0.10}$  &  $80.3^{+18.9}_{-18.2}$  &  $0.38^{+0.12}_{-0.11}$  \\\\ GRB 060418  &  8.81  &  1.49  &  $57.4^{+6.2}_{-5.9}$  &  $1.129^{+0.041}_{-0.040}$  &  $37.5^{+25.4}_{-22.6}$  &  $0.96^{+0.15}_{-0.14}$  \\\\ GRB 060421  &  90.0  &    &    &  &  $34.2^{+42.2}_{-34.2}$  &  $0.58^{+0.15}_{-0.46}$  \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lcccc|cc} \\hline &                  &            & \\multicolumn{2}{c}{WT Mode} & \\multicolumn{2}{c}{PC Mode} \\\\ GRB  &  Galactic $N_{\\rm H}$  &  redshift  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$ \\\\ \\hline GRB 060427  &  3.85  &    &  $10.3^{+3.0}_{-2.8}$  &  $1.78^{+0.18}_{-0.17}$  &  $12.0^{+4.7}_{-4.9}$  &  $1.42^{+0.21}_{-0.22}$  \\\\ GRB 060428A  &  76.7  &    &  $<16.19$  &  $2.30^{+0.53}_{-0.26}$  &  $18.6^{+12.6}_{-11.7}$  &  $1.03^{+0.15}_{-0.15}$  \\\\ GRB 060428B  &  1.33  &    &  $6.5^{+1.5}_{-1.4}$  &  $1.86^{+0.11}_{-0.10}$  &  $<1.28$  &  $1.040^{+0.115}_{-0.090}$  \\\\ GRB 060501  &  22.8  &    &    &  &  $30.0^{+99.4}_{-30.0}$  &  $1.62^{+2.21}_{-0.57}$  \\\\ GRB 060502A  &  3.45  &  1.51  &  $53.8^{+16.8}_{-15.3}$  &  $2.48^{+0.27}_{-0.25}$  &  $24.1^{+13.5}_{-12.8}$  &  $1.034^{+0.115}_{-0.058}$  \\\\ GRB 060502B  &  4.24  &    &    &  &  $7.6^{+22.3}_{-7.6}$  &  $1.10^{+2.77}_{-0.81}$  \\\\ GRB 060505  &  1.72  &    &    &  &  $8.1^{+19.1}_{-8.1}$  &  $0.65^{+0.69}_{-0.55}$  \\\\ GRB 060507  &  8.84  &    &    &  &  $4.8^{+2.1}_{-3.9}$  &  $1.08^{+0.14}_{-0.16}$  \\\\ GRB 060510A  &  36.1  &    &  $<5.19$  &  $2.35^{+0.38}_{-0.32}$  &  $<3.00$  &  $0.932^{+0.079}_{-0.043}$  \\\\ GRB 060510B  &  4.07  &  4.9  &  $321.2^{+40.3}_{-39.8}$  &  $0.239^{+0.021}_{-0.022}$  &  $145.5^{+120.7}_{-72.8}$  &  $1.19^{+0.17}_{-0.16}$  \\\\ GRB 060512  &  1.53  &  0.4428  &  $30.3^{+17.5}_{-15.1}$  &  $4.54^{+1.13}_{-0.96}$  &  $<2.31$  &  $1.80^{+0.11}_{-0.11}$  \\\\ GRB 060515  &  2.44  &    &    &  &  $<6.77$  &  $0.81^{+0.52}_{-0.33}$  \\\\ GRB 060522  &  4.10  &  5.11  &  $143.3^{+333.8}_{-143.3}$  &  $0.56^{+0.23}_{-0.20}$  &  $210.7^{+192.9}_{-138.5}$  &  $1.14^{+0.13}_{-0.11}$  \\\\ GRB 060526  &  5.02  &  3.21  &  $55.0^{+15.4}_{-14.8}$  &  $0.892^{+0.035}_{-0.035}$  &  $<28.22$  &  $0.931^{+0.086}_{-0.084}$  \\\\ GRB 060602A  &  2.48  &    &    &  &  $59.3^{+39.3}_{-31.8}$  &  $4.2^{+0.0}_{-2.4}$  \\\\ GRB 060602B  &  121.0  &    &    &  &  $396.6^{+138.7}_{-103.2}$  &  $2.29^{+0.67}_{-0.56}$  \\\\ GRB 060604  &  4.09  &  2.68  &  $159.3^{+18.0}_{-17.1}$  &  $1.172^{+0.050}_{-0.050}$  &  $71.1^{+38.6}_{-33.6}$  &  $1.17^{+0.13}_{-0.12}$  \\\\ GRB 060605  &  3.95  &  3.8  &  $0.053^{+150.957}_{-0.053}$  &  $0.63^{+0.41}_{-0.30}$  &  $31.5^{+37.6}_{-31.5}$  &  $1.000^{+0.054}_{-0.043}$  \\\\ GRB 060607A  &  2.42  &  3.082  &  $27.4^{+12.1}_{-11.7}$  &  $0.702^{+0.031}_{-0.031}$  &  $49.2^{+25.7}_{-24.7}$  &  $0.654^{+0.056}_{-0.054}$  \\\\ GRB 060614  &  1.87  &    &  $1.72^{+0.64}_{-0.62}$  &  $0.897^{+0.035}_{-0.035}$  &  $1.45^{+1.15}_{-0.61}$  &  $0.890^{+0.064}_{-0.043}$  \\\\ GRB 060707  &  1.44  &  3.43  &  $<56.85$  &  $0.89^{+0.23}_{-0.22}$  &  $3.1^{+51.6}_{-3.1}$  &  $1.071^{+0.132}_{-0.094}$  \\\\ GRB 060708  &  2.10  &    &  $11.9^{+2.8}_{-2.6}$  &  $2.27^{+0.19}_{-0.18}$  &  $3.0^{+1.1}_{-1.8}$  &  $1.218^{+0.079}_{-0.101}$  \\\\ GRB 060712  &  1.66  &    &    &  &  $15.3^{+2.3}_{-4.8}$  &  $1.81^{+0.13}_{-0.23}$  \\\\ GRB 060714  &  6.05  &  2.71  &  $138.8^{+19.8}_{-18.8}$  &  $0.810^{+0.043}_{-0.043}$  &  $93.0^{+37.5}_{-35.5}$  &  $1.024^{+0.101}_{-0.096}$  \\\\ GRB 060717  &  1.24  &    &    &  &  $24.6^{+29.7}_{-17.1}$  &  $2.36^{+1.65}_{-0.56}$  \\\\ GRB 060719  &  2.03  &    &  $26.3^{+6.8}_{-5.9}$  &  $2.22^{+0.38}_{-0.32}$  &  $38.7^{+5.7}_{-5.9}$  &  $1.78^{+0.19}_{-0.19}$  \\\\ GRB 060729  &  4.49  &  0.54  &  $20.6^{+2.0}_{-2.0}$  &  $2.206^{+0.075}_{-0.072}$  &  $13.2^{+1.7}_{-1.6}$  &  $1.067^{+0.038}_{-0.037}$  \\\\ GRB 060801  &  1.35  &  1.13  &    &  &  $62.0^{+48.7}_{-18.9}$  &  $0.94^{+0.27}_{-0.26}$  \\\\ GRB 060804  &  45.1  &    &    &  &  $<6.46$  &  $1.09^{+0.15}_{-0.12}$  \\\\ GRB 060805A  &  1.55  &    &    &  &  $8.0^{+7.8}_{-6.7}$  &  $1.29^{+0.41}_{-0.36}$  \\\\ GRB 060805B  &  3.16  &    &    &  &  $0.18^{+6.16}_{-0.18}$  &  $0.73^{+0.24}_{-0.21}$  \\\\ GRB 060807  &  2.45  &    &    &  &  $12.1^{+2.4}_{-2.2}$  &  $1.087^{+0.098}_{-0.096}$  \\\\ GRB 060813  &  30.5  &    &  $8.4^{+3.4}_{-3.2}$  &  $0.877^{+0.078}_{-0.077}$  &  $6.0^{+3.6}_{-3.3}$  &  $0.904^{+0.086}_{-0.076}$  \\\\ GRB 060814  &  2.33  &    &  $23.7^{+1.3}_{-1.3}$  &  $0.787^{+0.038}_{-0.038}$  &  $30.5^{+2.6}_{-2.5}$  &  $1.296^{+0.084}_{-0.082}$  \\\\ GRB 060825  &  29.2  &    &  $6.7^{+33.6}_{-6.7}$  &  $0.93^{+0.69}_{-0.47}$  &  $5.2^{+16.9}_{-5.2}$  &  $0.53^{+0.27}_{-0.28}$  \\\\ GRB 060901  &  21.3  &    &    &  &  $25.7^{+27.9}_{-18.8}$  &  $0.91^{+0.50}_{-0.44}$  \\\\ GRB 060904A  &  1.33  &    &  $14.0^{+1.2}_{-1.1}$  &  $1.100^{+0.051}_{-0.049}$  &  $<22.76$  &  $-0.00^{+0.91}_{-0.69}$  \\\\ GRB 060904B  &  11.3  &  0.7  &  $46.9^{+3.9}_{-3.7}$  &  $1.095^{+0.043}_{-0.041}$  &  $38.9^{+9.5}_{-6.3}$  &  $1.188^{+0.096}_{-0.142}$  \\\\ GRB 060906  &  9.81  &  3.69  &  $347.2^{+454.5}_{-281.2}$  &  $1.56^{+0.46}_{-0.40}$  &  $245.0^{+138.6}_{-130.5}$  &  $1.14^{+0.13}_{-0.16}$  \\\\ GRB 060908  &  2.34  &  2.43  &  $43.7^{+39.9}_{-36.9}$  &  $1.24^{+0.20}_{-0.19}$  &  $95.1^{+51.1}_{-45.7}$  &  $1.13^{+0.18}_{-0.12}$  \\\\ GRB 060912A  &  3.85  &  0.94  &  $17.1^{+54.2}_{-17.1}$  &  $0.78^{+0.85}_{-0.67}$  &  $18.6^{+7.8}_{-8.6}$  &  $0.88^{+0.11}_{-0.11}$  \\\\ GRB 060919  &  5.90  &    &    &  &  $92.2^{+38.9}_{-36.7}$  &  $1.08^{+0.42}_{-0.51}$  \\\\ GRB 060923A  &  4.52  &    &  $8.9^{+8.1}_{-6.4}$  &  $0.95^{+0.32}_{-0.28}$  &  $9.5^{+3.1}_{-6.1}$  &  $0.92^{+0.13}_{-0.23}$  \\\\ GRB 060923B  &  8.44  &    &    &  &  $29.8^{+12.9}_{-10.9}$  &  $1.16^{+0.33}_{-0.30}$  \\\\ GRB 060923C  &  5.14  &    &  $0.81^{+4.49}_{-0.81}$  &  $0.87^{+0.23}_{-0.12}$  &  $7.9^{+4.8}_{-3.7}$  &  $1.04^{+0.16}_{-0.18}$  \\\\ GRB 060926  &  7.58  &  3.2  &  $205.1^{+175.3}_{-129.3}$  &  $1.01^{+0.23}_{-0.22}$  &  $341.2^{+234.4}_{-195.1}$  &  $1.04^{+0.21}_{-0.28}$  \\\\ GRB 060927  &  4.61  &  5.6  &    &  &  $122.1^{+251.0}_{-122.1}$  &  $0.89^{+0.16}_{-0.16}$  \\\\ GRB 060928  &  94.3  &    &    &  &  $613.0^{+399.6}_{-523.5}$  &  $2.7^{+1.6}_{-1.1}$  \\\\ GRB 060929  &  3.61  &    &  $11.9^{+1.4}_{-1.3}$  &  $0.773^{+0.049}_{-0.049}$  &  $10.2^{+4.6}_{-5.9}$  &  $1.22^{+0.25}_{-0.25}$  \\\\ GRB 061002  &  3.71  &    &  $15.3^{+8.0}_{-7.0}$  &  $1.30^{+0.35}_{-0.30}$  &  $20.6^{+13.9}_{-15.3}$  &  $1.51^{+0.67}_{-0.67}$  \\\\ GRB 061004  &  4.69  &    &  $<10.84$  &  $0.70^{+0.52}_{-0.25}$  &  $3.3^{+1.6}_{-3.3}$  &  $0.976^{+0.092}_{-0.167}$  \\\\ GRB 061006  &  14.1  &  0.4377  &    &  &  $17.9^{+16.8}_{-17.9}$  &  $0.98^{+0.25}_{-0.32}$  \\\\ GRB 061007  &  1.77  &  1.26  &  $44.7^{+2.6}_{-2.5}$  &  $0.884^{+0.023}_{-0.023}$  &  $52.7^{+10.4}_{-9.7}$  &  $1.018^{+0.087}_{-0.083}$  \\\\ GRB 061019  &  37.8  &    &    &  &  $68.8^{+14.5}_{-11.4}$  &  $1.07^{+0.14}_{-0.22}$  \\\\ GRB 061021  &  4.52  &    &  $<0.20$  &  $0.963^{+0.050}_{-0.050}$  &  $4.63^{+1.22}_{-0.72}$  &  $1.029^{+0.055}_{-0.051}$  \\\\ GRB 061025  &  4.06  &    &    &  &  $2.6^{+13.0}_{-2.6}$  &  $0.72^{+0.63}_{-0.37}$  \\\\ GRB 061028  &  13.1  &    &  $9.7^{+6.3}_{-5.7}$  &  $0.39^{+0.15}_{-0.15}$  &  $10.7^{+10.4}_{-10.7}$  &  $1.46^{+0.36}_{-0.43}$  \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lcccc|cc} \\hline &                  &            & \\multicolumn{2}{c}{WT Mode} & \\multicolumn{2}{c}{PC Mode} \\\\ GRB  &  Galactic $N_{\\rm H}$  &  redshift  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$ \\\\ \\hline GRB 061102  &  4.57  &    &  $8.5^{+5.4}_{-4.7}$  &  $2.20^{+0.40}_{-0.35}$  &  $3.0^{+34.6}_{-3.0}$  &  $0.36^{+0.92}_{-0.70}$  \\\\ GRB 061110A  &  4.26  &    &  $10.3^{+1.2}_{-1.1}$  &  $1.989^{+0.074}_{-0.068}$  &  $<0.70$  &  $0.86^{+0.15}_{-0.10}$  \\\\ GRB 061110B  &  3.37  &  3.44  &    &  &  $342.2^{+250.0}_{-257.8}$  &  $1.24^{+0.31}_{-0.27}$  \\\\ GRB 061121  &  3.99  &  1.314  &  $42.9^{+6.1}_{-5.8}$  &  $0.617^{+0.042}_{-0.041}$  &  $53.5^{+8.5}_{-8.0}$  &  $0.914^{+0.061}_{-0.059}$  \\\\ GRB 061122  &  12.5  &    &    &  &  $5.1^{+4.7}_{-4.7}$  &  $1.04^{+0.17}_{-0.18}$  \\\\ GRB 061126  &  10.2  &    &  $9.5^{+4.6}_{-4.2}$  &  $0.86^{+0.16}_{-0.15}$  &  $14.1^{+2.1}_{-2.0}$  &  $0.936^{+0.065}_{-0.064}$  \\\\ GRB 061201  &  5.31  &    &  $<19.12$  &  $-0.03^{+0.62}_{-0.26}$  &  $7.9^{+4.3}_{-5.1}$  &  $0.68^{+0.17}_{-0.18}$  \\\\ GRB 061202  &  11.8  &    &  $45.4^{+4.8}_{-4.5}$  &  $0.917^{+0.085}_{-0.082}$  &  $51.8^{+5.5}_{-5.2}$  &  $1.208^{+0.103}_{-0.100}$  \\\\ GRB 061210  &  3.20  &    &    &  &  $62.8^{+92.0}_{-55.5}$  &  $1.9^{+1.7}_{-1.3}$  \\\\ GRB 061217  &  3.41  &    &    &  &  $58.7^{+200.7}_{-49.0}$  &  $0.40^{+1.13}_{-0.86}$  \\\\ GRB 061222A  &  9.01  &    &  $32.4^{+2.9}_{-2.8}$  &  $1.196^{+0.078}_{-0.076}$  &  $27.1^{+2.7}_{-2.6}$  &  $1.071^{+0.071}_{-0.069}$  \\\\ GRB 061222B  &  26.5  &  3.36  &  $395.7^{+143.6}_{-125.0}$  &  $1.60^{+0.15}_{-0.14}$  &  $448.1^{+338.5}_{-329.4}$  &  $0.97^{+0.18}_{-0.23}$  \\\\ GRB 070103  &  4.43  &    &    &  &  $27.1^{+3.1}_{-6.3}$  &  $1.35^{+0.18}_{-0.21}$  \\\\ GRB 070107  &  29.9  &    &  $3.4^{+2.3}_{-2.2}$  &  $0.903^{+0.060}_{-0.059}$  &  $7.0^{+2.9}_{-2.7}$  &  $1.102^{+0.077}_{-0.073}$  \\\\ GRB 070110  &  1.61  &  2.352  &  $10.3^{+17.4}_{-10.3}$  &  $0.904^{+0.092}_{-0.085}$  &  $10.1^{+11.7}_{-10.1}$  &  $1.107^{+0.063}_{-0.061}$  \\\\ GRB 070125  &  4.28  &    &    &  &  $7.1^{+2.9}_{-5.1}$  &  $1.11^{+0.21}_{-0.23}$  \\\\ GRB 070129  &  6.96  &    &  $11.59^{+0.77}_{-0.75}$  &  $0.846^{+0.027}_{-0.026}$  &  $7.6^{+2.2}_{-1.8}$  &  $1.270^{+0.062}_{-0.103}$  \\\\ GRB 070208  &  1.75  &  1.17  &    &  &  $71.2^{+15.3}_{-20.8}$  &  $1.18^{+0.14}_{-0.18}$  \\\\ GRB 070219  &  3.95  &    &  $22.7^{+175.7}_{-22.7}$  &  $1.1^{+4.8}_{-2.0}$  &  $31.9^{+19.3}_{-9.8}$  &  $1.53^{+0.21}_{-0.20}$  \\\\ GRB 070220  &  34.8  &    &  $10.4^{+4.6}_{-4.4}$  &  $0.435^{+0.077}_{-0.076}$  &  $18.2^{+7.2}_{-6.7}$  &  $0.61^{+0.11}_{-0.11}$  \\\\ GRB 070223  &  1.39  &    &  $52.5^{+5.7}_{-5.4}$  &  $0.748^{+0.098}_{-0.096}$  &  $47.6^{+15.8}_{-14.5}$  &  $1.23^{+0.33}_{-0.36}$  \\\\ GRB 070224  &  3.76  &    &  $<4.78$  &  $1.58^{+0.40}_{-0.19}$  &  $<4.11$  &  $1.20^{+0.24}_{-0.22}$  \\\\ GRB 070227  &  23.9  &    &    &  &  $10.0^{+17.0}_{-10.0}$  &  $1.25^{+0.48}_{-0.42}$  \\\\ GRB 070306  &  2.85  &    &  $39.2^{+2.4}_{-2.3}$  &  $1.248^{+0.065}_{-0.062}$  &  $33.2^{+3.2}_{-3.1}$  &  $1.115^{+0.087}_{-0.083}$  \\\\ GRB 070309  &  52.9  &    &    &  &  $0.48^{+19.90}_{-0.48}$  &  $5.4^{+0.0}_{-1.2}$  \\\\ GRB 070311  &  23.6  &    &    &  &  $23.7^{+8.9}_{-10.2}$  &  $1.13^{+0.21}_{-0.23}$  \\\\ GRB 070318  &  1.44  &  0.84  &  $44.8^{+5.4}_{-5.2}$  &  $0.621^{+0.054}_{-0.053}$  &  $65.1^{+6.1}_{-5.1}$  &  $1.004^{+0.077}_{-0.073}$  \\\\ GRB 070328  &  2.61  &    &  $20.37^{+1.01}_{-0.98}$  &  $1.237^{+0.039}_{-0.037}$  &  $20.4^{+2.5}_{-2.4}$  &  $1.171^{+0.087}_{-0.085}$  \\\\ GRB 070330  &  5.45  &    &    &  &  $8.8^{+7.0}_{-5.3}$  &  $0.976^{+0.084}_{-0.144}$  \\\\ GRB 070411  &  26.3  &  2.95  &    &  &  $155.3^{+164.1}_{-130.2}$  &  $1.29^{+0.21}_{-0.18}$  \\\\ GRB 070412  &  2.68  &    &  $19.6^{+10.9}_{-9.2}$  &  $1.23^{+0.43}_{-0.36}$  &  $12.4^{+5.2}_{-5.8}$  &  $1.08^{+0.13}_{-0.23}$  \\\\ GRB 070419A  &  2.42  &  0.97  &  $52.6^{+5.8}_{-5.5}$  &  $1.249^{+0.070}_{-0.067}$  &  $35.6^{+16.4}_{-19.2}$  &  $1.15^{+0.25}_{-0.26}$  \\\\ GRB 070419B  &  6.72  &    &  $13.75^{+0.94}_{-0.92}$  &  $0.822^{+0.031}_{-0.030}$  &  $6.4^{+2.2}_{-2.1}$  &  $0.601^{+0.073}_{-0.071}$  \\\\ GRB 070420  &  33.1  &    &  $5.7^{+3.6}_{-3.4}$  &  $1.52^{+0.12}_{-0.11}$  &  $7.4^{+4.3}_{-4.0}$  &  $1.017^{+0.092}_{-0.092}$  \\\\ GRB 070429A  &  8.39  &    &  $37.2^{+4.5}_{-4.1}$  &  $2.62^{+0.21}_{-0.19}$  &  $5.2^{+10.4}_{-5.2}$  &  $0.75^{+0.36}_{-0.29}$  \\\\ GRB 070429B  &  1.80  &    &    &  &  $60.7^{+44.9}_{-30.3}$  &  $2.1^{+1.0}_{-1.4}$  \\\\ GRB 070506  &  3.79  &  2.31  &    &  &  $29.6^{+48.1}_{-27.9}$  &  $1.02^{+0.23}_{-0.22}$  \\\\ GRB 070508  &  7.84  &    &  $27.0^{+1.5}_{-1.5}$  &  $0.907^{+0.038}_{-0.038}$  &  $22.1^{+4.7}_{-2.1}$  &  $0.92^{+0.12}_{-0.11}$  \\\\ GRB 070509  &  7.89  &    &    &  &  $41.8^{+21.7}_{-12.9}$  &  $1.19^{+0.36}_{-0.55}$  \\\\ GRB 070517  &  8.30  &    &    &  &  $3.9^{+5.1}_{-3.9}$  &  $1.11^{+0.23}_{-0.22}$  \\\\ GRB 070518  &  1.49  &    &  $8.9^{+2.2}_{-2.1}$  &  $1.62^{+0.13}_{-0.13}$  &  $1.6^{+4.0}_{-1.6}$  &  $1.21^{+0.29}_{-0.24}$  \\\\ GRB 070520A  &  2.15  &    &  $22.2^{+3.8}_{-3.5}$  &  $2.19^{+0.20}_{-0.19}$  &  $23.5^{+7.9}_{-17.6}$  &  $2.12^{+0.31}_{-0.37}$  \\\\ GRB 070520B  &  3.76  &    &  $19.4^{+1.5}_{-1.4}$  &  $1.517^{+0.065}_{-0.062}$  &  $17.0^{+12.4}_{-11.0}$  &  $1.43^{+0.32}_{-0.47}$  \\\\ GRB 070521  &  2.92  &    &  $56.8^{+79.4}_{-39.7}$  &  $1.18^{+1.03}_{-0.83}$  &  $58.6^{+6.9}_{-6.0}$  &  $1.05^{+0.12}_{-0.12}$  \\\\ GRB 070529  &  17.9  &  2.5  &  $534.1^{+543.1}_{-352.4}$  &  $1.94^{+0.83}_{-0.64}$  &  $247.6^{+82.3}_{-125.5}$  &  $1.02^{+0.13}_{-0.17}$  \\\\ GRB 070531  &  21.2  &    &    &  &  $2.2^{+13.2}_{-2.2}$  &  $0.07^{+0.25}_{-0.18}$  \\\\ GRB 070611  &  1.38  &  2.04  &    &  &  $17.9^{+38.3}_{-17.9}$  &  $0.92^{+0.24}_{-0.23}$  \\\\ GRB 070612B  &  15.7  &    &    &  &  $58.7^{+25.8}_{-29.4}$  &  $1.46^{+0.46}_{-0.80}$  \\\\ GRB 070615  &  4.93  &    &    &  &  $5.3^{+41.8}_{-5.3}$  &  $0.5^{+3.4}_{-4.2}$  \\\\ GRB 070616  &  34.0  &    &  $5.5^{+1.2}_{-1.2}$  &  $0.214^{+0.019}_{-0.019}$  &  $9.0^{+4.1}_{-6.0}$  &  $1.40^{+0.11}_{-0.18}$  \\\\ GRB 070621  &  3.81  &    &  $37.7^{+3.2}_{-3.0}$  &  $1.428^{+0.095}_{-0.092}$  &  $37.6^{+6.5}_{-5.7}$  &  $1.68^{+0.19}_{-0.20}$  \\\\ GRB 070628  &  57.8  &    &  $39.8^{+74.5}_{-39.8}$  &  $3.2^{+2.2}_{-1.2}$  &  $10.1^{+5.2}_{-5.0}$  &  $1.025^{+0.082}_{-0.079}$  \\\\ GRB 070704  &  71.0  &    &  $33.8^{+5.9}_{-5.7}$  &  $0.734^{+0.061}_{-0.061}$  &  $49.8^{+37.3}_{-20.8}$  &  $0.98^{+0.18}_{-0.33}$  \\\\ GRB 070707  &  6.15  &    &    &  &  $0.64^{+16.62}_{-0.64}$  &  $1.94^{+2.03}_{-0.61}$  \\\\ GRB 070714A  &  9.22  &    &    &  &  $46.9^{+20.6}_{-17.5}$  &  $1.04^{+0.36}_{-0.33}$  \\\\ GRB 070714B  &  6.43  &  0.92  &  $30.8^{+12.8}_{-10.9}$  &  $0.266^{+0.086}_{-0.084}$  &  $48.6^{+14.3}_{-15.1}$  &  $1.065^{+0.081}_{-0.103}$  \\\\ GRB 070721A  &  1.61  &    &  $7.1^{+14.4}_{-7.1}$  &  $2.48^{+1.16}_{-0.45}$  &  $4.4^{+6.5}_{-2.7}$  &  $1.38^{+0.36}_{-0.20}$  \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lcccc|cc} \\hline &                  &            & \\multicolumn{2}{c}{WT Mode} & \\multicolumn{2}{c}{PC Mode} \\\\ GRB  &  Galactic $N_{\\rm H}$  &  redshift  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$ \\\\ \\hline GRB 070721B  &  2.33  &    &  $1.6^{+2.1}_{-1.6}$  &  $0.421^{+0.091}_{-0.085}$  &  $0.76^{+2.32}_{-0.76}$  &  $0.501^{+0.098}_{-0.050}$  \\\\ GRB 070724A  &  1.17  &    &  $21.0^{+6.5}_{-6.0}$  &  $0.60^{+0.18}_{-0.18}$  &  $11.6^{+10.4}_{-11.6}$  &  $0.69^{+0.28}_{-0.29}$  \\\\ GRB 070724B  &  30.7  &    &    &  &  $27.5^{+15.8}_{-17.4}$  &  $1.17^{+0.49}_{-0.40}$  \\\\ GRB 070729  &  2.58  &    &    &  &  $0.024^{+9.135}_{-0.024}$  &  $0.62^{+0.86}_{-0.43}$  \\\\ GRB 070802  &  2.90  &  2.45  &    &  &  $122.7^{+80.0}_{-100.7}$  &  $1.08^{+0.17}_{-0.18}$  \\\\ GRB 070805  &  19.2  &    &    &  &  $0.057^{+39.184}_{-0.057}$  &  $1.2^{+4.0}_{-1.3}$  \\\\ GRB 070808  &  2.83  &    &  $79.6^{+53.2}_{-39.6}$  &  $0.60^{+0.66}_{-0.56}$  &  $85.3^{+18.8}_{-20.0}$  &  $1.04^{+0.26}_{-0.27}$  \\\\ GRB 070809  &  6.40  &    &    &  &  $<3.87$  &  $0.39^{+0.18}_{-0.13}$  \\\\ GRB 070810A  &  1.75  &  2.17  &    &  &  $68.1^{+27.7}_{-15.2}$  &  $1.141^{+0.128}_{-0.094}$  \\\\ GRB 070911  &  2.39  &    &    &  &  $7.3^{+4.6}_{-5.3}$  &  $1.10^{+0.24}_{-0.24}$  \\\\ GRB 070913  &  8.10  &    &    &  &  $11.9^{+44.2}_{-11.9}$  &  $1.4^{+4.1}_{-1.2}$  \\\\ GRB 070917  &  17.3  &    &    &  &  $1.4^{+56.3}_{-1.4}$  &  $0.69^{+1.25}_{-0.61}$  \\\\ GRB 070925  &  12.9  &    &    &  &  $1.7^{+30.3}_{-1.7}$  &  $0.96^{+1.13}_{-0.54}$  \\\\ GRB 071001  &  76.2  &    &  $<5.08$  &  $1.95^{+0.23}_{-0.23}$  &  $<3.00$  &  $1.46^{+0.12}_{-0.18}$  \\\\ GRB 071003  &  10.7  &    &    &  &  $4.8^{+3.5}_{-3.7}$  &  $1.04^{+0.13}_{-0.12}$  \\\\ GRB 071008  &  8.78e-01  &    &    &  &  $<25.70$  &  $1.82^{+1.56}_{-0.90}$  \\\\ GRB 071010A  &  6.54  &    &    &  &  $40.1^{+10.3}_{-12.2}$  &  $1.67^{+0.27}_{-0.33}$  \\\\ GRB 071010B  &  1.01  &  0.947  &    &  &  $24.0^{+21.9}_{-18.3}$  &  $1.10^{+0.33}_{-0.17}$  \\\\ GRB 071011  &  51.8  &    &    &  &  $34.3^{+32.9}_{-17.5}$  &  $1.35^{+0.45}_{-0.41}$  \\\\ GRB 071017  &  119.0  &    &    &  &  $489.6^{+353.4}_{-288.2}$  &  $1.05^{+0.99}_{-0.85}$  \\\\ GRB 071020  &  5.12  &    &  $4.7^{+1.9}_{-1.8}$  &  $0.906^{+0.078}_{-0.074}$  &  $3.7^{+3.9}_{-3.6}$  &  $0.70^{+0.15}_{-0.15}$  \\\\ GRB 071021  &  4.80  &    &  $15.1^{+2.4}_{-2.3}$  &  $1.097^{+0.091}_{-0.089}$  &  $13.7^{+1.9}_{-3.2}$  &  $1.205^{+0.062}_{-0.125}$  \\\\ GRB 071025  &  5.13  &    &  $5.7^{+1.3}_{-1.3}$  &  $0.409^{+0.043}_{-0.043}$  &  $8.7^{+2.2}_{-1.8}$  &  $1.273^{+0.092}_{-0.098}$  \\\\ GRB 071028A  &  5.09  &    &  $5.5^{+5.5}_{-4.7}$  &  $1.14^{+0.26}_{-0.24}$  &  $1.7^{+3.5}_{-1.7}$  &  $1.05^{+0.14}_{-0.19}$  \\\\ GRB 071028B  &  1.22  &    &    &  &  $1.2^{+15.3}_{-1.2}$  &  $1.01^{+1.13}_{-0.78}$  \\\\ GRB 071031  &  1.22  &  2.69  &  $88.1^{+8.9}_{-8.7}$  &  $0.848^{+0.026}_{-0.026}$  &  $17.0^{+34.4}_{-17.0}$  &  $0.86^{+0.14}_{-0.15}$  \\\\ GRB 071101  &  56.3  &    &    &  &  $124.7^{+314.6}_{-124.7}$  &  $1.7^{+2.3}_{-1.1}$  \\\\ GRB 071104  &  27.5  &    &    &  &  $2.0^{+5.8}_{-2.0}$  &  $0.85^{+0.11}_{-0.19}$  \\\\ GRB 071109  &  24.9  &    &    &  &  $<29.04$  &  $0.41^{+1.36}_{-0.67}$  \\\\ GRB 071112C  &  7.44  &  0.82  &  $5.9^{+4.8}_{-4.6}$  &  $0.733^{+0.072}_{-0.070}$  &  $7.1^{+8.4}_{-4.4}$  &  $0.740^{+0.080}_{-0.112}$  \\\\ GRB 071117  &  2.33  &  1.33  &    &  &  $117.9^{+19.8}_{-30.4}$  &  $1.15^{+0.16}_{-0.10}$  \\\\ GRB 071118  &  9.70  &    &  $19.5^{+5.5}_{-5.2}$  &  $0.69^{+0.15}_{-0.14}$  &  $18.9^{+3.9}_{-3.7}$  &  $0.714^{+0.099}_{-0.096}$  \\\\ GRB 071122  &  4.75  &  1.14  &  $20.2^{+12.0}_{-11.1}$  &  $1.07^{+0.15}_{-0.14}$  &  $<11.07$  &  $0.64^{+0.19}_{-0.13}$  \\\\ GRB 071227  &  1.25  &    &  $15.0^{+4.3}_{-4.0}$  &  $0.68^{+0.14}_{-0.14}$  &  $2.3^{+11.3}_{-2.3}$  &  $0.78^{+0.54}_{-0.34}$  \\\\ GRB 080120  &  7.70  &    &    &  &  $4.9^{+5.7}_{-3.1}$  &  $1.33^{+0.26}_{-0.27}$  \\\\ GRB 080121  &  1.23  &    &    &  &  $<6.81$  &  $1.01^{+0.71}_{-0.41}$  \\\\ GRB 080123  &  2.30  &    &  $9.2^{+2.7}_{-2.6}$  &  $0.74^{+0.11}_{-0.10}$  &  $14.0^{+10.7}_{-9.3}$  &  $1.62^{+0.47}_{-0.33}$  \\\\ GRB 080129  &  64.2  &    &    &  &  $11.1^{+11.3}_{-11.1}$  &  $1.053^{+0.100}_{-0.181}$  \\\\ GRB 080130  &  12.4  &    &    &  &  $<28.39$  &  $0.30^{+0.99}_{-0.50}$  \\\\ GRB 080205  &  7.14  &    &  $12.0^{+3.5}_{-3.3}$  &  $1.11^{+0.14}_{-0.14}$  &  $16.4^{+3.0}_{-5.7}$  &  $0.94^{+0.15}_{-0.19}$  \\\\ GRB 080207  &  1.95  &    &  $94.8^{+10.0}_{-9.5}$  &  $0.333^{+0.093}_{-0.093}$  &  $69.0^{+7.9}_{-8.0}$  &  $1.51^{+0.13}_{-0.15}$  \\\\ GRB 080210  &  5.47  &  2.64  &  $215.3^{+61.8}_{-53.7}$  &  $1.71^{+0.19}_{-0.18}$  &  $148.7^{+72.5}_{-66.8}$  &  $1.22^{+0.12}_{-0.10}$  \\\\ GRB 080212  &  7.26  &    &  $21.6^{+1.1}_{-1.1}$  &  $1.907^{+0.050}_{-0.049}$  &  $12.8^{+2.3}_{-2.1}$  &  $1.290^{+0.059}_{-0.060}$  \\\\ GRB 080218B  &  13.7  &    &    &  &  $32.9^{+12.1}_{-16.1}$  &  $1.54^{+0.30}_{-0.45}$  \\\\ GRB 080229A  &  8.97  &    &  $24.9^{+2.2}_{-2.2}$  &  $1.991^{+0.098}_{-0.093}$  &  $17.4^{+2.5}_{-2.4}$  &  $0.865^{+0.070}_{-0.068}$  \\\\ GRB 080229B  &  87.8  &    &    &  &  $49.3^{+37.6}_{-13.7}$  &  $1.73^{+0.34}_{-0.17}$  \\\\ GRB 080303  &  12.6  &    &  $3.6^{+5.9}_{-3.6}$  &  $1.01^{+0.22}_{-0.20}$  &  $1.3^{+3.0}_{-1.3}$  &  $0.99^{+0.15}_{-0.16}$  \\\\ GRB 080307  &  2.37  &    &  $2.5^{+1.4}_{-1.4}$  &  $0.445^{+0.057}_{-0.056}$  &  $8.7^{+4.4}_{-5.0}$  &  $0.888^{+0.094}_{-0.208}$  \\\\ GRB 080310  &  3.28  &  2.43  &  $50.0^{+5.4}_{-5.3}$  &  $0.677^{+0.017}_{-0.017}$  &  $<9.25$  &  $0.996^{+0.064}_{-0.065}$  \\\\ GRB 080319A  &  1.45  &    &    &  &  $10.0^{+4.6}_{-4.4}$  &  $1.16^{+0.20}_{-0.18}$  \\\\ GRB 080319B  &  1.12  &  0.94  &  $11.89^{+0.92}_{-0.92}$  &  $0.770^{+0.014}_{-0.014}$  &  $9.5^{+4.1}_{-3.9}$  &  $0.803^{+0.064}_{-0.063}$  \\\\ GRB 080319C  &  2.21  &  1.95  &    &  &  $71.7^{+19.9}_{-18.7}$  &  $0.690^{+0.081}_{-0.079}$  \\\\ GRB 080319D  &  21.7  &    &  $<143.82$  &  $0.42^{+3.12}_{-0.80}$  &  $8.0^{+7.5}_{-6.7}$  &  $1.02^{+0.21}_{-0.20}$  \\\\ GRB 080320  &  1.35  &    &    &  &  $5.1^{+2.1}_{-1.8}$  &  $1.060^{+0.103}_{-0.099}$  \\\\ GRB 080325  &  3.81  &    &  $28.0^{+1.3}_{-1.3}$  &  $1.360^{+0.043}_{-0.042}$  &  $23.6^{+9.1}_{-8.3}$  &  $1.46^{+0.27}_{-0.30}$  \\\\ GRB 080328  &  33.8  &    &  $10.2^{+3.8}_{-3.6}$  &  $1.139^{+0.094}_{-0.092}$  &  $10.2^{+3.4}_{-3.3}$  &  $1.029^{+0.077}_{-0.075}$  \\\\ GRB 080330  &  1.22  &  1.51  &  $16.7^{+15.3}_{-14.1}$  &  $1.14^{+0.16}_{-0.15}$  &  $12.0^{+13.8}_{-12.0}$  &  $0.89^{+0.13}_{-0.12}$  \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lcccc|cc} \\hline &                  &            & \\multicolumn{2}{c}{WT Mode} & \\multicolumn{2}{c}{PC Mode} \\\\ GRB  &  Galactic $N_{\\rm H}$  &  redshift  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$  &  intrinsic $N_{\\rm H}$  &  Spectral index ($\\beta)$ \\\\ \\hline GRB 080405  &  3.31  &    &    &  &  $<12.47$  &  $0.89^{+0.52}_{-0.36}$  \\\\ GRB 080409  &  23.4  &    &    &  &  $37.3^{+14.2}_{-19.1}$  &  $1.28^{+0.30}_{-0.22}$  \\\\ GRB 080411  &  5.77  &  1.03  &    &  &  $44.4^{+4.8}_{-4.5}$  &  $1.024^{+0.044}_{-0.044}$  \\\\ GRB 080413A  &  8.71  &  2.43  &  $187.2^{+42.1}_{-37.9}$  &  $2.17^{+0.17}_{-0.16}$  &  $53.4^{+75.8}_{-53.4}$  &  $1.13^{+0.21}_{-0.24}$  \\\\ GRB 080413B  &  3.06  &  1.1  &  $27.5^{+9.2}_{-8.5}$  &  $1.032^{+0.103}_{-0.100}$  &  $26.0^{+5.4}_{-5.4}$  &  $0.989^{+0.065}_{-0.040}$  \\\\ GRB 080426  &  36.5  &    &    &  &  $16.5^{+5.1}_{-11.5}$  &  $1.10^{+0.12}_{-0.25}$  \\\\ GRB 080430  &  9.59e-01  &  0.767  &  $15.4^{+7.7}_{-7.1}$  &  $1.44^{+0.21}_{-0.19}$  &  $34.8^{+4.3}_{-3.5}$  &  $1.102^{+0.074}_{-0.066}$  \\\\ GRB 080503  &  5.57  &    &  $5.1^{+1.6}_{-1.6}$  &  $0.346^{+0.051}_{-0.051}$  &  $<2.22$  &  $1.53^{+0.35}_{-0.20}$  \\\\ GRB 080506  &  16.6  &    &  $0.10^{+1.73}_{-0.10}$  &  $0.623^{+0.055}_{-0.030}$  &  $<2.50$  &  $0.990^{+0.122}_{-0.077}$  \\\\ GRB 080507  &  1.31  &    &    &  &  $18.6^{+22.3}_{-16.2}$  &  $1.69^{+1.05}_{-0.83}$  \\\\ GRB 080514B  &  3.75  &    &    &  &  $5.5^{+3.7}_{-5.5}$  &  $1.00^{+0.24}_{-0.26}$  \\\\ GRB 080515  &  4.65  &    &    &  &  $1.3^{+19.5}_{-1.3}$  &  $0.84^{+0.81}_{-0.39}$  \\\\ GRB 080516  &  32.2  &    &    &  &  $58.3^{+21.9}_{-15.9}$  &  $1.26^{+0.18}_{-0.24}$  \\\\ GRB 080517  &  8.12  &    &  $16.3^{+28.3}_{-16.3}$  &  $1.21^{+1.03}_{-0.80}$  &  $18.4^{+8.4}_{-8.2}$  &  $0.95^{+0.24}_{-0.27}$  \\\\ GRB 080520  &  6.84  &  1.55  &    &  &  $138.2^{+39.8}_{-59.3}$  &  $1.30^{+0.22}_{-0.28}$  \\\\ GRB 080523  &  2.10  &    &  $6.0^{+1.6}_{-1.5}$  &  $1.461^{+0.092}_{-0.088}$  &  $<2.39$  &  $0.83^{+0.18}_{-0.12}$  \\\\ GRB 080602  &  3.51  &    &  $12.7^{+6.3}_{-5.7}$  &  $0.90^{+0.25}_{-0.24}$  &  $13.5^{+3.4}_{-3.3}$  &  $1.01^{+0.13}_{-0.12}$  \\\\ GRB 080603A  &  4.69  &  1.69  &    &  &  $68.1^{+27.4}_{-25.2}$  &  $1.43^{+0.20}_{-0.19}$  \\\\ GRB 080603B  &  1.23  &  2.69  &  $93.2^{+22.4}_{-21.5}$  &  $0.678^{+0.060}_{-0.060}$  &  $79.8^{+75.7}_{-62.5}$  &  $0.99^{+0.23}_{-0.21}$  \\\\ GRB 080604  &  3.81  &  1.42  &  $<3.44$  &  $1.066^{+0.069}_{-0.066}$  &  $5.9^{+12.1}_{-5.9}$  &  $0.65^{+0.13}_{-0.13}$  \\\\ GRB 080605  &  6.67  &  1.64  &  $71.3^{+8.1}_{-7.9}$  &  $0.755^{+0.036}_{-0.035}$  &  $70.3^{+30.0}_{-25.1}$  &  $0.829^{+0.121}_{-0.073}$  \\\\ GRB 080607  &  1.69  &  3.04  &  $377.4^{+23.1}_{-21.9}$  &  $1.044^{+0.036}_{-0.035}$  &  $260.5^{+51.1}_{-29.9}$  &  $1.101^{+0.098}_{-0.092}$  \\\\ GRB 080613A  &  2.03  &    &    &  &  $<4.56$  &  $0.25^{+0.37}_{-0.20}$  \\\\ GRB 080613B  &  3.17  &    &  $2.9^{+2.4}_{-2.3}$  &  $0.173^{+0.079}_{-0.079}$  &  $8.0^{+12.0}_{-8.0}$  &  $1.26^{+0.66}_{-0.53}$  \\\\ GRB 080623  &  25.9  &    &    &  &  $14.2^{+6.4}_{-7.0}$  &  $1.14^{+0.16}_{-0.19}$  \\\\ GRB 080625  &  14.5  &    &    &  &  $8.8^{+6.4}_{-5.9}$  &  $1.28^{+0.23}_{-0.22}$  \\\\ GRB 080701  &  17.8  &    &    &  &  $57.9^{+21.8}_{-18.3}$  &  $1.87^{+0.44}_{-0.40}$  \\\\ GRB 080702A  &  15.3  &    &    &  &  $61.6^{+44.0}_{-50.0}$  &  $1.08^{+0.55}_{-0.79}$  \\\\ GRB 080702B  &  2.94  &    &    &  &  $3.9^{+37.3}_{-3.9}$  &  $1.01^{+1.87}_{-0.84}$  \\\\ GRB 080703  &  5.13  &    &    &  &  $2.3^{+1.6}_{-1.3}$  &  $0.585^{+0.053}_{-0.056}$  \\\\ GRB 080707  &  6.99  &  1.23  &  $20.8^{+55.5}_{-20.8}$  &  $1.44^{+0.83}_{-0.52}$  &  $38.1^{+18.8}_{-19.2}$  &  $1.11^{+0.17}_{-0.18}$  \\\\ GRB 080710  &  4.10  &  0.85  &    &  &  $12.5^{+4.9}_{-4.4}$  &  $1.050^{+0.068}_{-0.093}$  \\\\ GRB 080714  &  54.4  &    &  $<4.95$  &  $0.52^{+0.13}_{-0.13}$  &  $5.1^{+22.2}_{-5.1}$  &  $0.67^{+0.37}_{-0.20}$  \\\\ GRB 080721  &  6.94  &  2.6  &  $70.9^{+6.0}_{-5.9}$  &  $0.833^{+0.016}_{-0.016}$  &  $86.4^{+21.5}_{-20.4}$  &  $0.991^{+0.059}_{-0.057}$  \\\\ GRB 080723A  &  39.4  &    &    &  &  $62.1^{+16.3}_{-14.7}$  &  $1.19^{+0.20}_{-0.10}$  \\\\ GRB 080723B  &  74.3  &    &    &  &  $10.5^{+16.8}_{-10.5}$  &  $0.87^{+0.21}_{-0.18}$  \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  \\clearpage       \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $\\beta_1$ & $\\beta_2$ & $\\beta_3$ & $\\beta_4$ & $\\beta_5$ & $\\beta_6$  \\\\ \\hline GRB 050126  &  PC  &  1.29$^{+0.20}_{-0.20}$  &  1.04$^{+0.52}_{-0.36}$   \\\\ GRB 050128  &  WT    &  &  0.35$^{+1.74}_{-0.63}$   \\\\ &  PC  &  0.98$^{+0.13}_{-0.12}$  &  0.966$^{+0.088}_{-0.092}$   \\\\ GRB 050219A  &  WT  &  1.08$^{+0.18}_{-0.17}$  &  1.09$^{+0.29}_{-0.13}$   \\\\ &  PC    &  &  1.16$^{+0.48}_{-0.53}$   \\\\ GRB 050315  &  WT    &  &  1.09$^{+0.34}_{-0.29}$   \\\\ &  PC  &  1.07$^{+0.11}_{-0.10}$  &  0.979$^{+0.085}_{-0.059}$  &  1.088$^{+0.060}_{-0.104}$   \\\\ GRB 050318  &  WT  &  1.51$^{+3.51}_{-0.73}$   \\\\ &  PC  &  0.99$^{+0.10}_{-0.10}$  &  0.99$^{+0.12}_{-0.12}$   \\\\ GRB 050319  &  WT    &    &  &  1.19$^{+0.53}_{-0.28}$   \\\\ &  PC  &  1.49$^{+0.17}_{-0.16}$  &  0.908$^{+0.075}_{-0.069}$  &  1.09$^{+0.11}_{-0.14}$  &  14.0$^{+8.0}_{-4.7}$   \\\\ GRB 050401  &  WT  &  0.895$^{+0.041}_{-0.040}$  &  1.05$^{+0.26}_{-0.25}$   \\\\ &  PC  &  0.70$^{+0.16}_{-0.15}$  &  0.94$^{+0.12}_{-0.14}$   \\\\ GRB 050406  &  WT  &  1.28$^{+0.28}_{-0.17}$   \\\\ &  PC  &  1.13$^{+1.55}_{-0.44}$  &  -0.962$^{+unconstr.}_{-unconstr.}$   \\\\ GRB 050410  &  WT  &  1.28$^{+0.69}_{-0.32}$   \\\\ &  PC  &  0.89$^{+0.44}_{-0.44}$  &  1.2$^{+unconstr.}_{-unconstr.}$   \\\\ GRB 050416A  &  PC  &  0.80$^{+0.26}_{-0.25}$  &  1.08$^{+0.21}_{-0.20}$  &  1.356$^{+0.094}_{-0.069}$   \\\\ GRB 050421  &  WT  &  0.44$^{+0.32}_{-0.30}$  &  0.12$^{+0.70}_{-0.58}$   \\\\ &  PC    &  &  1.21$^{+0.14}_{-0.28}$   \\\\ GRB 050422  &  PC  &  1.69$^{+0.31}_{-0.27}$  &  1.02$^{+0.80}_{-0.31}$   \\\\ GRB 050502B  &  WT  &  1.246$^{+0.034}_{-0.034}$   \\\\ &  PC  &  0.942$^{+0.058}_{-0.091}$  &  1.22$^{+0.51}_{-0.33}$   \\\\ GRB 050505  &  PC  &  0.93$^{+0.16}_{-0.15}$  &  1.032$^{+0.070}_{-0.063}$  &  1.106$^{+0.126}_{-0.088}$   \\\\ GRB 050525A  &  PC  &  0.94$^{+0.30}_{-0.28}$  &  2.25$^{+0.31}_{-0.32}$   \\\\ GRB 050607  &  PC  &  1.20$^{+0.14}_{-0.14}$  &  1.208$^{+0.096}_{-0.105}$  &  0.90$^{+0.23}_{-0.22}$  &  1.47$^{+0.46}_{-0.27}$   \\\\ GRB 050712  &  WT  &  1.27$^{+0.14}_{-0.13}$   \\\\ &  PC  &  1.13$^{+0.10}_{-0.12}$  &  1.32$^{+0.19}_{-0.42}$   \\\\ GRB 050713A  &  WT  &  1.66$^{+0.25}_{-0.23}$  &  1.111$^{+0.065}_{-0.064}$  &  1.10$^{+0.61}_{-0.55}$  &  1.58$^{+0.86}_{-0.68}$   \\\\ &  PC    &    &    &  &  1.216$^{+0.080}_{-0.139}$   \\\\ GRB 050713B  &  PC  &  0.63$^{+0.21}_{-0.20}$  &  1.12$^{+0.18}_{-0.17}$  &  0.988$^{+0.060}_{-0.111}$   \\\\ GRB 050714B  &  WT  &  5.21$^{+0.69}_{-0.61}$   \\\\ &  PC  &  4.1$^{+1.6}_{-1.1}$  &  2.60$^{+0.46}_{-0.36}$   \\\\ GRB 050716  &  WT  &  0.126$^{+0.094}_{-0.093}$  &  0.536$^{+0.080}_{-0.078}$   \\\\ &  PC    &  &  0.86$^{+0.26}_{-0.25}$  &  1.18$^{+0.11}_{-0.10}$   \\\\ GRB 050717  &  WT  &  0.478$^{+0.097}_{-0.095}$  &  0.07$^{+1.00}_{-0.96}$   \\\\ &  PC    &  &  0.65$^{+0.13}_{-0.11}$  &  0.19$^{+0.64}_{-0.77}$   \\\\ GRB 050721  &  WT  &  0.61$^{+0.14}_{-0.13}$   \\\\ &  PC  &  0.577$^{+0.102}_{-0.098}$  &  1.45$^{+0.36}_{-0.24}$   \\\\ GRB 050724  &  WT  &  0.590$^{+0.049}_{-0.048}$  &  1.54$^{+0.13}_{-0.12}$   \\\\ &  PC    &    &  &  0.78$^{+0.11}_{-0.16}$   \\\\ GRB 050726  &  WT  &  1.19$^{+0.14}_{-0.13}$   \\\\ &  PC  &  0.851$^{+0.100}_{-0.067}$  &  1.36$^{+0.16}_{-0.14}$   \\\\ GRB 050730  &  WT  &  0.597$^{+0.033}_{-0.032}$  &  1.06$^{+0.98}_{-0.67}$   \\\\ &  PC  &  0.480$^{+0.063}_{-0.041}$  &  0.691$^{+0.035}_{-0.022}$   \\\\ GRB 050801  &  PC  &  0.946$^{+0.121}_{-0.083}$  &  0.85$^{+0.40}_{-0.28}$   \\\\ GRB 050802  &  WT  &  0.99$^{+0.19}_{-0.18}$   \\\\ &  PC  &  0.86$^{+0.12}_{-0.11}$  &  1.021$^{+0.043}_{-0.085}$   \\\\ GRB 050803  &  WT  &  0.81$^{+0.17}_{-0.15}$  &  2.16$^{+1.26}_{-0.88}$   \\\\ &  PC  &  0.39$^{+0.39}_{-0.18}$  &  0.850$^{+0.083}_{-0.081}$  &  2.01$^{+0.18}_{-0.32}$   \\\\ GRB 050814  &  WT  &  1.092$^{+0.068}_{-0.066}$  &  1.6$^{+1.6}_{-1.2}$   \\\\ &  PC  &  1.097$^{+0.100}_{-0.094}$  &  1.036$^{+0.092}_{-0.065}$  &  0.81$^{+0.21}_{-0.19}$   \\\\ GRB 050815  &  PC  &  0.65$^{+0.56}_{-0.28}$  &  0.62$^{+0.23}_{-0.33}$   \\\\ \\hline \\end{tabular} \\caption{Spectral indices for the time-resolved spectra. $\\beta_1$ corresponds to the time during which the decay followed $\\alpha_1$ in Table~6. Column densities for these spectra are given in Table~10.} \\label{tab:resbeta} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $\\beta_1$ & $\\beta_2$ & $\\beta_3$ & $\\beta_4$ & $\\beta_5$ & $\\beta_6$  \\\\ \\hline GRB 050819  &  WT  &  1.29$^{+0.27}_{-0.15}$   \\\\ &  PC  &  1.33$^{+0.32}_{-0.22}$  &  1.40$^{+0.50}_{-0.47}$   \\\\ GRB 050820A  &  WT  &  1.04$^{+0.20}_{-0.18}$  &  -0.197$^{+0.050}_{-0.038}$   \\\\ &  PC    &  &  0.853$^{+0.080}_{-0.077}$  &  1.084$^{+0.035}_{-0.034}$   \\\\ GRB 050822  &  WT  &  1.583$^{+0.074}_{-0.070}$  &  3.42$^{+2.18}_{-0.91}$  &  0.40$^{+1.25}_{-0.54}$   \\\\ &  PC  &  1.86$^{+0.70}_{-0.44}$  &  1.00$^{+0.17}_{-0.15}$  &  1.54$^{+0.17}_{-0.14}$   \\\\ GRB 050824  &  PC  &  0.90$^{+0.15}_{-0.15}$  &  1.76$^{+0.25}_{-0.41}$   \\\\ GRB 050826  &  PC  &  0.96$^{+0.34}_{-0.32}$  &  1.23$^{+0.43}_{-0.33}$  &  1.24$^{+0.55}_{-0.48}$   \\\\ GRB 050904  &  WT  &  0.365$^{+0.034}_{-0.034}$   \\\\ &  PC  &  0.652$^{+0.086}_{-0.082}$  &  0.890$^{+0.046}_{-0.038}$   \\\\ GRB050915B  &  WT  &  1.66$^{+0.20}_{-0.19}$  &  -1.76$^{+3.47}_{-0.76}$   \\\\ &  PC  &  1.08$^{+0.65}_{-0.48}$  &  1.424$^{+0.073}_{-0.192}$   \\\\ GRB050916  &  PC  &  1.47$^{+1.55}_{-0.61}$  &  0.72$^{+0.26}_{-0.23}$   \\\\ GRB 050922B  &  WT  &  1.706$^{+0.087}_{-0.084}$   \\\\ &  PC  &  1.314$^{+0.070}_{-0.067}$  &  1.37$^{+0.36}_{-0.35}$  &  1.92$^{+0.46}_{-0.69}$   \\\\ GRB 050922C  &  WT  &  1.032$^{+0.073}_{-0.072}$   \\\\ &  PC  &  0.969$^{+0.122}_{-0.075}$  &  1.319$^{+0.087}_{-0.087}$   \\\\ GRB 051001  &  WT  &  0.103$^{+0.096}_{-0.093}$  &  0.716$^{+0.082}_{-0.081}$   \\\\ &  PC    &  &  1.05$^{+0.42}_{-0.39}$  &  1.62$^{+0.40}_{-0.39}$   \\\\ GRB 051008  &  PC  &  0.90$^{+0.11}_{-0.11}$  &  1.25$^{+0.12}_{-0.33}$   \\\\ GRB 051016A  &  WT  &  2.00$^{+0.85}_{-0.72}$   \\\\ &  PC  &  3.7$^{+1.8}_{-1.4}$  &  2.12$^{+0.27}_{-0.20}$   \\\\ GRB 051016B  &  WT  &  4.00$^{+0.96}_{-0.79}$   \\\\ &  PC  &  1.14$^{+1.18}_{-0.54}$  &  1.03$^{+0.21}_{-0.21}$  &  1.047$^{+0.098}_{-0.080}$   \\\\ GRB 051021B  &  WT  &  0.29$^{+0.42}_{-0.26}$   \\\\ &  PC  &  0.51$^{+0.40}_{-0.24}$  &  1.51$^{+0.85}_{-0.72}$   \\\\ GRB 051022  &  PC  &  1.04$^{+0.11}_{-0.11}$  &  1.10$^{+0.37}_{-0.30}$   \\\\ GRB 051109A  &  WT  &  1.10$^{+0.17}_{-0.14}$  &  1.45$^{+1.13}_{-0.48}$   \\\\ &  PC    &  &  1.15$^{+0.15}_{-0.14}$  &  1.088$^{+0.030}_{-0.058}$   \\\\ GRB 051109B  &  PC  &  1.25$^{+0.28}_{-0.29}$  &  1.75$^{+0.52}_{-0.46}$  &  1.24$^{+0.35}_{-0.16}$   \\\\ GRB 051117A  &  WT  &  1.064$^{+0.018}_{-0.018}$  &  2.2$^{+2.6}_{-1.4}$   \\\\ &  PC  &  1.24$^{+0.13}_{-0.13}$    &  &  1.32$^{+0.17}_{-0.10}$   \\\\ GRB 051221A  &  PC  &  1.01$^{+0.18}_{-0.17}$  &  0.92$^{+0.18}_{-0.23}$  &  1.19$^{+0.13}_{-0.21}$   \\\\ GRB 051227  &  WT  &  0.46$^{+0.21}_{-0.21}$   \\\\ &  PC  &  0.78$^{+0.27}_{-0.25}$  &  1.27$^{+0.56}_{-0.50}$   \\\\ GRB 060105  &  WT  &  1.098$^{+0.057}_{-0.055}$  &  0.900$^{+0.031}_{-0.030}$   \\\\ &  PC    &    &  &  0.878$^{+0.065}_{-0.063}$  &  1.22$^{+0.17}_{-0.21}$  &  1.49$^{+0.15}_{-0.24}$   \\\\ GRB 060108  &  WT    &  &  0.38$^{+0.52}_{-0.26}$   \\\\ &  PC  &  0.98$^{+0.20}_{-0.18}$  &  0.96$^{+0.24}_{-0.30}$  &  1.17$^{+0.21}_{-0.24}$   \\\\ GRB 060109  &  WT  &  1.23$^{+0.18}_{-0.17}$   \\\\ &  PC  &  0.94$^{+0.45}_{-0.25}$  &  1.22$^{+0.22}_{-0.21}$  &  1.55$^{+0.16}_{-0.11}$   \\\\ GRB 060111A  &  WT  &  1.268$^{+0.039}_{-0.038}$   \\\\ &  PC  &  1.29$^{+0.34}_{-0.31}$  &  1.22$^{+0.10}_{-0.11}$   \\\\ GRB 060111B  &  WT  &  1.17$^{+0.24}_{-0.22}$  &  1.55$^{+0.62}_{-0.54}$   \\\\ &  PC    &  &  1.22$^{+0.21}_{-0.18}$  &  1.29$^{+0.16}_{-0.47}$   \\\\ GRB 060115  &  WT  &  0.852$^{+0.088}_{-0.084}$   \\\\ &  PC  &  1.002$^{+0.085}_{-0.082}$  &  1.21$^{+0.13}_{-0.11}$   \\\\ GRB 060124  &  PC  &  1.01$^{+0.82}_{-0.69}$  &  1.030$^{+0.061}_{-0.058}$  &  1.072$^{+0.050}_{-0.067}$   \\\\ GRB 060202  &  WT  &  1.372$^{+0.051}_{-0.050}$  &  1.052$^{+0.043}_{-0.043}$  &  0.992$^{+0.087}_{-0.085}$   \\\\ &  PC    &    &  &  0.72$^{+0.15}_{-0.14}$  &  2.26$^{+0.10}_{-0.18}$   \\\\ GRB 060204B  &  WT  &  0.863$^{+0.075}_{-0.074}$  &  1.70$^{+0.18}_{-0.17}$   \\\\ &  PC  &  0.76$^{+0.28}_{-0.27}$  &  1.03$^{+0.15}_{-0.14}$  &  1.81$^{+0.13}_{-0.15}$   \\\\ GRB 060206  &  WT  &  1.44$^{+1.51}_{-0.91}$   \\\\ &  PC  &  0.95$^{+0.21}_{-0.14}$  &  0.98$^{+0.17}_{-0.17}$   \\\\ GRB 060210  &  WT  &  0.934$^{+0.031}_{-0.031}$  &  1.33$^{+2.66}_{-0.66}$   \\\\ &  PC  &  1.103$^{+0.058}_{-0.056}$  &  1.047$^{+0.070}_{-0.080}$   \\\\ GRB 060211A  &  WT  &  0.942$^{+0.072}_{-0.069}$   \\\\ &  PC  &  1.10$^{+0.52}_{-0.23}$  &  0.64$^{+0.22}_{-0.13}$  &  1.44$^{+0.34}_{-0.16}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $\\beta_1$ & $\\beta_2$ & $\\beta_3$ & $\\beta_4$ & $\\beta_5$ & $\\beta_6$  \\\\ \\hline GRB 060211B  &  WT  &  1.23$^{+0.84}_{-0.38}$   \\\\ &  PC  &  0.70$^{+0.21}_{-0.25}$  &  0.88$^{+0.83}_{-0.61}$   \\\\ GRB 060218  &  WT  &  0.7474$^{+0.0102}_{-0.0099}$   \\\\ &  PC  &  1.91$^{+0.22}_{-0.21}$  &  4.13$^{+0.78}_{-0.17}$   \\\\ GRB 060219  &  PC  &  2.09$^{+1.10}_{-0.87}$  &  2.14$^{+0.26}_{-0.43}$  &  2.33$^{+0.81}_{-0.59}$   \\\\ GRB 060223A  &  WT    &  &  0.88$^{+0.42}_{-0.36}$   \\\\ &  PC  &  0.86$^{+0.32}_{-0.14}$  &  1.12$^{+0.18}_{-0.24}$   \\\\ GRB 060306  &  WT  &  1.17$^{+0.29}_{-0.27}$  &  0.91$^{+2.10}_{-1.00}$   \\\\ &  PC  &  0.66$^{+0.82}_{-0.37}$  &  1.21$^{+0.18}_{-0.17}$  &  1.301$^{+0.075}_{-0.174}$   \\\\ GRB 060312  &  WT  &  1.168$^{+0.097}_{-0.093}$   \\\\ &  PC  &  1.19$^{+0.34}_{-0.29}$  &  0.89$^{+0.20}_{-0.18}$   \\\\ GRB 060313  &  WT  &  0.74$^{+0.40}_{-0.19}$  &  1.5$^{+2.5}_{-1.2}$   \\\\ &  PC  &  0.59$^{+0.38}_{-0.17}$  &  0.71$^{+0.16}_{-0.24}$  &  0.832$^{+0.116}_{-0.067}$  &  1.38$^{+0.28}_{-0.14}$   \\\\ GRB 060319  &  PC  &  1.11$^{+0.11}_{-0.11}$  &  1.70$^{+0.16}_{-0.24}$   \\\\ GRB 060403  &  WT  &  0.39$^{+0.19}_{-0.18}$   \\\\ &  PC  &  0.78$^{+0.69}_{-0.48}$  &  1.17$^{+0.77}_{-1.16}$   \\\\ GRB 060413  &  WT  &  0.88$^{+0.10}_{-0.10}$   \\\\ &  PC  &  0.54$^{+0.18}_{-0.17}$  &  0.41$^{+0.13}_{-0.12}$  &  0.73$^{+0.32}_{-0.21}$   \\\\ GRB 060418  &  WT  &  1.143$^{+0.043}_{-0.041}$  &  1.04$^{+0.11}_{-0.11}$  &  1.03$^{+0.41}_{-0.29}$   \\\\ &  PC    &  &  0.88$^{+0.11}_{-0.11}$  &  1.31$^{+0.23}_{-0.24}$   \\\\ GRB 060427  &  WT  &  1.78$^{+0.18}_{-0.17}$   \\\\ &  PC  &  1.17$^{+0.31}_{-0.27}$  &  1.36$^{+0.13}_{-0.16}$   \\\\ GRB 060428A  &  WT  &  2.30$^{+0.50}_{-0.25}$   \\\\ &  PC  &  1.7$^{+1.8}_{-1.1}$  &  1.02$^{+0.14}_{-0.14}$  &  0.87$^{+0.11}_{-0.11}$  &  1.42$^{+0.31}_{-0.51}$   \\\\ GRB 060428B  &  WT  &  1.86$^{+0.11}_{-0.10}$   \\\\ &  PC  &  1.13$^{+0.40}_{-0.21}$  &  1.15$^{+0.21}_{-0.20}$   \\\\ GRB 060502A  &  WT  &  2.48$^{+0.27}_{-0.25}$   \\\\ &  PC  &  1.34$^{+0.17}_{-0.15}$  &  1.02$^{+0.13}_{-0.12}$  &  0.95$^{+0.15}_{-0.17}$   \\\\ GRB 060507  &  WT    &  &  1.6$^{+2.8}_{-5.4}$   \\\\ &  PC  &  0.83$^{+0.27}_{-0.16}$  &  1.16$^{+0.12}_{-0.19}$   \\\\ GRB 060510A  &  WT  &  2.30$^{+0.20}_{-0.28}$   \\\\ &  PC    &  &  0.871$^{+0.044}_{-0.043}$  &  0.975$^{+0.040}_{-0.085}$   \\\\ GRB 060510B  &  WT  &  0.190$^{+0.023}_{-0.023}$  &  0.814$^{+0.073}_{-0.073}$   \\\\ &  PC    &  &  0.92$^{+0.21}_{-0.19}$  &  1.243$^{+0.131}_{-0.099}$   \\\\ GRB 060512  &  WT  &  4.53$^{+1.13}_{-0.96}$   \\\\ &  PC  &  1.85$^{+0.42}_{-0.23}$  &  1.80$^{+0.12}_{-0.12}$   \\\\ GRB 060522  &  WT  &  0.51$^{+0.21}_{-0.19}$   \\\\ &  PC  &  0.65$^{+0.28}_{-0.23}$  &  1.17$^{+0.21}_{-0.14}$  &  1.29$^{+0.37}_{-0.29}$   \\\\ GRB 060526  &  WT  &  0.892$^{+0.035}_{-0.035}$   \\\\ &  PC  &  0.77$^{+0.13}_{-0.13}$  &  1.09$^{+0.20}_{-0.14}$  &  1.15$^{+0.21}_{-0.25}$  &  1.04$^{+0.63}_{-0.36}$   \\\\ GRB 060604  &  WT  &  1.173$^{+0.050}_{-0.050}$   \\\\ &  PC    &  &  1.14$^{+0.18}_{-0.18}$  &  1.24$^{+0.14}_{-0.16}$   \\\\ GRB 060605  &  WT  &  0.63$^{+0.37}_{-0.32}$   \\\\ &  PC  &  0.57$^{+0.13}_{-0.12}$  &  1.12$^{+0.14}_{-0.13}$  &  1.23$^{+0.12}_{-0.10}$   \\\\ GRB 060607A  &  WT  &  0.705$^{+0.032}_{-0.032}$  &  0.71$^{+0.25}_{-0.16}$  &  0.60$^{+0.84}_{-0.61}$   \\\\ &  PC    &  &  0.483$^{+0.036}_{-0.031}$  &  0.613$^{+0.035}_{-0.050}$   \\\\ GRB 060614  &  WT  &  0.254$^{+0.052}_{-0.051}$  &  1.332$^{+0.041}_{-0.040}$  &  1.07$^{+0.35}_{-0.21}$   \\\\ &  PC    &    &  &  0.822$^{+0.093}_{-0.057}$  &  0.98$^{+0.13}_{-0.14}$   \\\\ GRB 060707  &  WT  &  0.89$^{+0.23}_{-0.22}$   \\\\ &  PC  &  1.01$^{+0.14}_{-0.11}$  &  1.31$^{+0.22}_{-0.22}$  &  0.99$^{+0.22}_{-0.21}$   \\\\ GRB 060708  &  WT  &  2.26$^{+0.19}_{-0.18}$   \\\\ &  PC    &  &  1.22$^{+0.16}_{-0.14}$  &  1.19$^{+0.11}_{-0.11}$   \\\\ GRB 060712  &  PC  &  1.90$^{+0.51}_{-0.44}$  &  1.67$^{+0.62}_{-0.55}$  &  1.72$^{+0.17}_{-0.18}$   \\\\ GRB 060714  &  WT  &  0.816$^{+0.044}_{-0.043}$   \\\\ &  PC  &  1.38$^{+0.42}_{-0.30}$  &  0.93$^{+0.13}_{-0.13}$  &  1.32$^{+0.16}_{-0.17}$   \\\\ GRB 060719  &  WT  &  2.20$^{+0.37}_{-0.32}$   \\\\ &  PC  &  1.19$^{+0.74}_{-0.61}$  &  2.08$^{+0.27}_{-0.25}$  &  1.85$^{+0.13}_{-0.19}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $\\beta_1$ & $\\beta_2$ & $\\beta_3$ & $\\beta_4$ & $\\beta_5$ & $\\beta_6$  \\\\ \\hline GRB 060729  &  WT  &  2.06$^{+0.10}_{-0.10}$  &  2.96$^{+0.18}_{-0.16}$   \\\\ &  PC    &  &  1.96$^{+0.28}_{-0.28}$  &  1.060$^{+0.045}_{-0.044}$  &  1.088$^{+0.030}_{-0.041}$  &  0.867$^{+0.087}_{-0.120}$   \\\\ GRB 060801  &  PC  &  0.59$^{+0.30}_{-0.29}$  &  0.70$^{+0.26}_{-0.23}$  &  5.3$^{+2.4}_{-2.6}$   \\\\ GRB 060804  &  PC  &  1.09$^{+0.21}_{-0.19}$  &  1.27$^{+0.18}_{-0.34}$   \\\\ GRB 060805A  &  PC  &  1.28$^{+0.45}_{-0.39}$  &  1.17$^{+0.36}_{-0.59}$   \\\\ GRB 060807  &  PC  &  0.95$^{+0.56}_{-0.34}$  &  0.91$^{+0.21}_{-0.20}$  &  1.17$^{+0.12}_{-0.11}$  &  1.80$^{+0.27}_{-0.35}$   \\\\ GRB 060813  &  WT  &  0.884$^{+0.078}_{-0.078}$   \\\\ &  PC    &  &  0.882$^{+0.085}_{-0.081}$  &  1.11$^{+0.27}_{-0.40}$   \\\\ GRB 060814  &  WT  &  0.36$^{+0.14}_{-0.14}$  &  0.784$^{+0.039}_{-0.038}$  &  1.16$^{+0.13}_{-0.13}$   \\\\ &  PC    &    &  &  0.95$^{+0.21}_{-0.20}$  &  1.200$^{+0.092}_{-0.090}$  &  1.404$^{+0.067}_{-0.137}$   \\\\ GRB 060904A  &  WT  &  0.962$^{+0.058}_{-0.056}$  &  2.04$^{+0.14}_{-0.13}$   \\\\ &  PC    &    &  &  1.63$^{+0.47}_{-0.25}$   \\\\ GRB 060904B  &  WT  &  1.095$^{+0.043}_{-0.041}$   \\\\ &  PC  &  1.00$^{+0.25}_{-0.22}$  &  1.33$^{+0.13}_{-0.14}$   \\\\ GRB 060906  &  WT  &  1.36$^{+0.64}_{-0.50}$  &  1.57$^{+0.68}_{-0.38}$   \\\\ &  PC    &  &  1.20$^{+0.22}_{-0.22}$  &  1.04$^{+0.17}_{-0.14}$   \\\\ GRB 060908  &  WT  &  1.24$^{+0.20}_{-0.19}$  &  1.87$^{+1.21}_{-0.87}$   \\\\ &  PC  &  1.14$^{+0.27}_{-0.26}$  &  0.95$^{+0.20}_{-0.19}$   \\\\ GRB 060912A  &  PC  &  0.87$^{+0.15}_{-0.15}$  &  1.13$^{+0.15}_{-0.10}$   \\\\ GRB 060923A  &  WT  &  0.91$^{+0.34}_{-0.32}$  &  1.20$^{+1.37}_{-0.59}$  &  1.20$^{+0.31}_{-0.25}$   \\\\ &  PC  &  0.55$^{+0.40}_{-0.35}$  &  0.62$^{+0.52}_{-0.22}$  &  1.42$^{+0.26}_{-0.16}$   \\\\ GRB 060923C  &  WT  &  0.87$^{+0.22}_{-0.12}$   \\\\ &  PC  &  0.752$^{+0.101}_{-0.089}$  &  1.93$^{+0.28}_{-0.32}$  &  1.32$^{+0.70}_{-0.28}$   \\\\ GRB 060926  &  WT  &  0.78$^{+0.78}_{-0.42}$  &  1.03$^{+0.26}_{-0.25}$   \\\\ &  PC    &    &  &  1.18$^{+0.27}_{-0.30}$   \\\\ GRB 060927  &  PC  &  0.81$^{+0.22}_{-0.17}$  &  1.04$^{+0.38}_{-0.31}$   \\\\ GRB 060929  &  WT    &  &  0.777$^{+0.050}_{-0.049}$   \\\\ &  PC  &  1.48$^{+0.69}_{-0.59}$  &  1.00$^{+0.17}_{-0.33}$   \\\\ GRB 061002  &  WT  &  1.35$^{+0.39}_{-0.36}$  &  1.25$^{+0.91}_{-0.39}$   \\\\ &  PC    &  &  1.48$^{+0.71}_{-0.64}$   \\\\ GRB 061004  &  WT  &  0.70$^{+0.54}_{-0.25}$   \\\\ &  PC  &  0.37$^{+0.29}_{-0.14}$  &  1.35$^{+0.26}_{-0.31}$  &  1.23$^{+0.24}_{-0.23}$   \\\\ GRB 061007  &  WT  &  0.832$^{+0.064}_{-0.062}$  &  0.893$^{+0.025}_{-0.025}$   \\\\ &  PC    &  &  0.99$^{+0.33}_{-0.29}$  &  0.97$^{+0.15}_{-0.15}$  &  1.30$^{+0.13}_{-0.12}$   \\\\ GRB 061019  &  PC  &  1.035$^{+0.064}_{-0.239}$  &  1.44$^{+1.10}_{-0.68}$   \\\\ GRB 061021  &  WT  &  0.965$^{+0.052}_{-0.047}$   \\\\ &  PC  &  0.53$^{+0.21}_{-0.11}$  &  1.08$^{+0.11}_{-0.10}$  &  1.052$^{+0.066}_{-0.071}$   \\\\ GRB 061028  &  WT  &  0.39$^{+0.16}_{-0.15}$   \\\\ &  PC  &  0.73$^{+0.18}_{-0.17}$  &  9.0$^{+2.3}_{-1.5}$   \\\\ GRB 061102  &  WT  &  2.21$^{+0.40}_{-0.35}$   \\\\ &  PC    &  &  0.58$^{+0.55}_{-0.51}$   \\\\ GRB 061110A  &  WT  &  2.071$^{+0.094}_{-0.087}$  &  2.24$^{+0.14}_{-0.13}$   \\\\ &  PC    &  &  1.05$^{+0.25}_{-0.20}$  &  0.937$^{+0.110}_{-0.098}$  &  1.00$^{+0.32}_{-0.32}$   \\\\ GRB 061121  &  WT  &  -0.035$^{+0.078}_{-0.075}$  &  1.255$^{+0.101}_{-0.097}$  &  1.11$^{+0.14}_{-0.13}$  &  0.960$^{+0.084}_{-0.081}$  &  1.14$^{+0.29}_{-0.28}$   \\\\ &  PC    &    &    &  &  0.86$^{+0.18}_{-0.16}$  &  0.935$^{+0.077}_{-0.075}$  &  0.899$^{+0.065}_{-0.061}$   \\\\ GRB 061126  &  WT  &  0.86$^{+0.16}_{-0.15}$   \\\\ &  PC  &  0.952$^{+0.078}_{-0.075}$  &  0.936$^{+0.123}_{-0.075}$   \\\\ GRB 061201  &  WT  &  -0.03$^{+0.64}_{-0.26}$  &  0.08$^{+1.23}_{-0.54}$   \\\\ &  PC  &  0.33$^{+0.11}_{-0.10}$  &  1.26$^{+0.40}_{-0.22}$   \\\\ GRB 061202  &  WT  &  0.912$^{+0.084}_{-0.082}$   \\\\ &  PC  &  1.21$^{+0.21}_{-0.20}$  &  1.11$^{+0.15}_{-0.15}$  &  1.463$^{+0.097}_{-0.163}$   \\\\ GRB 061222A  &  WT  &  1.203$^{+0.079}_{-0.077}$  &  1.24$^{+0.39}_{-0.35}$  &  1.49$^{+1.27}_{-0.84}$   \\\\ &  PC    &  &  0.871$^{+0.070}_{-0.067}$  &  1.059$^{+0.083}_{-0.080}$  &  1.352$^{+0.039}_{-0.123}$   \\\\ GRB 070103  &  WT    &  &  1.81$^{+1.21}_{-0.88}$   \\\\ &  PC  &  1.36$^{+0.43}_{-0.39}$  &  1.05$^{+0.50}_{-0.44}$  &  1.43$^{+0.33}_{-0.31}$   \\\\ GRB 070107  &  WT  &  0.904$^{+0.060}_{-0.059}$   \\\\ &  PC  &  1.088$^{+0.102}_{-0.099}$  &  1.168$^{+0.076}_{-0.055}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $\\beta_1$ & $\\beta_2$ & $\\beta_3$ & $\\beta_4$ & $\\beta_5$ & $\\beta_6$  \\\\ \\hline GRB 070110  &  WT  &  0.906$^{+0.090}_{-0.086}$  &  1.25$^{+0.21}_{-0.17}$   \\\\ &  PC  &  1.08$^{+0.23}_{-0.19}$  &  0.81$^{+0.12}_{-0.11}$  &  1.124$^{+0.076}_{-0.073}$  &  1.13$^{+0.18}_{-0.15}$  &  1.16$^{+0.20}_{-0.16}$   \\\\ GRB 070125  &  PC  &  1.14$^{+0.21}_{-0.23}$  &  1.17$^{+0.25}_{-0.14}$   \\\\ GRB 070129  &  WT  &  0.862$^{+0.027}_{-0.027}$   \\\\ &  PC  &  1.46$^{+0.11}_{-0.11}$  &  1.53$^{+0.28}_{-0.27}$  &  1.29$^{+0.13}_{-0.16}$  &  1.65$^{+1.26}_{-0.59}$   \\\\ GRB 070208  &  PC  &  1.35$^{+0.44}_{-0.26}$  &  1.26$^{+0.32}_{-0.31}$  &  1.17$^{+0.26}_{-0.37}$   \\\\ GRB 070219  &  WT  &  1.1$^{+4.4}_{-2.0}$   \\\\ &  PC  &  0.61$^{+0.72}_{-0.56}$  &  1.4$^{+1.7}_{-1.3}$  &  2.06$^{+0.47}_{-0.31}$   \\\\ GRB 070220  &  WT  &  0.446$^{+0.098}_{-0.096}$  &  0.44$^{+0.13}_{-0.12}$   \\\\ &  PC    &    &  &  0.460$^{+0.081}_{-0.080}$  &  0.65$^{+0.22}_{-0.22}$   \\\\ GRB 070223  &  WT  &  0.56$^{+0.14}_{-0.14}$  &  1.08$^{+0.17}_{-0.16}$   \\\\ &  PC    &    &  &  1.68$^{+0.28}_{-0.19}$   \\\\ GRB 070224  &  WT  &  1.43$^{+0.20}_{-0.16}$   \\\\ &  PC  &  1.34$^{+0.15}_{-0.15}$  &  0.96$^{+0.37}_{-0.20}$   \\\\ GRB 070306  &  WT  &  1.137$^{+0.092}_{-0.089}$  &  1.366$^{+0.091}_{-0.087}$   \\\\ &  PC    &  &  1.00$^{+0.53}_{-0.24}$  &  1.11$^{+0.11}_{-0.11}$  &  1.16$^{+0.19}_{-0.18}$  &  1.42$^{+0.17}_{-0.12}$   \\\\ GRB 070311  &  PC  &  1.01$^{+0.23}_{-0.24}$  &  1.22$^{+0.19}_{-0.39}$  &  1.14$^{+0.98}_{-0.67}$   \\\\ GRB 070328  &  WT  &  1.107$^{+0.089}_{-0.084}$  &  1.271$^{+0.056}_{-0.055}$  &  1.238$^{+0.065}_{-0.064}$   \\\\ &  PC    &    &  &  1.22$^{+0.61}_{-0.56}$  &  1.219$^{+0.037}_{-0.083}$   \\\\ GRB 070412  &  WT  &  1.16$^{+0.40}_{-0.34}$   \\\\ &  PC    &  &  0.86$^{+0.19}_{-0.30}$  &  1.74$^{+0.55}_{-0.50}$   \\\\ GRB 070419A  &  WT    &  &  1.15$^{+0.13}_{-0.13}$   \\\\ &  PC    &  &  1.11$^{+0.20}_{-0.19}$  &  1.50$^{+2.59}_{-0.58}$   \\\\ GRB 070419B  &  WT  &  0.801$^{+0.030}_{-0.029}$   \\\\ &  PC  &  0.520$^{+0.044}_{-0.043}$  &  0.77$^{+0.14}_{-0.18}$   \\\\ GRB 070420  &  WT  &  1.50$^{+0.12}_{-0.11}$  &  0.63$^{+0.82}_{-0.40}$   \\\\ &  PC    &  &  1.023$^{+0.081}_{-0.079}$  &  0.91$^{+0.12}_{-0.11}$  &  1.07$^{+0.18}_{-0.17}$  &  1.25$^{+0.20}_{-0.26}$   \\\\ GRB 070429A  &  WT  &  2.83$^{+0.25}_{-0.23}$  &  2.08$^{+0.67}_{-0.42}$   \\\\ &  PC    &  &  0.77$^{+0.39}_{-0.17}$  &  1.21$^{+0.24}_{-0.26}$  &  0.48$^{+0.62}_{-0.40}$   \\\\ GRB 070508  &  WT  &  0.892$^{+0.047}_{-0.047}$  &  0.905$^{+0.063}_{-0.063}$   \\\\ &  PC    &  &  0.84$^{+0.15}_{-0.14}$  &  1.42$^{+0.13}_{-0.30}$   \\\\ GRB 070517  &  PC  &  1.15$^{+0.26}_{-0.25}$  &  0.98$^{+0.49}_{-0.36}$   \\\\ GRB 070518  &  WT  &  1.56$^{+0.12}_{-0.12}$   \\\\ &  PC  &  1.17$^{+0.27}_{-0.17}$  &  2.09$^{+0.31}_{-0.39}$   \\\\ GRB 070520A  &  WT  &  2.14$^{+0.20}_{-0.19}$   \\\\ &  PC    &  &  2.36$^{+0.52}_{-0.33}$   \\\\ GRB 070520B  &  WT  &  1.517$^{+0.065}_{-0.062}$   \\\\ &  PC  &  1.05$^{+0.63}_{-0.56}$  &  1.97$^{+0.56}_{-0.47}$   \\\\ GRB 070521  &  WT  &  1.18$^{+1.03}_{-0.83}$   \\\\ &  PC  &  0.86$^{+0.40}_{-0.37}$  &  0.83$^{+0.13}_{-0.12}$  &  1.52$^{+0.11}_{-0.20}$   \\\\ GRB 070529  &  WT  &  1.93$^{+0.79}_{-0.62}$   \\\\ &  PC  &  1.08$^{+0.92}_{-0.56}$  &  0.81$^{+0.12}_{-0.11}$  &  1.38$^{+0.26}_{-0.14}$   \\\\ GRB 070616  &  WT  &  0.079$^{+0.023}_{-0.023}$  &  0.704$^{+0.039}_{-0.039}$   \\\\ &  PC    &  &  0.90$^{+0.25}_{-0.23}$  &  1.56$^{+0.15}_{-0.21}$   \\\\ GRB 070621  &  WT  &  1.398$^{+0.093}_{-0.091}$   \\\\ &  PC  &  1.39$^{+0.37}_{-0.35}$  &  1.69$^{+0.26}_{-0.25}$  &  2.13$^{+0.30}_{-0.22}$   \\\\ GRB 070628  &  WT  &  3.2$^{+2.1}_{-1.3}$   \\\\ &  PC  &  0.94$^{+0.15}_{-0.14}$  &  1.076$^{+0.099}_{-0.093}$   \\\\ GRB 070704  &  WT  &  0.734$^{+0.061}_{-0.061}$   \\\\ &  PC  &  0.49$^{+0.92}_{-0.28}$  &  1.16$^{+0.31}_{-0.47}$   \\\\ GRB 070714B  &  WT  &  0.154$^{+0.100}_{-0.097}$  &  0.78$^{+0.21}_{-0.19}$   \\\\ &  PC    &  &  0.97$^{+0.16}_{-0.15}$  &  1.28$^{+0.29}_{-0.28}$  &  0.88$^{+0.13}_{-0.14}$   \\\\ GRB 070721A  &  WT  &  2.29$^{+0.95}_{-0.61}$   \\\\ &  PC  &  1.66$^{+0.23}_{-0.22}$  &  1.96$^{+0.34}_{-0.39}$   \\\\ GRB 070721B  &  WT  &  0.72$^{+0.17}_{-0.15}$  &  0.302$^{+0.107}_{-0.088}$   \\\\ &  PC    &  &  0.320$^{+0.113}_{-0.047}$  &  1.12$^{+0.15}_{-0.19}$   \\\\ GRB 070724A  &  WT  &  0.58$^{+0.18}_{-0.18}$   \\\\ &  PC  &  0.82$^{+0.27}_{-0.26}$  &  1.57$^{+0.60}_{-0.49}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $\\beta_1$ & $\\beta_2$ & $\\beta_3$ & $\\beta_4$ & $\\beta_5$ & $\\beta_6$  \\\\ \\hline GRB 070802  &  PC  &  0.79$^{+0.19}_{-0.17}$  &  1.22$^{+0.43}_{-0.41}$  &  1.43$^{+0.42}_{-0.43}$   \\\\ GRB 070808  &  WT  &  0.61$^{+0.66}_{-0.56}$   \\\\ &  PC  &  0.32$^{+0.44}_{-0.40}$  &  1.58$^{+0.42}_{-0.34}$   \\\\ GRB 070810A  &  PC  &  1.11$^{+0.36}_{-0.33}$  &  1.15$^{+0.12}_{-0.15}$   \\\\ GRB 071001  &  WT  &  1.95$^{+0.24}_{-0.22}$   \\\\ &  PC  &  1.52$^{+0.18}_{-0.24}$  &  1.35$^{+0.58}_{-0.29}$   \\\\ GRB 071010A  &  PC  &  1.96$^{+0.36}_{-0.38}$  &  1.40$^{+0.40}_{-0.33}$   \\\\ GRB 071011  &  PC  &  1.47$^{+0.13}_{-0.34}$  &  0.73$^{+0.50}_{-0.52}$   \\\\ GRB 071020  &  WT  &  0.98$^{+0.12}_{-0.11}$  &  0.818$^{+0.052}_{-0.047}$   \\\\ &  PC    &  &  0.742$^{+0.041}_{-0.036}$   \\\\ GRB 071021  &  WT  &  1.115$^{+0.090}_{-0.088}$  &  1.06$^{+0.69}_{-0.58}$   \\\\ &  PC    &  &  1.288$^{+0.052}_{-0.101}$   \\\\ GRB 071025  &  WT  &  0.398$^{+0.076}_{-0.074}$  &  0.363$^{+0.059}_{-0.058}$  &  0.67$^{+0.13}_{-0.12}$   \\\\ &  PC    &    &  &  1.28$^{+1.80}_{-0.75}$  &  1.278$^{+0.045}_{-0.076}$   \\\\ GRB 071028A  &  WT  &  1.10$^{+0.25}_{-0.22}$   \\\\ &  PC  &  0.97$^{+0.28}_{-0.19}$  &  1.03$^{+0.34}_{-0.23}$  &  0.99$^{+0.65}_{-0.33}$   \\\\ GRB 071031  &  WT  &  0.849$^{+0.026}_{-0.026}$   \\\\ &  PC  &  0.85$^{+0.20}_{-0.12}$  &  0.96$^{+0.20}_{-0.23}$  &  0.79$^{+0.71}_{-0.47}$   \\\\ GRB 071112C  &  WT  &  0.837$^{+0.087}_{-0.083}$  &  0.612$^{+0.103}_{-0.088}$   \\\\ &  PC    &  &  0.66$^{+0.17}_{-0.12}$  &  0.95$^{+0.18}_{-0.21}$   \\\\ GRB 071118  &  WT  &  0.69$^{+0.15}_{-0.14}$   \\\\ &  PC  &  0.81$^{+0.33}_{-0.31}$  &  0.654$^{+0.075}_{-0.073}$  &  0.67$^{+0.23}_{-0.28}$   \\\\ GRB 071227  &  WT  &  0.60$^{+0.14}_{-0.14}$   \\\\ &  PC    &  &  0.76$^{+0.30}_{-0.27}$  &  1.97$^{+1.91}_{-0.62}$   \\\\ GRB 080123  &  WT  &  0.72$^{+0.11}_{-0.10}$   \\\\ &  PC    &  &  1.37$^{+0.38}_{-0.35}$  &  2.28$^{+0.42}_{-0.99}$   \\\\ GRB 080129  &  PC  &  0.94$^{+0.23}_{-0.16}$  &  1.16$^{+0.30}_{-0.28}$   \\\\ GRB 080205  &  WT  &  1.09$^{+0.16}_{-0.15}$  &  1.49$^{+0.43}_{-0.36}$  &  0.77$^{+0.55}_{-0.29}$   \\\\ &  PC    &    &  &  0.66$^{+0.27}_{-0.26}$  &  0.856$^{+0.074}_{-0.153}$   \\\\ GRB 080207  &  WT  &  0.300$^{+0.090}_{-0.089}$   \\\\ &  PC    &  &  1.52$^{+0.11}_{-0.15}$   \\\\ GRB 080212  &  WT  &  2.371$^{+0.092}_{-0.090}$  &  1.689$^{+0.060}_{-0.058}$   \\\\ &  PC    &  &  1.089$^{+0.176}_{-0.087}$  &  1.485$^{+0.075}_{-0.185}$   \\\\ GRB 080229A  &  WT  &  1.932$^{+0.095}_{-0.090}$  &  1.34$^{+0.74}_{-0.64}$   \\\\ &  PC    &  &  0.86$^{+0.13}_{-0.12}$  &  0.769$^{+0.046}_{-0.045}$  &  0.82$^{+0.14}_{-0.27}$   \\\\ GRB 080303  &  WT  &  0.94$^{+0.20}_{-0.13}$   \\\\ &  PC  &  0.957$^{+0.076}_{-0.087}$  &  1.10$^{+0.27}_{-0.30}$   \\\\ GRB 080307  &  WT  &  0.386$^{+0.083}_{-0.079}$  &  0.528$^{+0.080}_{-0.078}$   \\\\ &  PC    &  &  0.647$^{+0.110}_{-0.091}$  &  1.15$^{+0.32}_{-0.60}$   \\\\ GRB 080310  &  WT  &  0.514$^{+0.020}_{-0.020}$   \\\\ &  PC    &  &  1.381$^{+0.076}_{-0.071}$  &  0.960$^{+0.141}_{-0.094}$   \\\\ GRB 080319B  &  WT  &  0.722$^{+0.035}_{-0.034}$  &  0.764$^{+0.013}_{-0.013}$   \\\\ &  PC    &  &  0.793$^{+0.079}_{-0.078}$  &  0.827$^{+0.105}_{-0.099}$  &  0.922$^{+0.087}_{-0.084}$  &  0.84$^{+0.32}_{-0.18}$   \\\\ GRB 080319C  &  PC  &  0.68$^{+0.26}_{-0.25}$  &  0.573$^{+0.044}_{-0.043}$  &  1.08$^{+0.31}_{-0.18}$   \\\\ GRB 080319D  &  WT  &  -3.36$^{+2.36}_{-0.63}$   \\\\ &  PC  &  0.90$^{+0.12}_{-0.11}$  &  2.5$^{+1.1}_{-1.7}$   \\\\ GRB 080320  &  PC  &  0.84$^{+0.13}_{-0.13}$  &  1.040$^{+0.076}_{-0.085}$  &  1.50$^{+0.24}_{-0.14}$   \\\\ GRB 080325  &  WT  &  1.338$^{+0.043}_{-0.042}$   \\\\ &  PC    &  &  1.45$^{+0.14}_{-0.26}$   \\\\ GRB 080328  &  WT  &  1.16$^{+0.12}_{-0.11}$  &  1.30$^{+0.19}_{-0.18}$   \\\\ &  PC    &  &  0.833$^{+0.068}_{-0.066}$  &  1.055$^{+0.098}_{-0.094}$   \\\\ GRB 080330  &  WT  &  1.19$^{+0.20}_{-0.18}$   \\\\ &  PC    &  &  0.87$^{+0.18}_{-0.14}$  &  1.19$^{+0.58}_{-0.48}$   \\\\ GRB 080409  &  PC  &  0.7$^{+4.4}_{-4.6}$  &  0.62$^{+0.13}_{-0.26}$   \\\\ GRB 080411  &  PC  &  1.00$^{+0.12}_{-0.12}$  &  1.027$^{+0.040}_{-0.023}$  &  0.62$^{+0.15}_{-0.18}$   \\\\ GRB 080413A  &  WT  &  2.13$^{+0.16}_{-0.16}$   \\\\ &  PC  &  1.52$^{+0.45}_{-0.36}$  &  1.09$^{+0.22}_{-0.26}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage   \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $\\beta_1$ & $\\beta_2$ & $\\beta_3$ & $\\beta_4$ & $\\beta_5$ & $\\beta_6$  \\\\ \\hline GRB 080413B  &  WT  &  1.03$^{+0.10}_{-0.10}$   \\\\ &  PC  &  0.94$^{+0.32}_{-0.30}$  &  0.889$^{+0.051}_{-0.050}$  &  1.16$^{+0.13}_{-0.17}$   \\\\ GRB 080426  &  PC  &  0.63$^{+0.40}_{-0.30}$  &  1.23$^{+0.18}_{-0.27}$   \\\\ GRB 080430  &  WT  &  1.39$^{+0.20}_{-0.17}$   \\\\ &  PC  &  0.64$^{+0.23}_{-0.22}$  &  1.109$^{+0.094}_{-0.063}$  &  1.11$^{+0.11}_{-0.10}$   \\\\ GRB 080503  &  WT  &  0.114$^{+0.062}_{-0.062}$  &  0.724$^{+0.099}_{-0.095}$  &  1.39$^{+0.42}_{-0.37}$   \\\\ &  PC    &    &  &  1.21$^{+0.17}_{-0.10}$   \\\\ GRB 080506  &  WT  &  0.285$^{+0.085}_{-0.084}$  &  0.85$^{+0.12}_{-0.12}$  &  1.22$^{+0.12}_{-0.11}$   \\\\ &  PC    &    &  &  0.881$^{+0.179}_{-0.083}$  &  1.28$^{+0.23}_{-0.13}$   \\\\ GRB 080516  &  PC  &  1.37$^{+0.50}_{-0.46}$  &  1.18$^{+0.26}_{-0.32}$   \\\\ GRB 080517  &  WT  &  1.22$^{+1.03}_{-0.79}$   \\\\ &  PC  &  0.72$^{+0.33}_{-0.29}$  &  1.54$^{+0.63}_{-0.67}$   \\\\ GRB 080523  &  WT  &  1.52$^{+0.11}_{-0.10}$  &  1.479$^{+0.166}_{-0.082}$   \\\\ &  PC    &  &  0.78$^{+0.17}_{-0.16}$  &  1.33$^{+0.43}_{-0.43}$   \\\\ GRB 080602  &  WT  &  0.60$^{+0.36}_{-0.33}$  &  1.25$^{+0.44}_{-0.38}$   \\\\ &  PC    &  &  0.90$^{+0.14}_{-0.13}$  &  0.775$^{+0.080}_{-0.077}$   \\\\ GRB 080603B  &  WT  &  0.49$^{+0.13}_{-0.13}$  &  0.788$^{+0.084}_{-0.081}$  &  0.95$^{+0.17}_{-0.16}$   \\\\ &  PC    &    &  &  0.89$^{+0.19}_{-0.18}$   \\\\ GRB 080604  &  WT  &  1.069$^{+0.068}_{-0.063}$   \\\\ &  PC  &  0.63$^{+0.11}_{-0.13}$  &  1.13$^{+0.56}_{-0.43}$   \\\\ GRB 080605  &  WT  &  0.737$^{+0.052}_{-0.051}$  &  0.772$^{+0.049}_{-0.048}$   \\\\ &  PC    &  &  0.69$^{+0.18}_{-0.18}$  &  1.32$^{+0.20}_{-0.22}$   \\\\ GRB 080607  &  WT  &  1.00$^{+0.13}_{-0.13}$  &  1.045$^{+0.041}_{-0.040}$  &  1.155$^{+0.098}_{-0.094}$   \\\\ &  PC    &    &  &  0.950$^{+0.066}_{-0.046}$  &  1.21$^{+0.55}_{-0.50}$   \\\\ GRB 080613B  &  WT  &  0.172$^{+0.081}_{-0.079}$  &  1.66$^{+1.50}_{-0.61}$   \\\\ &  PC    &    &  &  1.22$^{+0.77}_{-0.57}$   \\\\ GRB 080625  &  PC  &  1.31$^{+0.11}_{-0.22}$  &  1.65$^{+0.62}_{-0.65}$   \\\\ GRB 080703  &  PC  &  0.618$^{+0.098}_{-0.073}$  &  0.38$^{+0.19}_{-0.23}$   \\\\ GRB 080707  &  WT  &  1.29$^{+0.77}_{-0.47}$  &  3.7$^{+4.2}_{-3.4}$   \\\\ &  PC  &  0.79$^{+0.41}_{-0.34}$  &  1.44$^{+0.38}_{-0.36}$  &  1.00$^{+0.15}_{-0.21}$   \\\\ GRB 080710  &  PC  &  0.969$^{+0.113}_{-0.081}$  &  1.23$^{+0.16}_{-0.23}$   \\\\ GRB 080714  &  WT  &  0.54$^{+0.22}_{-0.14}$  &  0.54$^{+0.17}_{-0.11}$   \\\\ &  PC    &  &  0.73$^{+0.53}_{-0.47}$  &  0.762$^{+0.099}_{-0.208}$   \\\\ GRB 080721  &  WT  &  0.834$^{+0.028}_{-0.028}$  &  0.87$^{+0.12}_{-0.11}$  &  0.831$^{+0.019}_{-0.019}$  &  1.04$^{+0.52}_{-0.43}$   \\\\ &  PC    &    &    &  &  1.013$^{+0.048}_{-0.035}$   \\\\ GRB 080723A  &  PC  &  1.40$^{+0.59}_{-0.52}$  &  1.15$^{+0.30}_{-0.29}$  &  1.11$^{+0.12}_{-0.24}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage     \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $N_{\\rm H,1}$ & $N_{\\rm H,2}$ & $N_{\\rm H,3}$ & $N_{\\rm H,4}$ & $N_{\\rm H,5}$ & $N_{\\rm H,6}$  \\\\ \\hline GRB 050126  &  PC  &  $<10.43$  &  5.5$^{+37.1}_{-5.5}$   \\\\ GRB 050128  &  WT    &  &  $<123.27$   \\\\ &  PC  &  6.3$^{+3.0}_{-2.9}$  &  4.0$^{+1.6}_{-1.7}$   \\\\ GRB 050219A  &  WT  &  15.3$^{+5.2}_{-4.8}$  &  15.2$^{+8.9}_{-7.5}$   \\\\ &  PC    &  &  29.9$^{+28.2}_{-20.7}$   \\\\ GRB 050315  &  WT    &  &  112.5$^{+85.2}_{-65.7}$   \\\\ &  PC  &  35.8$^{+18.9}_{-17.8}$  &  79.1$^{+19.7}_{-17.1}$  &  100.0$^{+19.9}_{-14.3}$   \\\\ GRB 050318  &  WT  &  15.9$^{+84.8}_{-15.9}$   \\\\ &  PC  &  11.5$^{+10.0}_{-9.6}$  &  16.8$^{+12.8}_{-10.8}$   \\\\ GRB 050319  &  WT    &    &  &  97.3$^{+167.0}_{-97.3}$   \\\\ &  PC  &  $<16.50$  &  18.5$^{+28.0}_{-18.5}$  &  33.6$^{+41.7}_{-23.3}$  &  $(7.0^{+11.1}_{-2.1}) \\tim{3}$   \\\\ GRB 050401  &  WT  &  155.6$^{+19.2}_{-18.3}$  &  120.7$^{+115.2}_{-94.6}$   \\\\ &  PC  &  141.4$^{+91.1}_{-78.5}$  &  236.2$^{+209.0}_{-97.6}$   \\\\ GRB 050406  &  WT  &  $<3.29$   \\\\ &  PC  &  18.1$^{+121.6}_{-9.0}$  &  $(3.2^{+unconstr.}_{-unconstr.}) \\tim{4}$   \\\\ GRB 050410  &  WT  &  10.1$^{+19.3}_{-10.1}$   \\\\ &  PC  &  15.2$^{+39.7}_{-15.2}$  &  63.8$^{+unconstr.}_{-unconstr.}$   \\\\ GRB 050416A  &  PC  &  12.9$^{+13.6}_{-12.2}$  &  56.1$^{+18.7}_{-16.0}$  &  76.0$^{+10.2}_{-10.6}$   \\\\ GRB 050421  &  WT  &  80.0$^{+35.3}_{-27.0}$  &  19.4$^{+45.9}_{-19.4}$   \\\\ &  PC    &  &  47.5$^{+6.1}_{-14.6}$   \\\\ GRB 050422  &  PC  &  $<10.13$  &  $<19.43$   \\\\ GRB 050502B  &  WT  &  7.05$^{+0.67}_{-0.66}$   \\\\ &  PC  &  0.68$^{+1.78}_{-0.68}$  &  11.4$^{+11.0}_{-11.3}$   \\\\ GRB 050505  &  PC  &  119.7$^{+117.7}_{-107.1}$  &  212.7$^{+54.6}_{-51.0}$  &  136.7$^{+55.8}_{-67.2}$   \\\\ GRB 050525A  &  PC  &  20.8$^{+20.4}_{-17.0}$  &  97.1$^{+36.5}_{-22.9}$   \\\\ GRB 050607  &  PC  &  10.0$^{+3.9}_{-3.7}$  &  $<2.18$  &  3.9$^{+8.1}_{-3.9}$  &  14.5$^{+6.3}_{-14.0}$   \\\\ GRB 050712  &  WT  &  8.2$^{+3.2}_{-3.0}$   \\\\ &  PC  &  5.3$^{+3.1}_{-2.0}$  &  11.9$^{+8.5}_{-11.9}$   \\\\ GRB 050713A  &  WT  &  43.0$^{+8.9}_{-8.1}$  &  22.6$^{+2.2}_{-2.1}$  &  15.8$^{+16.6}_{-14.2}$  &  38.7$^{+33.8}_{-23.8}$   \\\\ &  PC    &    &    &  &  26.3$^{+2.9}_{-5.0}$   \\\\ GRB 050713B  &  PC  &  18.4$^{+10.4}_{-9.2}$  &  26.6$^{+8.1}_{-7.4}$  &  19.2$^{+4.1}_{-4.5}$   \\\\ GRB 050714B  &  WT  &  77.9$^{+12.6}_{-11.1}$   \\\\ &  PC  &  45.1$^{+26.5}_{-18.7}$  &  31.7$^{+6.1}_{-6.7}$   \\\\ GRB 050716  &  WT  &  8.8$^{+4.0}_{-3.7}$  &  2.0$^{+2.1}_{-2.0}$   \\\\ &  PC    &  &  5.7$^{+7.1}_{-5.7}$  &  9.5$^{+4.1}_{-2.4}$   \\\\ GRB 050717  &  WT  &  19.7$^{+5.4}_{-5.0}$  &  $<78.34$   \\\\ &  PC    &  &  11.2$^{+5.6}_{-5.1}$  &  3.4$^{+37.7}_{-3.4}$   \\\\ GRB 050721  &  WT  &  11.2$^{+5.2}_{-4.7}$   \\\\ &  PC  &  9.8$^{+3.9}_{-3.7}$  &  14.6$^{+12.1}_{-9.5}$   \\\\ GRB 050724  &  WT  &  51.1$^{+4.8}_{-4.7}$  &  43.6$^{+6.8}_{-6.1}$   \\\\ &  PC    &    &  &  26.8$^{+13.0}_{-15.1}$   \\\\ GRB 050726  &  WT  &  4.7$^{+2.6}_{-2.4}$   \\\\ &  PC  &  0.48$^{+2.14}_{-0.48}$  &  6.5$^{+3.9}_{-5.4}$   \\\\ GRB 050730  &  WT  &  52.7$^{+22.6}_{-21.8}$  &  275.2$^{+640.1}_{-275.2}$   \\\\ &  PC  &  $<45.59$  &  43.0$^{+31.5}_{-24.0}$   \\\\ GRB 050801  &  PC  &  $<1.98$  &  $<6.33$   \\\\ GRB 050802  &  WT  &  4.0$^{+3.4}_{-3.1}$   \\\\ &  PC  &  3.0$^{+2.4}_{-2.3}$  &  5.77$^{+0.95}_{-1.42}$   \\\\ GRB 050803  &  WT  &  13.5$^{+4.5}_{-4.2}$  &  40.0$^{+33.5}_{-24.0}$   \\\\ &  PC  &  $<9.71$  &  12.9$^{+2.4}_{-2.3}$  &  30.3$^{+5.2}_{-7.5}$   \\\\ GRB 050814  &  WT  &  73.0$^{+61.8}_{-58.9}$  &  432.7$^{+1457.4}_{-432.7}$   \\\\ &  PC  &  $<17.69$  &  183.0$^{+113.4}_{-119.6}$  &  27.2$^{+256.7}_{-27.2}$   \\\\ GRB 050815  &  PC  &  $<18.21$  &  11.0$^{+14.8}_{-11.0}$   \\\\ \\hline \\end{tabular} \\caption{Column densities (in excess of the Galactic value) for the time-resolved spectra, in units of $10^{20}$ \\cms. $N_{\\rm H,1}$ corresponds to the time during which the decay followed $\\alpha_1$ in Table~6. Spectral indices for these spectra are given in Table~9.} \\label{tab:resnh} \\end{center} \\end{table*}  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $N_{\\rm H,1}$ & $N_{\\rm H,2}$ & $N_{\\rm H,3}$ & $N_{\\rm H,4}$ & $N_{\\rm H,5}$ & $N_{\\rm H,6}$  \\\\ \\hline GRB 050819  &  WT  &  $<4.12$   \\\\ &  PC  &  $<4.11$  &  14.0$^{+10.1}_{-6.5}$   \\\\ GRB 050820A  &  WT  &  2.7$^{+3.7}_{-2.7}$  &  $<1.29$   \\\\ &  PC    &  &  2.2$^{+1.8}_{-1.7}$  &  52.7$^{+9.3}_{-12.1}$   \\\\ GRB 050822  &  WT  &  13.0$^{+1.4}_{-1.3}$  &  $<18.67$  &  $<44.94$   \\\\ &  PC  &  3.6$^{+9.7}_{-3.6}$  &  3.5$^{+3.5}_{-1.6}$  &  18.8$^{+3.6}_{-1.7}$   \\\\ GRB 050824  &  PC  &  3.3$^{+6.8}_{-3.3}$  &  40.5$^{+13.6}_{-11.4}$   \\\\ GRB 050826  &  PC  &  99.9$^{+46.8}_{-39.1}$  &  39.1$^{+39.1}_{-22.9}$  &  51.0$^{+50.0}_{-45.3}$   \\\\ GRB 050904  &  WT  &  298.0$^{+90.1}_{-86.9}$   \\\\ &  PC  &  $<70.58$  &  135.7$^{+52.7}_{-82.6}$   \\\\ GRB050915B  &  WT  &  9.5$^{+4.7}_{-4.4}$  &  455.2$^{+191056.8}_{-398.3}$   \\\\ &  PC  &  7.0$^{+22.0}_{-7.0}$  &  20.5$^{+10.4}_{-15.3}$   \\\\ GRB050916  &  PC  &  $<104.78$  &  26.0$^{+25.8}_{-14.1}$   \\\\ GRB 050922B  &  WT  &  13.4$^{+1.6}_{-1.5}$   \\\\ &  PC  &  9.0$^{+1.5}_{-1.4}$  &  7.2$^{+4.6}_{-6.8}$  &  18.1$^{+11.0}_{-7.6}$   \\\\ GRB 050922C  &  WT  &  17.7$^{+14.1}_{-13.7}$   \\\\ &  PC  &  2.7$^{+24.2}_{-2.7}$  &  49.5$^{+18.3}_{-17.8}$   \\\\ GRB 051001  &  WT  &  12.5$^{+3.6}_{-3.4}$  &  14.2$^{+2.3}_{-2.2}$   \\\\ &  PC    &  &  16.4$^{+11.7}_{-9.8}$  &  20.5$^{+8.5}_{-8.3}$   \\\\ GRB 051008  &  PC  &  32.4$^{+4.3}_{-3.9}$  &  36.2$^{+5.1}_{-8.0}$   \\\\ GRB 051016A  &  WT  &  27.4$^{+21.7}_{-17.9}$   \\\\ &  PC  &  50.5$^{+37.9}_{-29.7}$  &  37.4$^{+10.6}_{-12.2}$   \\\\ GRB 051016B  &  WT  &  85.0$^{+32.9}_{-28.3}$   \\\\ &  PC  &  $<27.07$  &  55.7$^{+23.5}_{-20.6}$  &  87.6$^{+17.5}_{-27.4}$   \\\\ GRB 051021B  &  WT  &  $<31.12$   \\\\ &  PC  &  0.75$^{+33.63}_{-0.75}$  &  24.3$^{+65.7}_{-24.3}$   \\\\ GRB 051022  &  PC  &  315.2$^{+35.0}_{-33.5}$  &  275.0$^{+166.2}_{-70.1}$   \\\\ GRB 051109A  &  WT  &  15.3$^{+49.5}_{-15.3}$  &  118.3$^{+543.8}_{-118.3}$   \\\\ &  PC    &  &  158.0$^{+70.0}_{-62.1}$  &  74.0$^{+27.2}_{-18.8}$   \\\\ GRB 051109B  &  PC  &  $<4.50$  &  27.4$^{+15.1}_{-12.9}$  &  16.0$^{+10.6}_{-9.5}$   \\\\ GRB 051117A  &  WT  &  10.07$^{+0.40}_{-0.39}$  &  9.3$^{+39.9}_{-9.3}$   \\\\ &  PC  &  8.4$^{+2.7}_{-2.6}$    &  &  10.4$^{+1.6}_{-4.0}$   \\\\ GRB 051221A  &  PC  &  16.6$^{+9.2}_{-4.4}$  &  7.4$^{+9.8}_{-6.4}$  &  35.2$^{+10.2}_{-9.5}$   \\\\ GRB 051227  &  WT  &  17.7$^{+8.2}_{-7.3}$   \\\\ &  PC  &  6.3$^{+6.9}_{-6.1}$  &  9.1$^{+7.5}_{-5.4}$   \\\\ GRB 060105  &  WT  &  19.4$^{+2.0}_{-1.9}$  &  18.1$^{+1.2}_{-1.1}$   \\\\ &  PC    &    &  &  18.8$^{+2.6}_{-2.5}$  &  16.5$^{+3.9}_{-6.6}$  &  16.5$^{+2.9}_{-6.5}$   \\\\ GRB 060108  &  WT    &  &  $<69.31$   \\\\ &  PC  &  $<26.61$  &  54.8$^{+58.4}_{-54.8}$  &  94.3$^{+39.1}_{-49.1}$   \\\\ GRB 060109  &  WT  &  17.5$^{+4.9}_{-4.5}$   \\\\ &  PC  &  2.9$^{+11.4}_{-2.9}$  &  22.8$^{+7.6}_{-6.9}$  &  28.2$^{+6.8}_{-8.4}$   \\\\ GRB 060111A  &  WT  &  18.64$^{+0.97}_{-0.95}$   \\\\ &  PC  &  15.0$^{+8.1}_{-7.4}$  &  15.9$^{+5.5}_{-5.7}$   \\\\ GRB 060111B  &  WT  &  11.6$^{+5.6}_{-5.1}$  &  20.1$^{+16.3}_{-14.7}$   \\\\ &  PC    &  &  18.4$^{+6.4}_{-5.6}$  &  25.4$^{+6.3}_{-16.4}$   \\\\ GRB 060115  &  WT  &  135.9$^{+59.2}_{-55.7}$   \\\\ &  PC  &  $<28.41$  &  0.99$^{+54.33}_{-0.99}$   \\\\ GRB 060124  &  PC  &  193.3$^{+573.1}_{-192.9}$  &  66.2$^{+18.2}_{-17.2}$  &  90.0$^{+16.0}_{-15.5}$   \\\\ GRB 060202  &  WT  &  43.1$^{+1.9}_{-1.9}$  &  38.4$^{+1.8}_{-1.7}$  &  37.1$^{+3.6}_{-3.4}$   \\\\ &  PC    &    &  &  37.6$^{+7.3}_{-6.7}$  &  49.5$^{+2.0}_{-4.5}$   \\\\ GRB 060204B  &  WT  &  19.5$^{+2.2}_{-2.2}$  &  24.4$^{+4.0}_{-3.8}$   \\\\ &  PC  &  12.9$^{+8.2}_{-7.3}$  &  10.0$^{+3.4}_{-3.2}$  &  20.7$^{+5.5}_{-7.1}$   \\\\ GRB 060206  &  WT  &  $<265.06$   \\\\ &  PC  &  20.5$^{+111.6}_{-20.5}$  &  49.2$^{+56.1}_{-49.2}$   \\\\ GRB 060210  &  WT  &  227.3$^{+24.6}_{-24.0}$  &  66.4$^{+1909.0}_{-66.4}$   \\\\ &  PC  &  195.0$^{+43.6}_{-41.0}$  &  164.9$^{+44.7}_{-44.3}$   \\\\ GRB 060211A  &  WT  &  7.4$^{+1.9}_{-1.8}$   \\\\ &  PC  &  1.9$^{+12.0}_{-1.9}$  &  $<5.82$  &  20.4$^{+9.1}_{-8.1}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $N_{\\rm H,1}$ & $N_{\\rm H,2}$ & $N_{\\rm H,3}$ & $N_{\\rm H,4}$ & $N_{\\rm H,5}$ & $N_{\\rm H,6}$  \\\\ \\hline GRB 060211B  &  WT  &  2.0$^{+17.6}_{-2.0}$   \\\\ &  PC  &  2.6$^{+10.8}_{-2.6}$  &  8.3$^{+29.2}_{-8.3}$   \\\\ GRB 060218  &  WT  &  26.34$^{+0.44}_{-0.44}$   \\\\ &  PC  &  22.1$^{+5.4}_{-5.0}$  &  62.0$^{+15.0}_{-7.4}$   \\\\ GRB 060219  &  PC  &  14.6$^{+19.3}_{-14.3}$  &  39.8$^{+12.2}_{-11.5}$  &  34.6$^{+15.7}_{-9.5}$   \\\\ GRB 060223A  &  WT    &  &  605.4$^{+1022.8}_{-506.2}$   \\\\ &  PC  &  56.1$^{+313.1}_{-56.1}$  &  319.2$^{+284.3}_{-244.8}$   \\\\ GRB 060306  &  WT  &  38.3$^{+11.0}_{-9.8}$  &  9.3$^{+110.8}_{-9.3}$   \\\\ &  PC  &  9.2$^{+23.4}_{-9.2}$  &  35.3$^{+6.6}_{-6.1}$  &  37.7$^{+4.2}_{-6.5}$   \\\\ GRB 060312  &  WT  &  10.2$^{+2.5}_{-2.4}$   \\\\ &  PC  &  $<2.55$  &  1.9$^{+2.6}_{-1.9}$   \\\\ GRB 060313  &  WT  &  $<7.84$  &  $<20.30$   \\\\ &  PC  &  0.52$^{+9.34}_{-0.52}$  &  $<5.61$  &  $<1.27$  &  5.4$^{+2.4}_{-5.3}$   \\\\ GRB 060319  &  PC  &  41.4$^{+4.8}_{-4.3}$  &  51.0$^{+7.3}_{-10.4}$   \\\\ GRB 060403  &  WT  &  19.2$^{+15.9}_{-14.1}$   \\\\ &  PC  &  $<26.92$  &  28.0$^{+63.2}_{-28.0}$   \\\\ GRB 060413  &  WT  &  126.9$^{+16.0}_{-15.6}$   \\\\ &  PC  &  43.0$^{+24.1}_{-22.0}$  &  87.3$^{+21.1}_{-20.2}$  &  162.2$^{+44.7}_{-37.5}$   \\\\ GRB 060418  &  WT  &  61.6$^{+6.7}_{-6.3}$  &  42.8$^{+15.9}_{-14.8}$  &  2.8$^{+59.3}_{-2.8}$   \\\\ &  PC    &  &  40.7$^{+19.2}_{-17.7}$  &  149.4$^{+92.8}_{-95.3}$   \\\\ GRB 060427  &  WT  &  10.4$^{+3.0}_{-2.8}$   \\\\ &  PC  &  3.8$^{+6.1}_{-3.8}$  &  12.3$^{+3.1}_{-1.5}$   \\\\ GRB 060428A  &  WT  &  $<15.50$   \\\\ &  PC  &  35.5$^{+142.8}_{-35.5}$  &  16.8$^{+11.7}_{-11.0}$  &  0.013$^{+10.099}_{-0.013}$  &  70.3$^{+32.1}_{-70.3}$   \\\\ GRB 060428B  &  WT  &  6.5$^{+1.5}_{-1.4}$   \\\\ &  PC  &  $<6.52$  &  1.5$^{+1.6}_{-1.5}$   \\\\ GRB 060502A  &  WT  &  53.8$^{+16.8}_{-15.3}$   \\\\ &  PC  &  $<6.72$  &  41.9$^{+17.9}_{-16.4}$  &  39.0$^{+14.2}_{-20.6}$   \\\\ GRB 060507  &  WT    &  &  101.2$^{+447.8}_{-101.2}$   \\\\ &  PC  &  $<6.06$  &  7.5$^{+3.3}_{-4.7}$   \\\\ GRB 060510A  &  WT  &  $<4.31$   \\\\ &  PC    &  &  $<0.76$  &  1.6$^{+3.6}_{-1.6}$   \\\\ GRB 060510B  &  WT  &  393.0$^{+48.1}_{-47.3}$  &  272.6$^{+86.2}_{-82.1}$   \\\\ &  PC    &  &  122.3$^{+219.2}_{-122.3}$  &  240.6$^{+166.7}_{-84.6}$   \\\\ GRB 060512  &  WT  &  30.2$^{+17.6}_{-15.1}$   \\\\ &  PC  &  2.3$^{+9.1}_{-2.3}$  &  $<3.55$   \\\\ GRB 060522  &  WT  &  120.4$^{+288.3}_{-120.4}$   \\\\ &  PC  &  $<187.92$  &  96.9$^{+202.5}_{-96.9}$  &  944.0$^{+2022.1}_{-434.9}$   \\\\ GRB 060526  &  WT  &  55.0$^{+15.4}_{-14.8}$   \\\\ &  PC  &  $<38.97$  &  $<42.72$  &  92.2$^{+80.5}_{-92.2}$  &  0.055$^{+137.888}_{-0.055}$   \\\\ GRB 060604  &  WT  &  159.1$^{+18.0}_{-17.0}$   \\\\ &  PC    &  &  20.0$^{+42.0}_{-20.0}$  &  149.4$^{+59.9}_{-96.6}$   \\\\ GRB 060605  &  WT  &  $<183.29$   \\\\ &  PC  &  $<44.44$  &  44.8$^{+72.0}_{-44.8}$  &  112.7$^{+58.9}_{-58.6}$   \\\\ GRB 060607A  &  WT  &  28.5$^{+12.3}_{-11.8}$  &  $<82.41$  &  101.5$^{+439.6}_{-101.5}$   \\\\ &  PC    &  &  4.9$^{+16.0}_{-4.9}$  &  59.0$^{+26.7}_{-17.2}$   \\\\ GRB 060614  &  WT  &  4.0$^{+1.4}_{-1.4}$  &  5.80$^{+0.71}_{-0.69}$  &  $<1.26$   \\\\ &  PC    &    &  &  $<1.56$  &  6.5$^{+3.1}_{-3.1}$   \\\\ GRB 060707  &  WT  &  $<56.85$   \\\\ &  PC  &  $<46.13$  &  52.0$^{+84.3}_{-44.2}$  &  209.3$^{+370.5}_{-102.9}$   \\\\ GRB 060708  &  WT  &  11.8$^{+2.8}_{-2.6}$   \\\\ &  PC    &  &  2.7$^{+2.7}_{-2.3}$  &  5.2$^{+3.2}_{-4.9}$   \\\\ GRB 060712  &  PC  &  10.2$^{+8.4}_{-7.0}$  &  15.7$^{+14.8}_{-10.7}$  &  15.6$^{+7.0}_{-6.7}$   \\\\ GRB 060714  &  WT  &  136.9$^{+19.8}_{-18.8}$   \\\\ &  PC  &  33.0$^{+111.0}_{-33.0}$  &  91.7$^{+51.6}_{-47.4}$  &  301.0$^{+107.1}_{-117.5}$   \\\\ GRB 060719  &  WT  &  26.2$^{+6.7}_{-5.9}$   \\\\ &  PC  &  21.6$^{+19.3}_{-15.6}$  &  45.0$^{+8.0}_{-7.4}$  &  45.5$^{+6.5}_{-11.6}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $N_{\\rm H,1}$ & $N_{\\rm H,2}$ & $N_{\\rm H,3}$ & $N_{\\rm H,4}$ & $N_{\\rm H,5}$ & $N_{\\rm H,6}$  \\\\ \\hline GRB 060729  &  WT  &  40.1$^{+4.0}_{-3.7}$  &  18.8$^{+3.7}_{-3.4}$   \\\\ &  PC    &  &  $<3.76$  &  12.8$^{+2.0}_{-1.9}$  &  13.0$^{+1.3}_{-1.6}$  &  16.7$^{+5.8}_{-8.2}$   \\\\ GRB 060801  &  PC  &  48.0$^{+42.0}_{-33.3}$  &  7.8$^{+22.0}_{-7.8}$  &  $(1.28^{+2.97}_{-0.64}) \\tim{3}$   \\\\ GRB 060804  &  PC  &  $<7.21$  &  8.2$^{+9.5}_{-8.2}$   \\\\ GRB 060805A  &  PC  &  5.1$^{+7.5}_{-2.9}$  &  4.1$^{+9.8}_{-4.1}$   \\\\ GRB 060807  &  PC  &  1.6$^{+10.9}_{-1.6}$  &  10.9$^{+5.4}_{-4.9}$  &  13.3$^{+2.9}_{-2.6}$  &  17.0$^{+8.0}_{-11.9}$   \\\\ GRB 060813  &  WT  &  8.6$^{+3.4}_{-3.3}$   \\\\ &  PC    &  &  5.8$^{+3.6}_{-3.4}$  &  13.8$^{+14.8}_{-13.8}$   \\\\ GRB 060814  &  WT  &  39.6$^{+8.2}_{-7.5}$  &  25.7$^{+1.4}_{-1.4}$  &  19.2$^{+3.4}_{-3.3}$   \\\\ &  PC    &    &  &  22.4$^{+6.9}_{-6.3}$  &  30.8$^{+3.1}_{-3.0}$  &  29.3$^{+3.0}_{-3.7}$   \\\\ GRB 060904A  &  WT  &  17.8$^{+1.6}_{-1.5}$  &  18.1$^{+2.5}_{-2.3}$   \\\\ &  PC    &    &  &  19.6$^{+12.0}_{-9.5}$   \\\\ GRB 060904B  &  WT  &  46.9$^{+3.9}_{-3.7}$   \\\\ &  PC  &  31.8$^{+22.0}_{-17.4}$  &  49.6$^{+13.9}_{-13.0}$   \\\\ GRB 060906  &  WT  &  166.2$^{+408.2}_{-166.2}$  &  325.1$^{+728.8}_{-270.9}$   \\\\ &  PC    &  &  312.6$^{+189.8}_{-171.0}$  &  263.8$^{+212.5}_{-186.1}$   \\\\ GRB 060908  &  WT  &  43.9$^{+40.1}_{-36.9}$  &  169.9$^{+226.9}_{-165.3}$   \\\\ &  PC  &  92.7$^{+72.1}_{-62.8}$  &  60.7$^{+37.2}_{-48.1}$   \\\\ GRB 060912A  &  PC  &  20.0$^{+11.8}_{-10.9}$  &  40.5$^{+8.1}_{-14.2}$   \\\\ GRB 060923A  &  WT  &  10.8$^{+8.9}_{-7.5}$  &  3.7$^{+27.7}_{-3.7}$  &  10.8$^{+8.9}_{-7.4}$   \\\\ &  PC  &  14.2$^{+14.8}_{-11.2}$  &  $<9.31$  &  32.4$^{+8.8}_{-14.0}$   \\\\ GRB 060923C  &  WT  &  0.78$^{+4.40}_{-0.78}$   \\\\ &  PC  &  $<1.68$  &  27.8$^{+11.9}_{-10.1}$  &  5.7$^{+13.8}_{-5.7}$   \\\\ GRB 060926  &  WT  &  0.11$^{+294.31}_{-0.11}$  &  246.7$^{+205.7}_{-155.3}$   \\\\ &  PC    &    &  &  453.6$^{+274.2}_{-254.6}$   \\\\ GRB 060927  &  PC  &  52.5$^{+320.0}_{-52.5}$  &  176.5$^{+518.5}_{-176.5}$   \\\\ GRB 060929  &  WT    &  &  12.1$^{+1.4}_{-1.4}$   \\\\ &  PC  &  16.2$^{+16.2}_{-12.4}$  &  11.2$^{+4.4}_{-3.6}$   \\\\ GRB 061002  &  WT  &  16.2$^{+9.4}_{-8.5}$  &  8.7$^{+17.1}_{-8.7}$   \\\\ &  PC    &  &  20.2$^{+16.0}_{-14.8}$   \\\\ GRB 061004  &  WT  &  $<11.41$   \\\\ &  PC  &  $<6.83$  &  6.3$^{+6.7}_{-5.9}$  &  5.7$^{+7.5}_{-5.7}$   \\\\ GRB 061007  &  WT  &  44.8$^{+7.2}_{-6.9}$  &  45.0$^{+2.8}_{-2.7}$   \\\\ &  PC    &  &  71.4$^{+45.1}_{-36.8}$  &  57.2$^{+19.3}_{-17.3}$  &  64.0$^{+15.1}_{-8.6}$   \\\\ GRB 061019  &  PC  &  67.6$^{+11.0}_{-11.5}$  &  76.5$^{+55.5}_{-34.2}$   \\\\ GRB 061021  &  WT  &  $<0.20$   \\\\ &  PC  &  $<4.93$  &  5.4$^{+2.3}_{-2.1}$  &  6.0$^{+1.0}_{-1.3}$   \\\\ GRB 061028  &  WT  &  10.0$^{+6.3}_{-5.7}$   \\\\ &  PC  &  $<3.28$  &  114.5$^{+22.2}_{-17.7}$   \\\\ GRB 061102  &  WT  &  8.6$^{+5.4}_{-4.7}$   \\\\ &  PC    &  &  $<10.28$   \\\\ GRB 061110A  &  WT  &  16.2$^{+1.6}_{-1.5}$  &  5.9$^{+1.8}_{-1.7}$   \\\\ &  PC    &  &  $<1.20$  &  $<0.49$  &  9.0$^{+8.3}_{-9.0}$   \\\\ GRB 061121  &  WT  &  83.2$^{+26.0}_{-22.8}$  &  70.4$^{+12.3}_{-11.5}$  &  40.5$^{+14.6}_{-13.5}$  &  60.5$^{+11.3}_{-10.6}$  &  102.1$^{+56.9}_{-47.2}$   \\\\ &  PC    &    &    &  &  37.8$^{+22.6}_{-20.1}$  &  55.2$^{+10.9}_{-10.3}$  &  66.1$^{+13.5}_{-18.7}$   \\\\ GRB 061126  &  WT  &  9.5$^{+4.6}_{-4.2}$   \\\\ &  PC  &  14.5$^{+2.6}_{-2.4}$  &  12.0$^{+3.3}_{-4.7}$   \\\\ GRB 061201  &  WT  &  $<19.16$  &  $<71.80$   \\\\ &  PC  &  1.7$^{+3.2}_{-1.7}$  &  13.2$^{+16.0}_{-9.9}$   \\\\ GRB 061202  &  WT  &  45.3$^{+4.7}_{-4.5}$   \\\\ &  PC  &  52.7$^{+11.4}_{-10.3}$  &  47.7$^{+8.0}_{-7.4}$  &  48.9$^{+5.6}_{-7.4}$   \\\\ GRB 061222A  &  WT  &  32.6$^{+3.0}_{-2.9}$  &  27.9$^{+14.3}_{-11.9}$  &  18.7$^{+30.2}_{-18.7}$   \\\\ &  PC    &  &  27.3$^{+2.9}_{-2.8}$  &  25.9$^{+3.1}_{-2.9}$  &  36.0$^{+3.3}_{-5.0}$   \\\\ GRB 070103  &  WT    &  &  40.6$^{+55.5}_{-27.9}$   \\\\ &  PC  &  30.1$^{+13.9}_{-12.1}$  &  23.0$^{+16.4}_{-13.2}$  &  26.3$^{+8.2}_{-8.4}$   \\\\ GRB 070107  &  WT  &  3.5$^{+2.3}_{-2.2}$   \\\\ &  PC  &  5.7$^{+3.8}_{-3.6}$  &  9.8$^{+2.1}_{-5.5}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $N_{\\rm H,1}$ & $N_{\\rm H,2}$ & $N_{\\rm H,3}$ & $N_{\\rm H,4}$ & $N_{\\rm H,5}$ & $N_{\\rm H,6}$  \\\\ \\hline GRB 070110  &  WT  &  10.6$^{+17.4}_{-10.6}$  &  $<9.18$   \\\\ &  PC  &  $<12.20$  &  $<11.15$  &  13.8$^{+14.1}_{-13.3}$  &  7.8$^{+32.2}_{-7.8}$  &  17.3$^{+24.0}_{-14.3}$   \\\\ GRB 070125  &  PC  &  5.8$^{+2.8}_{-4.9}$  &  3.5$^{+3.6}_{-3.5}$   \\\\ GRB 070129  &  WT  &  11.66$^{+0.77}_{-0.76}$   \\\\ &  PC  &  $<2.31$  &  16.3$^{+6.9}_{-6.4}$  &  14.2$^{+1.9}_{-4.0}$  &  11.5$^{+12.5}_{-11.5}$   \\\\ GRB 070208  &  PC  &  48.9$^{+40.2}_{-24.5}$  &  66.1$^{+36.3}_{-30.9}$  &  89.6$^{+38.0}_{-43.9}$   \\\\ GRB 070219  &  WT  &  23.6$^{+110.6}_{-23.6}$   \\\\ &  PC  &  8.2$^{+21.2}_{-8.2}$  &  19.3$^{+40.1}_{-19.3}$  &  48.8$^{+22.2}_{-18.7}$   \\\\ GRB 070220  &  WT  &  11.7$^{+6.1}_{-5.7}$  &  8.5$^{+7.3}_{-6.7}$   \\\\ &  PC    &    &  &  15.8$^{+5.2}_{-4.9}$  &  13.1$^{+7.7}_{-8.4}$   \\\\ GRB 070223  &  WT  &  55.5$^{+9.8}_{-8.8}$  &  49.7$^{+7.7}_{-7.1}$   \\\\ &  PC    &    &  &  57.4$^{+14.3}_{-19.2}$   \\\\ GRB 070224  &  WT  &  $<2.16$   \\\\ &  PC  &  $<1.60$  &  3.2$^{+6.6}_{-3.2}$   \\\\ GRB 070306  &  WT  &  40.8$^{+3.7}_{-3.6}$  &  37.4$^{+3.1}_{-3.0}$   \\\\ &  PC    &  &  34.1$^{+20.8}_{-18.1}$  &  31.6$^{+4.0}_{-3.8}$  &  38.1$^{+7.6}_{-4.1}$  &  32.9$^{+5.1}_{-3.8}$   \\\\ GRB 070311  &  PC  &  19.9$^{+10.0}_{-11.0}$  &  11.5$^{+11.8}_{-11.5}$  &  18.5$^{+26.9}_{-18.5}$   \\\\ GRB 070328  &  WT  &  16.8$^{+2.2}_{-2.1}$  &  19.0$^{+1.4}_{-1.4}$  &  21.2$^{+1.7}_{-1.7}$   \\\\ &  PC    &    &  &  32.4$^{+21.2}_{-17.5}$  &  21.7$^{+1.2}_{-1.8}$   \\\\ GRB 070412  &  WT  &  12.9$^{+9.2}_{-7.4}$   \\\\ &  PC    &  &  10.5$^{+8.3}_{-7.6}$  &  21.6$^{+9.7}_{-10.9}$   \\\\ GRB 070419A  &  WT    &  &  22.2$^{+8.2}_{-7.8}$   \\\\ &  PC    &  &  26.7$^{+14.1}_{-12.9}$  &  23.5$^{+24.1}_{-23.5}$   \\\\ GRB 070419B  &  WT  &  12.30$^{+0.88}_{-0.86}$   \\\\ &  PC  &  5.7$^{+1.3}_{-1.3}$  &  8.4$^{+5.5}_{-5.6}$   \\\\ GRB 070420  &  WT  &  3.5$^{+3.4}_{-3.3}$  &  $<40.65$   \\\\ &  PC    &  &  5.5$^{+3.3}_{-3.2}$  &  9.2$^{+5.4}_{-5.1}$  &  10.1$^{+9.1}_{-9.2}$  &  18.1$^{+10.8}_{-16.7}$   \\\\ GRB 070429A  &  WT  &  41.9$^{+5.3}_{-4.9}$  &  17.4$^{+10.9}_{-7.9}$   \\\\ &  PC    &  &  2.2$^{+10.0}_{-2.2}$  &  19.6$^{+7.7}_{-8.5}$  &  3.2$^{+22.4}_{-3.2}$   \\\\ GRB 070508  &  WT  &  24.7$^{+1.8}_{-1.8}$  &  27.2$^{+2.5}_{-2.5}$   \\\\ &  PC    &  &  26.1$^{+6.0}_{-5.6}$  &  29.2$^{+6.6}_{-5.7}$   \\\\ GRB 070517  &  PC  &  4.3$^{+5.2}_{-2.4}$  &  2.2$^{+11.2}_{-2.2}$   \\\\ GRB 070518  &  WT  &  7.3$^{+1.9}_{-1.8}$   \\\\ &  PC  &  $<3.49$  &  24.5$^{+8.0}_{-16.3}$   \\\\ GRB 070520A  &  WT  &  20.3$^{+3.6}_{-3.4}$   \\\\ &  PC    &  &  33.7$^{+9.0}_{-8.5}$   \\\\ GRB 070520B  &  WT  &  19.4$^{+1.5}_{-1.4}$   \\\\ &  PC  &  12.0$^{+17.3}_{-12.0}$  &  37.8$^{+19.2}_{-20.2}$   \\\\ GRB 070521  &  WT  &  56.7$^{+79.4}_{-39.7}$   \\\\ &  PC  &  77.5$^{+32.4}_{-26.2}$  &  56.2$^{+7.7}_{-7.1}$  &  63.8$^{+9.5}_{-5.2}$   \\\\ GRB 070529  &  WT  &  481.5$^{+484.3}_{-314.3}$   \\\\ &  PC  &  $<409.06$  &  99.0$^{+69.8}_{-62.0}$  &  601.1$^{+349.8}_{-244.7}$   \\\\ GRB 070616  &  WT  &  6.8$^{+1.6}_{-1.6}$  &  3.2$^{+1.8}_{-1.8}$   \\\\ &  PC    &  &  18.4$^{+12.7}_{-11.1}$  &  9.8$^{+4.8}_{-7.6}$   \\\\ GRB 070621  &  WT  &  35.0$^{+3.0}_{-2.9}$   \\\\ &  PC  &  35.4$^{+12.7}_{-11.1}$  &  38.9$^{+8.3}_{-7.6}$  &  40.4$^{+8.8}_{-14.9}$   \\\\ GRB 070628  &  WT  &  40.5$^{+73.8}_{-40.5}$   \\\\ &  PC  &  6.4$^{+9.3}_{-6.4}$  &  12.1$^{+3.2}_{-5.7}$   \\\\ GRB 070704  &  WT  &  33.8$^{+5.9}_{-5.7}$   \\\\ &  PC  &  $<75.75$  &  80.7$^{+35.6}_{-40.6}$   \\\\ GRB 070714B  &  WT  &  26.2$^{+15.1}_{-12.6}$  &  41.8$^{+29.1}_{-20.1}$   \\\\ &  PC    &  &  46.4$^{+17.3}_{-15.2}$  &  24.3$^{+21.0}_{-18.5}$  &  31.1$^{+16.3}_{-19.5}$   \\\\ GRB 070721A  &  WT  &  4.5$^{+10.7}_{-4.5}$   \\\\ &  PC  &  $<2.65$  &  25.0$^{+8.5}_{-17.7}$   \\\\ GRB 070721B  &  WT  &  2.5$^{+3.3}_{-2.5}$  &  1.0$^{+2.5}_{-1.0}$   \\\\ &  PC    &  &  $<2.09$  &  20.5$^{+8.7}_{-8.6}$   \\\\ GRB 070724A  &  WT  &  19.0$^{+6.2}_{-5.7}$   \\\\ &  PC  &  25.5$^{+10.1}_{-8.9}$  &  56.3$^{+20.7}_{-25.8}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $N_{\\rm H,1}$ & $N_{\\rm H,2}$ & $N_{\\rm H,3}$ & $N_{\\rm H,4}$ & $N_{\\rm H,5}$ & $N_{\\rm H,6}$  \\\\ \\hline GRB 070802  &  PC  &  $<33.71$  &  161.6$^{+150.9}_{-116.6}$  &  162.5$^{+351.7}_{-81.5}$   \\\\ GRB 070808  &  WT  &  79.7$^{+53.2}_{-39.6}$   \\\\ &  PC  &  57.8$^{+38.9}_{-29.4}$  &  95.0$^{+24.5}_{-22.2}$   \\\\ GRB 070810A  &  PC  &  49.5$^{+72.6}_{-49.5}$  &  74.0$^{+32.9}_{-35.3}$   \\\\ GRB 071001  &  WT  &  $<5.06$   \\\\ &  PC  &  $<3.07$  &  $<30.90$   \\\\ GRB 071010A  &  PC  &  44.6$^{+14.0}_{-12.9}$  &  48.4$^{+13.2}_{-30.1}$   \\\\ GRB 071011  &  PC  &  49.8$^{+26.8}_{-15.5}$  &  21.5$^{+57.0}_{-21.5}$   \\\\ GRB 071020  &  WT  &  7.2$^{+2.6}_{-2.5}$  &  0.65$^{+2.26}_{-0.65}$   \\\\ &  PC    &  &  5.6$^{+1.0}_{-1.3}$   \\\\ GRB 071021  &  WT  &  14.4$^{+2.3}_{-2.2}$  &  12.4$^{+22.3}_{-12.4}$   \\\\ &  PC    &  &  15.4$^{+2.8}_{-3.3}$   \\\\ GRB 071025  &  WT  &  6.0$^{+2.2}_{-2.1}$  &  3.6$^{+1.7}_{-1.6}$  &  7.5$^{+3.4}_{-3.1}$   \\\\ &  PC    &    &  &  0.34$^{+29.63}_{-0.34}$  &  9.2$^{+1.3}_{-1.2}$   \\\\ GRB 071028A  &  WT  &  4.1$^{+4.9}_{-4.1}$   \\\\ &  PC  &  1.4$^{+6.0}_{-1.4}$  &  $<4.29$  &  $<11.74$   \\\\ GRB 071031  &  WT  &  88.8$^{+8.9}_{-8.7}$   \\\\ &  PC  &  0.093$^{+52.872}_{-0.093}$  &  48.5$^{+53.9}_{-41.0}$  &  128.3$^{+237.9}_{-128.3}$   \\\\ GRB 071112C  &  WT  &  10.0$^{+5.4}_{-5.2}$  &  $<4.25$   \\\\ &  PC    &  &  1.6$^{+12.0}_{-1.6}$  &  24.0$^{+16.6}_{-15.8}$   \\\\ GRB 071118  &  WT  &  19.5$^{+5.5}_{-5.2}$   \\\\ &  PC  &  19.8$^{+12.9}_{-11.0}$  &  13.9$^{+2.7}_{-2.6}$  &  20.9$^{+6.9}_{-10.1}$   \\\\ GRB 071227  &  WT  &  12.9$^{+4.2}_{-3.9}$   \\\\ &  PC    &  &  4.6$^{+6.7}_{-4.6}$  &  36.4$^{+607.5}_{-15.6}$   \\\\ GRB 080123  &  WT  &  7.9$^{+2.6}_{-2.5}$   \\\\ &  PC    &  &  7.1$^{+7.1}_{-6.1}$  &  30.9$^{+20.7}_{-24.8}$   \\\\ GRB 080129  &  PC  &  1.8$^{+15.1}_{-1.8}$  &  24.5$^{+14.8}_{-12.7}$   \\\\ GRB 080205  &  WT  &  14.6$^{+4.2}_{-3.9}$  &  6.6$^{+7.8}_{-6.6}$  &  $<11.28$   \\\\ &  PC    &    &  &  5.7$^{+7.9}_{-5.7}$  &  12.4$^{+3.6}_{-4.5}$   \\\\ GRB 080207  &  WT  &  88.2$^{+9.4}_{-8.8}$   \\\\ &  PC    &  &  69.7$^{+8.6}_{-7.6}$   \\\\ GRB 080212  &  WT  &  31.9$^{+2.1}_{-2.0}$  &  16.6$^{+1.3}_{-1.3}$   \\\\ &  PC    &  &  8.3$^{+4.4}_{-4.1}$  &  18.3$^{+3.3}_{-4.9}$   \\\\ GRB 080229A  &  WT  &  22.1$^{+2.1}_{-2.0}$  &  15.8$^{+19.3}_{-15.5}$   \\\\ &  PC    &  &  16.4$^{+4.4}_{-4.2}$  &  16.3$^{+1.7}_{-1.6}$  &  12.8$^{+8.7}_{-8.2}$   \\\\ GRB 080303  &  WT  &  1.1$^{+5.2}_{-1.1}$   \\\\ &  PC  &  $<1.92$  &  11.6$^{+8.1}_{-11.6}$   \\\\ GRB 080307  &  WT  &  2.2$^{+2.0}_{-1.8}$  &  2.1$^{+1.8}_{-1.7}$   \\\\ &  PC    &  &  6.4$^{+2.7}_{-2.5}$  &  7.4$^{+6.6}_{-7.4}$   \\\\ GRB 080310  &  WT  &  47.3$^{+6.4}_{-6.2}$   \\\\ &  PC    &  &  $<3.27$  &  46.7$^{+25.6}_{-38.3}$   \\\\ GRB 080319B  &  WT  &  5.8$^{+2.2}_{-2.1}$  &  12.73$^{+0.85}_{-0.85}$   \\\\ &  PC    &  &  9.5$^{+5.2}_{-4.9}$  &  9.6$^{+6.7}_{-6.3}$  &  13.6$^{+5.2}_{-7.8}$  &  $<7.88$   \\\\ GRB 080319C  &  PC  &  50.9$^{+59.9}_{-50.2}$  &  73.5$^{+11.8}_{-11.2}$  &  74.0$^{+41.5}_{-53.0}$   \\\\ GRB 080319D  &  WT  &  $(7.4^{+23.4}_{-4.0}) \\tim{3}$   \\\\ &  PC  &  6.5$^{+4.2}_{-4.0}$  &  41.4$^{+35.1}_{-39.7}$   \\\\ GRB 080320  &  PC  &  3.2$^{+2.7}_{-2.5}$  &  4.5$^{+1.7}_{-1.4}$  &  13.8$^{+2.0}_{-5.7}$   \\\\ GRB 080325  &  WT  &  26.0$^{+1.3}_{-1.2}$   \\\\ &  PC    &  &  24.3$^{+8.0}_{-8.7}$   \\\\ GRB 080328  &  WT  &  11.7$^{+4.7}_{-4.4}$  &  12.6$^{+7.6}_{-7.0}$   \\\\ &  PC    &  &  7.3$^{+3.1}_{-3.0}$  &  11.0$^{+4.4}_{-4.1}$   \\\\ GRB 080330  &  WT  &  3.7$^{+3.3}_{-3.0}$   \\\\ &  PC    &  &  1.5$^{+3.4}_{-1.5}$  &  9.0$^{+12.9}_{-9.0}$   \\\\ GRB 080409  &  PC  &  $<98.04$  &  14.1$^{+13.8}_{-12.3}$   \\\\ GRB 080411  &  PC  &  42.2$^{+12.9}_{-12.0}$  &  46.8$^{+4.5}_{-2.5}$  &  31.6$^{+11.9}_{-24.9}$   \\\\ GRB 080413A  &  WT  &  163.5$^{+38.5}_{-36.0}$   \\\\ &  PC  &  $<62.49$  &  121.5$^{+95.5}_{-105.5}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\clearpage    \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\hline GRB  & Mode & $N_{\\rm H,1}$ & $N_{\\rm H,2}$ & $N_{\\rm H,3}$ & $N_{\\rm H,4}$ & $N_{\\rm H,5}$ & $N_{\\rm H,6}$  \\\\ \\hline GRB 080413B  &  WT  &  27.7$^{+9.2}_{-8.6}$   \\\\ &  PC  &  21.4$^{+27.8}_{-21.4}$  &  23.4$^{+4.6}_{-4.5}$  &  38.2$^{+15.0}_{-19.5}$   \\\\ GRB 080426  &  PC  &  $<15.79$  &  21.4$^{+5.2}_{-12.5}$   \\\\ GRB 080430  &  WT  &  13.4$^{+6.9}_{-6.3}$   \\\\ &  PC  &  7.9$^{+13.6}_{-7.9}$  &  36.1$^{+6.4}_{-3.3}$  &  43.5$^{+7.8}_{-9.5}$   \\\\ GRB 080503  &  WT  &  2.5$^{+2.0}_{-1.9}$  &  8.6$^{+2.8}_{-2.6}$  &  11.5$^{+8.9}_{-7.5}$   \\\\ &  PC    &    &  &  $<1.09$   \\\\ GRB 080506  &  WT  &  5.0$^{+3.7}_{-3.5}$  &  4.6$^{+3.7}_{-3.5}$  &  2.6$^{+2.8}_{-2.6}$   \\\\ &  PC    &    &  &  0.098$^{+5.072}_{-0.098}$  &  $<2.68$   \\\\ GRB 080516  &  PC  &  71.4$^{+37.3}_{-30.9}$  &  51.9$^{+24.3}_{-12.5}$   \\\\ GRB 080517  &  WT  &  17.0$^{+28.0}_{-17.0}$   \\\\ &  PC  &  9.8$^{+10.4}_{-8.9}$  &  40.1$^{+22.6}_{-21.6}$   \\\\ GRB 080523  &  WT  &  9.0$^{+1.8}_{-1.7}$  &  $<2.13$   \\\\ &  PC    &  &  $<1.51$  &  7.1$^{+9.5}_{-7.1}$   \\\\ GRB 080602  &  WT  &  11.0$^{+10.2}_{-8.5}$  &  13.2$^{+10.3}_{-7.9}$   \\\\ &  PC    &  &  9.2$^{+3.4}_{-3.2}$  &  12.5$^{+2.3}_{-2.2}$   \\\\ GRB 080603B  &  WT  &  144.6$^{+70.4}_{-57.8}$  &  79.0$^{+25.3}_{-24.2}$  &  116.8$^{+56.7}_{-49.8}$   \\\\ &  PC    &    &  &  85.2$^{+66.7}_{-57.4}$   \\\\ GRB 080604  &  WT  &  $<2.62$   \\\\ &  PC  &  6.6$^{+14.6}_{-6.6}$  &  0.21$^{+36.53}_{-0.21}$   \\\\ GRB 080605  &  WT  &  51.2$^{+10.3}_{-9.8}$  &  82.4$^{+11.7}_{-11.2}$   \\\\ &  PC    &  &  49.7$^{+40.2}_{-34.7}$  &  474.2$^{+283.1}_{-156.4}$   \\\\ GRB 080607  &  WT  &  558.9$^{+115.1}_{-103.0}$  &  369.1$^{+26.2}_{-24.7}$  &  345.2$^{+53.7}_{-50.5}$   \\\\ &  PC    &    &  &  258.9$^{+38.6}_{-22.5}$  &  99.0$^{+206.4}_{-99.0}$   \\\\ GRB 080613B  &  WT  &  2.8$^{+2.4}_{-2.2}$  &  $<18.28$   \\\\ &  PC    &    &  &  9.8$^{+16.0}_{-9.8}$   \\\\ GRB 080625  &  PC  &  10.2$^{+5.0}_{-5.9}$  &  44.4$^{+53.9}_{-15.4}$   \\\\ GRB 080703  &  PC  &  0.93$^{+2.39}_{-0.93}$  &  3.4$^{+9.5}_{-3.4}$   \\\\ GRB 080707  &  WT  &  13.8$^{+51.7}_{-13.8}$  &  101.1$^{+1444.1}_{-101.1}$   \\\\ &  PC  &  $<32.12$  &  54.0$^{+46.5}_{-36.7}$  &  39.3$^{+28.0}_{-24.4}$   \\\\ GRB 080710  &  PC  &  7.4$^{+6.9}_{-3.5}$  &  28.5$^{+14.4}_{-17.3}$   \\\\ GRB 080714  &  WT  &  $<11.72$  &  $<6.55$   \\\\ &  PC    &  &  30.1$^{+49.2}_{-30.1}$  &  20.0$^{+16.5}_{-16.3}$   \\\\ GRB 080721  &  WT  &  51.3$^{+9.5}_{-9.3}$  &  28.4$^{+35.5}_{-28.4}$  &  75.2$^{+7.4}_{-7.3}$  &  141.5$^{+196.2}_{-141.5}$   \\\\ &  PC    &    &    &  &  99.7$^{+14.1}_{-16.5}$   \\\\ GRB 080723A  &  PC  &  48.5$^{+42.5}_{-32.2}$  &  55.7$^{+24.3}_{-20.9}$  &  68.2$^{+12.3}_{-13.6}$   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{ll} \\hline Type of burst  & Bursts \\\\ \\hline No breaks  &  050408, 050520, 051021A, 051028, 051211B, 060121, 060123, 060805B, 060901, 060928\\\\ &  061025, 061122, 070309, 070724B, 070925, 071104, 080120, 080229B, 080405, 080507\\\\ &  080514B, 080603A, 080613A, 080723B, 041223, 050124, 050215B, 050219B, 050223, 050326\\\\ &  050412, 050509A, 050603, 050701, 050827,  {\\bf 050908},  {\\bf 050915A}, 051006, 051111\\\\ &  051117B, 051210, 051221B, 060110,  {\\bf 060116}, 060203, 060323, 060421, 060505, 060515\\\\ &  060602A, 060602B, 060717, 060825, 060919, 060923B, 061006, 061110B, 061210,  {\\bf 061222B}\\\\ &  070227,  {\\bf 070318},  {\\bf 070330}, 070411, 070506, 070509, 070531, 070611, 070612B, 070714A\\\\ &  070809, 070911, 070917, 071003, 071010B, 071028B, 071101, 071117, 071122,  {\\bf 080210}\\\\ &  080218B, 080319A, 080515, 080520, 080623, 080701, 080702A \\\\ &   \\\\ One break  &  051022, 070125, 080625, 050126, 050128, 050219A, 050318,  {\\bf 050401}, 050406, 050410\\\\ &  050421, 050422,  {\\bf 050502B},  {\\bf 050525A},  {\\bf 050712},  {\\bf 050714B},  {\\bf 050721},  {\\bf 050726},  {\\bf 050730},  {\\bf 050801}\\\\ &  050802, 050815,  {\\bf 050819}, 050824,  {\\bf 050904},  {\\bf 050915B}, 050916,  {\\bf 050922C}, 051008,  {\\bf 051016A}\\\\ &  051021B,  {\\bf 051227},  {\\bf 060111A},  {\\bf 060115},  {\\bf 060206},  {\\bf 060210}, 060211B,  {\\bf 060218}, 060223A,  {\\bf 060312}\\\\ &   {\\bf 060319},  {\\bf 060403},  {\\bf 060427}, 060428B, 060507,  {\\bf 060512},  {\\bf 060804}, 060805A,  {\\bf 060904B},  {\\bf 060908}\\\\ &   {\\bf 060912A}, 060927,  {\\bf 060929}, 061002, 061019, 061028, 061102, 061126,  {\\bf 061201},  {\\bf 070107}\\\\ &   {\\bf 070224},  {\\bf 070419B},  {\\bf 070517},  {\\bf 070518},  {\\bf 070520A}, 070520B,  {\\bf 070628},  {\\bf 070704}, 070721A,  {\\bf 070724A}\\\\ &   {\\bf 070808}, 070810A, 071001, 071010A, 071011,  {\\bf 071020},  {\\bf 071021}, 080129,  {\\bf 080207},  {\\bf 080303}\\\\ &   {\\bf 080319D},  {\\bf 080325}, 080409, 080413A, 080426, 080516, 080517,  {\\bf 080604},  {\\bf 080703}, 080710\\\\ &   \\\\ Canonical  &  070311,  {\\bf 050315}, 050319,  {\\bf 050416A},  {\\bf 050607},  {\\bf 050713A},  {\\bf 050713B},  {\\bf 050803},  {\\bf 050814},  {\\bf 050820A}\\\\ &   {\\bf 050822}, 050922B,  {\\bf 051016B},  {\\bf 051109A},  {\\bf 051109B},  {\\bf 051221A},  {\\bf 060105},  {\\bf 060108},  {\\bf 060109},  {\\bf 060204B}\\\\ &   {\\bf 050826} {\\bf 060219},  {\\bf 060313},  {\\bf 060413},  {\\bf 060418},  {\\bf 060428A},  {\\bf 060502A},  {\\bf 060510A},  {\\bf 060526},  {\\bf 060604},  {\\bf 060605}\\\\ &   {\\bf 060607A},  {\\bf 060614},  {\\bf 060707},  {\\bf 060708},  {\\bf 060712},  {\\bf 060714},  {\\bf 060719},  {\\bf 060729},  {\\bf 060807},  {\\bf 060814}\\\\ &   {\\bf 060904A},  {\\bf 060906},  {\\bf 060923A}, 060923C,  {\\bf 060926},  {\\bf 061004},  {\\bf 061021},  {\\bf 061121},  {\\bf 061202},  {\\bf 061222A}\\\\ &   {\\bf 070129}, 070219,  {\\bf 070220},  {\\bf 070306},  {\\bf 070328},  {\\bf 070420},  {\\bf 070429A},  {\\bf 070529},  {\\bf 070714B},  {\\bf 070721B}\\\\ &  070802,  {\\bf 071028A},  {\\bf 071112C},  {\\bf 071118},  {\\bf 080205},  {\\bf 080212},  {\\bf 080229A},  {\\bf 080310},  {\\bf 080320}, 080328\\\\ &   {\\bf 080330},  {\\bf 080430}, 080602,  {\\bf 080707},  {\\bf 080723A} \\\\ &   \\\\ Oddball  &  050505,  {\\bf 050716},  {\\bf 050717},  {\\bf 050724},  {\\bf 051001},  {\\bf 051117A},  {\\bf 060111B},  {\\bf 060124},  {\\bf 060202},  {\\bf 060211A}\\\\ &   {\\bf 060306},  {\\bf 060510B},  {\\bf 060522}, 060801,  {\\bf 060813},  {\\bf 061007},  {\\bf 061110A},  {\\bf 070103},  {\\bf 070110},  {\\bf 070208}\\\\ &   {\\bf 070223},  {\\bf 070412},  {\\bf 070419A},  {\\bf 070508},  {\\bf 070521},  {\\bf 070616},  {\\bf 070621},  {\\bf 071025},  {\\bf 071031},  {\\bf 071227}\\\\ &  080123,  {\\bf 080307},  {\\bf 080319B}, 080319C, 080411,  {\\bf 080413B},  {\\bf 080503}, 080506,  {\\bf 080523}, 080603B\\\\ &   {\\bf 080605},  {\\bf 080607},  {\\bf 080613B},  {\\bf 080714},  {\\bf 080721} \\\\  \\hline \\end{tabular} \\caption{Classification of GRB light curves. Bursts in bold are those for which a canonical light curve would be discernable. See text for details.} \\label{tab:types} \\end{center} \\end{table*}  \\clearpage"
    },
    "0812/0812.2980_arXiv.txt": {
        "abstract": "{For microquasars, the one time when these systems exhibit steady and powerful jets is when they are in the hard state.  Thus, our understanding of this state is key to learning about the disk/jet connection.  Recent observational and theoretical results have led to questions about whether we really understand the physical properties of this state, and even our basic picture of this state is uncertain.  Here, I discuss some of the recent developments and possible problems with our understanding of this state.  Overall, it appears that the strongest challenge to the standard truncated disk picture is the detection of broad iron features in the X-ray spectra, and it seems that either there is a problem with the truncated disk picture or there is a problem with the relativistic reflection models used to explain the broad iron features.}  \\FullConference{VII Microquasar Workshop: Microquasars and Beyond\\\\ September 1-5 2008\\\\ Foca, Izmir, Turkey}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "{\\em The most important point is that it really seems like the observations are telling us that there is a problem either with the truncated disk picture or with our interpretation of the broad iron features for black holes in the hard state.}  While thus far, the observations only indicate a problem down to a luminosity of $\\sim$1\\% $L_{\\rm Edd}$, there are strong implications for our entire understanding of the differences between the most fundamental black hole spectral states: Soft (or TD) and hard. Although the thermal disk component may also indicate a problem down to 0.1\\%, the evidence that this component requires a non-truncated disk is weaker than the broad iron features.  Thus, it is very important to determine if the relativistic interpretation for the broad iron features is the only viable one.  It has been suggested that the red wing of the iron line could be explained by Compton down-scattering of the line emission \\cite{lt07}; however, it is still not clear if this model can simultaneously explain the other properties of the hard state.  Also, it has been mentioned that it is not clear if the relativistic models are adequate for the case of stellar mass black holes.  Although they are often applied to the stellar mass case, the models mentioned above (pexriv and CDID) were originally developed for Active Galactic Nuclei (AGN). Some inaccuracies may be present when using pexriv for hotter accretion disks because of the way it deals with higher ionizations. While the CDID model deals with hot disks (thermal motions and higher ionization states) in a much improved manner, it is still true that the matter densities that are assumed are appropriate for AGN.  Developing a relativistic reflection model specifically for stellar mass black holes would be a positive step.  Also, more can be done on the observational side.  We currently have programs to try to extend the hard state studies to lower luminosities with {\\em XMM-Newton}, {\\em Suzaku}, and {\\em Swift}. In particular, we recently (2008 September) carried out a 100~ks {\\em Suzaku} observation of GX~339--4 at a significantly lower luminosity than the \\cite{tomsick08} {\\em Swift} and {\\em RXTE} study.  It will be very interesting to see if we find any change in the broad iron features.  \\vspace{5mm} \\emph{Acknowledgements} I would especially like to thank Emrah Kalemci for all his efforts in hosting this workshop.  I would also like to thank the Local and Scientific Organizing Committees for helping to make the meeting a great success.  I acknowledge many contributions from my collaborators, including Stephane Corbel, Emrah Kalemci, Phil Kaaret, Sera Markoff, Simone Migliari, Rob Fender, Charles Bailyn, and Michelle Buxton. I also acknowledge useful discussions with participants of the workshop.  I acknowledge partial support from NASA {\\em RXTE} Guest Observer grants NNG06GA81G and NNX06AG83G and from NASA {\\em Swift} Guest Observer grant NNX08AQ60G."
    },
    "0812/0812.0457_arXiv.txt": {
        "abstract": "We present Chandra detections of x-ray emission from the AGN in two giant Low Surface Brightness (LSB) galaxies, UGC~2936 and UGC~1455. Their x-ray luminosities are $1.8\\times10^{42}~ergs~s^{-1}$ and $1.1\\times10^{40}~ergs~s^{-1}$ respectively. Of the two galaxies, UGC~2936 is radio loud.  Together with another LSB galaxy UGC~6614 (XMM archival data) both appear to lie above the X-ray-Radio fundamental plane and their AGN have black hole masses that are low compared to similar galaxies lying on the correlation. However, the bulges in these galaxies are well developed and we detect diffuse x-ray emission from four of the eight galaxies in our sample. Our results suggest that the bulges of giant LSB galaxies evolve independently of their halo dominated disks which are low in star formation and disk dynamics. The centers follow an evolutionary path similar to that of bulge dominated normal galaxies on the Hubble Sequence but the LSB disks remain unevolved. Thus the bulge and disk evolution are decoupled and so whatever star formation processes produced the bulges did not affect the disks. ",
        "introduction": "Low Surface Brightness (LSB) galaxies are poorly evolved systems that have diffuse stellar disks, large HI disks and massive dark halos \\citep{ImpeyBothun1997}. The dark halo inhibits the growth of disk instabilities such as bars and spiral arms, this leads to an overall poor star formation rate over these galaxies \\citep{Boissier.etal.2008,O'Neil.etal.2007}. As a result LSB galaxies are metal poor \\citep{deNaray.etal.2004} and optically dim. Optical studies show that they span a wide range of morphologies from dwarfs to giant spirals \\citep{Beijersbergen.etal.1999}. Our paper focuses on giant LSB spirals which are characterised by prominent central bulges, optically dim disks and extended HI gas disks, a good example being Malin~1 \\citep{Barth.2007,ImpeyBothun1989}. These galaxies are also fairly isolated systems \\citep{Rosenbaum.Bomans}. The origin and evolution of these galaxies is still unclear; one possibility is that they form in low density enviroments and hence remain unevolved \\citep{Hoffman.etal.1992}.  Although LSB galaxy disks have been studied at length at optical and infrared wavelengths \\citep{Burkholder.etal.2001,Rahman.etal.2007,Hinz.etal.2007} not much is known about their nuclear properties. In particular, their nuclear black hole (BH) masses and AGN-bulge evolution history remain largely unconstrained. Optical spectroscopy indicates that a significant fraction have Active Galactic Nuclei (AGN) \\citep{Sprayberry.etal.1995} and large bulges \\citep{Schombert.1998}. Nuclear acivity has also been detected at radio wavelengths (Das et al. 2008, in preparation). However, the best way to detect and study AGN activity is using x-ray emission which will give rise to a compact x-ray bright source at the galaxy nucleus. There may also be a primordial x-ray emitting gaseous halo which will appear as diffuse x-ray emission associated with the galaxy or its bulge. Very little is known about BH masses in bulge dominated LSB galaxies but indirect studies suggest that it is low compared to normal galaxies \\citep{Pizzella.etal.2005}. This raises interesting questions regarding the evolution of BHs, AGN and bulges in poorly evolved and isolated galaxies.  In this paper we examine these issues with observations made using the Chandra X-ray Observatory.  Essentially nothing is known about the x-ray properties of LSB galaxies (possibly because the lack of star formation and hence the low number of high-mass x-ray binaries suggests low x-ray emission).  In the following sections we present the results of a pilot study of eight giant LSB galaxies with Chandra and discuss the implications of our observations. ",
        "conclusions": "We present Chandra observations of eight giant, LSB galaxies all of which have a sizable bulge. Our main results are the following.  \\noindent {\\bf (i)}~We have detected compact x-ray emission from the nuclei of two LSB galaxies, UGC~1455 ($L_{X}=1.1\\times10^{40}~erg^{-s}$) and UGC~2936 ($L_{X}=1.8\\times10^{42}~erg^{-s}$); it is due to AGN activity in these galaxies. The AGN emission is similar to that observed from the centers of nearby Seyfert and LINER galaxies.  \\noindent {\\bf (ii)}~For the galaxies UGC~2936 (our sample) and UGC~6614 (XMM archival data), we combined x-ray luminosities with radio luminosities and black hole masses to determine the location of these galaxies on the radio-x-ray fundamental plane. We find that both galaxies lie above the plane which suggests that their nuclei harbor less massive black holes compared to normal galaxies on the plane.  \\noindent {\\bf (iii)}~Diffuse x-ray emission was detected from the bulges of four galaxies. The luminosity is similar to that observed from the centers of nearby star forming galaxies. These results combined with AGN emission suggests that the AGN and bulges of LSB galaxies have followed an evolutionary path similar to bulge dominated bright galaxies even though their LSB disks are poorly evolved.  \\noindent {\\bf (iv)}~The detection of AGN and diffuse emission from the bulges of LSB galaxies shows that galaxies with unevolved disks can have normal bulges. Thus, whatever star formation processes made the bulge did not make the disk component in these galaxies. In fact the bulge and disk evolution appears to be distinctly decoupled in these galaxies. Giant LSB galaxies are thus good sites to study decoupled bulge-disk evolutionary processes."
    },
    "0812/0812.2452_arXiv.txt": {
        "abstract": "The power spectrum of number density perturbations of free electrons is obtained for the epoch of cosmological recombination of hydrogen. It is shown that amplitude of the electron perturbations power spectrum of scales larger than acoustic horizon exceeds  by factor of 17 the amplitude of baryon matter density ones (atoms and ions of hydrogen and helium). In the range of the first and second acoustic peaks such relation is 18, in the range of the third one  16. The dependence of such relations on cosmological parameters is analysed too. ",
        "introduction": "Important feature of the cosmological recombination of helium and hydrogen in the $\\Lambda$CDM-model is that it occurs when the total energy density of thermal radiation approximately is comparable to baryon matter one: $\\epsilon_{\\gamma}\\sim c^2\\rho_b$. Henceforth we will use the following definitions: $n_{\\rm HI}$ and $n_{\\rm HII}$ are number density of neutral and ionized hydrogen atoms; $n_{\\rm HeI}$, $n_{\\rm HeII}$ and $n_{\\rm HeIII}$ are number densities of neutral, singly and double ionized helium; $n_e=n_{\\rm HII}+n_{\\rm HeII}+2n_{\\rm HeIII}$ is number density of free electrons; $n_{\\rm H}=n_{\\rm HI}+n_{\\rm HII}$ is total number density of hydrogen nuclei; $n_{\\rm He}=n_{\\rm HeI}+n_{\\rm HeII}+n_{\\rm HeIII}$ is total number density of helium nuclei. We use the relative number densities (ionization fractions): $x_{\\rm HI}\\equiv n_{\\rm HI}/n_{\\rm H}$ is relative abundance of neutral hydrogen, $x_{\\rm HII}\\equiv n_{\\rm HII}/n_{\\rm H}$ are relative abundances of ionized hydrogen,  $x_{\\rm HeI}\\equiv n_{\\rm HeI}/n_{\\rm He}$, $x_{\\rm HeII}\\equiv n_{\\rm HeII}/n_{\\rm He}$ and  $x_{\\rm HeIII}\\equiv n_{\\rm HeIII}/n_{\\rm He}$-- relative abundances of neutral, singly and double ionized helium, $x_e\\equiv n_e/n_{\\rm H}$ -- relative number density of electrons. The ratio of total number densities of helium and hydrogen nuclei we define as $f_{\\rm He}\\equiv n_{\\rm He}/n_{\\rm H}$, which can be expressed via mass fraction of primordial helium $Y_P$, so that $f_{\\rm He}=Y_P/4(1-Y_P)$ (further we assume $Y_P=0.24$ \\cite{schramm1998}). These quantities obey obvious relationships: $x_e=x_{\\rm HII}+f_{\\rm He}x_{\\rm HeII}+2f_{\\rm He}x_{\\rm HeIII}$, $x_{\\rm HI}+x_{\\rm HII}=1$, $x_{\\rm HeI}+x_{\\rm HeII}+x_{\\rm HeIII}=1$. The total mass density of baryons can be expressed via hydrogen number density and mass fraction of primordial helium in the following way: $\\rho_{ b}\\approx m_p n_{b}=m_p(n_{\\rm H}+4n_{\\rm He})=m_pn_{\\rm H}(1+4f_{\\rm He})$, where $n_{b}$ is mean number density of baryons (protons and neutrons), where $m_{ p}$ is mass of proton.  At early stages of evolution of the Universe ($z>10^4$) all hydrogen and helium atoms were ionized completely by thermal  photons, so $x_{\\rm HII}=1$, $x_{\\rm HI}=0$, $x_{\\rm HeIII}=1$, $x_{\\rm HeI}=x_{\\rm HeII}=0$  and $x_e=1+2f_{\\rm He}$ \\cite{seager1999,seager2000}.  At $z\\sim 8000$ thermal photons with energies higher than ionization potential of ${\\rm HeII}$ from ground level reside in the short-wave tail of Planck function and their number density becomes too low to keep all helium in the ionization state of ${\\rm HeIII}$. It begins to recombine and $\\rm HeII$ ions  appear. Recombination of $\\rm HeII$ occurs in the conditions of local thermodynamic equilibrium (LTE), so, the ionization fraction of helium, $x_{\\rm HeIII}$, is described by Saha equation: \\begin{equation} \\label{SahaHeIII} {x_ex_{\\rm HeIII}\\over x_{\\rm HeII}} = {(2\\pi m_e k T_m)^{3/2}\\over h^3 n_{\\rm H}} e^{-\\chi_{\\rm HeII}/kT_m}, \\end{equation} where $T_m$ is matter temperature, $m_{e}$ is mass of electron, $h$ is Planck constant, $k$ is Boltzmann constant, $\\chi _{\\rm HeII}$ is ionization potential of HeII. Since at this epoch both hydrogen and helium are completely ionized  ($x_{\\rm HI}=0$, $x_{\\rm HII}=1$, $x_{\\rm HeI}=0$) we have $x_{\\rm HeII}=1-x_{\\rm HeIII}$ and $x_e=1+f_{\\rm He}(1+x_{\\rm HeIII})$, so equation (\\ref{SahaHeIII}) can be easily solved for $x_e$. Using it one can easily check that  already at $z\\sim 5000$ all helium atoms become singly ionized. Such state is kept up to $z\\sim 3500$ when $\\rm HeI$ begins to recombine. At this time the conditions are close to LTE yet. The metastable $2s$ level plays insignificant role in deviation of radiative recombination rate of ${\\rm HeI}$ from LTE one until the part of HeI is less than 1\\% of  total  helium content and ionized fraction $x_{\\rm HeII}$ is described yet enough accurately by Saha equation: \\begin{equation} \\label{SahaHeII} {x_ex_{\\rm HeII}\\over x_{\\rm HeI}} = 4{(2\\pi m_e k T_m)^{3/2}\\over h^3 n_{\\rm H}} e^{-\\chi_{\\rm He I}/kT_m}, \\end{equation} where $\\chi _{\\rm HeI}$ is ionization potential of HeI. Now $x_{\\rm HeIII}=0$ and $x_{\\rm HeI}= 1-x_{\\rm HeII}$. For accurate calculation of $x_{\\rm HeII}$ we must have the exact value of $x_e=x_{\\rm HII}+f_{\\rm He}x_{\\rm HeII}$. Despite $x_{\\rm HII}\\approx1$, the decrease of $n_{\\rm HII}$ in 0.1\\% caused by the hydrogen recombination leads to the variation of $n_e$ comparable to one caused by ${\\rm HeI}$ recombination because of the domination of hydrogen ($f_{\\rm H}=n_{\\rm H}/(n_{\\rm H}+n_{\\rm He})=0.921$). So, at this step the hydrogen recombination already must be taken into account too. The ionized fraction $x_{\\rm HII}$ is described yet enough accurately by Saha equation: \\begin{equation} \\label{SahaHII} {x_ex_{\\rm HII}\\over x_{\\rm HI}} = {(2\\pi m_e k T_m)^{3/2}\\over h^3 n_{\\rm H}} e^{-\\chi_{\\rm HI}/kT_m}, \\end{equation} where $\\chi _{\\rm HI}$ is ionization potential of HI. The system of these two equations can be reduced to the single cubic equation for $x_e$, which has one real root: \\begin{equation} \\label{xe3000} x_e=2\\sqrt{-A/3}\\cos{(\\alpha/3)}-B/3, \\end{equation} where $B=R_{\\rm HI}+R_{\\rm HeI}$, $R_{\\rm HeI}$ and $R_{\\rm HI}$ is right hand of equations (\\ref{SahaHeII}) and (\\ref{SahaHII}), $\\cos{\\alpha}=C/2\\sqrt{-A^3/27}$, $A=D-B^2/3$, $D=R_{\\rm HI}R_{\\rm HeI}-R_{\\rm HI}-f_{\\rm He}R_{\\rm HeI}$, $C=2B^3/27-BD/3-E$, $E=-R_{\\rm HI}R_{\\rm HeI}(1-f_{\\rm He})$. The code RECFAST was complemented by this solution in order to achieve more accurate computations of number density perturbations of ions. However, it does not change the results of calculations of unperturbed $x$'s noticeably \\cite{novosyadlyj2006a,novosyadlyj2006b}.  Metastable levels 2s HeI and HI cause a delay of recombination of HeII$\\rightarrow$HeI and HII$\\rightarrow$HI, violation of LTE population of levels and equilibrium of ionization-recombination processes (''bottleneck'' effect). Saha equation does not already describe adequately the recombination and equations of detailed balance must be used \\cite{seager1999}: \\begin{eqnarray} \\label{eqHeII} {dx_{\\rm HeII}\\over dz} = \\left(x_{\\rm HeII}x_e n_{\\rm H} \\alpha_{\\rm HeI} - \\beta_{\\rm HeI} (1-x_{\\rm HeII}) e^{-h\\nu_{\\rm HeI2^1s}/kT_m}\\right)\\nonumber \\\\ \\times {1 + K_{\\rm HeI} \\Lambda_{\\rm He} n_{\\rm H} (1-x_{\\rm HeII})e^{-h\\nu_{ps}/kT_m} \\over H(z)(1+z)\\left(1+K_{\\rm HeI} (\\Lambda_{\\rm He} + \\beta_{\\rm HeI}) n_{\\rm H} (1-x_{\\rm HeII}) e^{-h\\nu_{ps}/kT_m}\\right)}, \\end{eqnarray} and \\begin{eqnarray} \\label{eqHII} {dx_{\\rm HII}\\over dz} = \\left(x_ex_{\\rm HII} n_{\\rm H} \\alpha_{\\rm H} - \\beta_{\\rm H} (1-x_{\\rm HII}) e^{-h\\nu_{\\rm HI 2s}/kT_m}\\right)\\nonumber \\\\ \\times{1 + K_{\\rm H} \\Lambda_{\\rm H} n_{\\rm H}(1-x_{\\rm HII}) \\over H(z)(1+z)\\left(1+K_{\\rm H} (\\Lambda_{\\rm H} + \\beta_{\\rm H}) n_{\\rm H} (1-x_{\\rm HII}) \\right)}, \\end{eqnarray} where \\begin{eqnarray} \\label{alphaHeI} \\alpha_{\\rm HeI} =q\\left[\\sqrt{T_m\\over T_2}\\left(1+\\sqrt{T_m\\over T_2}\\right)^{1-p} \\left(1+\\sqrt{T_m\\over T_1}\\right)^{1+p}\\right]^{-1}\\!   \\mathrm {m^{3}c^{-1}}, \\end{eqnarray} \\begin{eqnarray} \\label{alphaHI} \\alpha_{\\rm H} = F\\cdot 10^{-19}at^{b}/(1 + ct^{d}),\\, \\mathrm{m^{3}c^{-1}} \\;\\;\\;  t=T_m/10^4. \\end{eqnarray} $\\alpha_{\\rm HeI}$ and $\\alpha_{\\rm H}$ are effective recombination coefficients of helium \\cite{hummer1998} and hydrogen \\cite{pequignot1991}, respectively. $K_{\\rm HeI}\\equiv \\lambda^3_{\\rm HeI2^1p}/[8\\pi H(z)]$ is the factor taking into account the cosmological redshifting of HeI  $2^1p-1^1s$  photons and $K_{\\rm H}\\equiv \\lambda^3_{\\rm H2p}/[8\\pi H(z)]$ is the factor taking into account the cosmological redshifting of Ly$\\alpha$ photons. The effective photoionization coefficients in  (\\ref{eqHeII}) and  (\\ref{eqHII}) are calculated via effective recombination coefficients as follows: \\begin{eqnarray} \\label{beta} \\beta=\\alpha (2\\pi m_e k T_m/h^2)^{3/2}e^{-h\\nu_{2s-1s}/kT_m}. \\end{eqnarray}  Before $z\\sim 800$ the matter temperature $T_m$ practically equals radiation temperature $T_{\\rm R}$ since until the time-scale of Thomson scattering remains essentially lower than the time-scale of expansion of the Universe, $t_{\\rm T}/t_{\\rm Hubble}<10^{-3}$. Therefore, the rate of temperature decreasing is governed by adiabatic cooling of radiation ($\\gamma=4/3$) caused by expansion of the Universe: \\begin{equation} \\label{Tm1} \\frac{dT_m}{dz} =\\frac{T_m}{(1+z)}. \\end{equation} And only after recombination at $z<800$ adiabatic cooling of ideal gas ($\\gamma=5/3$) begins to dominate over the heating caused by Compton effect which is the main process of energy transfer between electrons and photons. Cooling of plasma via free-free, free-bound and bound-bound transitions and collisional ionization as well as heating via photoionization and collisional recombination gives insignificant contribution into the rate of temperature change, it does not  exceed the  0.01$\\%$ of main processes -- adiabatic cooling and heating by Compton effect  \\cite{seager2000}. So, at this epoch the following equation for rate of temperature decreasing is enough accurate \\cite{seager1999}: \\begin{equation} \\label{Tm2} \\frac{dT_m}{dz} = \\frac{8\\sigma_{\\rm Th}a_{\\rm R} T_{\\rm R}^4}{3H(z)(1+z)m_ec}\\, \\frac{x_e}{1+f_{\\rm He}+x_e}\\,(T_m - T_{\\rm R}) + \\frac{2T_m}{(1+z)}, \\nonumber \\end{equation} The Table\\ref{tab} lists the values of all atomic constants and coefficients used in the equations (\\ref{SahaHeIII})-(\\ref{Tm2}).  In Fig.\\ref{xe} the relative number density of electrons $x_{e}$ at range $400\\le z\\le 10000$ is presented. The visibility function $d\\tau /dz e^{-\\tau}$ (dotted line) shows the region of the largest cosmological recombination rate and decoupling epoch -- here $\\tau$ is optical depth caused by Thomson scattering by electrons, $z=(a^{-1}-1)$ is redshift.  \\begin{figure} \\centerline{\\includegraphics[height=10cm]{xe.eps}} \\caption{Top panel: the dependence of relative number density of free electrons on redshift in $\\Lambda$CDM-model with parameters $\\Omega_{\\Lambda}=0.736$, $\\Omega_m=0.278$, $\\Omega_{b}=0.05$, ${\\rm h}=0.68$ \\cite{apunevych2007} (solid line) and  $\\Omega_{\\Lambda}=0.76$, $\\Omega_m=0.24$, $\\Omega_{b}=0.042$, ${\\rm h}=0.73$ \\cite{spergel2007} (dashed line overlaps with solid one). Dotted line represents visibility function $d\\tau /dz e^{-\\tau}$ ($\\times 270$). The relative difference $x_e$ for two sets of parameters of $\\Lambda$CDM-model is shown in bottom panel.} \\label{xe} \\end{figure} The calculations shown in Fig.\\ref{xe} have been done for $\\Lambda$CDM model with parameters  $\\Omega_{\\Lambda}=0.736$, $\\Omega_m=0.278$, $\\Omega_{b}=0.05$, ${\\rm h}=0.68$ \\cite{apunevych2007} and $\\Omega_{\\Lambda}=0.76$, $\\Omega_m=0.24$, $\\Omega_{b}=0.042$, ${\\rm h}=0.73$ \\cite{spergel2007}. The curves  $x_e(a)$ for two sets of parameters practically overlap: the difference $\\sim4\\%$ is in the range of maximum of visibility function ($z_{dec}\\approx1080$) and increases up to $\\sim6\\%$ for residual ionization $x_e\\sim 10^{-3}$ at $z\\le 900$. ",
        "conclusions": ""
    },
    "0812/0812.3568_arXiv.txt": {
        "abstract": "In the next years, several cosmological surveys will rely on imaging data to estimate the redshift of galaxies, using traditional filter systems with $4-5$ optical broad bands; narrower filters improve the spectral resolution, but strongly reduce the total system throughput. We explore how photometric redshift performance depends on the number of filters $n_f$, characterizing the survey depth by the fraction of galaxies with unambiguous redshift estimates. For a combination of total exposure time and telescope imaging area  of $270$ hrs m$^2$, $4-5$ filter systems perform  significantly worse, both in completeness depth and precision, than systems with $n_f\\gtrsim 8$ filters. Our results suggest that for low $n_f$, the color-redshift degeneracies overwhelm the improvements in photometric depth, and that even at higher $n_f$, the effective {\\it photometric redshift depth} decreases much more slowly with filter width than naively expected from the reduction in $S/N$. Adding near-IR observations improves the performance of low $n_f$ systems, but still the system which maximizes the photometric redshift completeness is formed by $9$ filters with logarithmically increasing bandwidth (constant resolution) and half-band overlap, reaching $\\sim 0.7$~mag deeper, with $10\\%$ better redshift precision, than $4-5$ filter systems. A system with $20$ constant-width, non-overlapping filters reaches only $\\sim 0.1$~mag shallower than $4-5$ filter systems, but has a precision almost 3 times better,  $\\delta z = 0.014(1+z)$ vs $\\delta z = 0.042(1+z)$. We briefly discuss a practical implementation of such a photometric system: the ALHAMBRA survey. ",
        "introduction": "Photometric redshift estimation is not a new technique (Baum 1962, Loh \\& Spillar 1986, Koo 1985, Connolly et al. 1995; see Koo 1999 for a history of the method), but it has considerably developed in the last decade, especially following the Hubble Deep Field (HDF) observations (Williams 1996, Casertano et al. 2000), which provided catalogs with excellent photometric quality and abundant spectroscopic redshift coverage. This allowed astronomers to thoroughly test standard photo-z techniques and try new approaches (Gwyn \\& Hartwick 1996, Lanzetta, Yahil \\& Fern\\'{a}ndez-Soto 1997, Sawicki, Lin \\& Yee 1997, Fern\\'{a}ndez-Soto, Lanzetta \\& Yahil 1999, Brunner et al. 1997, Ben\\'{\\i}tez et al. 1999, Ben\\'{\\i}tez 2000, Bolzonella, Miralles \\& Pell\\'{o} 2000).  As Hickson, Gibson \\& Callaghan (1994) first showed, multiband narrow filters can be much more efficient to obtain redshifts than spectroscopy if the large area of the imaging cameras is factored in. Several photometric surveys, using different filter systems, have been proposed or implemented in the last decade: the UBC-NASA survey (Hickson \\& Mulrooney 1998), CADIS (Wolf et al. 2001b), COMBO-17 (Wolf et al. 2001a), COSMOS-21 (Taniguchi 2004), ALHAMBRA (Moles et al. 2008), DES (DES collaboration 2005), LSST (Tyson 2006), PanStarrs (Kaiser 2007), VST (Arnaboldi et al. 2007),  and PAU (Ben\\'\\i tez et al. 2008). These surveys represent powerful  alternatives to deep spectroscopic surveys like DEEP2 (Davis et al. 2003), VVDS (Le F\\`{e}vre et al. 2003), or BOSS (Schlegel et al. 2007) at least for those scientific goals which only require limited redshift accuracy and low resolution spectral information.  However at least three of the imaging surveys (DES, LSST, PanStarrs) will work with photometric systems with $4-5$ optical broadband filters, similar to those traditionally used in Astronomy. It is obvious that using more, narrower filters, for a fixed exposure time, will significantly sacrifice photometric depth. However, photometric depth is not equivalent to {\\it photometric redshift} depth. The fewer the filters, the more prone the system is to color-redshift degeneracies; these make impossible to unambiguously determine the redshift for a galaxy, even if observed at relatively high $S/N$. The Hubble Ultra Deep Field (Beckwith et al. 2003) offers a good example. Despite the fact that the limiting magnitude in the HUDF is $0.9-1.4$ mag deeper than the HDF, the lack of a $U-$band filter in the HUDF makes the photometric redshift depth of both fields similar (Coe et al. 2006).  This letters explores the impact on photometric redshift performance of factors as the number of filters $n_{f}$, constant vs. logarithmically increasing bandwidth, half-band overlaps, and near-IR observations. We also briefly discuss a practical implementation of a medium band filter system: the ALHAMBRA survey. ",
        "conclusions": "We explore the performance of four different uniform filter systems with constant and logarithmically increasing ($\\Delta\\lambda\\propto\\lambda$) widths, and with half-width $\\Delta\\lambda/2$ overlaps or just a minimal overlap corresponding to the filter wings, and use, as a measure of survey effective depth, the fraction of galaxies with a compact, unimodal probability redshift distributions as a function of magnitude. Our simulations employ a realistic input catalog, based on HDF photometric redshifts, and correspond to a combination of total exposure time and telescope area of $270$ hrs m$^2$. We find that traditional $4-5$ optical filter systems clearly underperform, both in terms of completeness magnitude limit and precision, systems with $n_f\\gtrsim 8$ filters.  Our results suggest that for low $n_f$, the effect of color-redshift degeneracies dominates the advantages of increased photometric depth, and that even at higher $n_f$, the effective {\\it photometric redshift depth} decreases much more slowly with filter width than naively expected from the reduction in $S/N$. Adding near-IR observations increases the overall depth, alleviating color-redshift degeneracies and improving the relative performance of low $n_f$ systems. However the optimum performance still comes from a system with $9$ filters with logarithmically increasing bandwidth (constant resolution) and half-band overlap, which reaches $\\sim 0.7$mag deeper, with $10\\%$ better redshift precision, than $4-5$ filter systems. For many scientific applications, which require both precision and depth, the use of $>15$ medium band filters is clearly advantageous. A system with $20$ constant-width, non-overlapping filters reaches only $\\sim 0.1$mag shallower than $4-5$ filter systems, but has a precision almost 3 times better,  $\\delta z = 0.014(1+z)$ vs $\\delta z = 0.042(1+z)$, as a practical implementation of such a system, the ALHAMBRA survey, shows.  Since it is well known that color-redshift degeneracies worsen with magnitude depth, it can be expected that the relative decoupling between photometric depth and photometric redshift depth described here will be more significant for surveys which reach fainter limits than those considered in our simulations, and less important for shallower observations, where the color/redshift degeneracies are less of a problem. In any case, future projects will have to seek an optimum number of filters based on their particular observing parameters and science goals."
    },
    "0812/0812.1792_arXiv.txt": {
        "abstract": "{}% {We study turbulent transport coefficients that describe the evolution of large-scale magnetic fields in turbulent convection.}% {We use the test field method, together with three-dimensional numerical simulations of turbulent convection with shear and rotation, to compute turbulent transport coefficients describing the evolution of large-scale magnetic fields in mean-field theory in the kinematic regime. We employ one-dimensional mean-field models with the derived turbulent transport coefficients to examine whether they give results that are compatible with direct simulations.}% {The results for the $\\alpha$-effect as a function of rotation rate are consistent with earlier numerical studies, i.e.\\ increasing magnitude as rotation increases and approximately $\\cos \\theta$ latitude profile for moderate rotation. Turbulent diffusivity, $\\etat$, is proportional to the square of the turbulent vertical velocity in all cases. Whereas $\\etat$ decreases approximately inversely proportional to the wavenumber of the field, the $\\alpha$-effect and turbulent pumping show a more complex behaviour with partial or full sign changes and the magnitude staying roughly constant. In the presence of shear and no rotation, a weak $\\alpha$-effect is induced which does not seem to show any consistent trend as a function of shear rate. Provided that the shear is large enough, this small $\\alpha$-effect is able to excite a dynamo in the mean-field model. The coefficient responsible for driving the shear-current effect shows several sign changes as a function of depth but is also able to contribute to dynamo action in the mean-field model. The growth rates in these cases are, however, well below those in direct simulations, suggesting that an incoherent $\\alpha$-shear dynamo may also act in the simulations. If both rotation and shear are present, the $\\alpha$-effect is more pronounced. At the same time, the combination of the shear-current and $\\bm{\\Omega}\\times{\\bm J}$-effects is also stronger than in the case of shear alone, but subdominant to the $\\alpha$-shear dynamo. The results of direct simulations are consistent with mean-field models where all of these effects are taken into account without the need to invoke incoherent effects.}% {}% ",
        "introduction": "The solar magnetic field is thought to arise from a complicated interplay of turbulence, rotation, and large-scale shear flows (e.g.\\ Ossendrijver \\cite{O03} and references therein). Whilst numerical simulations of simple systems using fully periodic boxes and externally forced idealised flows exhibiting large-scale dynamos have been around for some time (e.g.\\ Brandenburg \\cite{B01,B05a}; Brandenburg et al.\\ \\cite{BBS01}; Mininni et al.\\ \\cite{MGM05}; Brandenburg \\& K\\\"apyl\\\"a \\cite{BK07}; Yousef et al.\\ \\cite{YHSea08a,YHSea08b}; K\\\"apyl\\\"a \\& Brandenburg \\cite{KB08}) and dynamos driven by the magnetorotational instability exhibit large-scale dynamos (e.g.\\ Brandenburg et al.\\ \\cite{BNST95}; Hawley et al.\\ \\cite{HGB96}), convection simulations have not been able to produce appreciable large-scale magnetic fields until recently (Rotvig \\& Jones \\cite{RJ02}; Browning et al.\\ \\cite{BMBT06}; Brown et al.\\ \\cite{BBBMNT07}; K\\\"apyl\\\"a et al.\\ \\cite{KKB08}, hereafter Paper I; Hughes \\& Proctor \\cite{HP08}). The main ingredient missing in many earlier simulations was a large-scale shear flow and boundary conditions which allow magnetic helicity fluxes out of the system. Indeed, the shear flow plays a dual role in dynamos: it not only generates new magnetic fields by stretching, but it also drives magnetic helicity fluxes along constant isocontours of shear which can allow efficient dynamo action (Vishniac \\& Cho \\cite{VC01}; Brandenburg \\& Subramanian \\cite{BS05}; Paper I). Recently, however, large-scale dynamos have also been found from rigidly rotating convection simulations without shear (K\\\"apyl\\\"a et al.\\ \\cite{KKB08b}).  Although large-scale magnetic fields can clearly be obtained from simulations, the origin of these fields in many cases (e.g.\\ Yousef et al.\\ \\cite{YHSea08a,YHSea08b}; Paper I; Hughes \\& Proctor \\cite{HP08}) is still uncertain. In the mean-field framework (e.g.\\ Moffatt \\cite{M78}; Parker \\cite{P79}; Krause \\& R\\\"adler \\cite{KR80}; R\\\"udiger \\& Hollerbach \\cite{RH04}), the dynamo process is described by turbulent transport coefficients that govern the evolution of large-scale magnetic field. The evolution equation for the large-scale part is obtained from the standard induction equation by decomposing magnetic and velocity fields into their mean and fluctuating parts, i.e.\\ ${\\bm B}= \\mean{\\bm B} + {\\bm b}$, ${\\bm U}= \\mean{\\bm U} + {\\bm u}$, which leads to \\begin{equation} \\frac{\\pd \\mean{\\bm B}}{\\pd t} = \\CURL({\\mean{\\bm U} \\times \\mean{\\bm B} + \\bm{\\mathcal{\\mean{E}}}} - \\eta \\mu_0 \\mean{\\bm J})\\;, \\label{equ:meanind} \\end{equation} where $\\eta$ is the molecular magnetic diffusivity, $\\mean{\\bm J} = \\mu_0^{-1}\\CURL{\\mean{\\bm B}}$ is the current density, and $\\mu_0$ is the vacuum permeability. The remaining term, $\\mean\\emf \\equiv \\overline{\\bm{u} \\times \\bm{b}}$, is the electromotive force describing the effects of small-scale turbulence on the evolution of mean fields and can be represented in terms of the mean fields and their derivatives \\begin{equation} \\mathcal{\\mean{E}}_i = \\alpha_{ij} \\mean{B}_j + \\eta_{ijk} \\mean{B}_{j,k} + \\ldots, \\label{equ:emf} \\end{equation} where $\\alpha_{ij}$ and $\\eta_{ijk}$ are tensorial coefficients, commas denote partial derivatives, and summation over repeated indices is assumed. Expression (\\ref{equ:emf}) is valid if the mean fields vary slowly in space and time.  Whilst mean-field models have been quite successful in reproducing many aspects of the solar magnetism (e.g.\\ Ossendrijver \\cite{O03}), they have often been hampered by the poor knowledge of the turbulent transport coefficients which could only be computed analytically using unrealistic or unjustified approximations, such as first order smoothing (FOSA). More recently, numerical models of convection in local Cartesian geometry have been employed to compute some of these coefficients in more realistic setups (Brandenburg et al.\\ \\cite{BTNPS90}; Ossendrijver et al.\\ \\cite{OSB01}, \\cite{OSRB02}; Giesecke et al.\\ \\cite{GZR05}; K\\\"apyl\\\"a et al.\\ \\cite{KKOS06}; Cattaneo \\& Hughes \\cite{CH06}; Hughes \\& Cattaneo \\cite{HC08}). To date, however, only coefficients relevant for the $\\alpha_{ij}$ term in Eq.~(\\ref{equ:emf}) have been determined from convection simulations. This is due to the limitations of the method used where a uniform magnetic field is imposed and the resulting electromotive force is measured. Furthermore, if the Lorentz force is retained in the simulations, dynamo-generated magnetic fields may grow to saturation, leading to quenching even if the imposed field is weak. At large magnetic Reynolds numbers such quenching can be very strong if there are no magnetic helicity fluxes, suggesting therefore small values of $\\alpha$ even for weak imposed fields.  During recent years an improved scheme of extracting turbulent transport coefficients has appeared which is referred to as the test field method (Schrinner et al.\\ \\cite{SRSRC05,Shea07}). In the test field method the velocity field of the simulation is used in a number of induction equations, which all correspond to a given set of large-scale test fields which do neither evolve nor react back onto the velocity field. The test fields are orthogonal so the coefficients can be obtained by matrix inversion. This method has been used successfully in setups where the turbulence is due to isotropic forcing without shear (Sur et al.\\ \\cite{SBS08}, Brandenburg et al.\\ \\cite{BRS08}) and with shear (Brandenburg et al.\\ \\cite{BRRK08}; Mitra et al.\\ \\cite{MKTB08}), respectively. Moreover, the method has been used to extract dynamo coefficients from more realistic setups where the turbulence is driven by supernovae (Gressel et al.\\ \\cite{GZER08}) and the magnetorotational instability (Brandenburg \\cite{B05b,B08}).  In the present paper we apply the method for the first time to convection simulations. We also seek to understand the dynamos reported in Paper I by applying the derived coefficients in a one-dimensional mean-field model. In the case of convection with rigid rotation it is likely that the large-scale fields are due to the turbulent $\\alpha$-effect that is present in helical flows (K\\\"apyl\\\"a et al.\\ \\cite{KKB08b}). However, when shear is present, there are various mechanisms that can generate large-scale fields: in helical flows a finite $\\alpha$-effect (e.g.\\ R\\\"adler et al.\\ \\cite{RKR03}; R\\\"adler \\& Stepanov \\cite{RS06}; R\\\"udiger \\& Kitchatinov \\cite{RK06}) with shear can excite a classical $\\alpha\\Omega$ or $\\alpha$-shear-dynamo (e.g.\\ Brandenburg \\& K\\\"apyl\\\"a \\cite{BK07}; K\\\"apyl\\\"a \\& Brandenburg \\cite{KB08}). Even if the mean value of $\\alpha$ is zero, strong enough fluctuations about zero in combination with shear can drive an incoherent $\\alpha$-shear dynamo (e.g.\\ Vishniac \\& Brandenburg \\cite{VB97}; Proctor \\cite{P07}). Finally, the shear-current (Rogachevskii \\& Kleeorin \\cite{RK03,RK04}; Kleeorin \\& Rogachevskii \\cite{KR08}) and $\\bm{\\Omega}\\times{\\bm J}$ (R\\\"adler \\cite{R69}; R\\\"adler et al. \\cite{RKR03}; Pipin \\cite{P08}) effects may operate even in nonhelical turbulence. If both rotation and shear are present in the system it is not obvious how to distinguish between the shear-current and $\\bm{\\Omega}\\times{\\bm J}$ effects. In the present paper we are able to extract the relevant turbulent transport coefficients responsible for most of these processes and determine which one of them is dominant in the different cases with the help of a one-dimensional mean-field model. In order to facilitate comparisons between the mean-field models and the direct simulations presented in Paper I, we use identical setups and overlapping parameter regimes as those used in Paper I in the determination of the transport coefficients. ",
        "conclusions": "We obtain turbulent transport coefficients governing the evolution of large-scale magnetic fields from turbulent convection simulations with the test field method. We study the system size and magnetic Reynolds number dependences of the coefficients. This is important because spurious results can be expected for small Reynolds numbers or when the aspect ratio of the domain is too small (Hughes \\& Cattaneo \\cite{HC08}). We find that for our standard system size, $L_{\\rm H}/d=4$, the coefficients are essentially identical to those obtained with a horizontal extent that is twice as large. As a function of $\\Rm$, all the coefficients are essentially constant for $\\Rm\\ga8$. This is in accordance with the theory but at odds with results from certain imposed field calculations (e.g.\\ Cattaneo \\& Hughes \\cite{CH06}). In these calculations the magnetic field is allowed to evolve until saturation which can cause strong quenching even if the imposed field itself is weak. This is particularly important if closed boundary conditions for the magnetic field are imposed, in which case magnetic helicity conservation can lead to catastrophic quenching (Vainshtein \\& Cattaneo \\cite{VC92}). More reliable results for the kinematic $\\alpha$-effect with the imposed field method can be obtained by resetting the magnetic field before it grows too large or develops substantial gradients (e.g.\\ Ossendrijver at al.\\ \\cite{OSRB02}; K\\\"apyl\\\"a et al.\\ \\cite{KKOS06}). More detailed comparison of the imposed field and test field methods is important, but beyond the scope of the present study.  The earlier determinations of transport coefficients from convection simulations have used the imposed field method (e.g.\\ Ossendrijver et al.\\ \\cite{OSB01,OSRB02}; K\\\"apyl\\\"a et al.\\ \\cite{KKOS06}) which yields the components of $\\alpha_{ij}$ but does not deliver $\\eta_{ij}$ because the imposed field is uniform. The test field method does not suffer from this restriction and $\\eta_{ij}$ and the $k$-dependence of the coefficients can be extracted. We find that for $k/k_1=0$, i.e.\\ for a uniform field, the results for $\\alpha$ and $\\gamma$ are consistent with those obtained from imposed field calculations, provided the magnetic field is reset before it grows too large and substantial gradients develop. As $k$ is increased, however, the qualitative behaviour of the coefficients changes. This is indicated by a partial sign change of $\\alpha$ and a complete sign change of $\\gamma$; see Figs.~\\ref{fig:kdep_128a} and \\ref{fig:kdep_128b}.  The turbulent diffusivity shows a robust behaviour regardless of the parameters of the simulations: the profile is proportional to the vertical velocity squared, $\\mean{u_z^2}$, as predicted by FOSA (e.g.\\ R\\\"adler \\cite{R80}). The value of $\\etat$ decreases almost proportional to $k^{-1}$, and shows a declining trend as a function of rotation and shear.  For the present parameters, the $\\alpha$-effect increases monotonically as rotation is increased. As a function of latitude, the diagonal components of $\\alpha_{ij}$ have a similar magnitude and peak near the pole with declining values towards the equator. The $\\alpha$-effect induced by shear is highly anisotropic: the $\\alpha_{xx}$ component has a similar profile and magnitude, but opposite sign, as $\\alpha_{xx}$ and $\\alpha_{yy}$ in the case of only rotation. This component also increases monotonically as a function of shear, whereas the shear-induced $\\alpha_{yy}$ remains small regardless of the strength of the shear. In the runs where rotation and shear are present, the diagonal components of $\\alpha_{ij}$ are roughly the sums of the corresponding coefficients in the cases with rotation and shear alone.  In addition to the $\\alpha$-effect, the $\\eta_{yx}$ component can contribute to a shear-current dynamo when $\\eta_{yx}S>0$. In our case, where $S<0$, such dynamo action is possible if $\\eta_{yx}<0$. We find that this coefficient shows negative regions near the base and near the top of the convectively unstable region, but the errors are of the same order of magnitude as the negative mean values in most cases.  In order to connect to earlier work, we use the test field results in a one-dimensional mean-field model in order to understand the excitation of dynamos using identical setups as in direct simulations (Paper I). We study here only the cases with shear and consider large-scale dynamos in the rigidly rotating case elsewhere (K\\\"apyl\\\"a et al.\\ \\cite{KKB08b}). The presently used dynamo model ignores $k$-dependence and is therefore likely to be too simple to fully describe the large-scale fields in the direct simulations. Nevertheless, the present results, taken at face value, seem to indicate that in the case with shear alone the derived dynamo coefficients are not sufficient to explain the dynamo but that an additional incoherent $\\alpha$-shear dynamo might be needed. This conjecture is based on the fact that mean-field $\\alpha$-shear and shear-current dynamos are both excited with similar growth rates which, however, are significantly smaller than those obtained from direct simulations in Paper I. Furthermore, a large-scale dynamo was marginal for $\\Sh=-0.03$ in Paper I, whereas for $\\Sh=-0.06$ it was found to be slightly subcritical in the present study. On the other hand, for the case with both shear and rotation, no additional incoherent effects seem to be needed. We find that in this case the regular $\\alpha$-shear dynamo produces larger growth rates than the combined shear-current and $\\bm{\\Omega}\\times\\bm{J}$ dynamo but neither effect alone seems to be strong enough to explain the dynamos in Paper I.  On a more general level, mean-field dynamo models of the Sun and other stars rely on parameterisations of turbulent transport coefficients. Even today, the majority of solar dynamo models bypass this problem and ignore most of the turbulent effects and rely on phenomenological descriptions of the $\\alpha$-effect and turbulent diffusion that are not without problems theoretically. On the other hand, some attempts have been made to incorporate the results for the transport coefficients from imposed field studies in mean-field models of the solar magnetism (e.g.\\ K\\\"apyl\\\"a et al.\\ \\cite{KKT06}; Guerrero \\& de Gouveia Dal Pino \\cite{GG08}) and models employing more general turbulence models have recently appeared (e.g.\\ Pipin \\& Seehafer \\cite{PS08}). We feel that this is a worthy cause to follow further with the present results."
    },
    "0812/0812.2791_arXiv.txt": {
        "abstract": "The Friedmann equations of general relativity can be derived from the first law of thermodynamics when the entropy of the apparent horizon of a spatially isotropic universe is given by the Bekenstein-Hawking entropy. We point out that if the entropy of the apparent horizon receives a logarithmic correction, the first law of thermodynamics leads to a modified Friedmann equation which corresponds precisely to the time-time component of the semi-classical Einstein field equations sourced by the trace anomaly of ${\\cal{N}}=4$ $U(N)$ super-Yang-Mills theory. This correspondence allows for a thermodynamic description of the dynamics of the Randall-Sundrum braneworld scenario. ",
        "introduction": "It is well known that the entropy and temperature of the event horizon of a Schwarzschild black hole are given by the Bekenstein-Hawking entropy \\cite{bekhawk} \\begin{equation} \\label{bekhawk} S_{\\rm BH}=\\frac{A}{4G} \\end{equation} and the Hawking temperature \\cite{hawk} \\begin{equation} \\label{hawktemp} T_H =\\frac{1}{2\\pi \\tilde{r}_A} , \\end{equation} respectively, where $A= 4\\pi \\tilde{r}_A^2$ is the proper area of the horizon and $G$ is Newton's constant. Similar expressions arise for the cosmological event horizon of de Sitter space \\cite{gibhawk}. More generally, if the entropy of a local Rindler causal horizon is proportional to its area, the Einstein field equations can be derived from the fundamental thermodynamic relation $\\delta Q = TdS $, where the heat flow, $\\delta Q$, is interpreted in terms of the energy flux across the horizon and the temperature is the Unruh temperature associated with a uniformly accelerating observer just within the horizon \\cite{jac}. Similar ideas have been employed in a cosmological context \\cite{frokof,dan,bousso,wang,gongwang,caikim,wwy}. It can be shown that if the entropy of the apparent horizon of a spatially isotropic universe is related to its area by Eq. (\\ref{bekhawk}), the Friedmann equations of general relativity follow from the first law of thermodynamics \\cite{caikim}.  These developments indicate that there exists a direct correspondence between Einstein gravity and the Bekenstein-Hawking entropy. However, Eq. (\\ref{bekhawk}) is a semi-classical relation and becomes modified when quantum corrections are taken into account. One such correction arises when one-loop effects break the conformal invariance of classical massless fields and induce an anomalous trace for the energy-momentum tensor. (For reviews, see, e.g., Refs. \\cite{duff,birrell}). In particular, the entropy of a Schwarzschild black hole that is emitting zero-mass particles receives a logarithmic correction of the form \\cite{furs,bm1,bm2}: \\begin{equation} \\label{quantumentropy} S= \\frac{A}{4G} + \\alpha \\ln \\left( \\frac{A}{4G} \\right)  , \\end{equation} where the dimensionless parameter, $\\alpha$, is determined by the conformal anomaly of the fields.  In a cosmological context, it was recently shown that a logarithmic correction to the entropy of the apparent horizon generates a modified Friedmann equation when the first law of thermodynamics is applied \\cite{cch}. On the other hand, the conformal anomaly is purely geometrical in nature and also modifies the dependence of the Friedmann equation on the energy density of the universe. In view of the above correspondence between gravity and thermodynamics, it is of interest to investigate whether the modifications due to the trace anomaly can be reproduced by a logarithmically-corrected Bekenstein-Hawking entropy. In the present work, we find that the corrections to the spatially flat Friedmann equation derived from the thermodynamic and field-theoretic perspectives are identical when the conformal field theory (CFT) is ${\\cal{N}} =4$, $U(N)$ super-Yang-Mills theory. As in the black hole background, the logarithmic correction to the entropy is proportional to the conformal anomaly of the fields. In this sense, therefore, the conformal anomaly can be interpreted as a quantum correction to the entropy of the apparent horizon.  The structure of the paper is as follows. We begin in Section 2 with the thermodynamic derivation of the Friedmann equation and establish a direct correspondence with anomaly-driven cosmology in Section 3. We employ this correspondence in Section 4 to present a thermodynamic description of the Randall-Sundrum braneworld scenario. We conclude with a summary in Section 5. ",
        "conclusions": "Cosmology driven by the trace anomaly of conformally invariant matter fields represents an important framework for investigating the evolution of the very early universe. The correspondence highlighted in the present work allows us to interpret the effect of the anomaly from a thermodynamic perspective as a logarithmic correction to the Bekenstein-Hawking entropy associated with the apparent horizon of the spatially flat FRW universe. The Friedmann equation derived by applying the first law of thermodynamics to the apparent horizon with such a corrected entropy has the same form as that derived directly from field-theoretic techniques. This is interesting given that a similar logarithmic correction to the Bekenstein-Hawking entropy is generated by one-loop effects near the event horizon of a Schwarzschild black hole \\cite{furs}. In both the cosmological and black hole backgrounds, the correction is determined by the number of zero-mass fields: for the FRW universe, $\\alpha \\propto -b$, whereas $\\alpha \\propto (c-b )$ for the Schwarzschild black hole \\cite{furs}.  We have focused on the trace anomaly of ${\\cal{N}} =4 $ $U(N)$ super-Yang-Mills theory and our conclusions rely on the condition that the anomaly coefficient defined in Eq. (\\ref{defd}) vanishes. A non-zero value would lead to terms involving higher derivatives of the scale factor in the expression for the effective energy density of the anomaly, Eq. (\\ref{tracesolution}). Nonetheless, we expect this correspondence between thermodynamics and gravity to hold for other theories where $d \\ne 0$. This follows since the value of this parameter can always be renormalized to zero by introducing a finite local counterterm (proportional to $R^2$) into the non-local action.  In conclusion, therefore, the present work provides a further example of how thermodynamics emerges from a combination of quantum physics and gravitation."
    },
    "0812/0812.2102_arXiv.txt": {
        "abstract": "Through detailed numerical simulations, we demonstrate that relativistic outflows (Lorentz factor $\\Gamma \\sim 7$) of electron-positron pairs can be produced by radiative acceleration even when the flow starts from a nearly pair equilibrium state at subrelativistic temperatures. Contrary to the expectation that pairs annihilate during an expansion stage for such low temperatures, we find that most pairs can survive for the situations obtained in our previous work. This is because in the outflow-generating region the dynamical timescale is short enough even though the fireball is optically thick to scattering. Several problems that should be solved to apply to actual active galactic nucleus jets are discussed. ",
        "introduction": "\\label{sec:intro}  One of the most challenging problems in astrophysics is the acceleration mechanism of relativistic jets from active galactic nuclei (AGNs) and Galactic black hole candidates. The bulk Lorentz factor of these jets is above 10 and the kinetic power is almost comparable to the Eddington luminosity. Although various jet models have been proposed, we do not yet have a satisfactory solution. Recent remarkable progress in magnetohydrodynamical (MHD) simulations includes Poynting-dominated jet formation by a rapidly rotating black hole with an accretion disk \\citep[e.g.][]{mck06,haw06}. Along the spin axis a centrifugal funnel filled with magnetic field is formed. The magnetic field in the funnel is considered to be amplified by the differential rotation of the accretion disk and the frame-dragging effect of the black hole. However, in the simulation of \\citet{mck06}, the total jet luminosity within the funnel is only 0.2\\% of the rest-mass energy accretion rate. While continuous developments in numerical studies are expected to probe MHD acceleration, other types of acceleration mechanisms are worth considering.  One of the alternative acceleration mechanisms is the thermal expansion of optically thick fireballs. If the matter content of jets is electron-positron pairs, thermal radiation pressure from accretion disks may be sufficient to accelerate the jets. The original model of the fireball \\citep{ree92} relevant to gamma-ray bursts (GRBs) assumes that its initial stage is radiation-dominated plasmas that are in complete thermal equilibrium at high temperatures comparable to the electron rest-mass energy $m_{\\rm e} c^2$. However, the characteristic size of AGNs is too large to achieve complete thermal equilibrium. The novel model proposed to overcome this problem is electron-positron pair outflow from a ``Wien fireball'' \\citep{iwa02,iwa04}, which is a photon-dominated, optically thick plasma, but the densities of photons and pairs are less than those in the complete thermal equilibrium fireballs. Pairs coupled with photons by scattering are thermally accelerated, expending the internal energy of the fireball. As is seen in the original model of the fireball, the Lorentz factor $\\Gamma=1/\\sqrt{1-\\beta^2}$ increases with the radius $R$, and the temperature drops as $T \\propto R^{-1}$. \\citet{iwa02,iwa04} showed that a sufficient amount of electron-positron pairs survive as a relativistic outflow, if the temperature of the fireball at the photosphere is relativistic ($\\gtrsim m_{\\rm e} c^2$).  When the temperature is relativistic, all the cross-sections of pair creation, annihilation and Compton scattering are of the same order. Therefore, a balance between pair creation and annihilation is realized, and the number densities of pairs and photons are of the same order as long as photons and pairs are coupled with each other. On the other hand, the condition to make the pair-annihilation timescale shorter than the dynamical timescale, \\begin{equation} n_+ \\sigma_{\\rm T} \\frac{R}{\\Gamma \\beta}> \\frac{3}{8} \\left[ 1+\\frac{2 \\theta^2}{\\ln{(1.12\\theta+1.3)}}\\right], \\end{equation} is almost the same as the condition of being optically thick to scattering, where $n_+$, $\\sigma_{\\rm T}$, and $\\theta=T/m_{\\rm e} c^2$ are the positron density in the comoving frame, Thomson cross section, and the normalized temperature, respectively. This implies that the number of pairs is almost conserved outside the photosphere, where photons and pairs are decoupled. Therefore, a relativistic temperature at the photosphere is required to obtain a sufficient amount of pairs avoiding annihilation.    The temperature of fireballs is determined by microscopic processes in their formation sites, presumably accretion disks, such as radiative cooling, Compton scattering, and $\\gamma-\\gamma$ pair production \\citep[see e.g.][]{sve84}. Assuming that pairs are confined with protons, several authors \\citep[see e.g.][and references there in]{kus87,bjo92} investigated pair equilibrium plasmas. However, a series of studies of pair equilibrium plasmas implies that the equilibrium temperature is too low for the Wien fireball model. The plasma temperature typically becomes $\\theta \\sim 0.1$ for plasmas of size $\\sim 10^{14}$ cm and luminosity $\\gtrsim 10^{45} {\\rm erg}$ ${\\rm s}^{-1}$, while the Wien fireball model requires a super-Eddington luminosity $\\gtrsim 10^{47}$ erg ${\\rm s^{-1}}$ (under the assumption of spherical symmetry) and a high temperature $\\theta > 0.5$.  The previous studies on pair equilibrium plasmas have not taken into account the dynamical effects of outflowing electron-positron pairs on the plasma temperature. It is not trivial to estimate the motional (or advective) effects of fireballs on radiative transfer, cooling processes, pair density and temperature. In our previous paper \\citep[][hereafter AT07]{asa07}, we simultaneously solved the dynamics of outflowing pairs, microphysics, and radiative transfer to obtain the internal structure of fireballs. Even in this simulation, we obtained only $\\theta \\sim 0.2$ at $R \\sim$ several times $10^{14}$ cm for the luminosity $L = 10^{47}$ erg ${\\rm s^{-1}}$, though a considerable amount of pairs outflow with a mildly relativistic velocity. In the simulations of AT07, the acceleration of outflows is not finished within the simulation scale. We pessimistically supposed that a rapid extinction of pairs inside the photosphere may occur, since the obtained temperature is so low.  However, the detailed situations in AT07 are different from the simplified inner boundary conditions of the outflow in the simulations of \\citet{iwa04}. The pair-photon ratio at the outer boundary of the simulation region is smaller than that of the Wien equilibrium plasma by a factor of $\\sim 2$. The photon spectrum obtained numerically is also different from the simple Wien spectrum. It may be valuable to simulate subsequent evolution of the outflows outside  such a tepid fireball obtained in the simulations in AT07.  In this Letter, we demonstrate that the tepid fireball obtained by the simulations of AT07 can be accelerated to a relativistic velocity conserving the pair number, even though the fireball temperature is not high enough. In \\S \\ref{sec:method}, we describe our simulation method, and show our results in \\S \\ref{sec:results}. \\S \\ref{sec:disc} is devoted to discussion. ",
        "conclusions": "\\label{sec:disc}  Our simulation shows that an energy release of $10^{47}$ erg ${\\rm s^{-1}}$ within a size of $\\sim 10^{14}$ cm produces a pair outflow of $\\Gamma \\sim 7$. Even though the fireball temperature is tepid, rapid annihilation of pairs is avoided. If the entire system we have considered is moving with a Lorentz factor $\\Gamma_{\\rm s} \\geq 1.06$ ($\\beta_{\\rm s} \\geq 0.34$), the Lorentz factor of the pair plasma seems to be larger than 10. Such a nonrelativistic outflow of protons may be easily achieved, as various MHD simulations show. Therefore, similar simulations to ours including the outflow of protons are valuable, while we have assumed that the background protons are static.    The reason why a rapid annihilation of pairs does not occur may be a slight difference in pair density from the prediction of the Wien equilibrium. This difference comes from the finite size of the energy-release region. In the energy-release region the outflowing material keeps on being heated and is evacuated with a mildly relativistic velocity. The rapid expansion makes the pair plasma escape smoothly from the optically thick region, avoiding pair annihilation. For such a marginal optical depth, the Wien equilibrium is no longer a reasonable approximation. The pair densities obtained numerically are suitable for avoiding pair annihilation and keeping the plasma optically thick to scattering, which are opposing requirements in general.  However, resultant densities are too low to convert the radiation energy into the kinetic energy of the pair plasma efficiently. The efficiency in our simulation is only 1/60 of the total luminosity. The results may depend on the profiles of the plasma heating rate. Simulations for other types of profiles of the plasma heating rate are worth studying. We do not simulate for $L>10^{47}$ erg ${\\rm s^{-1}}$, because the higher pair density requires high computational cost, while the temperature tends to decrease with $L$. This is another challenge for us, though there is a possibility of rapid extinction of pairs due to the lower temperature.   Of course, the most idealized method to efficiently produce relativistic outflows is to create fireballs with relativistic temperatures. Relativistic temperatures for optically thick plasmas assure a higher density of pairs, which is preferable for both the energy efficiency and the final Lorentz factor. However, as AT07 showed for steady solutions, even if the effects of mildly relativistic motion are taken into account, the temperature of fireballs is close to that of static pair equilibrium plasmas obtained by past studies. It is valuable to seek a much faster solution ($U > 1$) than AT07. As is usually seen in solutions of the accretion disk, by varying the boundary conditions we may obtain another type of solution with high temperature and high outward advection. However, it is difficult to make $U > 1$ within the energy-release region, because the effect of Compton drag may be crucial in such a compact region. The photon field is produced from the plasma itself so that the photon distribution is not sufficiently beamed in the inner region, where the plasma is produced. We expect plasma acceleration only in the outer region, where the photon distribution is highly beamed.  Another possibility is nonsteady solutions of outflows. We should note that the timescale of radiative transfer is much longer than the dynamical timescale. The lifetime of accretion plasma may not be long enough to achieve the steady-state radiation field obtained here and in AT07. As in the original idea of the Wien fireball model, an instantaneous heating may produce a runaway production of relativistic pairs because of slow annihilation in a hot plasma. Therefore, time-dependent simulations should be examined in a future work."
    },
    "0812/0812.2012_arXiv.txt": {
        "abstract": "This paper reports the discovery of a Very Low Luminosity Object (VeLLO) in the ``starless'' dense core L328, using the {\\it Spitzer Space Telescope} and ground based observations from near-infrared to millimeter wavelengths. The {\\it Spitzer} 8 {\\micron} image indicates that  L328 consists of three subcores of which the smallest one may harbor a source, L328-IRS while two other subcores remain starless. L328-IRS is a Class 0 protostar according to its bolometric temperature (44 K) and the high fraction ($\\sim 72 \\%$) of its luminosity emitted at sub-millimeter wavelengths.  Its inferred ``internal luminosity'' ($0.04 - 0.06 \\lsun$) using a radiative transfer model under the most plausible assumption of its distance as 200 pc is much fainter than for a typical protostar, and even fainter than other VeLLOs studied previously. Note, however, that its inferred luminosity may be uncertain by a factor of $2-3$ if we consider two extreme values of the distance of L328-IRS (125 or 310 pc). Low angular resolution observations of CO do not show any clear evidence of a molecular outflow activity. But broad line widths toward L328, and {\\it Spitzer} and near-infrared images showing nebulosity possibly tracing an outflow cavity, strongly suggest the existence of outflow activity.  Provided that an envelope of at most $\\sim 0.1 \\msun$ is the only mass accretion reservoir for L328-IRS, and the star formation efficiency is close to the canonical value $\\sim 30\\%$, L328-IRS has not yet accreted more than $\\rm 0.05 \\msun$. At the assumed distance  of 200 pc, L328-IRS is destined to be a brown dwarf. ",
        "introduction": "Knowledge of how low mass stars form has been significantly improved thanks to the recent availability of large telescopes and sensitive detectors, especially at millimeter and infrared wavelengths.  Molecular cloud condensations are understood to gradually increase in density via several processes such as ambipolar diffusion or turbulent dissipation with some contribution of magnetic fields.  When the condensations are dense enough to be detected with high density tracers such as the CS and $\\rm NH_3$ transitional lines, but have not yet formed a protostar, they are called ``starless'' cores (e.g., Benson \\& Myers 1989, Lee, Myers, \\& Tafalla 1999).  Some starless cores are believed to be gravitationally bound, and are called pre-stellar cores under the assumption that protostars are likely to form in the future (e.g., Ward-Thompson et al. 2007, Di Francesco et al.  2007).  Once a core become gravitationally unstable and collapses, it forms a protostar (the main accretion stage) with a disk and infalling envelope.  The protostar then evolves to the stage where the envelope has mostly dispersed and a young stellar object (YSO) with an optically thick disk is revealed, and eventually the disk becomes optically thin to finally reveal a normal main-sequence (MS) star (e.g., Andr\\'e et al.  2000, Robitaille et al.  2006).  Even though we have a good consensus on the global scenario of evolution from the cores to MS stars, we still have very little knowledge about which conditions decide how stars form and evolve in the dense cores.  Probably the initial stage of star formation may play the most important role in its further evolution.  This is why we study ``starless'' dense cores.  Dense molecular cores have typically been classified as ``starless'' if they contain no {\\it Infrared Astronomical Satellite} (IRAS) point source whose flux increases in the longer wavelengths (Beichmann et al. 1986; Benson \\& Myers 1989; Lee \\& Myers 1999). Therefore the definition of ``starless'' has been limited to the sensitivity of the IRAS of $\\rm \\sim 0.1(d/140)^2 L_\\odot$ where d is the distance to the source (Myers et al. 1987).  This implies that studying protostars embedded in dense cores has been limited only to sources brighter than 0.1 $\\rm L_\\odot$ at the distance of nearby star-forming regions such as Taurus, resulting in fewer studies of young stars fainter than 0.1 $\\rm L_\\odot$ forming in the cores.  The sensitivity of the {\\it Spitzer Space Telescope} (hereafter {\\it Spitzer}) at infrared wavelengths allows, for the first time, sensitive explorations to detect deeply embedded faint objects in dense cores that were not detected by IRAS or the Infrared Space Observatory (ISO) (Young et al. 2004; Dunham et al.  2006; Bourke et al. 2006).  Such objects may become protostars or brown dwarfs depending on how they obtain matter as time goes on.   The objects fainter than 0.1 $\\rm L_\\odot$ are called Very Low Luminosity Objects (VeLLOs) (Young et al. 2004, Di Francesco et al. 2007; Dunham et al.\\ 2008).  The key feature of these sources is that their luminosity is an order of magnitude fainter than the accretion luminosity ($\\rm \\sim 1.6 L_\\odot$) that the lowest mass protostar of $\\rm \\sim 0.08 M_\\odot$ can produce, under the standard star formation theory (e.g., Shu et al. 1987), with a typical accretion rate ($\\rm \\sim 2\\times 10^{-6} M_\\odot~yr^{-1}$) and protostellar radius of $\\rm \\sim 3R_\\odot$ (Dunham et al. 2006).  These may simply be either very faint protostars that will grow to become normal young stars, or proto-brown dwarfs that eventually become brown dwarfs.  So far there is no clear conclusion about what these objects are and will be in the future because of the limited number of VeLLOs studied.  There are only three VeLLOs studied in detail: L1014-IRS (Young et al. 2004), L1521F-IRS (Bourke et al. 2006), and IRAM 04191+1522 (Dunham et al. 2006).  Their properties are similar in terms of their faintness ($\\rm 0.06 - 0.09~L_\\odot$), but found to be different in their outflow activity and relations with physical properties of their parent cloud cores (Bourke et al. 2006, Dunham et al. 2008).  Detailed studies on more VeLLOs are needed to better understand these objects and their role in star formation.  This paper presents a study of another VeLLO, L328-IRS in the dark cloud Lynds 328, using observations performed with a part of {\\it Spitzer} c2d Legacy project (Evans et al. 2003, 2007) together with several complementary ground-based observations.  L328-IRS has been listed as one of 15 VeLLO candidates that were found in a search of c2d cores and clouds by Dunham et al. (2008).  L328 is a dense molecular cloud located in projection between two nearby large clouds, the Ophiuchus Molecular Cloud (hereafter the Ophiuchus) and the Aquila Rift. In L328, neither an IRAS point source nor any PMS stars had been found, and thus it was classified as ``starless'' (Lee \\& Myers 1999). Therefore this object has been studied  as one of the starless dense cores by many researchers (Visser et al.\\ 2001, 2002; Bacmann et al.\\ 2002, 2003; Crapsi et al.\\ 2005b; L\\\"ohr et al.\\ 2007). The one thing to note is that the line width of $\\rm N_2H^+$ 1-0 is broad ($\\rm \\Delta V=0.438~km~s^{-1}$) compared with other starless cores, indicating the existence of the significant contribution of turbulent gaseous motion (Crapsi et al.\\ 2005b).  Crapsi et al. (2005b) have classified L328 as ``not chemically and dynamically evolved'' because it does not show any evolved feature such as infall asymmetry, significant CO depletion, or deuteration.  The distance to L328 is an important issue for this study. We will explain how we adopt the distance of L328 in $\\S$2.  In $\\S$3 we describe {\\it Spitzer} observations as well as other ground-based observations.  The images of L328, photometric properties of L328-IRS, association of L328-IRS with L328, and outflow activity of L328-IRS are described in $\\S$4. The Spectral Energy Distribution (SED) of L328-IRS is modeled with a radiative transfer code and the fate of L328-IRS is discussed in $\\S$5. The summary can be found in the last section. ",
        "conclusions": "\\subsection{Radiative transfer modeling of the SED of L328-IRS}  The observed luminosity (i.e., bolometric luminosity) of L328-IRS was calculated to be about 0.19 $\\rm L_\\odot$ by using fluxes integrated through the observed wavelength range and its assumed distance of 200 pc. This luminosity ($\\rm L_{bol}$) is the sum of an internal luminosity ($\\rm L_{int}$) due to the central source and a luminosity ($\\rm L_{isrf}$) supplied by the interstellar radiation field (ISRF) from the outside.  The emission from the envelope is partly caused by heating by the internal luminosity source and partly by the ISRF.  Therefore it is necessary to separate the emission of the central source from other contributions to estimate its internal luminosity $\\rm L_{int}$.  This can be done through a detailed study of the radiative process of emission from the central source through the dusty envelope.  To do this we used a one dimensional radiative transfer package DUSTY (Ivezic. 1999) to calculate the SED of L328-IRS with assumed physical properties of the central source and its dusty envelope. Similar studies of other VeLLOs have recently been undertaken (Young et al. 2004; Dunham et al. 2006; Bourke et al. 2006); all the details of the input parameters to this modeling are described in those studies and Young \\& Evans (2005).  Our model consists of a stellar black body, a cooler disk, dusty envelope, and the Interstellar Radiation Field (ISRF).  Starting with an input spectrum from the central star and disk, the model calculates the emission reprocessed by the model envelope with an educated guess of the contribution of the ISRF.  By adjusting the free parameters in the model, the fitting process is iterated to get the $\\chi^2$ best fit of the modeled SED to the observed SED, and thus the best fit internal luminosity $\\rm L_{int}$.  The internal luminosity of the central source in the model is the luminosity from the stellar component plus the luminosity from the disk. For the stellar component, we assume the temperature and luminosity of a blackbody that provide a reasonable fit to the flux densities between 3.6 - 8.0 $\\micron$. The flux densities at these emergent from the central source, however, are absorbed by the dusty envelope surrounding L328-IRS and reprocessed to longer wavelengths. Especially the flux density at 8 $\\micron$ is highly affected by absorption which may be caused mainly by 9.7 $\\micron$ silicate features in the dusty envelope. These need to be taken into account.  The flux densities at 24 and 70 $\\micron$ are not fitted at all by the stellar component, and a dusty disk is included to fit those wavelengths.  We employed a model of a disk with surface density and temperature profiles described as $\\rm \\Sigma (r)= \\Sigma_0 (r/r_0)^{-p}$ and $\\rm T_r=T_0(r/r_0)^{-q}$, respectively (Butner et al. 1994).  We adopt the ``spherical disk'' approach of Butner et al. (1994) in the sense that in the code the emission from the model disk is averaged over all possible viewing angles and then added to the emission from the stellar blackbody to produce an input spectrum for the internal source.  In this modeling we tried to fix all free parameters whose variation has small effect on the results.  We assume a flared disk with a power index p=1.5 in the surface density distribution, but with a slightly less steep index (q=0.4) in the temperature distribution than in Butner et al. (1994) which helps to better fit the 24 $\\micron$ flux density than a model with q=0.5.  The inner radius of the disk was set to be the radius at which dust reaches its sublimation temperature (assumed to be 1500 K). It varied from 0.02 to 0.04 AU to be consistent with the internal luminosity of the central source.  On the other hand the outer radii were simply fixed to be 50 AU.  Usually the model was insensitive with the disk size.  We run the models with outer disk radii of 10 AU and 100 AU and found similar results for the inferred internal luminosity to the model with a disk size of 50 AU.  The code also includes the contribution from intrinsic disk emission which can arise from accretion of the material on the disk.  The input emission emerging at mainly short wavelengths is then reprocessed in the dusty envelope to produce emission at longer wavelengths. We adopted an approximate isothermal equilibrium sphere for the envelope density distribution (Tafalla et al 2002), $n(r)=n_0/[1+(r/r_0)^{\\alpha}]$, where $n$ is the density, $r$ is the distance from the center, and $n_0$, $r_0$, and $\\alpha$ are scaling parameters of the density distribution (note that this is {\\it not} the same $\\alpha$ from \\S~4.2.1).  Alternatively, a power-law density distribution for the envelope [$n(r) \\propto r^{-\\alpha}$] was also tested and it was found that a shallow power index such as $\\alpha=1.1$ resulted in a very similar fit to the flux densities at sub-millimeters to that of the approximate isothermal sphere .  The envelope is assumed to be heated by the emergent emission from the central source inside and the ISRF outside.  The inner radius of the envelope is set to be 50 AU and its outer radius is 10,000 AU which roughly matches the size of the envelope seen in the DSS image. The outer radius of the envelope is large enough to include the other sub-cores in L328. However, this radius was set to better match the sub-millimeter and its long-ward flux densities of the envelope with the model SED. Our simple model of the envelope considers the flux contribution from the smm2 core only, but not that from other sub-cores.  The dust opacities are those of ``OH5 dust'' (grains with thin ice mantles) by Ossenkopf \\& Henning (1994) which is known to be appropriate for cold dense cores (Evans et al.  2001; Shirley et al. 2005).  We adopted the Black-Draine ISRF (Black 1994, Draine 1978; see Evans et al. 2001) and assumed it to be attenuated by dust grains having optical constants of Draine \\& Lee (1984) with $\\rm A_V$ = 3.0.  This extinction was chosen to simulate the situation that L328 is somewhat embedded in diffuse dust between the Ophiuchus and the Aquila Rift clouds.  We tried to fix most of free parameters regarding the envelope so that the estimated mass of the envelope is consistent with the mass ($\\rm M_{env}$) of the envelope determined from the 350 $\\micron$ continuum.  The mass was calculated from $\\rm M_{env} = {{S_{\\nu} D^2} \\over {B_{\\nu}(T) \\kappa_{\\nu}}}$ where $S_{\\nu}$ is the flux density at 350 $\\micron$, $\\rm B_{\\nu} (T)$ is the Planck function at temperature (T), and $\\kappa_{\\nu}$ is dust opacity.  The estimated mass is about 0.05 - 0.09 $\\rm M_\\odot$ in the range of temperatures of 15 - 20 K with a measured flux density of 3.2 Jy ($20\\arcsec$ beam) with $\\rm \\kappa_{\\nu}=0.101~cm^{-2}~g^{-1}$ (OH5 dust opacities in Ossenkopf \\& Henning 1994).  The mass was also estimated using the 850 and 1200 $\\micron$ data and found to be consistent within a factor of 2 with the mass derived from the 350 $\\micron$ flux.  We chose the parameters of $\\rm n_0=6\\times 10^5~cm^{-3}$, $r_0=1500$ AU, and $\\alpha=2.0$ as reasonable parameters resulting in an envelope mass within a 2000 AU radius of 0.07 $\\rm M_\\odot$, consistent with the observed mass of the envelope of 20\\arcsec\\ aperture.  With this combination of parameters, we tried to obtain a best fit with stellar temperatures $\\rm T_{star}$=1500, 1800, 2000, 2500, 3000, and 3500 K by adjusting $\\rm L_{int}$.  Among these temperatures a central source of 0.04 - 0.06 $\\rm L_\\odot$ is found to produce the best fits to the observed SED.  Figure 8 shows a result of the best fit of the observed SED for $\\rm T_{star}$=2000 K and $\\rm L_{int}=0.05~L_\\odot$ which has a minimum value of $\\chi^2$. This result implies that most of the observed luminosity arises from emission at sub-millimeter wavelength regions heated by the ISRF, as modeled by Evans et al. (2001) for starless cores.  Dunham et al (2008) have estimated the internal luminosity of L328-IRS by using an empirical correlation between fluxes at 70 $\\micron$ and internal luminosities in the VeLLOs that were obtained from the calculation of the SED of embedded low luminosity protostars with the two dimensional radiative transfer code RADMC (Dullemond \\& Dominik 2004; Crapsi et al.\\ 2008). Their value of $\\rm L_{int}$ estimated from the 70 $\\micron$ flux is 0.07 $\\rm L_\\odot$ which is higher than our $\\rm L_{int}$.  However, their assumed distance of 270 pc is larger than our adopted value. If they had used 200 pc, then their luminosity would be $\\sim 0.04 \\rm L_\\odot$, consistent with our value.  Therefore our study suggests that, if the distance of L328 is correct, the luminosity of the central source is well below 0.1 $\\rm L_\\odot$ and qualifies L328-IRS as a VeLLO.  What if the assumed distance of L328 is different from 200 pc ? We calculated the radiative transfer model for two possible extreme distances 125 pc (the distance of the Ophiuchus) and 310 pc (the distance of the far edge of the Aquila), to examine other possible values of the internal luminosity of L328-IRS due to the uncertainty of the distance, with proper scaling of the mass and the size of its envelope at those distances. The internal luminosities for $\\rm T_{star}$=1500 to 3500 K are found to be $0.12 - 0.18 \\rm L_\\odot$ for the distance of 310 pc and $0.02 - 0.03 \\rm L_\\odot$  for the distance of 125 pc. This indicates that in the extreme scenario that L328 is possibly located at the far edge of Aquila Rift cloud at 310 pc (Strai\\v{z}ys et al. 2003), the internal luminosity of L328-IRS can be slightly higher than the borderline for a VeLLO. Therefore a good knowledge of the distance is a key issue in deciding whether faint sources like L328-IRS are of very low luminosity.  There are several pieces of evidence to suggest that our assumed distance of 200 pc for L328 is more reasonable than 310 pc or even 270 pc (the mean distance to the Aquila Rift).  First, the velocity of L328 is intermediate between that of Ophiuchus (125 pc) and the Aquila Rift, suggesting an intermediate distance.  Second, L328 is seen in absorption against the background emission at the longer {\\it Spitzer} wavelengths (5.8 -- 70 \\micron).  This background emission is likely due to PAH and warm dust directly associated with the Aquila Rift clouds, and strongly suggests that L328 is in the foreground or on the near side of the Rift\\footnote[1] {Our recent measurement of the distance of L328 using near-IR photometry of 2MASS data around L328 gives a distance of $220\\pm 40$ which is very consistent with our adopted distance (Maheswar \\& Lee 2008 in prep.).}.  To summarize, with our best estimate of the distance to L328 of 200 pc and the modeling of the SED of L328-IRS with a radiative transfer code, its inferred internal luminosity is 0.04 - 0.06 $\\rm L_\\odot$ and qualifies it as a VeLLO. At the distance extremes which seem to be unlikely, $\\rm L_{int}$ can be lower to 0.02 $\\rm L_\\odot$ and higher up to $\\rm 0.18 L_\\odot$ In any case, even if $\\rm L_{int} = 0.18 L_\\odot$ for L328-IRS, its luminosity would still be close to the VeLLO cutoff, which is not strongly based on physical properties, and would still be significantly lower than typical protostars.  Accurate determination of the distance of L328-IRS and higher dimensional radiative transfer modeling together with a more complete SED are required to better determine the luminosity of L328-IRS.  \\subsection{Evolutionary fate of L328-IRS }  Here we discuss the properties of L328-IRS and compare those with the properties of other VeLLOs.  First of all, we note that the internal luminosity of L328-IRS is fainter than that of other VeLLOs studied with a similar level of detail, given that all the distances are fully correct (L1014-IRS, 0.09 $\\rm L_\\odot$; L1521F-IRS, 0.06 $\\rm L_\\odot$; IRAM 04191+1522, 0.08 $\\rm L_\\odot$).  We saw several possible hints for the outflow of L328-IRS. Nebulosity of a few arc-seconds in length and a probable cavity structure a few tens of arc-seconds long are seen in L328.  The projected length of the NE cavity from L328-IRS is about 6,000 AU ($30\\arcsec$).  If we assume that the nebulosity and tenuous cavity structure represent something related to outflow activity, then the total size of the outflow from L328-IRS would be at least 12,000 AU.  This is larger than the outflow from L1014-IRS (a few hundred AU; Bourke et al. 2005), but comparable to the projected size ($\\sim 10,000$ AU) of nebulosity seen in the {\\it Spitzer} image of L1521F-IRS (Bourke et al. 2005), and much smaller than the pc scale CO outflow in IRAM 04191+1522.  We did not find any clear evidence that the outflow of L328-IRS may occur in an episodic fashion.  We looked for any possibility of an episodic event of the outflow by searching an area of $30\\arcmin \\times 30\\arcmin$ around L328 in DSS images and $5\\arcmin \\times 5\\arcmin$ in {\\it Spitzer} images and found no clear evidence of Herbig-Haro objects or shocked region indicating possible episodic events within the outflow activity.  Toward L328 no infall signatures are observed in molecular lines, no significant CO depletion exists (there is CO depletion, but it is not classified as significant compared to many other cores; Bacmann et al.\\ 2002; Crapsi et al.\\ 2005), and no enhanced deuteration is observed, implying that it is dynamically and chemically young (Crapsi et al. 2005)\\footnote[1] {The molecular line observations referred in this discussion have been made at ($\\alpha, \\delta$)$_{J2000}$=($\\rm 18^h16^m59\\fs50, -18^{\\circ} 02\\arcmin30\\farcs5$) which is about $40\\arcsec$ offset from the position of L328-IRS and closer to smm1 and smm3. Therefore the discussion may rely on  the assumption that similarities of the line characteristics between the observational position of the molecular lines and the position of L328-IRS exist.}. This is somewhat similar to what is found in L1014, but very different from the cores hosting the VeLLOs,  IRAM 04191+1522 and L1521F-IRS, that show many of these properties (Andr\\'e et al. 1999; Belloche et al. 2002; Crapsi et al. 2004).  The mass (at most of order 0.1 $\\rm M_\\odot$) of the parent core smm2 associated with L328-IRS as a reservoir providing gaseous material to it is much smaller than other cores with VeLLOs (L1014-IRS, 1.7 $\\rm M_\\odot$, Young et al. 2004; L1521F-IRS, 4.8 $\\rm M_\\odot$, Crapsi et al.\\ 2005; IRAM 04191+1522, 2.5 $\\rm M_\\odot$, Dunham et al. 2006).  Considering these points, L328-IRS is most likely a very low luminosity object, possibly having outflow activity, embedded in an envelope of small mass that is not dynamically or chemically evolved.  Is L328-IRS just a small version of a protostar which will eventually grow to a protostar by accreting material from its envelope, or is L328-IRS already on a different pathway to become a brown dwarf?  We find that its observed envelope mass is too small as a provider of material for L328-IRS to become a star.  If we assume a 30 percent star formation efficiency in L328 (Lada et al. 2007), only a further $\\rm \\sim0.03  M_\\odot$ of the observed envelope mass will accrete onto L328-IRS.  Therefore, unless other sub-cores in L328 are  involving in the accretion process, the star formation efficiency is unusually higher than the canonical value, and it has already accreted more than $\\sim 0.05 \\rm M_\\odot$, L328-IRS seems to have no hope to become a star in the future, and is probably destined to be a brown dwarf. However, if any of these assumptions is not true, alternative scenario in which  L328-IRS will grow as a faint protostar and evolve to a normal star is not ruled out either."
    },
    "0812/0812.2154_arXiv.txt": {
        "abstract": "The correlation function of the distribution of matter in the universe shows, at large scales, baryon acoustic oscillations, which were imprinted prior to recombination. This feature was first detected in the correlation function of the luminous red galaxies (LRG) of the Sloan Digital Sky Survey (SDSS). The final release (DR7) of the SDSS has been recently made available, and the useful volume is about two times bigger than in the old sample. We present here, for the first time, the redshift space correlation function of this sample at large scales together with that for  one shallower, but denser volume-limited subsample drawn from the 2dF redshift survey. We test the reliability of the detection of the acoustic peak at about $100 \\, \\mpch$ and the behaviour of the correlation function at larger scales by means of careful estimation of errors. We confirm the presence of the peak in the latest data although broader than in previous detections. ",
        "introduction": "\\label{sec:intro} About 380,000 years after the Big Bang, when the temperature falls low enough for recombination to occur, the matter in the Universe becomes neutral. At this epoch, the sound speed drops off abruptly and acoustic oscillations in the fluid become frozen. Their signature can be detected in both the Cosmic Microwave Background (CMB) radiation and the large-scale distribution of galaxies.  The characteristic length scale of the acoustic oscillations can be used as a ``standard ruler\" to probe the expansion history of the universe.  The imprints in the matter distribution of this feature, called baryon acoustic oscillations (BAO), should be detectable in both the correlation function and the matter power spectrum. Moreover, this feature should manifest itself as a single peak in the correlation function at about $100 \\, \\mpch$\\footnote{$h$ is the Hubble constant in units of 100 km s$^{-1}$ Mpc$^{-1}$.}. The first unambiguous detection of this feature in the correlation function (redshift space) of the Sloan Digital Sky Survey Luminous Red Galaxies (SDSS-LRG) was reported by \\citet{2005ApJ...633..560E}. Detection of these oscillations in the power spectrum of the galaxy distribution was first reported by \\citet{2005MNRAS.362..505C} for the 2-degree Field Galaxy Redshift Survey (2dFGRS). Cosmological parameters can be determined from the position of the baryon acoustic peak \\citep[see, e.g.,][]{2007ApJ...657...51P}.  In this letter we study the correlation functions of the largest available redshift surveys (2dFGRS and SDSS).  We focus on the two-point correlation function, and study the reliability of the detection of the BAO feature at $100 \\, \\mpch$ in the samples drawn from those surveys. In our calculations, we adopted a cosmological model with $\\Omega_M = 0.27$, and $\\Omega_{\\Lambda} = 0.73$. ",
        "conclusions": "\\label{sec:dis} The results show, first, that the baryon acoustic peak is a stable feature in the large-distance galaxy correlation function. We have presented the analysis of the LRG sample drawn from the last data release (DR7) of the SDSS. This survey is much larger and more compact than the data used before, so edge-correction effects are smaller now and results are therefore more robust. While the peak is most prominent in the distribution of luminous red galaxies (because of the large spatial extent of the sample), it can be seen, if we know what to look for, in smaller galaxy samples. We have reported the first detection of the baryon acoustic peak in a volume-limited sample of very luminous galaxies drawn from the 2dFGRS.  Second, the baryon peak is much wider than found before and than expected; the distortions responsible for that demand careful analysis.  Third, the minimum in the large-distance correlation functions of some samples demands explanation -- is it really the signature of voids?  And, fourth, a new internal method to estimate the errors of the correlation function has been introduced. The method is based on a generalization of the application of the bootstrap resampling techniques to smooth functions of sample means.   \\smallskip \\textbf{Acknowledgements}  We thank Darren Croton for the 2dfGRS samples and the mask data, Rien van de Weygaert and Dietrich Stoyan for their advise regarding the Voronoi-Poisson process, Robert Lupton for useful discussions and Jun Pan and Jan Hamann for insightful comments on a preliminary draft of this letter.  This work has been supported by the University of Valencia through a visiting professorship for Enn Saar and by the Spanish MEC projects AYA2006-14056 and CSD2007-00060, including FEDER contributions. PAM acknowledges support from the Spanish MEC through a FPU grant. ES and ET acknowledge support by the Estonian Science Foundation, grants No. 6104, 6106, 7146 and by the Estonian Ministry for Education and Science, grant SF0060067s08.  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 \\texttt{http://www.sdss.org/}."
    },
    "0812/0812.3152_arXiv.txt": {
        "abstract": "{ Narrow-band surveys to detect Ly$\\alpha$ emitters are powerful tools for identifying high, and very high, redshift galaxies. Although samples are increasing at redshifts $z = 3 - 6$, the nature of these galaxies is still poorly known. The number of galaxies detected at redshifts below $z \\sim 3$ are also small.} {% We study the properties of $z = 2.25$ Ly$\\alpha$ emitters and compare them with those of $z > 3$ Ly$\\alpha$ emitters.} {% We present narrow-band imaging made with the MPG/ESO 2.2m telescope and the WFI (Wide Field Imager) detector. Using this data, we have searched for emission-line objects. We find 170 candidate typical Ly$\\alpha$ emitters and 17 candidates that we regard as high UV-transmission Ly$\\alpha$ emitters. We have derived the magnitudes of these objects in 8 photometric bands from $u^*$ to $K_s$, and studied whether they have X-ray and/or radio counterparts.   } {% We demonstrate that there has been significant evolution in the properties of Ly$\\alpha$ emitters between redshift $z \\sim 3$ and $z = 2.25$. The spread in spectral energy distributions (SEDs) at the lower redshift is larger and we detect a significant AGN contribution in the sample. The distribution of the equivalent widths is narrower than at $z \\sim 3$, with only a few candidates with rest-frame equivalent width above the predicted limit of 240~{\\AA}. The star formation rates derived from the Ly$\\alpha$ emission compared to those derived from the UV emission are lower by on average a factor of $\\sim 1.8$, indicative of a significant absorption by dust. } {% Ly$\\alpha$ emitters at redshift $z = 2.25$ may be more evolved than Ly$\\alpha$ emitters at higher redshift. The red SEDs imply more massive, older and/or dustier galaxies at lower redshift than observed at higher redshifts. The decrease in equivalent widths and star formation rates indicate more quiescent galaxies, with in general less star formation than in higher redshift galaxies. At $z = 2.25$, AGN appear to be more abundant and also to contribute more to the Ly$\\alpha$ emitting population. } ",
        "introduction": "Over the past decade, a large number of so-called Lyman-$\\alpha$ (Ly$\\alpha$) emitters have been discovered at high redshift. Several techniques have been employed in the search for these galaxies, but the by far most common method is that of narrow-band imaging, where a narrow-band filter is tuned to Ly$\\alpha$ within a particular narrow redshift range. Objects with large equivalent widths (EWs) are thus selected by comparing the colours in the narrow-band image and complementary broad-band images. Spectroscopically confirmed Ly$\\alpha$ emitters now include several hundreds of sources at redshifts $z \\sim 3$ (e.g., M{\\o}ller \\& Warren 1993; Cowie \\& Hu 1998; Steidel et al. 2000; Fynbo et al. 2001, 2003; Matsuda et al. 2005; Venemans et al. 2007; Nilsson et al. 2007; Ouchi et al. 2008), $z \\sim 4.5$ (Finkelstein et al. 2007), $z \\sim 5.7$ (Malhotra et al. 2005; Shimasaku et al. 2006; Tapken et al. 2006) and $z \\sim 6.5$ (Taniguchi et al. 2005; Kashikawa et al. 2006). However, in the low redshift range, between $z \\approx 1.6$ corresponding to the atmospheric cut-off in the UV and $z \\sim 3$, little progress has been made. In practice, the lower redshift limit is in fact higher than $z \\approx 1.6$ because of the drop-off in CCD sensitivity and a more typical lower redshift limit is $z \\sim 2.0$. Eight narrow-band surveys have so far been published below $z \\sim 3$; Fynbo et al. (1999), Pentericci et al. (2000), Stiavelli et al. (2001), Fynbo et al. (2002, 2003a, 2003b), Francis et al. (2004) and Venemans et al. (2007). Furthermore, Van Breukelen et al. (2005) published a sample of Ly$\\alpha$ emitters (LAEs) at $z \\sim 2.5$ using integral-field spectroscopy. The difficulty in observing the Ly$\\alpha$ line between redshifts $2 < z < 3$ lies in the low throughput of optical systems, the low efficiency of CCDs in this wavelength range, and higher extinction in the atmosphere. Even so, the advantage of a smaller luminosity distance is rewarding in that a higher \\emph{flux} limit equals the same \\emph{luminosity} limit as surveys at higher redshift. It also facilitates follow-up observations into the nature of these objects.  At higher redshifts, Ly$\\alpha$ emitters have been observed to be increasingly bluer, younger and smaller with increasing redshift. At $z \\sim 3$, Gawiser et al. (2006) inferred stellar masses of a few~$\\times 10^8$~M$_{\\odot}$, almost no dust extinction, and ages of the order of $100$~Myr. In a follow-up paper, Gawiser et al. (2007) studied a stacked sample of 52 Ly$\\alpha$ emitters without Spitzer detections and confirmed their results of young ages and that no dust appears to be present in these systems, although they inferred slightly higher masses. The galaxies in their sample with Spitzer detections are presented in Lai et al. (2008). For these galaxies, older ages and higher masses are reported and the authors propose that $z \\sim 3$ Ly$\\alpha$ emitters may have a wide range of properties. In Nilsson et al. (2007), a stacked sample of Ly$\\alpha$ emitters at $z = 3.15$ were studied. Here, small masses and low dust contents were inferred, but the ages were unconstrained. At even higher redshift, Pirzkal et al. (2007) showed that a sample of $4 < z < 5.7$ Ly$\\alpha$ emitters had very young ages of a few Myr and small stellar masses, in the range $10^6 - 10^8$~M$_{\\odot}$. These results agreed well with those of Finkelstein et al. (2007), who studied almost 100 Ly$\\alpha$ emitters at $z = 4.5$. Finkelstein et al. (2008) reported, studying a different sample, finding very dusty, massive galaxies with $A_{1200}$ as high as 4.5 magnitudes and masses as high as several $\\times 10^{10}$~M$_{\\odot}$, at $z = 4.5$. They argued that they observed a bimodality in the properties of Ly$\\alpha$ emitters; young and blue galaxies versus old and dusty. Thus, the properties of $z \\geq 3$ Ly$\\alpha$ emitters are poorly constrained, but these galaxies tend to be considered young, blue, small and dust-free. With the data presented here, we aim to extend the study of the properties of Ly$\\alpha$ emitters to lower redshifts, where a wider range of the SED can be studied in the optical/infrared and the luminosity distance is smaller, allowing a more detailed analysis.  This paper is organised as follows. Section~\\ref{sec:obs} presents the observations leading to our sample, as well as the data reduction. Section~\\ref{sec:sel} includes the object detection and candidate selection and a discussion about possible interlopers in the sample. In Sect.~\\ref{sec:basic} we present all the basic characteristics of the sample, including photometry, AGN contribution, surface density, and sizes, and the equivalent width distribution of the candidates. We summarise the results in Sect.~\\ref{sec:disc}.  \\vskip 5mm Throughout this paper, we assume a cosmology with $H_0=72$ km s$^{-1}$ Mpc$^{-1}$ (Freedman et al. 2001), $\\Omega _{\\rm m}=0.3$ and $\\Omega _\\Lambda=0.7$. Magnitudes are given in the AB system. ",
        "conclusions": "\\label{sec:disc} In this paper, the observations and selection leading to a sample of 170 robust LAE candidates at $z = 2.25$ have been presented. This is the first large-scale survey of LAEs at redshifts around two, and several indications of evolution are seen in the properties of LAEs between redshift $z \\sim 3$ and $z \\sim 2$.  This is however unsurprising, since the separation in these redshifts corresponds to longer than 1 Gyr. The main arguments for an evolution in the properties of this type of galaxy are: \\begin{itemize} \\item[-] \\emph{The large spread in SEDs} \\end{itemize} The most conspicuous evolution detected was that of the emergence of a large number of ``red'' LAEs. The 12 individual detections in the relatively shallow $K_s$ images and the red mean SED of the total stacked sample, presented in Fig.~\\ref{fig:redsed}, indicated that a significant subsample, of up to two-thirds of the total sample, of LAEs are massive, and potentially dusty, galaxies. We also showed that the rest-frame UV colours of these galaxies cannot be explained by a simple stellar population, and that even though the majority of the candidates are consistent with being young galaxies, a significant fraction of the sample displays more complicated colours. \\begin{itemize} \\item[-] \\emph{Potentially larger fraction of AGN} \\end{itemize} The result above is also related to the detection of at least nine AGN in the sample. Again, this is not a surprising result since redshift $z \\sim 2$ corresponds to the peak of the AGN number density distribution (see e.g. Wolf et al. 2003 and references therein). The number density of AGN in this sample is in principle consistent with surveys at higher redshift, but is likely to be higher by a factor of two. \\begin{itemize} \\item[-] \\emph{Change in SFR ratios} \\end{itemize} The ratio of UV to Ly$\\alpha$ derived SFRs has a higher median value in this survey than results at redshift $z \\sim 3$. The spread is large, but the trend is clear. This result, and the smaller EWs discovered in general for this sample (see also the next point) at high redshift may indicate that LAEs in this survey are expected to be more affected by dust than LAEs at $z \\sim 3$. \\begin{itemize} \\item[-] \\emph{Narrower EW distribution} \\end{itemize} Another intriguing and related result is that of the narrower equivalent width distribution. Firstly, as detailed in section~\\ref{ewsec}, the distribution is difficult to explain without invoking complicated arguments about the properties of the galaxies. A possible explanation would be that the star formation histories of Ly$\\alpha$ emitters are complex. Secondly, the difference in the distributions between redshifts three and two is further evidence of a higher dust quantity in the lower redshift galaxies.  In conclusion, by comparing observations of Ly$\\alpha$ emitters at redshift $z = 2.25$ with galaxies selected in the same manner at higher redshifts, several evolutionary signatures become evident in the properties of these galaxies. At lower redshifts, there appear to be fewer objects, with redder colours and higher dust contents, smaller equivalent widths, and a higher fraction of objects containing AGN. Future SED fitting of these galaxies will reveal more information into the properties such as dust, mass and age (Nilsson et al., in prep.)."
    },
    "0812/0812.4128_arXiv.txt": {
        "abstract": "Our position inside the Galaxy requires all-sky surveys to reveal its large-scale properties. The zero-level calibration of all-sky surveys differs from standard 'relative' measurements, where a source is measured in respect to its surroundings. All-sky surveys aim to include emission structures of all angular scales exceeding their angular resolution including isotropic emission components. Synchrotron radiation is the dominating emission process in the Galaxy up to frequencies of a few GHz, where numerous ground based surveys of the total intensity up to 1.4~GHz exist. Its polarization properties were just recently mapped for the entire sky at 1.4~GHz. All-sky total intensity and linear polarization maps from WMAP for frequencies of 23~GHz and higher became available and complement existing sky maps. Galactic plane surveys have higher angular resolution using large single-dish or synthesis telescopes. Polarized diffuse emission shows structures with no relation to total intensity emission resulting from Faraday rotation effects in the interstellar medium. The interpretation of these polarization structures critically depends on a correct setting of the absolute zero-level in Stokes U and Q. ",
        "introduction": "All-sky radio continuum surveys provide basic information on our local environment and the large-scale properties of the Galaxy. They are required to model the Galactic emission components in 3-D (Sun {\\it et al.} (2008)) and guide more sensitive higher-angular resolution observations of the Galactic plane or other regions or objects of interest. All-sky surveys are quite time consuming projects, which require special observing methods to accurately measure large-scale sky emission. They need similar telescopes in the northern and southern hemisphere and have to adapt calibration data from additional instruments to provide the absolute level of sky emission. Thus all-sky surveys are rare, in particular in linear polarization, where the signals are much weaker than in total intensities.  In the following we describe the basic methods to calibrate and adjust total intensity surveys as well as the more complex requirements being applied for polarization surveys. We do not discuss the instrumental corrections to be taken into account during the data reduction process, which can be found in the original publications. ",
        "conclusions": "\\label{sec:concl}  We have discussed calibration methods being applied to all-sky and Galactic plane continuum surveys in total intensity and linear polarization. The interpretation of polarized emission emerging from Faraday rotation effects in the interstellar medium requires special attention, in paparticularn absolute zero-level setting is essential."
    },
    "0812/0812.3708_arXiv.txt": {
        "abstract": "It has been well established that the $f$-mode of relativistic ordinary-fluid neutron stars displays a universal scaling behavior. Here we study whether the ``ordinary'' $f_{\\rm o}$- and ``superfluid'' $f_{\\rm s}$-modes of superfluid neutron stars also show similar universal behavior. We first consider a simple case where the neutron superfluid and normal fluid are decoupled, and with each fluid modeled by a polytropic equation of state. We find that the $f_{\\rm o}$-mode obeys the same scaling laws as established for the $f$-mode of orindary-fluid stars. However, the oscillation frequency of the $f_{\\rm s}$-mode obeys a different scaling law, which can be derived analytically from a homogenous two-fluid stellar model in Newtonian gravity. Next the coupling effect between the two fluids is studied via a parameterized model of entrainment. We find that the coupling in general breaks the universal behavior seen in the case of decoupled fluids. Based on a relativistic variational principle, an approximated expression is derived for the first-order shift of the $f_{\\rm s}$-mode squared frequency due to the entrainment. ",
        "introduction": "\\label{sec:intro}   Ground-based gravitational wave detectors are either already operating or will soon be operating. For example, the LIGO detectors have conducted five science runs since 2002. While the current detectors are still not sensitive enough to detect gravitational waves directly, interesting upper limits on the gravitational-wave strains from several potential astrophysical sources (e.g., isolated pulsars and stochastic background) have been placed \\citep{abbott-2007-1,abbott-2007-2}. It is very likely that the detection of gravitational waves will come within a decade or so after the upgrade of the detectors.   Pulsating neutron stars are one of the interesting potential sources of gravitational waves. They may also rotate rapidly enough for various kinds of rotational induced mode instabilities to develop and enhance the emitted gravitational waves \\citep{andersson2003gwi}. The pulsation modes of neutron stars are damped due to the emission of gravitational waves, and hence are called quasinormal modes \\citep[see, e.g.,][for reviews]{kokkotas1999qnm,ferrari08}, instead of normal modes as in Newtonian theory. The study of quasinormal modes of neutron stars has a long history dating back to the pioneering works of Thorne and his collaborators \\citep{thorne1967tnr,price1969tnr,thorne1969npg1,thorne1969npg2,campolattaro1970npg}. It is by now well established that the quasinormal mode spectrum of a neutron star is tremendously rich and the gravitational waves emitted from a pulsating star carry important information about the internal structure of the star \\citep[e.g.,][]{andersson96, andersson1998tgw,Benhar:1998au,kokkotas01, benhar04,tsui2005prl}. The gravitational-wave signals from pulsating neutron stars, should they be detected by ongoing or future detectors, can thus in principle be used to constrain the supranuclear equation of state (EOS), which is still poorly understood.   While the quasinormal modes of a pulsating neutron star in general depend sensitively on the stellar model and EOS, some empirical universal behaviors have also been observed with different EOSs. \\citet{andersson1998tgw} noted that the frequencies and damping times of the leading gravitational-wave $w$-modes and fluid $f$-mode of nonrotating neutron stars can be approximated by some empirical relations which depend only on the mass and radius of the star for most EOSs \\citep[see also][]{Benhar:1998au}. In particular, \\citet{andersson1998tgw} obtained the following formulas respectively for the real and imaginary parts of $f$-mode frequency: \\begin{equation} {\\rm Re}(\\omega M) = \\alpha_1 M + \\alpha_2 C^{3/2} , \\label{eq:nils_real_f} \\end{equation} \\begin{equation} {\\rm Im}(\\omega M) =  C^4 \\left( \\beta_1 C + \\beta_2 \\right) , \\label{eq:nils_img_f} \\end{equation} where $\\alpha_1$, $\\alpha_2$, $\\beta_1$, and $\\beta_2$ are model-independent constants determined by curve fitting. $M$ is the mass of the star and the dimensionless parameter $C \\equiv M/R$ (in units where $G=c=1$) is the compactness of the star. The real part of $\\omega$ is the mode oscillation frequency, while the imaginary part is inversely proportional to the damping time due to the emission of gravitational waves. Recently, the physical mechanism behind such universal behaviors has been investigated in detail by \\citet{tsui2005uqn}. In their notations, the (complex) frequency $\\omega $ of the $w$-mode or $f$-mode can be approximated by \\begin{equation} \\omega M = a C^2 + b C + c , \\label{eq:scaling_law} \\end{equation} where the complex parameters $(a,b,c)$ are also model-independent constants determined by curve fitting. Exploring such universality for the $w$-modes, \\citet{tsui2005prl} have developed an inversion scheme to determine the mass, radius, density distribution, and the EOS of a neutron star model from the frequencies of the first few $w$-modes of the star.   Mature neutron stars are known to be very cold on the nuclear temperature scale ($10^{10}$ K). It is believed that a newborn neutron star can cool down rather quickly (on a timescale of a few weeks to months) to the transition temperature ($\\sim 10^9$ K) for nucleon superfluidity and superconductivity to occur \\citep[see, e.g.,][]{lombardo-1999,lombardo-2001-578,andersson-2005-763}. For sufficiently high density, nuclear matter can transform to deconfined quark matter, which may also be in the so-called color-superconducting phase \\citep[see, e.g.,][]{alford-2003-67,alford-2004-30}. To date, there is still no direct evidence for the existence of nucleon superfluidity in neutron stars. However, the well-established pulsar glitch phenomenon is best explained by the pinning and unpinning of large numbers of superfluid vortices to the solid crust \\citep{Radhakrishnan-1969,lyne-1993}. So, how would nucleon superfluidity affect the dynamics and hence the oscillation mode spectrum of neutron stars? For the simplest model, neutron stars consist of neutrons, protons, and electrons. When neutrons become superfluid, they dynamically decouple from protons and electrons. The protons can also be in a superconducting state, but they are coupled to the ordinary electron fluid via electromagnetic intereaction on a very short timescale. In effect, the neutron star interior can be approximated by a two-fluid model: one fluid is the neutron superfluid; the other fluid is a conglomerate of all other charged constituents.  To a first approximation, the two interpenetrating fluids could be considered as independent. However, due to the strong interaction between neutrons and protons, the two fluids in general can couple via the so-called entrainment effect. Entrainment is a multifluid effect and it arises when the flow of one fluid induces a momentum in the other fluid. In this work, we will study the parameterized entrainment model used previously by \\citet{andersson-2002-66}.   Building on the formalism developed by Carter and his collaborators \\citep[e.g.,][]{carter1989rfd,comer1993hfm,comer1994hfr,carter-1995-454, langlois1998drr}, \\citet{comer-1999-60} and \\citet{andersson-2002-66} have calculated the quasinormal modes of such nonrotating two-fluid stellar model \\citep[see, e.g.,][for the corresponding study in Newtonian theory]{lindblom1994osn,lee1995non,prix2002aon}. These works show the existence of a new family of modes for a two-fluid star, the so-called superfluid modes. These modes have the distinguishing characteristic that the two fluids are essentially counter-moving, as a manifestation of an extra fluid degree of freedom which is missing in the study of pulsating ordinary, single fluid neutron stars \\citep{comer-1999-60}. Another characteristic of the superfluid modes is that they depend sensitively on the entrainment effects between the two fluids \\citep{andersson-2002-66}. In a two-fluid star, there is essentially a doubling of the fluid modes. For example, the single $f$-mode which exists in an ordinary fluid star is splitted into an ``ordinary'' $f$-mode (denoted by $f_{\\rm o}$), where the two fluids tend to move together, and a superfluid counterpart (denoted by $f_{\\rm s}$) where the fluids are counter-moving \\citep{comer-1999-60}. On the other hand, the $w$-modes do not show such kind of mode doubling effect \\citep{comer-1999-60}. The two fluids of the $w$-modes always move in ``lock-step'' due to the fact that $w$-modes are largely spacetime oscillations.  As mentioned above, the frequencies of the $f$- (or $w$-modes) of nonrotating ordinary-fluid neutron stars can be approximated by universal scaling laws (\\ref{eq:nils_real_f})-(\\ref{eq:scaling_law}) for different EOSs. It is thus natural to ask whether this universality also holds for two-fluid neutron star models. In this paper, we investigate whether the $f_{\\rm o}$ and $f_{\\rm s}$ modes of two-fluid stars establish any kind of universality. In order to provide the reader a better understanding of the main results obtained in this paper, we shall give a brief summary in the following.   In this work we first study the simplest case of two decoupled fluids, each with a polytropic EOS. We vary the polytropic indices of the two fluids to mimic the effects of different EOSs. We find that the $f_{\\rm o}$-mode (i.e., the ``ordinary'' $f$-mode) can still be approximated very well by the same scaling laws (\\ref{eq:nils_real_f})-(\\ref{eq:scaling_law}) as established for ordinary-fluid neutron stars. For the superfluid $f_{\\rm s}$-mode, we find that the real part of the mode frequency satisfies a different universal scaling law (Eq.~(\\ref{eq:newton_fs})), while the imaginary part in general does not have any universal behavior. We are also able to derive the universal scaling curve for the real part of the $f_{\\rm s}$-mode frequency analytically based on a homogeneous two-fluid model in Newtonian gravity.  We then study the effect of coupling between the two fluids via a parameterized model of entrainment. We analyze both numerically and analytically how the entrainment affects the $f_{\\rm s}$-mode frequency. Based on a relativistic variational principle, we have derived a general integral formula to calculate the first-order shift in the mode (squared) frequency (Eq.~(\\ref{eq:1st_order_rel})). Furthermore, for the particular class of background EOS and entrainment models we considered in the study, we are able to approximate the integral formula by an algebraic expression (\\ref{eq:1st_order_rel_final}) which depends explicitly only on the model parameters.  The main results obtained in this work (Eqs.~(\\ref{eq:newton_fs}) and (\\ref{eq:w2_over_w2f2})) can be used by other researchers to obtain a good approximation to the $f_{\\rm s}$-mode frequency (without the need of constructing their own numerical code) for the class of background EOS and entrainment models that have been used extensively to study superfluid neutron stars.   The plan of this paper is as follows. Section~\\ref{sec:formal} summarizes briefly the general relativistic two-fluid formalism. In section~\\ref{sec:eos} we describe the EOS models we use in the study. Section~\\ref{sec:results} presents our numerical results. Section~\\ref{sec:fs_mode_derive} presents our analytical analysis for a homogeneous two-fluid star model. The effects of entrainment on the superfluid mode frequency is studied analytically in Section~\\ref{sec:pert}. Finally, we summarize our results in section~\\ref{sec:conclude}. We use units where $G=c=1$ unless otherwise noted. ",
        "conclusions": "\\label{sec:conclude}  In this work we have studied whether the $f_{\\rm o}$- and $f_{\\rm s}$-modes of superfluid neutron stars exhibit any kind of universal scaling laws as have been seen in the $f$-mode of ordinary-fluid neutron star models. We first focus on the simplified case of two decoupled fluids, each with a polytropic EOS. We vary the polytropic indices to mimic the effects of different EOSs. Our numerical results show that the $f_{\\rm o}$-mode, where the two fluids move in ``lock-step'', obeys the same universal scaling laws as the $f$-mode of ordinary fluid stars.  On the other hand, we find that the oscillation frequency of the $f_{\\rm s}$-mode, which corresponds essentially to counter motion between the two fluids, obeys a different scaling law (Eq.~(\\ref{eq:newton_fs})). We have also derived the scaling law analytically based on a homogeneous two-fluid stellar model in Newtonian gravity. However, the damping time of the $f_{\\rm s}$-mode in general does not exhibit any kind of universality. While we have only used a generalized polytropic EOS in our study, we believe that our conclusion holds in general for superfluid neutron star models in which the two fluids exist throughout the whole star and are decoupled in the sense that the master function $\\Lambda$ (ie, the EOS) can be decomposed into two contributions corresponding to each fluid, i.e., $\\Lambda(n^2, p^2) = \\Lambda_{\\n} (n^2) + \\Lambda_{\\p} (p^2)$.  The inclusion of a coupling term in the master function will in general break the universal behavior. To illustrate the effect of coupling, we have studied the entrainment between the two fluids using a parameterized entrainment model. We show numerically how the $f_{\\rm s}$-mode frequency increases with the strength of the entrainment. Furthermore, based on a relativistic variational principle, we have carried out a perturbative analysis and have derived an expression for the first-order shift of the frequency due to the entrainment. If the superfluid $f_{\\rm s}$-modes could be detected by future gravitational wave detectors, then the derivation of the observed mode frequencies from the universal scaling curve for the decoupled fluids could then be a useful probe to the coupling effects between the neutron superfluid and normal fluids inside neutron stars. In summary, our main results (Eqs.~(\\ref{eq:newton_fs}) and (\\ref{eq:w2_over_w2f2})) can be used to obtain a good approximation to the oscillation frequency of the $f_{\\rm s}$-modes for the generalized polytropic EOS and entrainment models that have been used extensively to study superfluid neutron stars \\citep{comer-1999-60,andersson2001srg,andersson-2002-66,prix2002aon, yoshida03}.   While we have not yet detected the gravitational waves emitted by neutron stars, it is worthy to mention that the first gravitational-wave search sensitive to the $f$-modes has recently been carried out by the LIGO detectors and interesting upper limits on the wave strain (within the predicted range of some theoretical models) have also been placed \\citep{abbott-2008sgr}. In view of the fact that the advanced LIGO detectors will have more than a factor of 10 improvement on the sensitivity of the wave strain, gravitational-wave astroseismology may soon become a reality."
    },
    "0812/0812.1227_arXiv.txt": {
        "abstract": "{The abundance patterns of metal-poor stars provide us a wealth of chemical information about various stages of the chemical evolution of the Galaxy.  In particular, these stars allow us to study the formation and evolution of the elements and the involved nucleosynthesis processes. This knowledge is invaluable for our understanding of the cosmic chemical evolution and the onset of star- and galaxy formation. Metal-poor stars are the local equivalent of the high-redshift Universe, and thus offer crucial observational constraints on a variety of issues regarding the early Universe.  This review presents an introduction to metal-poor stars and their role as ``cosmic lab'' for the study of neutron-capture nucleosynthesis processes, particularly that of the r-process. The metal-poor star HE~1523$-$0901 serves as an example for this group of objects. It displays in its spectrum the strongest overabundance of neutron-capture elements associated with the r-process. Heavy neutron-capture elements such as Eu, Os, and Ir were measured, as well as the radioactive elements Th and U. Abundance of Th and U, in conjunction with those of stable elements make possible nucleo-chronometry, i.e., the determination of stellar ages. HE~1523$-$0901 appears to be $\\sim13$\\,Gyr old. Age uncertainties range from 2 to 5\\,Gyr for individual chronometers, and are largly due to theoretical uncertainties in the initial production ratio of the employed chronometers. The decay product of the radioactive elements, lead, can be used to constrain r-process calculations. Only few such stars are currently known with detected U. These objects, however, are crucial for the study of this nucleosynthesis process. Once more objects are discovered, and assuming an old age for them (infered from their low metallicity), stars with measured Th \\textit{and} U abundances can become stellar age calibrators. This way, ages of stars in which only Th is measured (many more stars are available with a Th detection only), can be derived \\textit{indepedently} of model calculations.}    \\FullConference{10th Symposium on Nuclei in the Cosmos\\\\ July 27 - August 1 2008\\\\ Mackinac Island, Michigan, USA}  \\begin{document} ",
        "introduction": "The first stars that formed from the pristine gas left after the Big Bang were very massive, of the order of 100\\,M$_{\\odot}$ [1]. After a very short life time these co-called Population\\,III stars exploded as supernovae, which then provided the first metals to the interstellar medium. All subsequent generations of stars formed from chemically enriched material. Metal-poor stars are early Population\\,II objects and belong to the stellar generations that formed from the non-zero metallicity gas left behind by the first stars. Due to their low masses ($\\sim0.8$\\,M$_{\\odot}$) they have extremely long lifetimes that exceed the current age of the Universe of $\\sim14$\\,Gyr years. Hence, these stellar ``fossils'' of the early Universe are still observable today.  In their atmospheres these old objects preserve information about the chemical composition of their birth cloud. They thus provide archaeological evidence of the earliest times of the Universe. In particular, the chemical abundance patterns provide detailed information about the formation and evolution of the elements and the involved nucleosynthesis processes. This knowledge is invaluable for our understanding of the cosmic chemical evolution and the onset of star- and galaxy formation. Galactic metal-poor stars are the local equivalent of the high-redshift Universe. They allow us to derive observational constraints on the nature of the first stars and supernovae, and on various theoretical works on the early Universe in general [2].  Focusing on long-lived low-mass metal-poor main-sequence and giant stars, we are observing stellar chemical abundances that reflect the composition of the interstellar medium during their star formation processes [3]. Stars spend $\\sim90\\%$ of their lifetime on the main sequence before they evolve to become red giants. Main-sequence stars only have a shallow convection zone that preserves the stars' birth composition over billions of years. Stars on the red giant branch have deeper convection zones that lead to a successive mixing of the surface layers with nuclear burning products from the stellar interior. In the lesser evolved giants the surface composition has not yet been significantly altered by any such mixing processes and are useful as tracers of the chemical evolution. Further evolved stars (e.g., asymptotic giant branch stars) usually have surface compositions that have been altered through repeated events that dredge up events of nucleosynthetic burning products.  The main indicator used to determine stellar metallicity is the iron abundance, [Fe/H], which is defined as \\mbox{[A/B]}$ = \\log_{10}(N_{\\rm A}/N_{\\rm B})_\\star - \\log_{10}(N_{\\rm A}/N_{\\rm B})_\\odot$ for the number N of atoms of elements A and B, and $\\odot$ refers to the Sun. With few exceptions, [Fe/H] traces the overall metallicity of the objects fairly well. This review focuses on metal-poor stars that have around 1/1,000 of the solar Fe abundances, and are thus able to probe the earliest epochs of nucleosynthesis processes.  A detailed summary of the history and the different ``classes'' of metal-poor stars and their role in the early Universe can be found in ref. [4]. ",
        "conclusions": ""
    },
    "0812/0812.2297_arXiv.txt": {
        "abstract": "{\\small We study the gravitational collapse problem of rotating shells in three-dimensional Einstein gravity with and without a cosmological constant. Taking the exterior and interior metrics to be those of stationary metrics with asymptotically constant curvature, we solve the equations of motion for the shells from the Darmois-Israel junction conditions in the {\\it co-rotating} frame. We study various collapse scenarios with {\\it arbitrary} angular momentum for a variety of geometric configurations, including anti-de Sitter, de Sitter, and flat spaces. We find that the collapsing shells can form a BTZ black hole, a three-dimensional Kerr-dS spacetime, and an horizonless geometry of point masses under certain initial conditions. For pressureless dust shells, the curvature singularity is {\\it not} formed due to the angular momentum barrier near the origin. However when the shell pressure is nonvanishing, we find that for all  types of shells with polytropic-type equations of state (including the perfect fluid and the generalized Chaplygin gas), collapse to a naked singularity is {\\it possible} under generic initial conditions. We conclude that in three dimensions angular momentum does not in general guard against violation of cosmic censorship.} ",
        "introduction": "\\label{sec:intro}  In  three-dimensional Einstein gravity, there are no dynamical degrees of freedom, i.e., no gravitons that mediate the interactions between massive objects in the vacuum \\cite{Dese:84}, with or without the cosmological constant, unless some higher derivative terms, like the (gravitational) Chern-Simons term, are introduced \\cite{Dese:82}.   Although spacetime outside of matter is locally of constant curvature, some non-trivial spacetime solutions to the field equations exist that have  either black hole \\cite{btz} or cosmological event horizons \\cite{Dese:84b,Park:98} when  the cosmological constant is nonzero\\footnote{In general cosmological horizons in dS space do not necessarily imply the existence of black hole horizons in the associated AdS space. However, in three dimensions they are closely related, i.e, there is one-to-one correspondence between the ``event'' horizons --  cosmological horizons in dS and  outer black hole horizons in AdS -- for overlapping ranges of ADM mass and angular momentum. See Ref. \\cite{Park:98} for the details of the mapping.}. Their metrics have the general form \\begin{equation} \\label{eq:metbtz} (ds)^2 = -N^2(R) dT^2 + \\frac{dR^2}{N^2(R)} + R^2(d\\varphi+N^{\\varphi}(R)dT)^2 \\end{equation} with the lapse (squared) and shift functions \\begin{equation} N^2(R) = \\frac{(\\alpha_{o} R^2 +\\gamma_{o} )(\\alpha_{i} R^2 + \\gamma_{i} )}{\\ell^2R^2},~ {N^{\\varphi}(R)=-\\frac{\\mbox{sign}(\\alpha_{o} \\alpha_{i}) \\sqrt{\\gamma_{o} \\gamma_{i}}}{\\ell R^2}} \\end{equation} respectively. Here, $\\alpha_{o/i}$ and $\\gamma_{o/i}$ are constant parameters whose signs (and values) depend on whether we are considering the black hole solution in AdS space or the cosmological solution in dS space.  The ADM mass and  angular momentum of this class of spacetimes, with a cosmological constant $\\Lambda = \\pm 1/\\ell^2$, are given by \\begin{equation} \\label{BTZ} {M_{BTZ}}=\\frac{R_{o}^2+R_{i}^2}{ {8  G} \\ell^2},~~{J_{BTZ}} = \\frac{2R_{o}R_{i}}{ {8 G} \\ell} \\end{equation} for the (BTZ) black hole solution (where $\\alpha_o=\\alpha_i=+1,~ \\gamma_o=-R_o^2,~ \\gamma_i=-R_i^2)$ \\cite{btz} and \\begin{equation} \\label{KdS} {M_{KdS_3}}=\\frac{R_{o}^2-R_{(i)}^2}{ {8  G} \\ell^2},~~{J_{KdS_3}} = \\frac{2R_{o}R_{(i)}}{{8  G} \\ell} \\end{equation} for the (KdS$_3$) cosmological solution (where $\\alpha_o=+1,~ \\alpha_i=-1,~ \\gamma_o=+R_o^2,~ \\gamma_i=+R_{(i)}^2)$ \\cite{Dese:84b,Park:98}, respectively\\footnote{Here,  mass and angular momentum agree with the definitions in Ref. \\cite{Bala:01} but differ in sign from Ref. \\cite{Stro:01}.}, where the three-dimensional Newton constant $G$ is assumed positive\\footnote{ In three-dimensional gravity, there is no {\\it a priori} reason to fix the sign of Newton's constant. For some recent discussion  on the sign choice, see Ref. \\cite{Dese:82}.}.   For the black hole case of Eq. (\\ref{BTZ}), we have ${M_{BTZ}}\\ge {J_{BTZ}}/\\ell \\ge 0 $ in order that there is no naked conical singularity and $R_{o/i}$ denotes the {\\it outer/inner} horizon. On the other hand, for the cosmological solution, there is no constraint on $M_{KdS_3}$ and $J_{KdS_3}$ in order that the horizon exists, unless $J_{KdS_3}=0$ is considered \\cite{Park:98}: Even $M_{KdS_3}<0$ is allowed also. In this case, $R_{o}$ denotes the (cosmological) event horizon \\cite{Dese:84b} but $R_{(i)}$ does not signify an inner horizon; the real parameter $R_{(i)}$ is introduced just for convenience \\cite{Park:98}.  Over the years black hole and cosmological solutions in three dimensions have fascinated theorists because of the potential insights they afford into quantum gravity. Amongst the most intriguing applications \\cite{mr,Pele:95,Ilha:99,Gutt:05} have been black hole formation and the associated issues of gravitational collapse and the cosmic censorship conjecture \\cite{cch}. While there are numerous examples of initial conditions that form a naked singularity in the context of general relativity, none of them are generic as required by the terms of cosmic censorship conjecture. However it has been proposed that three dimensions might be an exception \\cite{Hube:04,mo}. A recent study of collapsing {\\it shells} in three dimensions with {\\it no} angular momentum showed that a naked singularity and a Cauchy horizon can form as the final result of shell collapse for a {\\it broad} variety of initial data \\cite{Pele:95,mo}. This feature is quite generic since the results are independent of the types of  collapsing shell matter such as pressureless dust, polytropic matter, and a generalized Chaplygin gas (GCG).  Furthermore, similar behaviour also appeared in previous studies of collapsing dust in three dimensions \\cite{mr} and the formation of topological black holes in four dimensions \\cite{ms}. (See refs. \\cite{Ilha:99,Ilha:97} for higher-dimensional extensions.) The formation of a naked singularity is of great importance in association with the AdS/CFT correspondence \\cite{Hube:04,Hert:04,MyHub} since one might ask whether or not the quantum correlation function at the AdS boundary is properly defined in the presence of a bulk singularity. This in turn raises additional issues associated with the collapse to a naked singularity, particularly the proper inclusion of quantum (gravity) effects in the bulk \\cite{Vaz:2007hn}.  Thus far the possible scenarios for the emergence of a naked singularity have been explicitly analyzed for collapse to final states (black holes or otherwise) with no angular momentum \\cite{mr,Pele:95,Ilha:99,Gutt:05,mo,ms,Ilha:97}. Indeed, relatively little is explicitly known about the gravitational collapse of matter with nonzero angular momentum in any dimension. It is therefore natural to take into account an extension of gravitational collapse to exterior black holes with rotation (or other possible configurations) and to investigate how rotational effects alter previous results  \\cite{mr,mo} concerning   cosmic censorship violation. While this problem is technically formidable in higher dimensions, a study of  gravitational collapse of shells in three dimensions is considerably more tractable. So we shall consider this problem in our paper\\footnote{ After completing this work, a paper \\cite{Vaz:0805} investigating a rotating dust {\\it cloud} appeared. Their  ``no singularity'' result is consistent with our result for dust shells in Sec. 3. However, our results for shells with pressure in Sec. 4 imply that the naked singularity would {\\it not} be avoided even for rotating clouds with pressure \\cite{Mann:08}.}.  To this end, we introduce a co-rotating coordinate system on the shell, which simplifies the matching procedure. The angular momentum produces a potential barrier around the origin, preventing the shell from contracting to zero size. We analyze the possible collapse scenarios by investigating the effective potential  for all types of equation of state. For a pressureless dust shell, we find that its angular momentum prevents the creation of a singularity at the origin, unlike the non-rotating case \\cite{mo}. Previous work on rotating dust shell collapse was carried out in the Hamiltonian formalism, an alternative to the Darmois-Israel formalism \\cite{Cris:04}.  For a more restricted class of scenarios, they found that singularities were also avoided.  Our work goes beyond this insofar as we consider shells with pressure. In this case, curvature singularities of finite spatial extent are formed before they meet the barrier, and the effective potential and surface stress-energy tensor diverge. Therefore under rather generic conditions it is possible for a naked singularity to form, violating cosmic censorship as in the non-rotating collapse scenario \\cite{mo}.  The outline of our paper is as follows. In Sec. \\ref{sec:coll}, we present a set-up for the two-dimensional hypersurfaces with arbitrary rotations and a spherical symmetry. Considering a co-rotating frame where the co-moving observer on the shell sees only the radial motion yields a simplified equation of motion for the shell. In Sec. \\ref{sec:dust}, we consider the evolution of a pressureless dust shell and compare it to the results for non-rotating shell collapse \\cite{mo}. In Sec. \\ref{sec:pressure}, we consider  shells with pressure whose equations of state are the polytrope type, which includes the perfect fluid and the GCG. Finally, we shall summarize and discuss our results in Sec. \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  Our investigation of the gravitational collapse of rotating shells in three dimensions has uncovered a number of interesting features. We have studied whether or not angular momentum can significantly change the collapse scenario and its resulting cosmic censorship violations in the non-rotating cases in the literature.  For asymptotically AdS boundary conditions, we find that the rotating shell collapses to either a black hole or a minimum value and then expands out to infinity. For shells composed of pressureless dust these are the only scenarios; the centrifugal barrier forbids a naked singularity from forming.  However for shells with pressure we have another scenario in which a naked singular ring can form, violating cosmic censorship. When the exterior spacetime is taken to be a  geometry with a point mass, one might expect that the collapsing shell could form a curvature singularity within finite time, since there is no event horizon as with the non-rotating system in Ref. \\cite{mo}. However we have shown that a naked singularity never forms in the collapse of a rotating dust shell due to a centrifugal barrier in the effective potential experienced by the shell. Collapse scenarios for shells with pressure show that a naked ring singularity of finite size can be formed, where the effective potential and the surface stress-energy tensor diverge. For asymptotically dS boundary conditions a collapsing shell with pressure can form a naked singularity, also. If the interior spacetime is (K)dS, then a cosmological horizon can form from an expanding shell. Which of these scenarios occurs depends on the choice of parameters and initial conditions.  For asymptotically flat boundary conditions the qualitative behaviours are similar, {\\it i.e.}, a rotating collapsing dust shell does not form a naked singularity, but a shell with pressure can form a naked singularity.  For a polytropic shell this will be a naked singular ring of finite size. The qualitative behaviour is more or less intermediate between the asymptotically AdS and dS cases as implied by the choice of $\\alpha^{\\pm}_{o/i}$ and illustrated in Figs. \\ref{fig:numpics2} and \\ref{fig:numpics8}, though there are some anomalous regions that do not reveal this simple trend.  The most intriguing lesson of this paper is that the centrifugal barrier in the effective potential governing the time evolution of a rotating dust shell can prevent formation of a naked singularity. However if the shell has pressure, a ring singularity may form  for a typical class of equations of state, i.e., $\\omega >1$, where the pressure dominates the angular momentum, with quite general initial data. So if the exterior spacetime is assumed to be a geometry without a (covered) black hole event horizon, then the singularity may be naked and the violation of cosmic censorship is possible.  We have found that a collapsing shell  can form either a black hole/cosmological horizon or a naked singularity or bounce to infinity, depending on the initial data. This suggests a set of phase transitions \\cite{Pele:95}  along with accompanying critical phenomena, similar to that discovered for scalar matter \\cite{crit}. It would be of great interest to study these phenomena explicitly as they could reveal universal features of the critical exponents in three dimensions.  Since we have confined our study to three dimensions (where there is no gravitational radiation), some caution is warranted in applying our results to higher dimensions. We have found that the angular momentum does not in general prevent violation of cosmic censorship in three dimensions, where local gravitational interactions vanish outside of matter.  An interesting extension of our work would be to include higher derivative terms, whose effect is to produce such interactions. Another interesting extension is the inclusion of quantum effects, since they might be expected to prevent singularity formation, thereby sidestepping the cosmic censorship issue.  What impact these modifications have on our results remains to be investigated.  \\vspace{5mm} {\\bf Acknowledgment}\\\\ {We would like to thank Sang Pyo Kim and Shin Nakamura for exciting discussions and the headquarter of APCTP for warm hospitality during the APCTP-TPI Joint Focus Program and Workshop. J. J. Oh would like to thank Wontae Kim, Seungjoon Hyun, Hongbin Kim, Jaehoon Jeong, Hyeong Chan Kim, Gungwon Kang, Inyong Cho, Constantinos Papageorgakis, Andrew Strominger, and Chi-Ok Hwang for useful and helpful discussions. R. B. Mann was supported by the Natural Sciences and Engineering Research Council of Canada. J. J. Oh was supported by the Korea Research Council of Fundamental Science \\& Technology (KRCF). M.-I. Park was supported by the Korea Research Foundation Grant funded by Korea Government(MOEHRD) (KRF-2007-359-C00011). }   \\nc{\\PR}[3]{Phys. Rev. {\\bf #1}, #2 (#3)} \\nc{\\NPB}[3]{Nucl. Phys. {\\bf B#1}, #2 (#3)} \\nc{\\PLB}[3]{Phys. Lett. {\\bf B#1}, #2 (#3)} \\nc{\\PRD}[3]{Phys. Rev. {\\bf D#1}, #2 (#3)} \\nc{\\PRL}[3]{Phys. Rev. Lett. {\\bf #1}, #2 (#3)} \\nc{\\PREP}[3]{Phys. Rep. {\\bf #1}, #2 (#3)} \\nc{\\EPJ}[3]{Eur. Phys. J. {\\bf #1}, #2 (#3)} \\nc{\\PTP}[3]{Prog. Theor. Phys. {\\bf #1}, #2 (#3)} \\nc{\\CMP}[3]{Comm. Math. Phys. {\\bf #1}, #2 (#3)} \\nc{\\MPLA}[3]{Mod. Phys. Lett. {\\bf A #1}, #2 (#3)} \\nc{\\CQG}[3]{Class. Quant. Grav. {\\bf #1}, #2 (#3)} \\nc{\\NCB}[3]{Nuovo Cimento {\\bf B#1}, #2 (#3)} \\nc{\\ANNP}[3]{Ann. Phys. (N.Y.) {\\bf #1}, #2 (#3)} \\nc{\\GRG}[3]{Gen. Rel. Grav. {\\bf #1}, #2 (#3)} \\nc{\\MNRAS}[3]{Mon. Not. Roy. Astron. Soc. {\\bf #1}, #2 (#3)} \\nc{\\JHEP}[3]{JHEP {\\bf #1}, #2 (#3)} \\nc{\\JCAP}[3]{JCAP {\\bf #1},, #2 {(#3)}} \\nc{\\ATMP}[3]{Adv. Theor. Math. Phys. {\\bf #1}, #2 (#3)} \\nc{\\AJP}[3]{Am. J. Phys. {\\bf #1}, #2 (#3)} \\nc{\\ibid}[3]{{\\it ibid.} {\\bf #1}, #2 (#3)} \\nc{\\ZP}[3]{Z. Physik {\\bf #1}, #2 (#3)} \\nc{\\PRSL}[3]{Proc. Roy. Soc. Lond. {\\bf A#1}, #2 (#3)} \\nc{\\LMP}[3]{Lett. Math. Phys. {\\bf #1}, #2 (#3)} \\nc{\\AM}[3]{Ann. Math. {\\bf #1}, #2 (#3)} \\nc{\\hepth}[1]{{\\tt [arxiv:hep-th/{#1},]}} \\nc{\\grqc}[1]{{\\tt [arxiv:gr-qc/{#1},]}} \\nc{\\astro}[1]{{\\tt [arxiv:astro-ph/{#1},]}} \\nc{\\hepph}[1]{{\\tt [arxiv:hep-ph/{#1},]}} \\nc{\\phys}[1]{{\\tt [arxiv:physics/{#1},]}}"
    },
    "0812/0812.2068_arXiv.txt": {
        "abstract": "We study and elucidate the mechanism of spiral density wave excitation in a differentially rotating flow with turbulence which could result from the magneto-rotational instability. We formulate a set of wave equations with sources that are only non-zero in the presence of turbulent fluctuations.  We solve these in a shearing box domain, subject to the boundary conditions of periodicity in shearing coordinates, using a WKBJ method. It is found that, for a particular azimuthal wave length, the wave excitation occurs through a sequence of regularly spaced swings during which the wave changes from leading to trailing form. This is a generic process that is expected to occur in shearing discs with turbulence. Trailing waves of equal amplitude propagating in opposite directions are produced, both of which produce an outward angular momentum flux that we give expressions for as functions of the disc parameters and azimuthal wave length.  By solving the wave amplitude equations numerically we justify the WKBJ approach for a Keplerian rotation law for all parameter regimes of interest. In order to quantify the wave excitation completely the important wave source terms need to be specified. Assuming conditions of weak nonlinearity, these can be identified and are associated with a quantity related to the potential vorticity, being the only survivors in the linear regime. Under the additional assumption that the source has a flat power spectrum at long azimuthal wave lengths, the optimal azimuthal wave length produced is found to be determined solely by the WKBJ response and is estimated to be $2\\pi H$, with $H$ being the nominal disc scale height. In a following paper by Heinemann \\& Papaloizou, we perform direct three dimensional simulations and compare results manifesting the wave excitation process and its source with the assumptions made and the theory developed here in detail, finding excellent agreement. ",
        "introduction": "\\label{sec:introduction} Accretion discs are ubiquitous in astrophysics, occurring in close binary systems, active galactic nuclei and around protostars \\citep[see e.g.][for reviews]{1995ARA&A..33..505P,1996ARA&A..34..703L}. Ever since their importance was first realized it has been clear that some form of turbulence is necessary to provide the anomalous angular momentum transport implied by observed luminosities and inferred accretion rates. This has usually been parametrised using the \\citet{1973A&A....24..337S} $\\alpha$-parametrisation.  The most likely source of turbulence is through the magneto-rotational instability (MRI) \\citep[see][]{1991ApJ...376..214B,1998RvMP...70....1B}. Both local and global simulations that display sustained MRI turbulence have been performed with and without net flux. In all cases prolific spiral density (SD) wave excitation has been noted (e.g.\\ \\citealt{2005AIPC..784..475G} in the local case and \\citealt{1998ApJ...501L.189A} in the global case). These waves may be crucial for explaining various phenomena in accretion disk systems. In the context of protoplanetary disks for instance, it has recently been suggested that stochastic gravitational forces derived from density variations due to SD waves may play an important role in driving the migration of low mass protoplanets \\citep{2004MNRAS.350..849N,2005A&A...443.1067N}. This possibly remains a viable mechanism even in so-called dead zones where the ionization fraction is too low for MHD turbulence to occur, see the recent simulations by \\citealt*{2007ApJ...670..805O}. In general, SD waves may lead to significantly enhanced angular momentum transport in magnetically inactive regions of accretion disks. It is therefore important to gain an understanding of the processes leading to the excitation of spiral density waves, how generic the phenomenon is, and how the wave amplitudes scale with physical parameters.  In order to do this, we assume weakly nonlinear conditions, under which the important source terms for exciting the SD waves are expected to be proportional to what we call the pseudo potential vorticity (PPV), which is equal to the potential vorticity to linear order but differs from it in the nonlinear regime. The  potentially important role that PPV plays for the excitation of SD waves in rotating shear flows has been recognised in earlier work \\citep[see e.g.][]{1997PhRvL..79.3178C}. These authors found by numerically integrating the linearised equations of motion of compressible, plane Couette flow that there exists a linear mode coupling between specified vortical perturbations and (free) SD waves that leads to efficient excitation of the latter (see also \\citealt{2005A&A...437....9B} for a similar study in the context of accretion disks). However, the possibility of the generation of vortical perturbations through the action of turbulence has not yet been fully assessed.  It is the purpose of this paper to develop a mathematically rigorous theory of the excitation of SD waves within the WKBJ framework. Our approach is inspired by the work of \\citet{2004JAtS...61..211V} who developed an analogous theory for small amplitude inertia-gravity waves in a local, quasi-geostrophic model of the Earth's atmosphere. Apart from illuminating the excitation process of SD waves in rotating shear flows by putting it on a firm mathematical basis, the theory also enables us to calculate the amplitude of the excited waves and the associated angular momentum transport explicitly.  In a following paper \\cite[][paper II]{2008arXiv0812.2471H} we perform numerical simulations which to study SD wave excitation in MRI driven turbulence \\citep[see also][]{2005AIPC..784..475G,2006ApJ...653..513S} and make detailed comparisons to the WKBJ theory presented here. In this scenario, magnetically dominated turbulent stresses cause vortical perturbations which then in turn lead to the excitation of the observed SD waves. The role that turbulent stresses play is thus indirect. In this aspect the excitation process differs from that considered in Lighthill's theory of aerodynamic noise generation \\citep{1952RSPSA.211..564L}.  The plan of this paper is as follows: In section \\ref{sec:shearing-box} we describe the shearing box model, giving the basic equations and defining the background shear flow in which the hydromagnetic turbulence responsible for the SD wave excitation, is generated. In section \\ref{sec:wave-equations} we derive equations describing the excitation of SD waves. These take the form of linear wave equations with both linear and nonlinear source terms that are determined by the turbulence. We focus on waves that are nearly independent of the vertical coordinate which have been found to dominate in simulations carried out in paper II and which can be dealt with using a vertical averaging procedure. We formulate the law of conservation of angular momentum for linear, non-dissipative SD waves and derive a convenient expression for the average radial flux which we later use to estimate the angular momentum flux arising due to the excited waves.  In section \\ref{sec:wkbj-theory} we go on to develop the WKBJ theory of wave excitation. A Fourier analysis is carried out enabling each azimuthal wave number $k_y$ to be considered separately. The WKBJ theory applies to a shearing box for which the boundary conditions are the imposition of periodicity in shearing coordinates and involves a sequence of excitations uniformly spaced and localized in time, during which the wave swings from leading to trailing. The wave amplitude and wave action produced in a swing are calculated from a WKBJ formalism involving the evaluation of integrals along anti-Stokes lines.  We compare the results derived from asymptotic theory to results obtained by numerically integrating the ordinary differential equations describing the evolution of the appropriate Fourier amplitude. We find excellent agreement between these approaches. This agreement persists under all conditions of interest. This is in spite of the fact that asymptotic theory formally requires a parameter depending on the azimuthal wave number to be small. Finally we discuss our results in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} In this paper we have developed a theory of SD wave excitation in a rotating shear flow with turbulence which may result from the magneto-rotational instability but under the assumption that the magnetic field is too weak to affect the form of the waves significantly. We considered the commonly adopted shearing box model for which the flow is subject to the boundary condition of periodicity in shearing coordinates \\citep{1965MNRAS.130..125G}.  The main feature resulting from the shear is that wave excitation occurs through a sequence of regularly spaced swings as the wave changes from leading to trailing form. For a fixed azimuthal wave number $k_y$, and a Keplerian rotation profile, the swings are separated by a time interval \\mbox{$\\delta{}t_\\mathrm{s}=T_\\mathrm{orb}/(qk_yL_x)$}, where the orbital period $T_\\mathrm{orb}=2\\pi/\\Omega$. For the optimal azimuthal wave number \\mbox{$k_y=\\kappa/c$} (see discussion below) and \\mbox{$L_x=H$} as an estimated radial correlation length of the turbulence, which should also be the minimum box size required to capture its essential properties, it follows that \\mbox{$\\delta{}t_\\mathrm{s}\\approx{}T_\\mathrm{orb}$}.  The wave equations governing the excitation during a particular swing were found to depend on time alone and under the assumption that the important source terms causing the wave excitation are associated with the pseudo potential vorticity, they could be solved to find the asymptotic wave form and net positive wave action produced. The form of the wave equations necessitated a WKBJ analysis in the complex plane.  In this respect the formalism differs from shearing box analyses that adopt rigid or free boundary conditions, or which assume strictly harmonic forcing with radial boundaries extended to infinity, rather than periodicity in shearing coordinates. In the former cases one can separate out a harmonic time dependence and solve a problem in space for the wave amplitude \\citep[e.g.][]{1987MNRAS.228....1N}.  The analysis of the wave excitation process driven by pseudo potential vorticity carried out in this paper has similarities to an analysis of inertia-gravity waves excited in the earth's atmosphere by \\citet{2004JAtS...61..211V} who perform an analogous WKBJ analysis in the complex plane.  The excitation process produces waves of equal amplitude propagating in opposite directions. As these waves are both trailing, by symmetry each produces an equal outward angular momentum flow. Even when the excitation process is linear, as the waves propagate away, the radial wave length shortens until shock dissipation eventually occurs \\citep[e.g.][]{2001ApJ...552..793G}. Thus waves are always likely to be seen to manifest nonlinear effects as the characteristic radial wave length shortens. When waves behave linearly during the initial excitation, but subsequently undergo significant but not complete dissipation between successive swings, the rate of angular momentum transport can be estimated as the wave action produced in single swing given by equation (\\ref{eq:flux-limit}). This situation is found to occur in the simulations presented in paper II.  \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{figures/epsilon.pdf} \\end{center} \\caption{Exponential dependence of the WKBJ amplitudes on the azimuthal wave number $k_y$ for a Keplerian rotation law.} \\label{fig:F-of-eps} \\end{figure} An important parameter is the value of the azimuthal/horizontal wave number, $k_y,$ for which the wave excitation is most favoured, which we define as being the value for which the wave amplitude produced is maximal. According to (\\ref{eq:pywave}) and (\\ref{eq:ppmwave}) the wave amplitude is a product of a known function of $k_y$ and the square of the Fourier amplitude of PPV at the time of swing. The latter quantity, being determined by the nonlinear hydromagnetic turbulence, cannot be found from the wave excitation calculations performed here. This aspect is discussed in paper II where relevant numerical simulations are performed and analysed. Here we shall anticipate results and assume that the PPV spectrum is relatively flat at small $k_y$ with the consequence that it does not affect the dependence of the wave amplitude on $k_y$ significantly.  The wave amplitude depends on the azimuthal wave number $k_y$ through the parameter \\begin{equation*} \\epsilon = \\frac{q\\Omega k_y c}{k_y^2 c^2 + \\kappa^2} \\end{equation*} in such a way that it is exponentially small for $\\epsilon\\ll{}1$, see (\\ref{eq:wave-amplitude}) and (\\ref{eq:Aplus}). Given the fact that $\\epsilon$ is small both in the small and the long azimuthal wave number limit, we deduce that wave excitation will be most effective near the optimal wave number \\begin{displaymath} k^\\mathrm{opt}_y = \\kappa/c \\end{displaymath} for which $\\epsilon$ takes its maximum value \\begin{displaymath} \\epsilon_\\mathrm{max} = q\\Omega/(2\\kappa). \\end{displaymath} For a Keplerian disc with $q=3/2$ and $\\kappa=1$ we have \\mbox{$k_y^\\mathrm{opt}=1/H$} and thus $\\epsilon_\\mathrm{max}=3/4$. We recall that WKBJ theory gives very accurate results for values of $\\epsilon$ as large as this. For illustrative purposes we plot the exponential dependence of the wave amplitudes as a function of azimuthal wave number in Fig.~\\ref{fig:F-of-eps}.  We see that amplitude of the excited wave falls off rapidly away from the optimal wave number.  The arguments given above suggest that SD wave production will be most effective for \\mbox{$k_y\\sim{}k_y^{\\rm{opt}}$}. This is the longest possible azimuthal wave length for a box with \\mbox{$L_y=2\\pi{}H$} as is commonly adopted. For boxes of this size and smaller wave production is expected to be most effective at the longest azimuthal wave length. On the other hand once $L_y$ exceeds $2\\pi{}H$ the longest wave length is expected to no longer be the most effective. This is fully supported by the simulation results presented in paper II. In this paper we confirm the main features of the excitation process described here and verify the dominance of the pseudo potential vorticity related source terms. Although the waves are observed to become nonlinear very soon after the initial excitation, the main features of the analysis presented here are confirmed. This suggests that useful extensions can be made to the analysis of wave excitation under more general conditions such as those that incorporate significant self-gravity. We plan to undertake these in the near future."
    },
    "0812/0812.2890_arXiv.txt": {
        "abstract": "The appearance and time variability of  accreting millisecond X-ray pulsars (hereafter AMXPs, e.g. Wijnands \\& van der Klis 1998) depends strongly on the accretion rate,  the effective viscosity and the effective magnetic diffusivity of  the disk-magnetosphere boundary. The accretion rate is the main parameter which determines the location of the magnetospheric radius of the star for a given stellar magnetic field. We introduce a classification of accreting neutron stars as a function of the accretion rate and show the corresponding stages obtained from our global 3D magnetohydrodynamic (MHD) simulations and from our axisymmetric MHD simulations. We discuss the expected variability features in each stage of accretion, both periodic and quasi-periodic (QPOs). We conclude that the periodicity may be suppressed at both very high and very low accretion rates. In addition the periodicity may disappear when  ordered funnel flow accretion is replaced by  disordered accretion through the interchange instability. ",
        "introduction": "\\begin{figure} \\centering \\includegraphics[height=.36\\textheight]{cases} \\caption{Main regimes of accretion to a rotating neutron star which occur for different accretion rates. Cases A to F show different regimes when the accretion rate gradually decreases. A is the Boundary Layer regime where the magnetic field is completely buried or screened, and matter accretes directly to the surface of the star; B is the Magnetic Boundary Layer regime where the magnetosphere is small but may influence  the dynamics of the matter flow; Regimes C and D are where the magnetosphere is large enough to channel the matter flow;  however, accretion may also occur through instabilities (case C); Regime E is the Propeller stage.  Regime F is the radio-pulsar stage.} \\end{figure}   \\noindent {\\bf A. BOUNDARY LAYER (BL).} If the accretion rate is sufficiently large  that $r_m < R_*$, then the star's magnetic field is completely buried by the accreting matter (e.g. Cumming, Zweibel, Bildsten 2001; Lovelace et al. 2005), and matter accretes to the star directly through a boundary layer, which is expected to be approximately axisymmetric because there is no channeling by the magnetic field. In this case, high-frequency QPOs may be associated with inhomogeneities at the boundary between inner radius of the disk and the star. Axisymmetric hydrodynamic simulations of accretion through the BL have been recently performed by Fisker \\& Balsara (2005).    Additional factors such as radiative pressure in the inner regions of the disk around NS may also determine the characteristic radius in the disk and corresponding high-frequency QPOs (e.g., Miller, Lamb, \\& Psaltis 1998; Miller \\& Lamb 2001).      \\begin{figure} \\centering \\includegraphics[height=.17\\textheight]{mbl-fun-4} \\caption{3D MHD simulations of accretion through the magnetic boundary layer of a star (in the stable regime)  with a misalignment angle $\\Theta=15^\\circ$. Left two panels show slices of density distribution and sample magnetic field lines in the case of a small magnetosphere. The two right-hand panels show similar slices but for a tiny magnetosphere.} \\end{figure}   \\begin{figure*} \\centering \\includegraphics[height=.16\\textheight]{mbl-3} \\caption{3D MHD simulations of accretion through an unusual magnetic boundary layer which formed in the case of a small misalignment angle $\\Theta=5^\\circ$. From left to right: 3D image of the accretion disk; an equatorial slice of the density distribution, and the Fourier power spectrum.} \\end{figure*}   \\begin{table} \\begin{tabular}{l@{\\extracolsep{0.5em}}l@{}lll}  \\hline &                                             & Type      & $\\dot M$ (in $M_\\odot/yr$) & Type of Variability                   \\\\ \\hline \\multicolumn{2}{l}{$r_m < R_*$}                 & Boundary Layer          & $\\dot M > 7.3\\times 10^{-8}$            & QPO                     \\\\ \\multicolumn{2}{l}{$R_* < r_m < 2 R_*$}       & Magnetic BL  & $6.5\\times 10^{-9} < \\dot M < 7.3\\times 10^{-8}$     & QPO ($P_*$)                  \\\\ \\multicolumn{2}{l}{$2 R_*< r_m < 3.1 R_* $}   & Funnel, Instabilities  & $1.4\\times 10^{-9} < \\dot M < 6.5\\times 10^{-9}$   & $P_*$, QPO \\\\ \\multicolumn{2}{l}{$3.1 R_*< r_m < 12 R_* $}  & ``Propeller\"  &   $1.3\\times 10^{-11} < \\dot M < 1.4\\times 10^{-9}$  & QPO ($P_*$)  \\\\ \\multicolumn{2}{l}{$r_m > 12 R_* $}           & Pulsar    & $\\dot M < 1.3\\times 10^{-11}$    & radio $P_*$   \\\\ \\hline \\end{tabular} \\caption{Ranges of magnetospheric radii, accretion rates and expected variability features in different regimes of accretion for a neutron star with $P_*=2.5 {\\rm ms}$.} \\label{tab:refval} \\end{table}  \\medskip  \\noindent{\\bf B. MAGNETIC BOUNDARY LAYER (MBL)}. At lower accretion rates, the star's magnetic field is only partially buried, so that it influences matter flow around the star. We will draw an approximate line between MBL case and case of larger magnetospheres at $r_m =2R_*$, so that the MBL cases correspond to $R_* < r_m < 2 R_*$. Even in this range there are a number of possible cases with different observational properties (Romanova \\& Kulkarni 2008). \\smallskip  \\noindent{\\it MBL with No Magnetosphere}. If the magnetospheric radius is very close to the stellar radius, then it cannot channel matter to the poles. However, some magnetic flux threads the inner regions of the disk and thus may influence the dynamics of the disk (e.g., Warner \\& Woudt 2002). The authors suggested that the inner region of the disk may form an equatorial belt, which may be responsible for the QPOs in dwarf novae,  a sub-class of cataclysmic variables. Such MBLs were proposed theoretically and should be investigated in greater detail in the future. First exploratory simulations have been done by Romanova \\& Kulkarni (2008).  \\smallskip  \\noindent{\\it  MBL with Small Magnetosphere}. We investigated accretion to a star with a small magnetosphere in global 3D MHD simulations (Romanova \\& Kulkarni 2008). We observed from simulations that matter may accrete in stable or unstable regime.  First we investigated accretion in stable regime to a star with misalignment angle $\\Theta=15^\\circ$ in two cases of small and tiny magnetosphere. We observed that small magnetosphere formed and matter accreted through funnel streams (see Figure 2, left panels). In the case of about twice as smaller dipole field of the star a tiny magnetosphere formed and also disrupted the disk (see Figure 2, right panels). Equatorial slice shows that the magnetic field lines are wrapped by the disk and may influence its structure. In particular, non-axisymmetric bending-type waves have been observed in the inner parts of the disk.  Another run at smaller misalignment angle $\\Theta=5^\\circ$ led to {\\it unstable accretion} but of unusual kind:  instead of sporadically forming unstable tongues as in the unstable regime (Kulkarni \\& Romanova 2008a,b; Romanova, Kulkarni \\& Lovelace 2008), we observed the formation of two symmetric equatorial tongues which connect the star with the disk and which rotate around the star with angular velocity of the inner disk (see Figure 3). The tongues form two equatorial spots which rotate faster than the star: each spot rotates about 4 times faster than the star. The frequency approximately corresponds to the frequency of the inner disk and will change when position of the inner disk change. Thus a star in this accretion regime will reveal high frequency QPOs with frequency drift determined by the accretion rate to the star. Such frequency drift and correlation with matter flux is observed in AMXPs (e.g., van der Klis 2000).  \\begin{figure} \\centering \\includegraphics[height=.3\\textheight]{stab-unstab} \\caption{ Stable (left) and unstable (right) regimes of accretion. The bottom panels show light-curves from the hot spots (from Romanova, Kulkarni \\& Lovelace 2008).} \\end{figure}   \\medskip \\noindent{\\bf C and D. UNSTABLE AND STABLE REGIMES OF ACCRETION}. For even lower accretion rates, the disk matter is stopped by the magnetosphere at a larger distance, a few radii of the star. Recent 3D MHD simulations have shown that in this case the matter may accrete either through symmetric {\\bf funnel streams} which are lifted above the magnetosphere (stable regime), or through an interchange instability, where temporary stochastic tongues of matter form at the disk-magnetosphere boundary and penetrate deeply into magnetosphere. These tongues form multiple hot spots on the surface of the star. Figure 4 shows both regimes of accretion and the corresponding light-curves associated with the hot spots on the surface of the star. \\smallskip  \\noindent {\\it Stable Regime of Accretion (Funnel Streams).} In the stable regime, matter accretes in two ordered funnel streams (Figure 4, left panel, see 3D MHD simulations by Romanova et al. 2003, 2004a; Kulkarni \\& Romanova 2005). These two funnels form two ordered spots at the surface of the star, which vary only slightly around their preferred positions, so that the stellar period is clearly observed in the Fourier power spectrum. This is  the regime where X-ray pulsations in AMXPs are expected. If the magnetic field is more complex than the dipole, then the hot spots have a more complex structure. However,  spots will have a definite preferred locations so that a definite periodicity is expected (Long et al. 2008). Our simulations  shown that a star may spin up or down due to interaction of the magnetosphere with the disk, so that the star's  period is expected to wander around the rotational equilibrium state, which corresponds to the condition $r_m\\approx 1.2 r_{cor}$, which has been derived from simulations by Long et al. (2005) (see also Ghosh \\& Lamb 1979). Such variations of the period have been recently observed, e.g., by Burderi et al. (2007) and Papitto et al. (2008). In addition to periodic variations associated with the hot spots, QPOs associated with oscillations of the inner disk regions are expected (Lovelace \\& Romanova 2007;  Lai \\& Zhang 2008).  Simulations have shown that at small misalignment angles $\\Theta < 5^\\circ$ a hot spot may rotate faster or slower than the star depending on the accretion rate (Romanova et al. 2003; Romanova et al. 2006). This may lead to the rapid phase variations and other interesting phenomena observed in AMXPs (Lamb et al. 2008a,b).   \\begin{figure} \\centering \\includegraphics[height=.25\\textheight]{sketch1-2} \\caption{Disk-magnetosphere interaction for different  magnetic Prandtl numbers. The left-hand panel is for  $Pr=1$, when the disk matter flows in and penetrates through the magnetosphere diffusively at the same rate. The right-hand panel is for $Pr>>1$, where the radial accretion speed is significantly larger than the diffusion speed of the field through the matter.} \\end{figure}  \\begin{figure} \\centering \\includegraphics[height=.2\\textheight]{prop-2} \\caption{The left panel shows outflow of matter (background, selected field lines (red) and streamlines (blue) in the strong propeller regime. The fluxes at the right panel show that the process of accretion and outflows is episodic (from Romanova et al. 2005).} \\end{figure}  \\smallskip  \\noindent The {\\it Unstable Regime of Accretion} has been observed in multiple 3D MHD simulations in cases where the misalignment angle of the dipole is small, $\\Theta=5^\\circ$ (Kulkarni \\& Romanova 2008a,b; Romanova et al. 2008).  In the unstable regime, matter accretes to the star through several ``tongues\" which form at the inner edge of the disk and accrete to the surface of the star on the dynamical time-scale, forming temporary hot spots on its surface. The light-curve from the hot spots looks stochastic (see Figure 4, right panel), and in many cases no period is detected in the wavelet or Fourier spectra. Thus, in the strongly unstable regime, a star with a full-sized magnetosphere may be in the stochastic regime of accretion. In less strongly unstable cases, both tongues and funnels are present, and periodicity can be detected. \\smallskip  \\noindent{\\it  Boundary Between Stable and Unstable Regimes:} We observed from numerous simulations that (1) the stable regime usually dominates at high misalignment angles, $\\Theta \\gtrsim 30^\\circ$; (2) enhancement of the accretion rate usually leads to the onset of the unstable regime (Kulkarni \\& Romanova, these Proceedings). The last condition can be interpreted as follows: for a NS with a given period of rotation $P_*$, the growth of the accretion rate leads the gravitational force to dominate over the centrifugal (or, $r_m << r_{cor}$),  and the instability appears more easily. At smaller accretion rates, the magnetospheric radius moves out, closer to $r_{cor}$, and the centrifugal force becomes strong enough to oppose instability. Detailed comparisons of stable/unstable regimes with instability criteria were also performed (Kulkarni \\& Romanova 2008a; Romanova et al. 2008).  From other side we observed that at larger accretion rate the density of matter at the inner edge of the disk is larger which is favorable for the onset of instability.  \\medskip    \\noindent{\\bf F. ``PROPELLER'' REGIME}. At even smaller accretion rates, the magnetospheric radius $r_m$ becomes larger than the corotation radius $r_{\\rm cor}$, and the star enters the ``propeller\" regime. Axisymmetric simulations performed in our group have shown that propeller regime may be either $\\it strong$ or $\\it weak$.  \\noindent{\\it In the Strong Propeller Regime}, disk matter acquires angular momentum from the rotating magnetosphere fast enough that most of it is ejected in conical outflows. At the same time a significant amount of angular momentum and energy flows along the open stellar field lines, giving an axially-directed Poynting-flux dominated jet (Romanova et al. 2005; Ustyugova et al. 2006). Multiple simulations have shown that the strong propeller operates when $Pr >> 1$. In this case the field lines are bunched into an X-type point (see Figure 5, right panel) which is favorable for outflows, as discussed by Shu and collaborators (e.g., Cai et al. 2008). In this regime, the cycles of matter accumulation at the inner disk, its diffusion through magnetosphere and outbursts to jets are observed with some accretion to the star (see also Goodson et al. 1997). Smaller ratios $r_m/r_{cor}$ and higher diffusivities (still at $Pr \\gg 1$) are also favorable for the strong propeller regime. Oscillations are often quasi-periodic, with frequency $\\nu_{QPO}=0.01-0.2 \\nu_*$, where the lowest frequency is observed at the lowest diffusivity, and vice-versa. \\smallskip  \\noindent{\\it In the Weak Propeller Regime}, when $\\Pr\\approx 1$, the magnetosphere interacts with the disk, and the star loses its angular momentum without outflows. The disk may oscillate due to such interaction, and most of the incoming disk matter may accrete to the star's surface (Romanova et al. 2004b). It is often the case that the rotating magnetosphere (modified by the azimuthal wrapping of the field) pushes the disk to larger distances, and matter may accrete from that point. In the limit of  a very weak propeller where $r_m\\approx r_{cor}$, stellar pulsations and different QPO frequencies are expected as in the case of funnel flow accretion. If magnetosphere is large enough then accretion may be stopped by the centrifugal barrier of the fast rotating magnetosphere in both ``propeller\" regimes.  \\medskip  \\noindent{\\bf G. PULSAR REGIME}. ~For even smaller accretion rates, the radius of the magnetosphere becomes larger than the light cylinder radius, $r_{L}=c/\\Omega_*$ . In this regime, accretion to the surface of the NS is suppressed, and the star becomes a radio-pulsar. This stage is least investigated numerically and it is not clear, for example, at what values of $r_m/r_{LC}$ the radio-pulsar emission dominates. ",
        "conclusions": "Analysis of different stages of evolution of accreting magnetized neutron stars have shown that they have different appearances for different accretion rates. In addition, their appearance may be different depending on other factors, such as the magnetic diffusivity and the Prandtl number (ratio of viscosity to diffusivity). We chose as an example a typical NS with period $P_*=2.5~{\\rm ms}$ ($\\nu_*=400 {\\rm Hz})$, and derived accretion rates from equation (2) corresponding to the boundaries between different regimes. The results are shown in the Table. As we see from simulations of different  regimes, periodic oscillations (with the star's period) are observed only in the regime of magnetospheric accretion which corresponds to a rather narrow range of accretion rates, $1.4\\times 10^{-9}~{M_\\odot/yr} < \\dot M < 6.5\\times 10^{-9}{M_\\odot/yr}$. This fact may explain the appearance of pulsations in intermittent AMXPs in a narrow range of $\\dot M$ as discussed by Altamirano et al. (2008) in the case of SAX J1748.9-2021. If the pulsations are also present in the magnetic boundary layer case and in the weak propeller regime, then these boundaries for  $\\dot M$ will be somewhat wider."
    },
    "0812/0812.3402_arXiv.txt": {
        "abstract": "The impact of environment on AGN activity up to $z\\sim1$ is assessed by utilizing a mass-selected sample of galaxies from the 10k catalog of the zCOSMOS spectroscopic redshift survey.  We identify 147 AGN by their X-ray emission as detected by $XMM$-Newton from a parent sample of 7234 galaxies.  We measure the fraction of galaxies with stellar mass $M_*>2.5\\times10^{10}$ M$_{\\sun}$ that host an AGN as a function of local overdensity using the 5th, 10th and 20th nearest neighbors that cover a range of physical scales ($\\sim1-4$ Mpc).  Overall, we find that AGNs prefer to reside in environments equivalent to massive galaxies with substantial levels of star formation.  Specifically, AGNs with host masses between $0.25-1\\times10^{11}$ M$_{\\sun}$ span the full range of environments (i.e., field-to-group) exhibited by galaxies of the same mass and rest-frame color or specific star formation rate.  Host galaxies having $M_*>10^{11}$ M$_{\\sun}$ clearly illustrate the association with star formation since they are predominantly bluer than the underlying galaxy population and exhibit a preference for lower density regions analogous to SDSS studies of narrow-line AGN.  To probe the environment on smaller physical scales, we determine the fraction of galaxies ($M_*>2.5\\times10^{10}$ M$_{\\sun}$) hosting AGNs inside optically-selected groups, and find no significant difference with field galaxies.  We interpret our results as evidence that AGN activity requires a sufficient fuel supply; the probability of a massive galaxy to have retained some sufficient amount of gas, as evidence by its ongoing star formation, is higher in underdense regions where disruptive processes (i.e., galaxy harrassment, tidal stripping) are lessened. ",
        "introduction": "The local environment of galaxies harboring Active Galactic Nuclei (AGN) and QSOs has long been thought to play a potential role in triggering mass accretion onto Supermassive Black Holes (SMBHs).  With many of the properties of galaxies (e.g., morphology, color, star formation rate) clearly dependent on environment \\citep[e.g., ][]{ka04,ba06,co06,cu09} and the possibility of a common history of mass assembly for SMBHs and their host bulges \\citep[e.g.,][]{gr04,bo06,ho08,so08}, we expect that AGNs may prefer to reside in specific environments most nuturing for their growth. Identifying environmental factors might allow us to determine the physical mechanism(s) responsible for driving accretion such as major mergers of galaxies that has been demonstrated through numerical simulations to be able to remove angular momentum from rotationally-supported gas thus transfering mass to the nuclear regions \\citep[e.g.][]{ba96,mi96}.  For example, there is evidence that AGNs reside in dark matter halos with masses $M_{halo}\\sim10^{12-13} M_{\\sun}$ \\citep{por04,hop07,ma08,pa08,bon08}, a mass regime comparable to the group-scale environments thought to be fertile ground for galaxies to coalesce.  As well, high density regions such as massive clusters of galaxies are expected to be inhospital environments for AGN \\citep{dr85} given the strong empirical association between AGN activity and concurrent star formation \\citep[e.g., ][]{ka03,jahnke04,si09} although counterevidence exists for the radio-loud population \\citep{hill91,be07}.  Environmental studies to date have presented seemingly disparate results most likely due to varying selection methods and physical scales used to characterize environment.  Using a large sample of narrow-line AGN from the SDSS, \\citet{mi03} find that there is no environmental dependence on AGN activity using a magnitude-limited sample of galaxies and meaning environment on scales of around a few megaparsecs.  Upon further investigation, \\citet{ka04} find that an environmental dependence, similar to star formation, emerges when implementing a selection based on stellar mass and considering the luminosity of the AGN.  Supportive of this scenario, \\citet{coil07} find that quasars in the DEEP2 survey fields have environments similar to blue galaxies.  On the other hand, luminous quasars from the SDSS have an overabundance of galaxies within their vicinity on smaller scales \\citep[$<0.1$ Mpc;][]{se06} that is in agreement with clustering analysis based on quasar pairs \\citep{he06} but may still have significant biases in the employed methods \\citep{pad08}. Furthermore, the enhancement in star formation attributed to galaxy mergers on scales of 0.01-0.1 kpc is not reflected in AGN accretion \\citep{li08} for which the authors conclude may be due to varying timescales with star formation preceeding AGN activity \\citep[e.g.,][]{sch07} occurring within a more relaxed host galaxy.  X-ray selected surveys with both $Chandra$ and XMM-$Newton$ now enable the study of the environment of AGNs including the obscured population at higher redshifts ($z\\gtrsim0.3$) where optical selection of narrow-line AGN is difficult.  \\citet{gr05} first looked into the environments of X-ray selected sources detected in the $Chandra$ Deep Fields and found that the near-neighbor counts were identical to the galaxies without AGN.  This led the authors to conclude that mergers were not the physical mechanism triggering mass accretion especially since the AGN host galaxies were no more asymmetric than the average galaxy of equivalent luminosity \\citep[see][for a morphological study of AGN hosts in COSMOS]{ga08}.  While \\citet{ge07} claim based on a large spectroscopic sample of galaxies from DEEP2 that X-ray selected AGN at $z\\sim1$ prefer higher density environments, their results are consistent with those of galaxies having similar host properties (i.e., absolute magnitude, rest-frame color).  From a complementary perspective, \\citet{ma07} has measured the AGN content of galaxy clusters and found no significant difference with the fraction of field galaxies hosting AGN.  Recently, \\citet{gi08} find clustering lengths of AGN in the COSMOS field comparable to massive galaxies and conclude that the clustering signal is reflective of SMBHs preferring to reside in galaxies with $M_*>3\\times10^{10}$ M$_{\\sun}$.  To date, an environmental analysis of X-ray selected AGN based on both a large spectroscopic survey of galaxies and careful consideration of selection to disentangle the degeneracy between mass and environment seen in galaxy studies \\citep[e.g.,][]{ba06,van08,cu09} has not yet been attempted.  Here, we utilize the rich multiwavelength observations of the COSMOS field \\citep{sc07} to determine the role of environment in triggering AGN at $0.1<z<1.0$.  The COSMOS survey is roughly a 2 square degree region of the sky selected to have full coverage with all major observatories both from the ground (i.e., Subaru, VLT) and space (e.g., $HST$, $Spitzer$, XMM-$Newton$).  The zCOSMOS survey \\citep{lilly07,lilly09} targets objects for optical spectroscopy with the VLT in two separate observing programs. A 'bright' sample ($i_{ACS}<22.5$) is observed with a red grism to provide a wavelength coverage of $5500-9500$ ${\\rm \\AA}$ ideal to identify galaxies ($L_*$) up to $z \\sim 1.2$.  A deeper program, not utilized in the present study, targets faint galaxies ($B<25$), selected to be in the redshift range $1.5\\lesssim z\\lesssim 2.5$ using a blue grism for an effective wavelength coverage of $3600<\\lambda<6700~{\\rm \\AA}$.  For the present study, we select galaxies based on reliable spectroscopic redshifts from the zCOSMOS 'bright' program and their stellar mass estimates based on broad-band photometry.  Those that host AGN are identified by their X-ray emission as detected by XMM-$Newton$ \\citep{ha07,cap07,cap08}.  We use the nearest neigbor approach to determine the projected local galaxy density using the three-dimensional galaxy distribution as fully described in \\citet{ko09a}.  In addition, the zCOSMOS optically-selected group catalog \\citep{kn09} offers a complementary perspective on the role of environment.  The low optical luminosity and obscured AGN further enable us to determine if the host galaxies of AGN exhibit a trend similar to the color-density or SFR-density relations found for non-active galaxies.  Finally, we refer the reader to \\citet{si09} that presents the properties (i.e., stellar mass, SFR, rest-frame color) of the hosts of X-ray selected AGN in the COSMOS field.  Throughout this work, we assume $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.75$, $\\Omega_{\\rm{M}}=0.25$ and use AB magnitudes. ",
        "conclusions": "We have assessed whether the environment plays any role in regulating mass accretion onto SMBHs, on scales of a few megaparsecs, using 7234 galaxies ($i_{acs}<22.5$; $0.1<z<1.0$) from the zCOSMOS survey ('Bright' program).  To do so, a parent sample of 2457 galaxies is carefully constructed based on stellar mass ($log~M_*>10.4$) with no obvious color ($U-V$) bias.  Galaxies undergoing substantial accretion onto a SMBH are identified by the presence of an AGN with X-ray emission detected by XMM-$Newton$; we find 147 AGN with 101 residing in host galaxies having $log~M_*>10.4$.  With two characterizations of the local environment, we measure the fraction of galaxies as a function of their galaxy overdensity based on the zCOSMOS density field \\citep{ko09a} and presence in optically-selected galaxy groups \\citep{kn09}.  Based on the zCOSMOS density field, we find that the AGNs, with exception of those residing in the most massive host galaxies ($log~M_*>11$), span a broad range of environments \\citep[see][for similar conclusions]{ge07}, from the field to massive groups, equivalent to galaxies of similar stellar mass.  Our findings do appear to be consistent with low redshift studies based on narrow-line AGN from the SDSS \\citep{mi03,ka04,co08} that taken together demonstrate that the environments depend on the stellar mass of their host galaxy.  The lack of an environmental dependence on AGN activity over the lower mass regime ($10.4<log~M_*<11$) in zCOSMOS is compatible with the results of \\citet{mi03} since their magnitude-limited sample, dominated by low luminosity AGN, most likely includes many lower mass galaxies.  An environmental dependence with AGN activity less common in higher density regions emerges for 'strong AGN' residing in galaxies of higher mass \\citep[][]{ka04,co08} and is also evident in our study based on zCOSMOS galaxies with $log~M_*>11$. We contend that the discrepant findings based on SDSS studies at low redshift may be due to the differences in the underlying mass distribution of galaxes rather than the differences of the AGN selection.  We find that the presence of young stars in the host galaxies of AGN \\citep[][]{ka03,jahnke04,si09} is also indicative of their environments. Galaxies hosting AGN in zCOSMOS follow a similar color-density relation to non-active galaxies \\citep{cu09} with blue hosts ($U-V<1.6$) having a lower mean overdensity than red galaxies.  An equivalent relation based on SFR or that weighted by stellar mass (specific) supports this assertion.  As fully discussed in \\citet{cu09}, the sSFR is seen to be more strongly dependent on environment than SFR due to the underlying mass dependency discussed above.  The majority of massive ($log~M_*>11$) AGN hosts having colors bluer than the underyling galaxy population agrees with their residence in lower density environments.  Therefore, we conclude that the incidence of AGN activity as a function of environment depends on the properties of its host galaxy namely its mass content both in the form of stars and gas in the interstellar medium.  Furthermore, we carry out a complementary analysis of the environments of AGN by utilizing the optically-selected catalog of galaxy groups \\citep{kn09} in zCOSMOS.  Such groups have halo masses $\\sim10^{12}-10^{14}$ M$_{\\sun}$ similar to those that host the more luminous quasars \\citep[e.g.,][]{por04,pa08,bon08} and less massive than clusters at low redshift.  We measure the incidence of AGN activity in galaxies in groups compared to that in the 'field'.  This method effectively enables us to probe smaller physical scales ($\\lesssim1$ Mpc) not possible with the density field due to the spatial resolution of $\\pm1000$ km s$^{-1}$ along the line-of-sight.  An excess signal around quasars, based on either near-neighbor counts \\citep{se06} or the quasar correlation function \\citep{he06}, highlights the importance of probing scales $\\sim100$ kpc.  Although, the environments of 31 X-ray selected AGN \\citep{wa05} with $0.4<z<0.6$ from the Canada-France-Redshift Survey fields are indistinguishable from a well-matched control sample of galaxies on scales around 30-500 kpc.  We find in zCOSMOS that AGN are essentially equally likely to reside in or out of a galaxy group irrespective of its properties (i.e., effective richness).  The fraction of group galaxies ($log~M_*>10.4$) hosting an AGN is similar to 'field' galaxies ($\\sim3\\%$).  This is in agreement with both our results based on the density field for galaxies spanning the full mass range $log~M_*>10.4$, and those of \\citet{ge08} which consider the host galaxy properties.  We note that the current AGN sample in zCOSMOS residing in galaxy groups is too small to investigate if an environmental effect for the most massive galaxies exists as exercised using the density field.  Our lack of a strong environmental effect is also consistent with the incidence of X-ray selected AGN in galaxy clusters \\citep{ma07} as compared to the field.  Also, an increase in the AGN fraction in the higher of two redshift bins is significantly smaller than the enhancement reported for galaxy clusters \\citep{ea07} and in our case most likely due to the luminosity evolution of the AGN population since both the group and field samples show an increase of equal magnitude.  We conclude that internal processes \\citep{ho_he06} play an important role in the growth of SMBHs during an AGN phase of moderate-luminosity ($L_X\\sim10^{43}$ erg s$^{-1}$; 'Seyfert mode') that accounts for the bulk of the Cosmic X-ray Background \\citep[e.g.,][]{gi07}.  Major mergers of galaxies, possibly relevant for the more luminous quasar phenomenon, may not be the primary mechanism for fueling these AGN due to the lack of any enhancement of activity in specific environments likely to be conducive for merging (i.e., galaxy overdensities comparable to the small group scale).  In the opposite sense, we do find that for the most massive galaxies ($log~M_*>11$) the environment plays a role possibly through various physical processes (e.g., tidal stripping, harrassment) in higher density environments that concurrently shuts down star formation and accretion onto supermassive black holes.  This highlights the requirements for a galaxy to host an accreting SMBH, as observed by their X-rays, a massive bulge \\citep[e.g.,][]{sa04,gr05,si08a,ga08} and a sufficient fuel supply as indicated by the young stars usually present in AGN host galaxies \\citep{ka03,si09}."
    },
    "0812/0812.1631_arXiv.txt": {
        "abstract": "{} {Properties of \\textit{transient horizontal magnetic fields} (THMFs) in both plage and quiet Sun regions are obtained and compared.} {Spectro-polarimetric observations with the Solar Optical Telescope (SOT) on the \\emph{Hinode} satellite were carried out with a cadence of about 30 seconds for both plage and quiet regions located near disk center. We select THMFs that have net linear polarization ($LP$) higher than 0.22\\%, and an area larger than or equal to 3 pixels, and compare their occurrence rates and distribution of magnetic field azimuth. We obtain probability density functions (PDFs) of magnetic field strength and inclination for both regions.} {The occurrence rate in the plage region is the same as for the quiet Sun. The vertical magnetic flux in the plage region is $\\sim$8 times larger than in the quiet Sun. There is essentially no preferred orientation for the THMFs in either region. However, THMFs in the plage region with higher $LP$ have a preferred direction consistent with that of the plage-region's large-scale vertical field pattern. PDFs show that there is no difference in the distribution of field strength of horizontal fields between the quiet Sun and the plage regions when we avoid the persistent large vertical flux concentrations for the plage region. } {The similarity of the PDFs and of the occurrence rates in plage and quiet regions suggests that a local dynamo process due to the granular motion may generate THMFs all over the sun. The preferred orientation for higher $LP$ in the plage indicates that the THMFs are somewhat influenced by the larger-scale magnetic field pattern of the plage.} ",
        "introduction": "Horizontal fields with high time variation occurring all over the Sun were observed with the SOLIS and GONG instruments \\citep{Harvey2007}. High resolution spectroscopic observations with the SOT spectropolarimeter (SP) \\citep{Tsuneta2008SoPh,Suematsu2008SoPh,Ichimoto2008SoPh,Shimizu2008SoPh} aboard \\emph{Hinode} \\citep{Kosugi2007} have confirmed this finding and extended the studies considerably \\citep{Lites2007,Orozco2007}. Recently \\citet{Ishikawa2008} have reported the presence of THMFs in a plage region. \\citet{Centeno2007} and \\citet{Ishikawa2008} show examples of the emergence of THMFs both in the quiet Sun and in a plage region. In addition, \\citet{Tsuneta2008ApJ} shows the ubiquitous occurrence of THMFs in polar regions. These horizontal magnetic fields are subject to the upflows and downflows of convection. Earlier reports of flux emergence in the quiet Sun \\citep{Lites1996,DePontieu2002,MartinezGonzalez2007} appear to be examples of THMFs.  Plage regions and the quiet Sun differ considerably in their amounts of vertical magnetic flux. The properties of THMFs in plage regions and in the quiet Sun, however, appear to be similar \\citep{Centeno2007,Ishikawa2008}. A fundamental question is whether there is any difference in the properties of THMFs in these regions. In this paper, we compare THMFs in a plage region with those in a quiet Sun region to analyze their similarities and differences, and to clarify any relationship between the granular-sized THMFs and the existing vertical magnetic field. Relationships between THMFs and the larger-scale predominately vertical fields associated with plage and network regions may shed light on whether local or global-scale dynamo processes generate these transient fields. At present, the evidence that a local dynamo is playing a significant role for the quiet Sun magnetism comes from the Hanle-effect investigation by \\citet{Trujillo2004}, who inferred a magnetic energy density which is the order of 20\\% of the kinetic energy density produced by the convective motions in the quiet solar photosphere, and showed that the observed scattering polarization signals do not seem to be modulated by the solar cycle. In this paper we try to provide more evidence in favor of a local dynamo by comparing \\emph{Hinode} observations of THMFs in a plage region and in the quiet Sun.  \\begin{figure} \\includegraphics[height=8cm]{QS_MDI.eps} \\caption{$SOHO$ MDI longitudinal magnetograph taken at 01:35:01~UT on 2007 November~30. The approximate extent of the area scanned by the SP is shown by the rectangular box. We did not get simultaneous magnetgorams with \\emph{Hinode}, and cannot tell the exact position of scan. The larger square shown with dashed lines shows the area used for estimate of total vertical flux. The three images to the right represent the continuum intensity, the net linear polarization $LP$, and circular polarization of the slot obtained almost simultaneously with the magnetograph.} \\label{QSmdi} \\end{figure}  \\begin{figure} \\includegraphics[height=8cm]{Plage_MDI2.eps} \\caption{$SOHO$ MDI longitudinal magnetograph taken at 15:59:02~UT on 2007 February~10. The caption is the same as Fig.\\ref{QSmdi}. Note that three rectangles with dotted line in slot images show the area which is dominated by persistent strong vertical magnetic flux concentrations for the plage region, and is masked for the analysis. } \\label{PRmdi} \\end{figure} ",
        "conclusions": "\\subsection{THMF and local dynamo process} The vertical magnetic flux in the plage is about 8 times larger than that of the quiet Sun. Note that we derive these magnetic flux estimates from pixels with field strength stronger than $50$~Gauss in the regions marked by the boxes with dashed lines in Figs.~\\ref{QSmdi} and \\ref{PRmdi}. Horizontal fields dominate the plage region without persistent large vertical magnetic flux concentrations as well as the quiet Sun, and there is no difference in magnetic field strength of horizontal fields between these regions. In addition the occurrence rates of THMFs are nearly identical. If the THMF occurrence rate was in any way directly related to the \\textit{global} vertical fields forming the plage region, then we would expect the occurrence rate in the plage region to be much larger than that of the quiet Sun. This suggests that the emergence of the THMFs does not have a direct causal relationship with the vertical magnetic fields in the plage region. The same THMF occurrence rate, no preferred orientation, and similar field strength distributions for both regions strongly suggest that a common local process that is not directly influenced by global magnetic fields produces THMFs. As ubiquitous THMFs are receptive to convective motion \\citep{Centeno2007,Ishikawa2008}, a reservoir for THMFs may be located near solar surface, and these magnetic fields are carried to the surface with convective flow.  Such reservoir can be maintained by a local dynamo process due to near-surface convective motion \\citep{Cattaneo1999,Vogler2007}. Indeed, numerical simulations \\citep{Abbett2007,Schussler2008} have shown that local dynamo action can generate horizontal magnetic structures in the quiet Sun. Such a local dynamo process could naturally explain the similarity in occurrence rates and field strength PDFs, including the fact that THMFs do not have a preferred orientation. Moreover, the horizontal field strength is distributed in a range smaller than the equipartition field strength of near-surface turbulent granules. This would suggest that local dynamo process takes place in a relatively shallow layer with the size of a granule. The similarity in field strength distribution also indicates that properties of THMFs do not depend on the seed field e.g. global fields.  Other possibilities for the origin of THMFs include debris from decaying active region, magnetic fields failed to emerge from the convection region to the photosphere \\citep{magara2001}, and extended weak magnetic fields in the upper convection zone generated by \\textit{explosion} \\citep{moreno1995}. If the reservoir is maintained by one of these processes the THMFs would be expected to be affected by global toroidal fields in terms of properties of THMFs described above. \\citet{Steiner2008} carried out numerical simulations of the surface layers of the Sun with two different sets of initial and boundary conditions, and compared their simulations with Hinode observations. The first set starts with vertical unipolar magnetic fields appropriate for plage regions, and the second set with horizontal uniform magnetic fields at the bottom boundary area. In both runs they found that spatial and temporal average for the horizontal fields surpass that for the vertical fields but the first run has more vertical fields than the other. As these simulations show, the environment of magnetic fields seems to affect the distribution of the magnetic field vector. Within the context of these simulations the observed properties of THMFs such as the similarity in the occurrence rates and magnetic field distributions, and no preferred orientation of THMFs would be difficult to explain.  A slight preferred orientation of THMFs with higher $LP$ toward the global plage polarity suggests that these THMFs may be influenced by the global plage field. However, because any strong vertical fields associated with the emergence of these THMFs \\citep{Ishikawa2008} are not observed, they are probably not directly created from the vertical magnetic fields forming the plage as suggested by \\citet{Isobe2008}. Thus, even if these THMFs with higher $LP$ are related to the global toroidal system, the relationship would be indirect --- these THMFs with high $LP$ may result from fragmented elements of plage flux tossed about by the convective motions below the photosphere.  \\subsection{THMF and chromospheric and coronal heating} Vertical magnetic elements on the solar surface are naturally connected with the upper atmosphere. It is essential for coronal and chromospheric heating problem to understand their properties and the connectivity with the chromosphere and the corona. This raises a parallel question for the case of the THMFs: do the ubiquitous THMFs in the quiet Sun's photosphere have any consequence on upper atmosphere? We here estimate the magnetic flux carried by THMFs as follows. \\begin{equation} E_{mag}=\\frac{B^{2}}{8\\pi}\\times V \\times f \\times \\frac{1}{\\tau} \\times \\frac{1}{A} \\label{fluxeq} \\end{equation} where $B$ is average magnetic field strength of THMFs ($\\sim310$ G), $f$ average filling factor ($\\sim 0.15$),  and $A$ the observing area  ($2\\arcsec\\times164\\arcsec$) for the quiet Sun, and $B$ of $\\sim370$ G, $f$ of $\\sim 0.18$, and $A$ of $2\\farcs4\\times164\\arcsec\\times0.84$ for the plage region. In order to obtain field strength and filling factor of THMFs, the Milne Eddington inversion was performed for pixels with Stokes $Q$, $U$, or $V$ amplitudes 4.5 times larger than their noise levels (section 3.2). Since the selected pixels would include vertical flux tubes, we further set the condition that $LP$ be higher than 0.22\\%, which is threshold used for detecting THMFs (section 3.1). Average field strength and filling factor of the pixels thus chosen are used for $B$ and $f$ in equation (\\ref{fluxeq}). THMFs are at most as large as a granule where they appear \\citep{Ishikawa2008}, and the size is comparable to that of a granule \\citep[][in prep.]{Ishikawa2008b}. We substitute $\\sim(1\\arcsec)^{3}$ for the volume ($V$). THMF appears in the quiet Sun and in the plage region every 54 s and 47 s as described in previous section, and we substitute 54 s and 47 s for $\\tau$, respectively. Then, we obtain $\\sim2\\times10^{6}$ erg cm$^{-2}$ s$^{-1}$ and $\\sim5\\times10^{6}$ erg cm$^{-2}$ s$^{-1}$ for the quiet Sun and the plage region, respectively. Note that the occurrence rate may be under-estimated due to our limited sensitivity. Thus, the estimated energy flux should be taken as a minimum value.  \\begin{table} \\begin{center} \\begin{tabular}{lccc} \\hline \\hline \\multicolumn{1}{c}{Energy loss} (erg cm$^{-2}$ sec$^{-1}$) &Quiet Sun&Active region\\\\ \\hline \\ \\ Corona &\\mbox{$3\\times10^{5}$}&\\mbox{$10^{7}$}\\\\ Chromosphere &\\mbox{$4\\times10^{6}$}&\\mbox{$2\\times10^{7}$}\\\\ \\hline Magnetic energy &Quiet Sun&Plage region&\\\\ supplied by THMFs  &\\mbox{$\\sim2\\times10^{6}$}&\\mbox{$\\sim5\\times10^{6}$}\\\\ \\hline \\end{tabular} \\end{center} \\caption{Chromospheric and coronal energy losses \\citep{withbroe&noyes1977} and magnetic energy flux (bottom) in the quiet Sun and in a plage deduced in this paper} \\label{EL} \\end{table}  Table \\ref{EL} shows energy losses of active regions and quiet regions obtained from the empirical models \\citep[][Table 1]{withbroe&noyes1977}. Derived estimates of energy losses in each region indicate the required energy input. The magnetic energy flux that THMFs provide to the quiet Sun ($\\sim2\\times10^{6}$ erg cm$^{-2}$ s$^{-1}$) is comparable to total chromospheric energy loss of 4$\\times10^{6}$ erg cm$^{-2}$ s$^{-1}$. The total energy loss of the corona is only one tenth of the energy loss in the chromosphere. The magnetic energy flux in a plage region of $\\sim5\\times10^{6}$ erg cm$^{-2}$ s$^{-1}$ associated with THMFs is comparable to the total chromospheric and coronal energy loss.  We demonstrated that the horizontal magnetic fields could contribute to the chromospheric and coronal heating \\citep[e.g.,][]{Abbett2007,Isobe2008}. Its transient nature may accompany an associated transient chromosheric heating and dynamic events such as many small-scale dynamical events as shown by the \\textit{Hinode} CaII H movies.  We point out a possible connection between the horizontal magnetic fields and high chromospheric activities such as jets discovered by \\textit{Hinode}/SOT \\citep{shibata2007}. We also indicate that the THMFs  we observe are probably the tip of the ice berg due to our limited sensitivity, and there may be more ubiquitous magnetic fields unresolved by \\textit{Hinode} \\citep{Trujillo2004}."
    },
    "0812/0812.2283_arXiv.txt": {
        "abstract": "We present high-resolution {\\it HST} images of all 35 AGNs with optical reverberation-mapping results, which we have modeled to create a nucleus-free image of each AGN host galaxy.  From the nucleus-free images, we determine the host-galaxy contribution to ground-based spectroscopic luminosity measurements at $\\lambda 5100$\\,\\AA.  After correcting the luminosities of the AGNs for the contribution from starlight, we re-examine the H$\\beta$ \\rl\\ relationship.  Our best fit for the relationship gives a powerlaw slope of 0.52 with a range of $0.45 - 0.59$ allowed by the uncertainties.  This is consistent with our previous findings, and thus still consistent with the naive assumption that all AGNs are simply luminosity-scaled versions of each other.  We discuss various consistency checks relating to the galaxy modeling and starlight contributions, as well as possible systematic errors in the current set of reverberation measurements from which we determine the form of the \\rl\\ relationship. ",
        "introduction": "One of the key developments in extragalactic astronomy over the past decade has been the discovery that supermassive black holes are present in most, if not all, galaxies having a stellar bulge. Remarkably, the mass of the black hole is tightly correlated with the stellar velocity dispersion of the host galaxy bulge (\\citealt{ferrarese00,gebhardt00}), pointing to a close link between the growth and evolution of galaxy stellar populations and the growth of nuclear black holes.  Determining the masses of black holes in active galaxies is a crucial step toward understanding this connection, and provides fundamental insight into the physics of accretion and emission processes in the black hole environment.  Unfortunately, most active galactic nuclei (AGNs) are too distant for black hole masses to be measured using spatially resolved stellar or gas dynamics.  The technique that has been most successful for the measurement of the black hole mass (\\mbh) in AGNs is reverberation mapping (\\citealt{blandford82,peterson93}).  With this technique, the AGN continuum (typically measured at 5100\\,\\AA) and broad emission lines (most notably, H$\\beta$) are monitored over an extended time period.  Since the emission-line regions are photoionized by the central source, changes in the AGN continuum strength are followed by changes in the emission-line fluxes, with a time lag that depends on the light-travel time across the broad-line region (BLR).  This time lag can be measured by cross-correlation of the continuum and emission-line light curves, and gives the radius of the BLR. Combining the BLR radius with the broad emission-line velocity width then gives the virial mass enclosed within the BLR, which is dominated by the black hole (e.g., \\citealt{peterson98a,kaspi00}).  The validity of reverberation masses has been upheld by the detection of virial behavior in the broad line region in a subset of objects (e.g., \\citealt{peterson99,peterson00a,onken02,kollatschny03}), as well as the consistency of reverberation masses with other dynamical mass methods, such as stellar dynamics (\\citealt{davies06,onken07}) and gas dynamics \\citep{hicks08}.  Due to the long-term nature of reverberation-mapping projects, these measurements have only been carried out for a relatively small sample of AGNs in the past: about 36 Seyferts 1s and low-luminosity quasars.  The BLR radius--luminosity correlation ($R_{\\rm BLR} \\propto L^{\\alpha}$) derived from this reverberation sample is the basis for \\emph{all} secondary techniques used to estimate black hole masses in distant AGNs (e.g., \\citealt{laor98,wandel99,mclure02,vestergaard06}) and is an essential tool used to search for cosmological evolution of the \\msigma\\ relationship (e.g., \\citealt{peng06,woo08}).  The power of the \\rl\\ relationship comes from the simplicity of using it to quickly estimate \\mbh\\ for large samples of objects, even at high redshift, using two simple measurements from a single spectrum of each object.  \\citet{peterson04} compiled and consistently reanalyzed the database of available reverberation-mapping data for 35 AGNs, with the goal of improving the measurements of the size of the BLR and thereby improving their mass measurements.  Subsequently, \\citet{kaspi05} reanalyzed the \\rl\\ relationship and found a power-law slope of $\\alpha = 0.665 \\pm 0.069$ using the optical continuum and broad H$\\beta$ line.  However, many of the AGNs in the sample reside in host galaxies that are comparable in luminosity to the AGN itself.  Even worse, with the typically large apertures employed in reverberation-mapping campaigns (i.e. 5$\\farcs$0 $\\times$ 7$\\farcs$5), the host-galaxy starlight contribution to the spectroscopic luminosity of any given source is substantial.  Failure to account for the enhancement of the luminosity by starlight results in an artificially steep slope for the \\rl\\ relationship, as a larger percentage of the luminosity in faint AGNs is contributed by the host galaxy.   A preliminary study by \\citet{bentz06a} presented two-dimensional fits to high-resolution \\hst\\ images of 14 reverberation-mapped AGNs, from which the host-galaxy contribution was determined.  The luminosities of the 14 sources were corrected and the \\rl\\ relationship was re-examined, resulting in a measured power-law slope of $\\alpha \\approx 0.5$, consistent with the naive prediction that all AGNs are simply luminosity-scaled versions of each other.  In this work, we present high-resolution \\hst\\ images of the rest of the H$\\beta$ reverberation-mapped AGNs, bringing the total to 35. We improve upon our previous two-dimensional fits and reanalyze the \\rl\\ relationship after correcting the luminosity of every AGN for the contribution from starlight.  We show that these new results are consistent with those from our preliminary study, and discuss the new measurements in light of known systematic errors that may affect the slope of the \\rl\\ relationship.  Throughout this work, we will assume a standard flat $\\Lambda$CDM cosmology with $\\Omega_{\\rm B} = 0.04$, $\\Omega_{\\rm DM} = 0.26$, $\\Omega_{\\Lambda} = 0.70$, and $H_0 = 70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "Throughout this work, we have improved upon our original methods by using more accurate and more conservative profiles to model the host galaxy components, as well as a more conservative color correction method, all resulting in conservative measurements of the host galaxy starlight as measured for every object contributing to the \\rl\\ relationship.  Even so, the best fit to the relationship has not changed significantly from that presented by \\citeauthor{bentz06a}: $0.519^{+0.063}_{-0.066}$ compared to the previous value of $0.518 \\pm 0.039$.  It would appear that given reasonably high-resolution, unsaturated images of AGN host galaxies and reasonable fits to the host galaxy surface brightness profiles, the AGN luminosity can be corrected for the contamination from starlight fairly accurately.  AGN BLRs can be modeled assuming the central radiation field is the only source of heating and ionization of the gas. The simplest, most naive assumption is that all BLRs are made of identical clouds, with the same density, column density, composition and ionization parameter.  In addition, the radiation from all AGNs can be assumed to have the same spectral energy distribution (SED).  This results in the prediction that $R_{\\rm BLR} \\propto L^{0.5}$.  Any changes in the BLR gas distribution or the AGN SED, especially those associated with source luminosity, will result in a different slope.  In fact, AGNs have been observed to have different SEDs as a function of luminosity (e.g., \\citealt{mushotzky89,zheng93}).  It is therefore interesting that our present work based on the results from several reverberation-mapping campaigns produces a slope for the \\rl\\ relationship that is consistent with the naively expected slope of 0.5, and that, to a first approximation, brighter AGNs are simply ``scaled up'' versions of fainter AGNs.  Removing the host-galaxy starlight component reduces the scatter\\footnote{Throughout this manuscript, ``scatter'' refers to the $1-\\sigma$ deviation from the best fit \\rl\\ relationship.} in the \\rl\\ relationship, from $39-44$\\% as found by the best-fit FITEXY results from \\citet{kaspi05} to $34-40$\\% as found here.  It also flattens the slope of the relationship considerably.  This has the overall effect of biasing samples that used previous versions of the \\rl\\ relationship to estimate black hole masses.  The host-galaxy starlight is often removed through spectroscopic decomposition of single-epoch spectra before the size of the BLR radius is estimated from the continuum luminosity.  However, it is crucial to use a \\rl\\ relationship in which the objects providing the calibration are {\\it also} corrected for host-galaxy starlight, which we have provided here.  Compared with the best-fit relationship of \\citet{kaspi05}, we find that the calibration of the \\rl\\ relationship presented here results in black hole masses that are $\\sim 30$\\% smaller at $L = 10^{46}$, $\\sim 50$\\% larger at $L = 10^{44}$, and a factor of $\\sim 3$ larger at $L = 10^{42}$.  At the low-luminosity end ($L \\la 10^{43}$), there may still be some uncertainty as to the behavior of the \\rl\\ relationship.  In the current sample of objects with reveberation-mapping results, there are somewhat fewer objects at lower-luminosity and they have larger uncertainties than the other, higher-luminosity objects in the sample. It should be kept in mind that the lower-luminosity objects were, in general, the first targets of ground-based monitoring campaigns due to their relatively low redshifts and high apparent brightnesses.  The larger uncertainties in their measurements is partially due to the less-rigorous control over observational factors in those early monitoring campaigns (such as observing cadence, spectral resolution, detector efficiency, etc.) simply because there was a lack of experience in this field at that time.  The problems of small sample size and relatively larger uncertainties for lower luminosity objects will soon be mitigated by two independent reverberation mapping campaigns which have been recently carried out at MDM Observatory and at Lick Observatory and targeted the low-luminosity end of the relationship.  Preliminary results from the MDM campaign promise to replace several measurements, as was done in the case of NGC\\,4593 \\citep{denney06} and NGC\\,4151 \\citep{bentz06b} in the 2005 MDM campaign. And preliminary results from the Lick campaign promise to add several new objects to the low-luminosity end of the \\rl\\ relationship (e.g., \\citealt{bentz08}).  Internal reddening is known to be a problem in some of the very nearest, and lower-luminosity, objects in the current reverberation sample.  For example, NGC\\,3227 is known to have substantial internal reddening compared to most of the other objects in the current sample and should therefore require one of the largest reddening corrections. \\citet{crenshaw01} determined a reddening curve for NGC\\,3227 and find that at 5100\\,\\AA, $A_{\\lambda} / E(B-V) = 3.6$ and $E(B-V) = 0.18$. This implies an extinction at 5100\\,\\AA\\ of 0.65\\,mag, which, if corrected for, would increase the luminosity here by a factor of 1.8, or 0.26\\,dex, moving the location of NGC\\,3227 in the bottom panel of Figure~5 from slightly left of the best fit \\rl\\ relationship to right on top of it.  As we do not have similar corrections for the other objects in this sample, we do not apply the reddening correction for NGC\\,3227 in our determination of the \\rl\\ relationship.  However, based on the magnitude of the correction determined for NGC\\,3227, there does not seem to be any reason to expect that correcting all the sources for internal reddening will have much effect on the slope of the \\rl\\ relationship.  An obvious issue that may also be addressed with the galaxy fits that we have presented in this work is the relationship between black hole mass and host galaxy bulge luminosity (or mass; \\citealt{kormendy95,magorrian98}).  We discuss this relationship for the AGNs in our sample in a related manuscript \\citep{bentz09}."
    },
    "0812/0812.2626_arXiv.txt": {
        "abstract": "{Evolution of galaxies through cosmic time has been widely studied at high redshift, but there are a few studies in this field at lower redshifts. However, low--redshifts studies will provide important clues to the evolution of galaxies, furnishing the required link between local and high--redshift universe.} % {In this work we focus on the metallicity of the gas in spiral galaxies at low redshift looking for signs of chemical evolution. We analyze the metallicity contents of star forming galaxies of similar luminosities at different redshifts, we studied the metallicity of star forming galaxies from SDSS--DR5 (Sloan Digital Sky Survey--Data Release 5), using different redshift intervals from 0.1 to 0.4.} {We used the public data of SDSS--DR5 processed with the STARLIGHT spectral synthesis code, correcting the fluxes for dust extinction, estimating metallicities using the $R_{23}$ method, and analyzing the samples with respect to the [{N\\,\\textsc{ii}}] $\\lambda$6583/[{O\\,\\textsc{ii}}] $\\lambda$3727 line ratio.} {From a final sample of 207 galaxies, we find a decrement in 12+log(O/H) corresponding to the redshift interval 0.3 $<$ z $<$ 0.4 of $\\sim$0.1 dex with respect to the rest of the sample, which can be interpreted as evidence of the metallicity evolution in low--z galaxies.} {} ",
        "introduction": "Determination of the chemical composition of the gas in galaxies versus cosmic time provides a very important tool for understanding galaxy evolution, due to its important impact on fields such as stellar evolution and nucleosynthesis, gas enrichment processes, and the primary or secondary nature of the different species.  Optical emission lines in galaxies have been widely used to estimate abundances in extragalactic {H\\,\\textsc{ii}} regions (e.g. Aller 1942, Searle 1971, Pagel 1986, Shields 1990). The direct method of estimating metallicity in galaxies is known as the $``$$T_e$ method$\"$ (Pagel et al. 1992, Skillman $\\&$ Kennicutt 1993), and consists of measuring the oxygen abundance from the electron temperature of the {H\\,\\textsc{ii}} region obtained using, the ratio of a high--excitation auroral line such as [{O\\,\\textsc{iii}}] $\\lambda$4363 to the lower excitation [{O\\,\\textsc{iii}}] $\\lambda\\lambda$ 4959, 5007 lines.  However, [{O\\,\\textsc{iii}}] $\\lambda$4363 is too weak, not only in metal--rich, but even in metal--poor galaxies. In the absence of this auroral line, metallicities have to be estimated from strong line ratios, such as the $R_{23}$ ratio introduced by Pagel et al. (1979). This is a commonly used and well--calibrated method of estimating metallicity abundances (see Alloin et al. 1979, Edmunds $\\&$ Pagel 1984, McCall et al. 1985, Dopita $\\&$ Evans 1986, Pilyugin 2000, 2001, Tremonti et al. 2004, Kewley $\\&$ Dopita 2002, Liang et al. 2006, hereafter L06). Nevertheless, this method has the disadvantage of being double--valued as a function of 12$+$log(O/H). Besides  $R_{23}$, there are other strong--line methods useful for determining abundances of high--metallicity, star--forming galaxies, for example the N2=[{N\\,\\textsc{ii}}] $\\lambda$6583/ {H$\\alpha$} method (Denicol\\'o et al. 2002, Pettini $\\&$ Pagel 2004), the  $S_{23}$=([{S\\,\\textsc{ii}}] $\\lambda\\lambda$6717, 6731 +[{S\\,\\textsc{iii}}] $\\lambda\\lambda$9069,9532)/{H$\\beta$} (V\\'{\\i}lchez  $\\&$  Esteban 1996), among others. For a detailed discussion of the different methods see Bresolin (2006) and Kewley  $\\&$  Ellison (2008).  Many studies of metallicity evolution exist as a function of cosmic time,  although many of them refer to damped Lyman $\\alpha$ systems at z $>$ 2 (e.g. Pettini et al. 2002, Henry $\\&$ Prochaska 2007). For the evolution of the mass--metallicity relation of star--forming galaxies, Erb et al. (2006) find evolution at z $\\gtrsim$ 2 and z $\\sim$ 3.5,  respectively. Also, Brooks et al. (2007) and Finlator $\\&$ Dav\\'e (2008), among others, derive cosmological models of the mass--metallicity relation, while at intermediate redshifts (1 $<$ z $<$ 2),  there are several important studies of the evolution of the chemical composition of the gas, such as the ones by Lilly et al. (2003), Kobulnicky et al (2003), and Maier et al. (2006).  Among the studies at z $<$ 1, mainly based on small samples,  Savaglio et al. (2005) have investigated the mass--metallicity relations using galaxies at 0.4 $<$ z $<$ 1, finding that the metallicity is lower at higher redshift, for the same stellar mass, by ~0.15 dex. Also, Maier at al. (2005), with a sample of 30 galaxies with 0.47 $<$ z $<$ 0.92, find that one--third have  metallicities lower than those of local galaxies with similar luminosities and star formation rates. In contrast, Carollo $\\&$ Lilly (2001) use emission--line ratios of 15 galaxies in a range of 0.5 $<$ z $<$ 1 to find that their metallicities appear to be remarkably similar to those of local galaxies selected with the same criteria. A similar result was found for the luminosity--metallicity relation by Lamareille et al. (2006), comparing star--forming galaxies at local and intermediate redshift (0.2 $<$ z $<$ 1, split into 0.2 redshift bins). However, Buat et al. (2008) and Kobayashi et al. (2007) have derived a model of metallicity evolution as a function of $z$, which show a progressive increase in metallicity with time, even at low redshift.  There is, then, a need to increase the lower redshift galaxy samples to ascertain whether or not at such low redshifts (i.e. small lookback times, of the order of 8.4 Gyr for z $\\sim$ 0.5, using a concordance cosmology) any evidence exists of metallicity evolution. The SDSS database provides an excellent opportunity for extending these studies up to z $\\sim$0.4 to explore a possible evolution of metallicity at low--redshift using larger samples, thus deriving more statistically significant results. ",
        "conclusions": "Although similar studies of metallicity using SDSS exist, they are either restricted to a redshift $\\sim$0.1 (e.g. Tremonti et al. 2004) or do not separate their samples as a function of redshift (e.g. L06), thus making it impossible to detect metallicity evolution. Finally, studies at high redshift are statistically limited.  After exploring possible biases on the sample, the results obtained in the present work suggest  \\emph{prima facie} evidence of an intrinsic metallicity evolution in the local Universe, showing a decrement of $\\sim$0.1 dex in 12+log(O/H) at a redshift interval 0.3 $<$ z $<$ 0.4, a result consistent with the models of metallicity evolution in Buat et al. (2008) and Kobayashi et al. (2007)."
    },
    "0812/0812.3831_arXiv.txt": {
        "abstract": "We present an updated version of MegaZ-LRG \\citep{2007MNRAS.375...68C} with photometric redshifts derived with the neural network method, ANNz as well as five other publicly available photo-z codes (HyperZ, SDSS, Le PHARE, BPZ and ZEBRA) for $\\sim$1.5 million Luminous Red Galaxies (LRGs) in SDSS DR6. This allows us to identify how reliable codes are relative to each other if used as described in their public release. We compare and contrast the relative merits of each code using $\\sim$13000 spectroscopic redshifts from the 2SLAQ sample. We find that the performance of each code depends on the figure of merit used to assess it. As expected, the availability of a complete training set means that the training method performs best in the intermediate redshift bins where there are plenty of training objects. Codes such as Le PHARE, which use new observed templates perform best in the lower redshift bins. All codes produce reasonable photometric redshifts, the 1-$\\sigma$ scatters ranging from 0.057 to 0.097 if averaged over the entire redshift range. We also perform tests to check whether a training set from a small region of the sky such as 2SLAQ produces biases if used to train over a larger area of the sky. We conclude that this is not likely to be a problem for future wide-field surveys. The complete photometric redshift catalogue including redshift estimates and errors on these from all six methods can be found at www.star.ucl.ac.uk/$\\sim$mbanerji/MegaZLRGDR6/megaz.html ",
        "introduction": "Photometric redshifts will be one of the key ingredients for us to improve our understanding of the Universe in the decade to come. Up to date, galaxy large scale structure surveys relied mainly on spectroscopic redshifts to produce high precision power spectrum measurements of the galaxy distribution \\citep[e.g.][]{2005MNRAS.362..505C,2007ApJ...657..645P}. Combined with CMB experiments these surveys have provided evidence that the Universe is flat and is likely to be dominated by a dark energy component \\citep{2008arXiv0803.0547K}.   \\begin{figure*} \\begin{center} \\includegraphics[width=15.5cm,angle=0]{./figs/aitoffdr6_v3.eps} \\caption{Map of the MegaZ-LRG sample (blue) covering the SDSS DR6 area as well as the 2SLAQ sample (red). For clarity only a random subsample of galaxies have been plotted.} \\end{center} \\end{figure*}  However having a considerable step up in the size of the spectroscopic surveys will be a hard task to achieve for technical reasons. Several Multi-fibre optical spectrographs are currently being built (FMOS) \\citep{2006SPIE.6269E.136D} or being designed (WFMOS), but it is unlikely that they will be able to survey a considerable part of the sky. On the other hand radio interferometers may be able to perform spectroscopic surveys of the sky reasonably quickly \\citep{2004NewAR..48.1063B,2005MNRAS.360...27A} but the timescale for the technical advances to allow for this will be relatively long.  The alternative to a full spectroscopic survey is to obtain multi-colour images of the sky and perform photometric redshift estimates for the galaxies we have available \\citep[e.g.][]{2003AJ....125..580C}. In a pilot study with high redshift Luminous Red Galaxies (LRGs) it has been shown \\citep{2006astro.ph..5302P,2007MNRAS.374.1527B} that producing large scale measurements with photometric redshifts is possible and competitive with a smaller spectroscopic redshift survey. Using the same dataset \\citet{2008MNRAS.385.1257B} have shown that photometric redshifts can also be used to study small scale halo model signatures.  On the other hand there are many caveats of photometric redshifts that have to be assessed in order for us to be completely confident that these measurements are reliable to the level of systematics that we expect in future surveys. For instance \\citet{2007MNRAS.374.1527B} have performed a detailed study of whether star-galaxy separation influences the cosmological measurements given that the LRGs that have been selected are contaminated at the per cent level by M-type stars which have similar colours. They have also assessed whether there is a significant contamination from dust corrections in the galaxy, by obtaining estimates of the power spectrum in different regions of obscuration in the sky.  We extend this analysis concentrating on the level of systematic effects that is introduced by the use of different photometric redshift techniques. We have selected the same sample as was selected in the MegaZ-LRG catalogue \\citep{2007MNRAS.375...68C} and used several different photometric redshift techniques on the same galaxies available from the literature, including artificial neural networks, template fitting techniques and Bayesian techniques. We note here that LRGs have well defined 4000 $\\AA$ break, hence this strong feature makes photometric redshift estimation an easier task. Here all codes compared produce reasonable photometric redshifts and we are comparing more subtle differences between codes.  In Sec.\\ref{sec:data} we describe the MegaZ-LRG data used. In Sec.\\ref{sec:methods} we describe all the methods we have used to estimate the photometric redshifts for the LRG sample. In Sec.\\ref{sec:results} we compare different statistics for the different photo-z results. We perform an analysis to check for gradients across the sky which could arise from training sets if they only belong to a small area of the sky in Sec.\\ref{sec:harmonics} and we present the catalogue in Sec.\\ref{sec:catalogue}. Our conclusions are drawn in Sec.\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We have presented an updated version of the MegaZ-LRG catalogue. This catalogue contains about 1.5 million objects with accurate photometric redshifts which can be used for a range of science applications. The catalogue is available on-line and contains SDSS ID information so all SDSS data can be retrieved for each object as well as the photometric redshifts from each of the six public codes.  We have run several comparisons of code and template libraries on the 2SLAQ LRG sample. We conclude that there are differences in the codes and stress that a more thorough comparison is needed where the effects of the codes and template libraries are disentangled. This will allow us to pinpoint where the discrepancies are arising. An approach based on first principles such as that presented in \\citet{2008arXiv0811.2600B} is also timely. We have used several figures of merit to assess which code + template library performed best for this set of galaxies. We conclude that different codes perform with different strengths depending on the figure of merit used. We outline more specifically the findings below:  \\begin{itemize}  \\item As expected, the availability of a complete training set means the training method, ANNz performs best in the intermediate redshift bins where there are plenty of spectroscopic redshifts.  \\item Le PHARE performs very well particularly in the lower redshift bins suggesting the Poggianti templates may be a better fit to LRGs at those redshifts compared to other templates used in this comparison.   \\item  HyperZ run with Bruzual \\& Charlot templates gives better results than using the same with CWW templates once again highlighting the importance of template choice.  \\item The SDSS template code gives very good results compared to other codes at the highest photo-z bins despite having only one evolving template for the LRGs. Given the narrow range in SEDs of our sample of Luminous Red Galaxies, the strengths of template-based codes with extensive template libraries are not adequately highlighted by this comparison.  \\item ZEBRA shows a small average bias indicating the importance of the template optimisation technique in removing biases.  \\end{itemize}  As expected the training code performs best where the training set is large and complete and the template sets overtake the training code if the training set starts to become sparse. The importance of template choice is highlighted by the fact that most figures of merit show codes used in conjunction with the CWW templates to perform worse than those using other training or synthetic templates. This suggests that the CWW templates are not a very good match to the SEDs of these LRGs.  There is a discrepancy between the scatters found for these codes ranging from 0.057 to 0.097. Both values are considered good results for photo-z estimates as one would expect from LRGs but there is a clear difference between the different code and template combinations that are run. Given that these differences will also depend on galaxy type and training set size, it is imperative that we carry out a more thorough comparison where the effects of codes and templates are deconstructed, in order to understand what factors affect the photo-z accuracy. We caution the reader that the results presented here are specific to a sub-sample of galaxies, namely LRGs, which have a narrow range in SEDs and a complete and representative spectroscopic training set available. The conclusions presented here could and probably would change if the comparison were made with different galaxies or a different training set size.  We have also produced a set of tests to assess whether the fact that the training sets are from a restricted area in the sky affects the photometric redshifts significantly. We conclude that there is little or no difference between the results from template methods and training set methods across the sky and that the difference found is not likely to be a source of the training set being restricted in area. This is promising for future wide-field photometric redshift surveys such as the Dark Energy Survey, PanStarrs, Euclid and JDEM."
    },
    "0812/0812.0003_arXiv.txt": {
        "abstract": "We use semi-analytic modeling on top of the Millennium simulation to study the joint formation of galaxies and their embedded supermassive black holes. Our goal is to test scenarios in which black hole accretion and quasar activity are triggered by galaxy mergers, and to constrain different models for the lightcurves associated with individual quasar events. In the present work we focus on studying the spatial distribution of simulated quasars.  At all luminosities, we find that the simulated quasar two-point correlation function is fit well by a single power-law in the range $0.5 \\la r \\la 20\\, h^{-1} \\rm{Mpc}$, but its normalization is a strong function of redshift.  When we select only quasars with luminosities within the range typically accessible by today's quasar surveys, their clustering strength depends only weakly on luminosity, in agreement with observations.  This holds independently of the assumed lightcurve model, since bright quasars are black holes accreting close to the Eddington limit, and are hosted by dark matter haloes with a narrow mass range of a few $10^{12} h^{-1} \\rm{M}_{\\odot}$. Therefore the clustering of bright quasars cannot be used to disentangle lightcurve models, but such a discrimination would become possible if the observational samples can be pushed to significantly fainter limits. Overall, our clustering results for the simulated quasar population agree rather well with observations, lending support to the conjecture that galaxy mergers could be the main physical process responsible for triggering black hole accretion and quasar activity. ",
        "introduction": "At the beginning of this century, very bright quasars powered by black holes (BHs) with masses of the order of $10^{9} \\rm{M}_{\\odot}$ were discovered at redshifts up to $z \\sim 6$ \\citep{fan00,fan01}.  At the same time, X-ray observations showed that the space density of Active Galactic Nuclei (AGN) peaks at $z \\sim 2-3$, and AGN with high X-ray luminosities are more common at higher redshift with respect to their low-luminosity counterpart \\citep{steffen03, cowie03, cattaneo03, ueda03, hasinger05}.  \\citet{heckman04}, using optical data from SDSS, found that at low redshift only BHs with a mass $\\la 10^{7} \\rm{M}_{\\odot}$ are actively growing.  Combined, these observations suggest that supermassive black holes (SMBHs) grow ``anti-hierarchically'': the more massive BHs were already in place at high redshift, and since then the accretion activity has shifted to smaller scales.  Understanding how this evolution of BH growth relates to cosmic structure formation, how BH accretion depends on the environment, and how BHs interact with their host galaxies, have become central questions of cosmology that need to be answered for a full understanding of galaxy formation. In fact, it has become clear in recent years that BHs and galaxies are linked and mutually influence each other. This co-evolution has been explored in recent years through analytic, semi-analytic and fully numerical approaches in numerous studies \\citep[e.g.,][]{silk98, kauffmann00, merloni04, diMatteo05, cattaneo05, croton06a, monaco07, malbon07, marulli06, marulli07}.  Starting from \\citet{soltan82}, many papers have combined the present-day BH mass function with the AGN luminosity density of quasars at various redshifts to conclude that most of the mass in todays BHs must have been accumulated during phases of bright AGN activity \\citep[see also][]{yu02,elvis02,marconi04,merloni08}. The duration of these highly-efficient accretion phases could range from a few $10^{7} \\rm{yr}$ \\citep{yu02} up to $10^{8} \\rm{yr}$ \\citep{marconi04}, values that strongly depend on the BH mass range considered and on the assumed radiation efficiency $\\epsilon$. In fact, the precise value of this quasar lifetime is still an open question \\citep{martini04}.  Estimates of the duration of individual accretion events using, for example, the {\\em proximity effect} \\citep{carswell82,bajtlik88}, have suggested lifetimes of the order of $1 ~ \\rm{Myr}$ \\citep{kirkman08}.  \\citet{haiman01} and \\citet{martini01} suggested to use quasar clustering to obtain estimates of the quasar lifetime \\citep[see also the seminal work of ][]{cole89}. The reasoning behind this conjecture is simple: if quasars are strongly clustered, their hosts must be rare objects, and therefore they must also be long events in order to account for the total quasar luminosity density observed. If, on the other hand, their clustering is comparable to the clustering of small dark matter haloes, their hosts must be much more common, and their luminous phases must therefore have short duration.  The advent of wide-field surveys like SDSS and 2dF quasi-stellar object (2dFQSO) \\citep{york00,croom04} with their observation of thousands of quasars has allowed a detailed investigation of the clustering properties of accreting BHs. \\citet{croom05} and \\citet{porciani04} calculated the correlation function of quasars observed in 2dF in the redshift range $0.5 \\la z \\la 2$.  Both groups found that the clustering strength is an increasing function of redshift, but that it does not depend significantly on quasar luminosity.  The inferred values of the bias would suggest that quasars of the observed luminosities are hosted by haloes of a few $10^{12} h^{-1} \\rm{M}_{\\odot}$, which remains approximately constant with redshift, since haloes of a fixed mass are progressively more clustered towards higher redshift \\citep[see also][]{grazian04}.  Following the approach of \\citet{haiman01} and \\citet{martini01}, the estimated quasar lifetime would be a few $10^{7} \\rm{yr}$, reaching $\\sim 10^{8} \\rm{yr}$ at the highest observed redshifts.  More recent studies on larger samples and at different redshifts have confirmed these results \\citep{shen07, myers07, coil07, daAngela08,padmanabhan08, ross09}. However, the magnitude range covered by these surveys is typically quite narrow, and this may explain the lack of evidence for a significant dependence of clustering on luminosity. When \\citet{shen08b} analyze the clustering of the $10 \\%$ brightest objects of their sample, they find that these quasars have a higher bias compared to the full sample.  Using hydrodynamical simulations of isolated galaxy mergers \\citep{springel05a, diMatteo05}, \\citet{hopkins05} studied the luminosity distribution of accreting BHs, whose activity is triggered by the merger event. \\citet{hopkins05} found that the luminosity distribution of the simulated AGN is equivalent to a highly efficient accretion phase (with very high Eddington ratios), followed by a decaying phase where AGN spend most of their life. During this extended period, they would appear as faint AGN, even though they may, in fact, contain quite massive BHs.  Based on these results, \\citet{lidz06} explored the dependence of quasar clustering on luminosity, using an analytic approach to connect quasars, black hole masses and halo masses.  In a quasar model in which the bright end of the luminosity function is populated by BHs accreting close to their peak luminosity, and the faint end is mainly populated by BHs accreting at low Eddington ratio, there should be no strong dependence of clustering on quasar luminosity, i.e., bright and faint AGN should actually be the same type of objects, but seen in different stages of their evolution. They should therefore be hosted by dark matter haloes of similar masses and hence exhibit similar clustering properties. Assuming a relation between the quasar B-band peak luminosity and the mass of the host haloes, \\citet{lidz06} tested this prediction, and indeed found that only a narrow range of halo masses should host active quasars, with a median characteristic mass of $M_{\\rm halo} \\sim 1.3 \\times 10^{13} \\rm{M}_{\\odot}$. As pointed out by the authors, only future surveys that will be able to observe the faint quasars in their quiescent stage will be able to test this picture, and to rule out the alternative hypothesis of luminosity-dependent quasar clustering.  In the present work we explore the properties of quasar clustering using a semi-analytic model for galaxy formation and BH accretion developed on the outputs of the Millennium Simulation \\citep{springel05b}. Compared to other theoretical work, we do not have to make assumptions about the halo population hosting BHs nor about the relation between the halo mass and the quasar luminosity (or BH mass), since they are a natural outcome of the simulation of the galaxy formation process. However, we have to make assumptions about the physics of BH accretion, and what triggers AGN activity. In this work we are especially interested in testing the assumption that galaxy mergers are the primary physical mechanism responsible for triggering accretion onto massive BHs. To this end we explore the simulation predictions for quasar clustering and the quasar luminosity function obtained with a pure Eddington-limited lifetime model and a model that includes a low-luminosity accretion mode as described by \\citet{hopkins05} \\citep[see also][hereafter Paper I]{marulli08}.  After an introduction to our methodology and a review of some basic properties of our simulated AGN population (Section 2), we show their clustering properties in Section 3, where we also compare our results with the most recent observational optical data available. In Section 4, we show the relation between luminous BHs, quiet BHs and their host haloes. Finally, we summarize and discuss our results in Section 5. ",
        "conclusions": "In this series of papers we investigate semi-analytic models for BH accretion and quasar emission in the context of a comprehensive galaxy formation model developed for the Millennium Simulation.  The physical scenario for BH growth we study is based on the model for BH accretion from \\citet{kauffmann00}, as revised by \\citet{croton06a}, which assumes that galaxy mergers are the primary physical mechanism responsible for efficiently feeding central BHs. In Paper I \\citep{marulli08} we used the most recent observations of the local BH population to test basic predictions of the model for the local BH demographics, testing also different theoretical models for the quasar lifetime with goal to reproduce the observed quasar luminosity function. We found an overall good agreement between the predicted and the observed BH properties, and that the faint-end of the observed luminosity function can be better reproduced when a quasar lightcurve model is adopted that includes long quiescent accretion after an Eddington-limited accretion phase.  In the present work we used the spatial distribution of active BHs as a further test of our model for BH accretion. Throughout the paper, we compared the results obtained adopting two different theoretical models for the quasar lifetime: pure Eddington-limited accretion (Mod I), and a model in which Eddington-limited accretion is followed by a long, weak accretion phase (Mod II), modeled after \\citet{hopkins05}.  The main difference between the predictions of the two models is in the faint-end of the luminosity function. The long low-luminosity accretion phase allowed by Mod II gives rise to a large population of massive BHs that at low redshifts are accreting at low Eddington ratios, in agreement with the observational results of, for example, \\citet{heckman04} and \\citet{netzer07}, who found that in the local universe only BHs with mass $\\la 10^{7} \\rm{M}_{\\odot}$ are experiencing high-efficient accretion. As also recently pointed out by \\citet{hopkins08}, it is only by studying the properties of the faint AGN population that the quiescent phase described by \\citet{hopkins05} can be tested.  Independent of the model adopted for the lightcurve, the two-point correlation function of our simulated AGN can be approximated by a single power-law in the range $0.5 \\la r \\la 20~ h^{-1} \\rm{Mpc}$.  The bias between AGN and the dark matter is a strong function of redshift, but, at a given epoch, it is approximately constant in the range $1.0 \\la r \\la 20 ~h^{-1} \\rm{Mpc}$.  As expected, the correlation lengths of AGN obtained with Mod I or Mod II differ only for the faint population: the correlation length of faint AGN obtained with Mod II is consistent with the correlation length of $10^{12}-10^{13} h^{-1} \\rm{M}_{\\odot}$ haloes, whereas faint AGN obtained with Mod I exhibit the same clustering as $10^{11}-10^{12} h^{-1} \\rm{M}_{\\odot}$ haloes.  Recent results from optical quasar surveys like SDSS and 2dFQSO have not found evidence for a strong dependence of clustering on luminosity \\citep[e.g.,]{porciani04, croom05, myers07, daAngela08}, except for \\citet{shen08b} who detect an excess of clustering for their $10\\%$ brightest quasars.  Our results are consistent with these observations if we consider only quasars with an intrinsic luminosity within the range probed by these surveys. However, if we compare the clustering properties of AGN over a very extended range of luminosity, then the correlation length becomes a moderately strong function of luminosity and the value of the correlation length of the faint population in particularly is seen to depend on the lightcurve model assumed. The fact that the clustering of the observed quasars does not depend on luminosity could be explained in two ways: quasars of different luminosities are powered by BHs of the same mass that are in different stages of their evolution, and/or the typical mass of haloes hosting quasars is approximately constant for the luminosity range probed by observations. From our results the second hypothesis seems to be clearly favoured. The mass range of haloes hosting $L_{\\ast}$ quasars is narrow enough that a significant luminosity dependence of clustering cannot be detected with the current observational samples, independent of the lightcurve model.  We also directly compared the clustering of our $L_{\\ast}$ quasars with the most recent observational data, and found very good agreement. Since quasars at these luminosities are objects very close to their peak luminosity, and therefore correspond to objects accreting at high Eddington ratios, we cannot, however, use these results as a sensitive test of our lightcurve models. Nevertheless, the good agreement with observations indicates that our merger-triggered BH accretion model predicts a spatial distribution of quasar that is consistent with observations. This non-trivial outcome can be viewed as a further success of the hierarchical galaxy formation paradigm, and the cold dark matter hypothesis.  We note that a similar result for the luminosity dependence of AGN clustering has been found in \\citet{marulli09}, who analyzed  mock AGN Chandra catalogues constructed with the same semi-analytic model adopted in  this work. Furthermore, \\citet{thacker08} have recently found very similar results modeling the AGN spatial properties in an hydrodynamical simulation.  In future work we will compare the merger-triggered quasar model with alternative suggestions for the physical triggering mechanism of quasar activity, such as disk-instabilities occuring in isolated galaxies. We expect that quasar clustering statistics can here be a potentially powerful discriminant to further constrain the viable physical models for the evolution of supermassive black holes, and their co-evolution with galaxies."
    },
    "0812/0812.1097_arXiv.txt": {
        "abstract": "{We study the clustering properties of galaxy clusters expected to be observed by various forthcoming surveys both in the X-ray and sub-mm regimes by the thermal Sunyaev-Zel'dovich effect. Several different background cosmological models are assumed, including the concordance $\\Lambda$CDM and various cosmologies with dynamical evolution of the dark energy. Particular attention is paid to models with a significant contribution of dark energy at early times which affects the process of structure formation. Past light cone and selection effects in cluster catalogs are carefully modeled by realistic scaling relations between cluster mass and observables and by properly taking into account the selection functions of the different instruments. The results show that early dark-energy models are expected to produce significantly lower values of effective bias and both spatial and angular correlation amplitudes with respect to the standard $\\Lambda$CDM model. Among the cluster catalogs studied in this work, it turns out that those based on \\emph{eRosita}, \\emph{Planck}, and South Pole Telescope observations are the most promising for distinguishing between various dark-energy models.} ",
        "introduction": "One of the main quests of contemporary astrophysics is the determination of the nature of dark energy, the mysterious component of the cosmic fluid responsible for the accelerated expansion of the Universe. While the evidence of its presence has became compelling in the past decade \\citep{AS06.1,RI07.1,DU08.1,KO08.1,RU08.1,KI08.1}, its nature remains entirely unexplained. In particular, while virtually all the present observational evidence is in concordance with a cosmological-constant interpretation of the dark energy, its possible dynamical evolution is not well constrained. The detection of this time evolution would hint at dark energy being different from vacuum energy, and would call for some more general explanation, such as minimally coupled scalar fields \\citep{RA88.1,WE88.2,BR00.1}. Since, in this standard generalization, dark energy is not supposed to clump on the scales of the largest cosmic structures, the only way its nature can be unveiled is by studying the expansion history of the Universe. The expansion rate of the Universe as a function of cosmic time in turn affects the process of structure formation, and consequently the many observable properties of cosmic structures that are accessible to observations.  The most immediate effect of cosmology on cosmic structures is on the number counts of objects, especially the most massive and extreme ones such as galaxy clusters. A generic dynamical dark-energy model that postulates a dark-energy density increasing with redshift will necessarily imply an earlier structure formation than in cosmological-constant models (if the linear amplitude of density fluctuations at present is fixed), and thus a higher abundance of objects at any time.  An alternative channel for the detection of possible effects of dynamical-dark energy is the study of clustering properties of cosmic structures. Models predicting a higher abundance of objects imply that high-mass clusters are less exceptional, and as a consequence the linear bias with respect to the underlying dark-matter density-field should be reduced. The linear bias, the structure abundance, and the linear correlation function of density fluctuations all affect the angular and spatial correlation functions that are observed in cluster catalogs, and all depend on the behavior of dark energy. We are therefore justified in exploring the effect of different quintessence models on the clustering properties of massive galaxy clusters, and understanding whether differences between them could be detected significantly in cluster catalogs produced by forthcoming experiments. This is the purpose of the present work.  We focus on blind cluster surveys based on the X-ray emission of the hot intra-cluster plasma and on the spectral distortion of the CMB radiation produced by the thermal Sunyaev-Zel'dovich (\\citealt{SU72.1}, SZ henceforth) effect. The cosmological models that we address range from the concordance cosmological-constant $\\Lambda$CDM cosmogony to early dark-energy models, to intermediate models with a constant equation-of-state parameter for dark energy $w_\\mathrm{x} \\ne -1$ or with a gentle evolution in $w_\\mathrm{x}$ with time. \\cite{SA07.1} performed a simple preliminary study of the two-point angular correlation function predicted to be observed by \\emph{Planck} in models with early quintessence. As will become evident from the results shown in this paper, our findings are qualitatively consistent with those, while a quantitative comparison is not directly possible because of the different catalog definitions. Observational results about the clustering properties of galaxy clusters measured in optical and infrared catalogs can be found in \\cite{BR07.1}, \\cite{PA08.1} and \\cite{ES08.1}.  This paper is structured as follows. In Sect.~\\ref{sct:cosmo}, a brief overview of the different cosmological models employed in this paper is presented, with a description of their main features and a summary of the various cosmological parameters. In Sect.~\\ref{sct:clus}, we review the formalism used to describe clustering of galaxy clusters in the past light cone of the observer. In Sect.~\\ref{sct:survey}, we detail the properties of the X-ray and SZ surveys analyzed in the present work and summarize the scaling laws used in linking the mass and redshift of objects to their observable properties in Sect.~\\ref{sct:scaling}. In Sect.~\\ref{sct:cat}, we describe the properties of the cluster catalogs obtained therefrom. In Sect.~\\ref{sct:res}, we present our results on the spatial and the angular correlation functions, and in Sect.~\\ref{sct:sum}, we summarize our conclusions. We shall use throughout the \\cite{SH02.1} prescription for the computation of both the cluster mass function and the linear bias. ",
        "conclusions": "\\label{sct:sum}  We have studied the clustering properties of galaxy clusters in various cosmological models with dynamical and non-dynamical evolution in the dark-energy component. In addition to the concordance $\\Lambda$CDM cosmology, we addressed a model with a constant equation-of-state parameter for dark energy, $w_\\mathrm{x} = -0.8$, a dynamical-dark energy model with the $w_\\mathrm{x}(z)$ parametrization proposed by \\cite{KO08.1}, and four models with a non-negligible amount of early dark energy. Cosmological models with dynamical dark energy are generically expected to form structures earlier, affecting both the abundance of massive galaxy clusters and their spatial distribution and clustering properties, this latter property having been the issue here.  To predict forthcoming observations, we computed the effective bias as a function of redshift and observed spatial and angular correlation functions as a function of (physical or apparent) separation that are expected to be measured in cluster catalogs produced by planned blind surveys both in the X-ray and in the sub-mm regimes by means of the thermal SZ effect. The X-ray surveys that we considered in this work are the \\emph{eRosita} wide survey, which is described in detail in the related mission definition document, and the XMM cluster survey, based on existing pointings of the XMM satellite. For the SZ surveys, we focused on the South Pole Telescope, the Atacama Cosmology Telescope, and the \\emph{Planck} all-sky survey of the Cosmic Microwave Background.  For computing the clustering properties of objects contained in each catalog as a function of cosmology, we employed a well-established formalism that takes past-light cone and selection effects on the cluster sample into account. To link the limiting flux of each survey to the minimum mass that enters the respective catalog at a given redshift, we adopted realistic scaling relations, based both on observations and numerical simulations, between the mass of a cluster and its X-ray luminosity or SZ flux density, or its integrated Compton-$y$ parameter. It turns out that the minimum mass entering a cluster catalog depends not only on the instrument considered, but also on the cosmological model adopted. This is so because the scaling relations mentioned above depend on the underlying world model through the expansion rate and cosmological distances.  As one could naively expect, the number of objects entering a given catalog at a given redshift depends heavily on the cosmology, the highest cluster abundances being present in models with EDE. This is caused by the well known higher mass function displayed by these models and the lower minimum mass entering the catalogs. The SPT catalog is the most extended in redshift due to the low limiting SZ flux density ($5$ mJy), while distributions for XCS and \\emph{Planck} are the most limited in redshift, because of the quite shallow limiting X-ray flux of the former and large beam of the latter, which tends to dilute the signal.  For all catalogs, the first two models with early quintessence, EDE1 and EDE2, display the smallest effective bias at all redshifts. This is due to the high abundance of structures present in these models. On the other hand, the concordance $\\Lambda$CDM model always displays the highest effective bias, for all catalogs and at all redshifts, while the other models lie somewhere in-between the two. Because of the width of their extent in redshift, SPT and ACT catalogs show the flattest trend for the effective bias, especially SPT, for which $b_\\mathrm{eff}(z)$ is at most $\\sim 10$ at $z = 3$. Other catalogs, particularly those based upon \\emph{Planck} and XCS have a much steeper trend, the effective bias reaching considerably higher values already at $z \\sim 2$.  Concerning the spatial correlation function, all the EDE models almost coincide on all scales, having a correlation amplitude smaller than that for the models $\\Lambda$CDM, K08 and constant $w_\\mathrm{x} = -0.8$, which are also very similar to each other. The largest difference is displayed by the catalogs produced with \\emph{eRosita} and SPT. In the former, the relative errors are also smallest, thus making \\emph{eRosita} the most promising instrument for the detection of EDE through the spatial correlation function, if the full cluster catalog is to be used. Detection might also be possible for SPT and \\emph{Planck} at intermediate scales $r \\sim r_0$, while it is completely out of question for XCS and ACT, whose error bars are too large. This is due to the shallow limiting fluxes for both of them, which is not adequately compensated for by the area covered (as is the case for \\emph{Planck}). The situation is similar for the angular correlation function.  There the smallest error bars are produced by \\emph{eRosita} as well, while the EDE4 model produces the smallest correlation amplitude for all the catalogs.  To explore the redshift evolution in the spatial and angular correlation functions, we also performed different cuts in redshift of the various cluster catalogs, in a way that the number of pairs of objects in each bin would remain approximately the same. For the number of spatial pairs, we adopted the double binning $z \\le 0.1$ and $z > 0.1$, with the exception of SPT which allows the inclusion of a third bin, $z > 0.3$, with the modification of the second bin to $0.1 < z \\le 0.3$. This same binning was also employed for the number of angular pairs.  The result of this analysis is that the spatial correlation function increases with increasing redshift, the correlation length for \\emph{eRosita} growing by $\\sim 80\\%$ between the low- and high-redshift bins. The relative errors show that the increment in the correlation function is significant for \\emph{eRosita}, SPT and \\emph{Planck}, for all cosmological models considered in this work. Also, comparing the amplitude of the ratio between spatial correlation functions in EDE models and $\\Lambda$CDM-like cosmologies, with the size of the error bars indicates that it is better to focus on the low-redshift cluster subsample in order to maximize significant differences between the concordance model and the EDE models in the \\emph{eRosita} and \\emph{Planck} catalogs. The high-redshift bin, $z > 0.3$, is better for the SPT catalog, giving results compatible with those from the full catalog. As for the angular correlation function, it tends to decrease with increasing redshift, a trend that is significant for all models and catalogs except the usual XCS and ACT. In the absence of errors, in the high-redshift bin there would be the highest chance of distinguishing a $\\Lambda$CDM model from models with an early quintessence contribution, since there the ratios of correlation amplitudes are at their highest. However, the errorbars are also large there, so that \\emph{eRosita} is the only survey expected to permit significant detection of deviations from the concordance model at all redshifts. In general, if one is interested in optimizing the differences between angular correlation functions measured in different models, it does not pay to subdivide catalogs according to redshift.  We also found that the same is true when cluster catalogs are binned according to the observable used to define them, an approach more directly motivated from the observational point of view. Specifically, binning the SPT catalog according to the flux density always produces ratios of angular correlation functions in different cosmologies that are similar to those for the full catalog, while the error bars are always larger.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=\\hsize]{Figures/angular_z_gt_0_3}\\hfill \\end{center} \\caption{As Fig.~\\ref{fig:angular}, but only for catalogs restricted to $z > 0.3$. Only three cosmologies are shown here for simplicity.} \\label{fig:angular_z_gt} \\end{figure}  Before concluding, two notes of caution are in order. First of all, some discussion arose in the literature on whether the semi-analytic calculations performed in \\cite{BA06.1} on the spherical collapse model in cosmologies with dynamical-dark energy are indeed correct, and how accurately they are reproduced by $N$-body simulations \\citep{FR08.1,FR08.2,GR08.1}. While the discussion is not yet settled, \\cite{SA07.1} performed the same calculations as \\cite{BA06.1} using a different approach, and found results that are perfectly consistent with the latter. In addition, \\cite{SC08.1} used the same approach as \\cite{BA06.1} to successfully evaluate the spherical collapse behavior in a modified gravity scenario. Hence, we are at least reasonably confident that the approach followed by \\cite{BA06.1}, also employed in this work, is fundamentally correct. Moreover, numerical simulations using the correct early-time behavior of the growth factor for scaling the initial conditions yield results other than those of \\cite{GR08.1} and \\cite{FR08.1}, which tend towards the expectation from the analytic work of \\cite{BA06.1}. Although precise direct integrations of the spherical collapse equations are difficult, further work avoiding approximations has so far confirmed our earlier results (Pace et al., in preparation). Even though definitive conclusions are not reached yet, these facts seem to justify our choice.  Secondly, estimates of the number counts of galaxy clusters detected by \\emph{Planck} seem to fall substantially below the estimates of \\cite{SC07.1}, and the redshift distribution is apparently shallower than that represented in Fig.~\\ref{fig:dist} (\\emph{Planck} SZ Challenge (in preparation), see also \\citealt{LE08.1}). While definitive new limits for \\emph{Planck} are not yet available, this part of the results should be read with caution. In particular, a decrease in the number of objects in the \\emph{Planck} catalog would produce an increase in the relative errors, which scale as the inverse of the square root of the number of pairs of objects.  The present work shows that, while mild modifications to the redshift evolution in the equation-of-state parameter for dark energy $w_\\mathrm{x}(z)$ have a negligible impact on the clustering properties of galaxy clusters, more exotic models such as early dark-energy cosmologies can change the effective bias of collapsed objects significantly, and thus also the spatial and angular correlation amplitudes of galaxy clusters. We have shown that at least some of the forthcoming blind surveys both in the X-ray and sub-mm regimes will be able to distinguish significantly between these models and more generally place constraints on the time evolution in the dark-energy density.  In general, we expect object clustering to be less effective than e.g.,\\ direct abundance data in determining the cosmological model. For instance, in the \\emph{eRosita} catalog the number of clusters at $z \\gtrsim 1$ increases by about one order of magnitude between the $\\Lambda$CDM and the EDE4 models. This variation is much larger than the corresponding variation in the correlation functions, while the size of the relative errors, assuming Poisson statistics, is comparable. Object counting at high-redshift is also expected to be capable of distinguishing cosmological models with a gentle variation in dark-energy density from standard cosmology. The number of high-$z$ clusters in the \\emph{eRosita} catalog is higher by a factor $\\sim 2.5$ in the model with constant $w_\\mathrm{x} = -0.8$ than for $\\Lambda$CDM.  Nevertheless, the results of this work show that clustering of massive clusters by itself remains a fundamental channel to unravel the effect of the expansion history of the Universe on the process of structure formation, although employing this information in addition to simple object number counts is certainly an interesting issue to be explored.  {\\small"
    },
    "0812/0812.2788_arXiv.txt": {
        "abstract": "\\chandra\\ data of the X-ray source \\xf47\\  were obtained in the ACIS Survey of \\m33\\ (\\cm33) in 2006.  During one of the observations, the source varied from a high state to a low state and back, in two other observations it varied from a low state to respectively intermediate states. These transitions are interpreted as eclipse ingress and egresses of a compact object in a high mass X-ray binary system. The phase of mid eclipse is given by HJD~245\\,3997.476$\\pm$0.006, the eclipse half angle is $30.6\\degr\\pm1.2\\degr$. Adding \\xmm\\ observations of \\xf47\\ in 2001 we determine the binary period to be 1.732479$\\pm$0.000027~d. This period is also consistent with \\ro\\ HRI observations of the source in 1994. No short term periodicity compatible with a rotation period of the compact object is detected. There are indications for a long term variability similar to that detected for Her X--1. During the high state the spectrum of the source is hard (power law spectrum with photon index $\\sim$0.85) with an unabsorbed luminosity of 2\\ergs{37} (0.2--4.5 keV). We identify as an optical counterpart a $V \\sim 21.0~{\\rm mag}$ star with $T_{\\rm eff} > 19000~{\\rm K}$, $\\log(g) > 2.5$. CFHT optical light curves for this star show an ellipsoidal variation with the same period as the X-ray light curve. The optical light curve together with the X-ray eclipse can be modeled by a compact object with a mass consistent with a neutron star or a black hole in a high mass X-ray binary. However, the hard power law X-ray spectrum favors a neutron star as the compact object in this second eclipsing X-ray binary in \\m33. Assuming a neutron star with a canonical mass of 1.4 \\mbox{$\\:M_{\\sun}$} and the best fit companion temperature of 33000~K, a system inclination $i = 72\\degr$ and a companion mass of 10.9 \\mbox{$\\:M_{\\sun}$} are implied. ",
        "introduction": "Only a few eclipsing X-ray binaries (XRBs) have been detected in the Milky Way and other nearby galaxies. However, eclipsing XRBs are very interesting sources as they provide (when X-ray and optical data are included) the most reliable detailed mass measurements for compact stellar remnants possible \\citep[with the exception of the few known radio pulsars in binaries, including double neutron stars, see e.g.][]{2006csxs.book....1P}.  Of specific interest is the first eclipsing XRB in \\m33, \\mx7\\ (hereafter \\xe7), which showed properties that were not known from XRBs in the Galaxy or the Magellanic Clouds. \\xe7\\ was already detected as a variable by the \\ein\\ observatory with a luminosity at maximum in excess of \\oergs{38} \\citep[][]{1981ApJ...246L..61L,1983ApJ...275..571M} and it stayed active in all following observations. Its variability was explained through the model of an eclipsing XRB with an orbital period of 3.45~d and an eclipse duration of $\\sim$0.4~d based on \\ro\\ and \\asca\\ data \\citep[][]{1997AJ....113..618L,1999MNRAS.302..731D}. With the help of the well-constrained \\chandra\\ position of \\xe7, \\citet[][]{2004A&A...413..879P}  identified a B0I to O7I star of 18.89 mag in $V$ as optical counterpart: this star shows the ellipsoidal light curve of a high mass XRB (HMXB) with the \\xe7\\ binary period. They argued that the compact object in the system is a black hole based on the mass of the compact object derived from orbital parameters and the optical companion mass, the lack of pulsations, and analysis of the extracted X-ray spectrum. Those authors concluded that \\xe7\\ would be the first eclipsing black hole HMXB.  The black hole nature of \\xe7\\ was firmly established by \\citet[][]{2006ApJ...646..420P} using observations of the \\chandra\\ Advanced CCD Imaging Spectrometer (ACIS) survey of \\m33\\ \\citep[\\cm33, see][]{2008ApJS..174..366P} which accumulated in seven ACIS-I pointing directions a total exposure of 200~ks each. During several of these pointings \\xe7\\ was in the field of view. The \\cm33\\ measurements of \\xe7\\ resolved for the first time the eclipse ingress and egress as well as constrained the light curve of \\xe7\\ for binary phases around eclipse. In addition, \\citet[][]{2006ApJ...646..420P} identified \\xe7\\  on archival HST WFPC2 images. Through detailed modeling of the optical light curve, the observed X-ray eclipse and improved parameters for the companion star, a lower mass limit for the compact object in the system of 9~\\MSOL\\ \\citep[see also][]{2007A&A...462.1091S} has been obtained. The black hole mass was further constrained as 15.65~\\MSOL\\ with the help of a radial velocity curve based on Gemini North spectra and careful modeling \\citep[][]{2007Natur.449..872O}. This -- at the time -- made \\xe7\\ the stellar black hole with the highest mass that had been determined with high precision.  In this paper we report on the detection and optical identification of the second eclipsing XRB in \\m33\\ within the \\cm33\\ project. The source was already detected in the \\xmm\\ survey of \\m33\\ of \\citet[][hereafter PMH2004]{2004A&A...426...11P} and is source  no. 47 in their catalogue (\\xf47, hereafter \\x47). During the \\xmm\\ observations it showed strong time variability and was classified as a transient XRB candidate by \\citet[][]{2006A&A...448.1247M}. It was first reported in the \\m33\\ \\ro\\ catalogue by \\citet[][]{2001A&A...373..438H} as source 19. \\x47\\ is listed as  no. 8 in the ``first look\" \\cm33\\ catalogue \\citep[][]{2008ApJS..174..366P}. In several \\cm33\\ observations the source was detected far off axis with strong time variability which -- in a preliminary analysis -- could be explained as the signature of a 1.7 d eclipsing XRB \\citep[][]{2006ATel..905....1P}. A massive star from the Local Group survey catalogue of \\m33\\ stars \\citep[][]{2006AJ....131.2478M} could be identified as the optical counterpart of the X-ray source. In a re-analysis of the deep 3.6m Canada-France-Hawaii Telescope (CFHT) photometric survey images of \\m33\\ \\citep[][]{2006MNRAS.371.1405H}, the proposed optical companion showed the ellipsoidal light curve of a HMXB in \\m33\\ \\citep[][]{2006ATel..913....1S}. A coherent analysis of the \\chandra\\ and \\xmm\\ X-ray and optical data allows us to further constrain the system parameters. Our follow-up analysis of \\x47\\ in the \\ro\\ data also indicates time variability consistent with the \\chandra\\ and \\xmm\\ ephemeris. ",
        "conclusions": "The time variability of \\x47\\  in the \\cm33\\ observations can be explained with a 1.7 d eclipsing X-ray binary with an eclipse half angle of $30.6\\degr$. The identification with an OB star in \\m33\\ that shows an ellipsoidal light curve, clearly establishes the system as the second HMXB in \\m33. We do not detect pulsations in the X-ray flux and therefore the compact object in the system  could be a neutron star or a black hole. The hard X-ray spectrum, however, may indicate a neutron star system \\citep[see e.g.][]{1983ApJ...270..711W}.  We can not exclude that the long term variability of the system in the \\chandra, \\xmm\\ and \\ro\\ observations can be of irregular nature. However, it also could reflect a long term periodicity similar to those known in Her X--1, LMC X--4 or SMC X--1. The \\chandra\\ bright state observations (ObsID 6387 and 6385) are separated by 84 days. With our sparse unsystematic monitoring, we can not decide if the source had another on-state in between as would be expected by extrapolating from the 35 d period of Her X--1, a system with a similar orbital period.  Inspecting the available \\citet[][]{2006AJ....131.2478M} colour and emission-line images we find that the object is located in a region with few (bright) stars. Also, there are no close-by emission nebulae, indicating that the object is older than \\mx7\\ which is still located in a \\HII\\ region.  \\x47\\ was serendipitously detected far off axis in the \\cm33\\ observations. Only with the help of dedicated deep observations with good time resolution it will be possible to characterize long term variability of \\x47\\ and to search for the expected pulsation period below 50 seconds."
    },
    "0812/0812.0602_arXiv.txt": {
        "abstract": "Simulations predict that shocks from large-scale structure formation and galactic winds have reduced the fraction of baryons in the warm, photoionized phase (the \\Lya\\ forest) from nearly 100\\% in the early universe to less than 50\\% today.  Some of the remaining baryons are predicted to lie in the warm-hot ionized medium (WHIM) phase at $T=10^5-10^7$~K, but the quantity remains a highly tunable parameter of the models.  Modern UV spectrographs have provided unprecedented access to both the \\Lya\\ forest and potential WHIM tracers at $z\\sim0$, and several independent groups have constructed large catalogs of far-UV IGM absorbers along $\\sim30$ AGN sight lines.  There is general agreement between the surveys that the warm, photoionized phase makes up $\\sim30$\\% of the baryon budget at $z\\sim0$.  Another $\\sim$10\\% can be accounted for in collapsed structures (stars, galaxies, etc.).  However, interpretation of the $\\sim100$ high-ion (O\\,VI, etc) absorbers at $z<0.5$ is more controversial.  These species are readily created in the shocks expected to exist in the IGM, but they can also be created by photoionization and thus not represent WHIM material.  Given several pieces of observational evidence and theoretical expectations, I argue that {\\it most} of the observed O\\,VI absorbers represent shocked gas at $T\\sim10^{5.5}$~K rather than photoionized gas at $T<10^{4.5}$~K, and they are consequently valid tracers of the WHIM phase.  Under this assumption, enriched gas at $T=10^{5-6}$~K can account for $\\sim$10\\% of the baryon budget at $z<0.5$, but this value may increase when bias and incompleteness are taken into account and help close the gap on the 50\\% of the baryons still ``missing''. ",
        "introduction": "One of the major legacies of space-based UV instruments, in particular {\\it HST} and {\\it FUSE}, is the opening of the low-redshift intergalactic medium (IGM) to spectroscopic investigation.  Perversely, far more is known about the IGM at high-redshift than in the local universe, simply because many of the key diagnostic lines are rest-frame UV transitions only accessible to optical telescopes at $z\\ge1.6$.  Indeed, before {\\it HST}/FOS observations of low-redshift AGN, there was little interest in the low-$z$ IGM; the observed absorber frequency per unit redshift $\\dndz$ of the \\Lya\\ forest lines in optical spectra drops as $(1+z)^{2.75}$ at redshifts of $z=2-4$ (Lu \\etal\\ 1991), thus predicting that \\Lya\\ systems at $z\\sim0$ would be exceedingly rare.  However, FOS revealed a significant break in the steep power-law slope at $z\\sim1.5$ (Weymann \\etal\\ 1998) due to the evolution of AGN ionizing radiation fields and the growth of large-scale structure.  The local IGM clearly provides astrophysical information of interest to cosmologists, and the sensitive UV spectrographs of the last decade have given us unprecedented opportunities.  Large-scale simulations (e.g., Dav\\'e \\etal\\ 1999, 2001; Cen \\& Ostriker 1999, 2006) predict that the baryon fraction $\\Omega_{\\rm Ly\\alpha}$ in photoionized IGM (traced by \\Lya\\ forest lines) should drop from nearly unity at $z\\sim4$ to a more modest fraction at local times as AGN radiation and large-scale structure shocks heat and ionize the IGM.  Concurrently, the mass in the warm-hot ionized medium (WHIM, $T=10^5-10^7$~K) is predicted to rise from nearly zero to a significant fraction of the total baryon budget at $z<0.5$.  The actual fractions in the different temperature and pressure phases depend sensitively on the methodology and parameters used in the simulations.  Therefore, establishing observational values for $\\Omega_{\\rm Ly\\alpha}$ and \\OmegaWHIM, as well as contributions from other phases at $z\\sim0$, are of crucial importance to our understanding of everything from large-scale structure formation to galactic metal feedback to the distribution of AGN throughout cosmic evolution.  Measuring the photoionized phase is relatively straightforward.  Despite the thinning of the \\Lya\\ forest at low $z$, \\HI\\ absorbers are still the strongest and most common in the IGM.  Several large surveys have established $\\Omega_{\\rm Ly\\alpha}\\approx 0.3\\,\\Omega_b$ at $z<0.5$ (e.g., Penton \\etal\\ 2000, 2004; Lehner \\etal\\ 2007; Danforth \\& Shull 2008).  Probing the WHIM phase is somewhat more complicated.  Hydrogen at these temperatures is almost entirely ionized, with no diagnostic spectral lines in any waveband.  Instead, we must rely on less direct tracers.  The first is to probe the vanishingly small neutral hydrogen fraction ($f_{\\rm HI}<10^{-5}$ at $T\\sim10^5$~K) through thermally-broadened \\Lya\\ lines in UV spectra (e.g., Richter \\etal\\ 2004, Lehner \\etal\\ 2007).  This method is fraught with many observational ambiguities and detection challenges (Danforth \\etal\\ 2009).  Tracing WHIM through X-ray transitions of common, highly-ionized metal ions such as \\OVII\\ and \\OVIII\\ (which peak in abundance at $T>10^6$~K) shows promise, but presents even more observational challenges.  Detections of a handful of \\OVII\\ and \\OVIII\\ absorbers in the IGM have been claimed toward several AGN (Nicastro \\etal\\ 2005, Fang \\etal\\ 2002, 2007), but none has yet been confirmed and much doubt exists as to their veracity (Kaastra \\etal\\ 2006, Rasmussen \\etal\\ 2007).  Current X-ray telescopes simply do not have the sensitivity and spectral resolution to detect the expected IGM absorbers.  \\begin{table}[b] \\begin{tabular}{lcrlc} \\hline \\tablehead{1}{c}{b}{O\\,VI Survey} &\\tablehead{1}{c}{b}{AGN} &\\tablehead{1}{c}{b}{${\\cal N}_{\\rm OVI}$} &\\tablehead{1}{c}{b}{$\\Delta z$} &\\tablehead{1}{c}{b}{$\\dndz$ ($>$30~m\\AA)} \\\\ \\hline Danforth \\& Shull (2005) & 31 & 40 & 2.2 & $17\\pm3$             \\\\ Danforth \\& Shull (2008) & 28 & 83 & 5.2 & $15^{+3}_{-2}$       \\\\ Tripp et al. (2008)      & 16 & 51 & 3.2 & $15.6^{+2.9}_{-2.4}$ \\\\ Thom \\& Chen (2008)      & 16 & 27 & 2.5 & $10.4\\pm2.2$\\tablenote{Blind survey, requiring {\\it both} O\\,VI lines for detection} \\\\ \\hline \\end{tabular} \\label{tab:a} \\end{table}  Far-UV transitions of highly-ionized metal ions (\\OVI\\ \\lam\\lam1032, 1038; \\NV\\ \\lam\\lam1238, 1242; \\NeVIII\\ \\lam\\lam770, 780, etc) offer the best current hope for WHIM measurement.  \\OVI\\ is the most promising of these with $\\sim100$ confirmed IGM detections to date.  \\OVI\\ reaches a peak fractional abundance under collisional ionization equilibrium (CIE, Sutherland \\& Dopita 1993) at $T\\sim300,000$~K, and nothing short of soft X-ray photons ($h\\nu>114$~eV) will produce it non-thermally.  However, as we will see, \\OVI\\ has a number of ambiguities as a WHIM tracer.  The idea of tracing WHIM with \\OVI\\ was pioneered in early work on individual AGN sight lines (Tripp \\etal\\ 2000, Savage \\etal\\ 2002, etc).  Each netted a few \\OVI\\ absorbers over a short redshift pathlength ($\\Delta z<0.6$).  The largest survey to date of absorption lines in the low-$z$ IGM (including \\OVI) was carried out by Danforth \\& Shull (2008, hereafter DS08): a survey of 28 AGN sight lines with both {\\it FUSE} and STIS/E140M data, following up the smaller Danforth \\etal\\ (2005, 2006) survey of 31 {\\it FUSE} sight lines ($z_{\\rm abs}<0.15$).  \\Lya\\ absorbers were identified interactively and checked for corresponding absorption in seven metal ion species (\\OVI, \\NV, \\CIV, \\CIII, \\SiIV, \\SiIII, and \\FeIII).  The survey netted $\\sim650$ \\HI\\ absorbers (many in \\Lya, \\Lyb, etc. enabling curve-of-growth analysis) over a redshift pathlength $\\Delta z=5.2$, with 83 \\OVI\\ absorbers and smaller numbers of detections in other metal ions.  The \\OVI\\ detections showed $\\dndz=15^{+3}_{-2}$ to a rest-frame equivalent width of $W_{1032}>30$ m\\AA, or $\\dndz=40^{+14}_{-8}$ to $W_{1032}>10$ m\\AA.  Two other ambitious IGM \\OVI\\ surveys were published concurrently with DS08 (see Table~1), featuring many of the same AGN sight lines.  Tripp \\etal\\ (2008) searched 16 STIS/E140M datasets blindly for \\OVI\\ and surveyed \\HI\\ systems for \\OVI\\ absorption, resulting in 51 \\OVI\\ detections.  Despite differences in redshift pathlength definition and other systematic effects, they found surprisingly compatible results: $\\dndz=15.6^{+2.9}_{-2.4}$ for $W_{1032}>30$ m\\AA.  Thom \\& Chen (2008) performed a true blind \\OVI\\ survey of 16 sight lines and found 27 \\OVI\\ absorbers for a frequency\\footnote{Thom \\& Chen (2008) required a detection in {\\it both} O\\,VI lines.  Their $\\dndz$ value is correspondingly lower than both DS08 and Tripp \\etal\\ (2008) and can be taken as a lower limit.} $\\dndz=10.4\\pm2.2$. ",
        "conclusions": ""
    },
    "0812/0812.2607_arXiv.txt": {
        "abstract": "Steady, spherically symmetric, adiabatic accretion and wind flows around non-rotating black holes were studied for fully ionized, multi-component fluids, which are described by a relativistic equation of state (EoS). We showed that the polytropic index depends on the temperature as well as on the composition of fluids, so the composition is important to the solutions of the flows. We demonstrated that fluids with different composition can produce dramatically different solutions, even if they have the same sonic point, or they start with the same specific energy or the same temperature. Then, we pointed that the Coulomb relaxation times can be longer than the dynamical time in the problem considered here, and discussed the implication. ",
        "introduction": "It is generally inferred from observations that the matter falling onto black holes is of very high temperature, both in microquasars \\citep{cetal03} as well as in AGNs \\citep{rc00}. The electron temperature around 10$^9$ K and/or the proton temperature around 10$^{12}$ K or more are accepted as typical values within few tens of the Schwarzschild radius, $r_s$, of the central black holes. Moreover, the general theory of relativity demands that the matter crosses the black hole horizon with the speed of light ($c$). In other words, close to black holes, the matter is relativistic in terms of its bulk speed and/or its temperature. On the other hand, at large distances away from black holes, the matter should be non-relativistic.  It is also inferred from observations that the astrophysical jets around black hole candidates have relativistic speeds \\citep{b03}. Since the jets originate from the accreting matter very close to black holes, their base could be very hot. At a few hundred Schwarzschild radii above the disc plane, they can expand to very low temperatures but very high speeds (Lorentz factor ${\\gamma}\\ {\\ga}$ a few). And as the fast moving matter of the jets hits the ambient medium and drastically slows down to form shocks and hot spots, once again the thermal energy increases to relativistic values though the bulk velocity becomes small. Relativistic flows are inferred for gamma-ray bursts (GRBs) too. In the so-called collapsar model scenario \\citep{w93}, the collimated bipolar outflows emerge from deep inside collapsars and propagate into the interstellar medium, producing GRBs and afterglows. In such model, these collimated outflows are supposed to achieve Lorentz factors ${\\gamma}\\ {\\ga}\\ 100$.  It is clear in the above examples that as a fluid flows onto a black hole or away from it, there are one or more transitions from the non-relativistic regime to the relativistic one or vice-versa. It has been shown by quite a few authors that to describe such trans-relativistic fluid, the equation of state (EoS) with a fixed adiabatic index $\\Gamma$ ($=c_p/c_v$, the ratio of specific heats) is inadequate and the relativistically correct EoS \\citep{c38,s57} should be used \\citep[\\egn][]{t48,metal05,rcc06}.  A fluid is said to be thermally relativistic, if its thermal energy is comparable to or greater than its rest mass energy, \\ie if $kT\\ \\ga\\ mc^2$. The thermally non-relativistic regime is $kT \\ll mc^2$. Here, $T$ is the temperature, $k$ is the Boltzmann constant, and $m$ is the mass of the particles that constitute the fluid. So it is not just the temperature that determines a fluid to be thermally relativistic, but it is the ratio, $T/m$, that determines it. Therefore, together with the temperature, the composition of the fluid (\\ie either the fluid is composed of electron-positron pairs, or electrons and protons, or some other combinations) will determine whether the fluid is in the thermally relativistic regime or not.  The study of relativistic flows around compact objects including black holes was started by \\citet{m72}. It was basically recasting the transonic accretion and wind solutions around Newtonian objects obtained by \\citet{b52} into the framework of the general theory of relativity. Since then, a number of authors have addressed the problem of relativistic flows around black holes, each focusing on its various aspects \\citep[\\egn][]{bm76, ft85, c96, d01, d02, metal04, betal06, fk07, metal07}. Barring a few exceptions \\citep[\\egn][]{bm76,metal04}, most of these studies used the EoS with a fixed ${\\Gamma}$, which, as we have noted, is incapable of describing a fluid from infinity to the horizon. \\citet{bm76} for the first time calculated the spherical accretion and wind solutions around Schwarzschild black holes, while using an approximate EoS for the single-component relativistic fluid \\citep{m71}. \\cite{metal04} modified the EoS used by \\citet{bm76} to obtain thermally driven spherical winds with relativistic terminal speeds. However, there has been no extensive study of the effects of fluid composition on the solutions of transonic flows around black holes. We in this paper investigate the effects.  The paper is organized as follows. In the next section, we present the governing equations including the EoS. In section 3, we present the sonic point properties. In section 4, we present the accretion and wind solutions. In section 5, we discuss the validity of our relativistic EoS. Discussion and concluding remarks are presented in the last section. ",
        "conclusions": "In this paper, we have investigated the effects of fluid composition on the solutions of accretion and wind flows onto black holes. In order to elucidate the effects, we have considered a very simple model of spherical flows onto Schwarzschild black holes, and non-conservative processes and magnetic fields have been ignored.  First, we have suggested an approximate EoS for multi-component fluids in equation (5e), and studied the thermal properties of fluids with the EoS. Three temperature ranges have been categorized; for $kT < m_ec^2$, any type of fluids are thermally non-relativistic, for $kT > m_pc^2$, any type of fluids are thermally relativistic, and for $m_ec^2 < kT < m_pc^2$, the degree to which fluids are relativistic is determined by the composition of the fluids as well as the temperature (Figure 1a). Then we have shown that although at the same temperature the $e^- - e^+$ fluid is most relativistic (Figure 1a), at the same sound speed it is least relativistic (Figure 1c), compared to the fluids with protons. It is because whether a fluid is relativistic or not depends on the competition between the thermal energy and the rest mass energy of the fluid.  The thermal properties of fluids carry to the sonic point properties. The sound speed at the sonic point, $a_c$, explicitly depends only on the sonic point location, $r_c$ (it implicitly depends on the specific energy, ${\\cal E}$, and the proton proportion, $\\xi$, through $r_c$). Therefore, comparing the thermodynamic quantities at the same $r_c$ is equivalent to comparing those quantities at the same $a_c$. We have shown that at the same $r_c$, the $e^- - e^+$ fluid is least relativistic, and a fluid with a finite $\\xi$ is most relativistic (Figures 2c and 2d).  Then, we have presented the global solutions of accretion and wind flows for the same $r_c$ but different $\\xi$, for the same ${\\cal E}$ but different $\\xi$, and for the same $T$ at large distances from black holes but different $\\xi$. In all the cases, the flows can be dramatically different, if the composition is different. This asserts that the effects of fluid composition are important in the solutions, and hence, incorporating them properly into the solutions through the EoS is important.  Lastly, we have noted that the EoS in equation (5e) is based on the assumptions that the distribution of the constituent particles is relativistically Maxwellian and the multi-components are of single temperature. However, at the same time, we have pointed out that while the Coulomb relaxation times are normally shorter than the dynamical time far away from black holes, close to black holes they can be longer. It means that close to black holes, the assumptions for the EoS can be potentially invalidated. The implication of it needs to be understood, and we leave further consideration of this issue for future studies."
    },
    "0812/0812.3956_arXiv.txt": {
        "abstract": "We present a sensitive search for the $^3\\!P_1\\rightarrow^3\\!P_0$ ground state fine structure line at 205 microns of ionized nitrogen (\\nii) in one of the highest redshift quasars (J1148+5251 at z=6.42) using the IRAM 30\\,m telescope. The line is not detected at a (3$\\sigma$) depth of 0.47 Jy\\,km\\,s$^{-1}$, corresponding to a \\nii\\ luminosity limit of L$_{\\rm [NII]}<$4.0$\\times10^8$\\,L$_{\\odot}$ and a L$_{\\rm [NII]}$/L$_{\\rm FIR}$ ratio of $<2\\times10^{-5}$. In parallel, we have observed the CO(J=6--5) line in J1148+5251, which is detected at a flux level consistent with earlier interferometric observations.  Using our earlier measurements of the [CII] 158\\,micron line strength, we derive an upper limit for the \\nii/[CII] line luminosity ratio of $\\sim$1/10 in J1148+5251. Our upper limit for the [CII]/\\nii ratio is similar to the value found for our Galaxy and M\\,82 (the only extragalactic system where the \\nii\\ line has been detected to date).  Given the non--detection of the \\nii\\ line we can only speculate whether or not high--z detections are within reach of currently operating observatories. However, \\nii\\ and other fine strucure lines will play a critical role in characterizing the interstellar medium at the highest redshifts (z$>$7) using the Atacama Large Millimeter/submillimeter Array (ALMA), for which the highly excited rotational transitions of CO will be shifted outside the accessible (sub--)millimeter bands. ",
        "introduction": "Forbidden atomic fine structure transitions are important cooling lines of the interstellar medium (ISM). They provide effective cooling in cold regions where allowed atomic transitions can not be excited, thus providing critical diagnostic tools to study the star--forming ISM.  Perhaps the most important cooling line is the forbidden $^2$P$_{3/2}\\rightarrow^2$P$_{1/2}$ fine-structure line of ionized carbon ([CII]) at 158 microns. Other main cooling lines are the oxygen [OI] (63 microns) and [OIII] (at 52 and 88 microns) lines, as well as the nitrogen [NII] lines at 122 and 205 microns (hereafter \\nii). Observations of a combination of these lines have proven to provide key diagnostics regarding the physical properties of the atomic interstellar medium in galaxies (e.g., probes for interstellar UV radiation field, photon dominated regions [PDRs], temperature, gas density and mass, e.g., Petuschowski \\& Bennett 1993, van der Werf 1999). As the ionization potential of carbon is 11.3\\,eV (hydrogen: 13.6\\,eV), [CII] is a tracer for both the neutral atomic and ionized medium, predominantly of PDRs.  The ionization potentials for oxygen and nitrogen, on the other hand, are 13.6\\,eV and 14.53\\,eV, respectively, implying that their forbidden ionized fine structure lines trace the ionized medium only.  The \\nii\\ line, which is the focus of this study, is of particular interest as it has a critical density that is very close to that of [CII], thus potentially providing complementary information on the origin of the [CII] emission (e.g.  Oberst et al.\\ 2006).  Given the poor atmospheric transmission at mid--/far--infrared wavelengths most of of these lines in the local universe can only be studied using airborne or space--based observatories. At sufficiently high redshift, however, some of the lines are shifted to the (sub--)millimeter bands that can be observed from the ground. Especially at redshifts approaching the Epoch of Reionization (z$>$7) these lines provide the only tools by which to study the properties of the star--forming ISM using (sub--)millimeter facilities, as the excited rotational lines of the standard tracer, CO, will be shifted outside the observable (sub--)millimeter bands (Walter \\& Carilli 2008).  Whereas the [CII] line has now been abundantly detected in the local universe (e.g., Stacey et al.\\ 1991, Malhotra et al.\\ 1997, Luhman et al.\\ 1998, Madden et al.\\ 1997), measurements of the \\nii\\ line are scarce.\\footnote{This is due to the fact that [CII] was accessible with the Infrared Space Observatory Long Wavelength Spectrometer (ISO/LWS, Clegg et al.\\ 1996) whereas \\nii\\ was not.} The line was first detected by {\\em FIRAS} aboard {\\em COBE} in the Milky Way (Wright et al.\\ 1991), where a global [CII]/\\nii\\ flux ratio of 10.4$\\pm$1 was observed (first laboratory measurements were obtained by Brown et al.\\ 1994). The \\nii\\ line was later detected in the Galactic HII region G333.6--0.2 (Colgan et al.\\ 1993) and subsequently in the Carina Nebula (Oberst et al. 2006); the first detection from the ground (using SPIFI on AST/RO). Again a flux ratio around 10 (9.2$\\pm$3.0) relative to the [CII] line was found.  The \\nii\\ line was detected in only one extragalactic system, M\\,82, using the Kuiper Airborne Observatory, with a [CII]/\\nii\\ flux ratio of 20.0$\\pm$3.4 (Petuchowski et al.\\ 1994). At high redshift, the ratio may be susceptible to N/C abundance variations (e.g., Matteucci \\& Padovani 1993). Given the potential importance of the \\nii\\ line in studies of the ISM in the very early universe, multiple attempts were performed to detect this line at high redshift (4C41.17 and PC2047+0123: Ivison \\& Harris 1996, the Cloverleaf: Benford 1999, APM\\,08279: Krips et al.  2007, see Table~1 for a summary).  Here we present a sensitive search for \\nii\\ emission in one of the highest redshift quasars, SDSS J114816.64+525150.3 (hereafter: J1148+5251) at z=6.42. It has a far--infrared (FIR) luminosity of 2.2$\\times10^{13}$ L$_\\odot$ (Bertoldi et al.\\ 2003a, Beelen et al.\\ 2006) and hosts a large reservoir of molecular gas (2\\,$\\times$\\,10$^{10}$\\,M$_\\odot$), the prerequisite for star formation, which has been detected through redshifted rotational transition lines of CO (Walter et al.\\ 2003, Bertoldi et al.\\ 2003b, Walter et al.\\ 2004).  Besides BR\\,1202--0725 (Iono et al.\\ 2006), it is the only system detected in [CII] at z$>$0 to date (Maiolino et al.\\ 2005, Walter et al.\\ 2008) and is thus an obvious target to search for emission from \\nii. We use a concordance, flat $\\Lambda$CDM cosmology in this letter, with $H_0$=71\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm M}$=0.27 and $\\Omega_{\\rm \\Lambda}$=0.73 (Spergel et al.\\ 2007), and a resulting luminosity distance of 63.8\\,Gpc at z=6.42 (Wright 2006).    \\begin{figure} \\epsscale{1} \\plotone{f1.eps} \\caption{The CO(6--5) line in J1148+5251 from the data that were taken in parallel to the \\nii\\ observations. Velocities are given offset to the tuned redshift of z=6.4189. The best Gaussian fit is overplotted as a solid curve (see Sec.~3.1 for fit parameters). \\label{f2}} \\end{figure} ",
        "conclusions": "\\subsection{Comparison to [CII] and Implications}  The ionization potential of carbon of 11.3\\,eV is below that of hydrogen (13.6\\,eV), whereas the one for nitrogen is above that value (14.53\\,eV). The \\nii\\ and [CII] transitions have nearly identical, low critical densities (\\nii: 44\\,cm$^{-3}$, [CII]: 46\\,cm$^{-3}$, assuming an electron temperature of 8000\\,K, e.g. Oberst et al. 2006), i.e.\\ their line ratio is given by the N$^+$/C$^+$ abundance ratio in the ionized medium. As pointed out by Oberst et al.\\ (2006), this ratio is insensitive to the hardness of the radiation field as the energy levels of the next ionization states (N$^{++}$ and C$^{++}$) are also similar (29.6 and 24.4 eV, respectively).  Our \\nii\\ observations are overplotted on the [CII] spectrum obtained at the PdBI (Walter et al.\\ 2009) in Fig.~3. The integrated flux of the [CII] measurement is I$_{\\rm [CII]}$=3.9$\\pm$0.3\\,Jy\\,km\\,s$^{-1}$, i.e. an order of magnitude higher than our \\nii\\ upper limit. In terms of luminosities, the [CII] luminosity is L$_{\\rm [CII]}$=4.69$\\pm$0.35$\\times10^9$\\,L$_\\odot$, a factor of 12 higher than the our upper limit for L$_{\\rm [NII]}$.  Our derived upper limit for the [CII]/\\nii\\ flux ratio thus is comparable to the values found in the Milky Way (Wright et al.\\ 1991) and M\\,82 (Petuchowski et al.\\ 1994) and sets tighter limits on the [CII]/\\nii\\ ratio than previous measurements at high redshifts (Table~1). Given the non--detection of the \\nii\\ line in J1148+5251 and in other sources at high redshift we can only speculate whether or not high--z detections are within reach of currently operating observatories.   \\subsection{Outlook}  Probing back into the Epoch of Reionization (z$>$7), forbidden atomic fine structure lines will likely be the only means by which to directly constrain the star--forming interstellar medium using (sub--)millimeter interferometers. This is due to the fact that only very high rotational transitions of CO can be observed at these redshifts using (sub--)millimeter facilities. However, strong emission from the highest (J$>6$) CO levels requires very highly excited gas (either due to high kinetic temperatures, high densities, or both) that is typically not present in large quantities in normal starforming environments (e.g. Daddi et al.\\ 2008).  Therefore high--J CO lines are faint and difficult to detect in such systems.  Recent studies of the CO excitation in high--redshift galaxies have shown that the CO excitation ladder (`CO SEDs') peak at around J$\\sim$6 for QSOs and J$\\sim$5 for submillimeter galaxies (Wei{\\ss} et al.\\ 2005, 2007a, 2007b, Riechers et al.\\ 2006).  \\begin{figure} \\epsscale{1} \\vspace*{6mm} \\plotone{f3.eps}  \\caption{Comparison of the \\nii\\ non--detection towards J1148+5251 and the [CII] emission (continuum subtracted) in the same source (Walter et al.\\ 2009). \\label{f3}} \\end{figure}   \\begin{figure*} \\epsscale{1.0} \\plotone{f4.eps} \\caption{Frequency coverage of redshifted fine structure lines for some current (and future) key telescopes as a function of redshift. The following lines are plotted (in order of increasing frequency, given in GHz]): CI$\\nu$492.2, CO(6--5)$\\nu$691.5, CI$\\nu$809.3, [NII]$\\nu$1461.1 (=\\nii), [CII]$\\nu$1900.5, [OI]$\\nu$2060, [NII]$\\nu$2460, [OIII]$\\nu$3393, [OI]$\\nu$4745, [OIII]$\\nu$5786). The {\\em left} panel shows the coverage for the PdBI and CARMA, the {\\em middle} panel the respective coverages for the EVLA and the IRAM 30\\,m telescope and the {\\em right} panel for ALMA (bands 2--9, band 5 is not fully funded and thus appears in lighter color). The three black points in the left and middle panels highlight the positions of the transitions in J1148+5251 discussed in this letter (CO(6--5), \\nii\\ and [CII]). \\label{f4}} \\end{figure*}  Figure 4 shows which atomic fine strucure lines can be observed in a given frequency band of some (present and future) telescopes as a function of redshift.  The CO(6--5) transition (which lies close to the peak of the highly excited QSOs) is also shown for comparison. It is obvious from this figure that at z$>$7, the atomic fine structure lines will be critical probes to constrain the properties of the interstellar medium using (sub--)millimeter instruments. Redshifted mm--wavelength molecular line observations have so far only probed the molecular medium of the target sources. Fine structure lines of various species, and in particular their luminosity ratios, will also yield information on the neutral and ionized gas that is more dilute.  Even though we have not detected the \\nii\\ line in our observations, we have reached a limit that is consistent with what is found in our Galaxy and in the starburst galaxy M\\,82. This may indicate that we are close to being able to detect atomic fine structure lines other than [CII] at high redshift --- however it is not clear if detections can be obtained given currently operating facilities. In the ALMA era, these lines will provide the fundamental tools needed to constrain the physical properties of the star--forming ISM in the earliest epoch of galaxy formation. These observations will then complement fine structure line measurements in the nearby universe using future airborne and space missions (e.g., {\\em SAFIRE} on {\\em SOFIA} and, e.g., {\\em PACS} onboard {\\em HERSCHEL}).  \\vspace{5mm}  The authors thank Roberto Maiolino, Alberto Bolatto and Helmut Wiesemeyer for useful comments on the manuscript. FW and DR acknowledge the hospitality of the Aspen Center for Physics, where parts of this manuscript were written.  DR acknowledges support from NASA through Hubble Fellowship grant HST-HF-01212.01-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. CC acknowledges support from the Max-Planck-Gesellschaft and the Alexander von Humboldt-Stiftung through the Max-Planck-Forschungspreis 2005.  \\vspace{-2mm}"
    },
    "0812/0812.1448_arXiv.txt": {
        "abstract": "%  We present recent advances in theoretical studies of the formation and evolution of dust in primordial supernovae (SNe) that are considered to be the main sources of dust in the early universe. Being combined with the results of calculations of dust formation in the ejecta of Population III SNe, the investigations of the evolution of newly formed dust within supernova remnants (SNRs) show that smaller grains are predominantly destroyed by sputtering in the shocked gas, while larger grains are injected into the ambient medium. The mass of dust grains surviving the destruction in SNRs reaches up to 0.1--15 $M_\\odot$, which is high enough to account for the content of dust observed for the host galaxies of quasars at $z > 5$. In addition, the transport of dust formed in the ejecta causes the formation of low-mass stars in the dense shell of primordial SNRs and affects the elemental composition of those stars. We also show that the flat extinction curve is expected in the high-redshift universe where SNe are the possible sources of dust. ",
        "introduction": "Recent submillimeter and millimeter observations have detected the continuum thermal emission of dust from a growing number of quasars at redshifts larger than $z=5$ (Bertoldi et al.\\ 2003; Priddey et al.\\ 2003; Robson et al.\\ 2004; Beelen et al.\\ 2006; Wang et al.\\ 2008a, 2008b). The estimated mass of dust in the host galaxies of these quasars is 10$^8$--10$^9$ $M_\\odot$, which implies a rapid enrichment of the interstellar medium (ISM) with dust in such young systems. In addition, the rest-frame near-infrared photometries of quasars at $z > 5$ with {\\it Spitzer } (Hines et al.\\ 2006; Jiang et al.\\ 2006) have confirmed the existence of hot dust heated by the central active nuclei. This indicates that the fundamental structures of quasars such as accretion disks and hot-dust tori had been already established during the first billion years of cosmic history. However, the origin and nature of dust in the high-redshift universe are not necessarily the same as those in our Galaxy. In fact Maiolino et al.\\ (2004a) have found that the rest-frame UV extinction curve observed for the quasar SDSS J1048+4637 at $z=6.2$ is quite different from those in local galaxies, suggesting that the source and evolution of dust at high redshifts ($z > 5$) are different from those at low redshifts ($z < 4$).  Dust grains in the early universe have great impacts on the formation processes of stars and galaxies. Dust grains control the chemistry of the ISM by locking up a large fraction of refractory elements and by serving as catalysts for chemical reactions of gas. Especially they foster the formation of hydrogen molecules on their surfaces (Cazaux \\& Tielens 2004; Cazaux \\& Spaans 2004), enhancing the star formation rate (SFR) in metal-poor star-forming gas clouds (Hirashita \\& Ferrara 2002). In dense metal-poor molecular clouds, dust grains provide efficient cooling pathways of gas through their thermal emission, which triggers the fragmentation of gas clouds into low-mass clumps of $\\sim$0.1--1 $M_\\odot$ for metallicity of $10^{-6}$--$10^{-4}$ $Z_\\odot$ (Schneider et al.\\ 2002, 2003, 2006a; Omukai et al.\\ 2005; Tsuribe \\& Omukai 2006) and affects the initial mass function (IMF) in the metal-poor universe (e.g., Schneider et al.\\ 2006b). Dust grains also influence the thermal balance in the interstellar space by producing photoelectrons that heat the gas and by reemitting the absorbed stellar light at infrared wavelengths.  The reprocessing of stellar light by dust at high redshift is also crucial in understanding the structure formation history of the universe from relevant observations. Obscuration of stellar light by dust leads to a serious misleading in interpreting the cosmic SFR and the IMF of the very early generation of stars from the cosmic infrared background radiation (e.g., Hauser \\& Dwek 2001). Thermal emission from high-$z$ dust dominates the far-infrared background radiation and can even distort the cosmic microwave background radiation (Loeb \\& Haiman 1997). Thus, the investigation of dust in the early universe is one of the most important subjects to reveal the nature of stellar populations and underlying physical processes at the dark age of the universe.  In the previous studies that assessed the effects of dust grains at high redshift, the composition and size distribution of dust grains were assumed to be the same as those in our Galaxy, and their amount was treated as a parameter (e.g., Loeb \\& Haiman 1997). However, the physical processes involving the interactions of dust with photons and gas are sensitive to the chemical composition and size of dust. The size distribution and amount of dust are determined by the competition between its production and destruction and vary with time according to the star formation activity. Therefore, in order to evaluate the impacts of high-$z$ dust, it is essential to clarify the evolution of dust in the early epoch of the universe by taking account of the formation and destruction of dust self-consistently.  At redshift larger than 5 when the cosmic age is less than 1.2 billion years, the formation of dust in mass-loss winds from asymptotic giant branch (AGB) stars is too inefficient to supply a copious amount of dust into the ISM because most of low-mass stars cannot evolve off the main sequence in such an early epoch. Hence, the main formation sites of dust are considered to be in the ejecta of Type II SNe (SNe II) evolving from short-lived massive stars with the progenitor masses $M_{\\rm pr}=$ 8--40 $M_\\odot$. Indeed the chemical evolution models point out that 0.1--1 $M_\\odot$ of dust per SN II event is required to form to explain a large content of dust observed in high-redshift galaxies (Morgan \\& Edmunds 2003; Maiolino et al.\\ 2006; Dwek et al.\\ 2007). On the other hand, numerical simulations (e.g., Bromm et al.\\ 2002; Abel et al.\\ 2002; Yoshida et al.\\ 2006) and semi-analytic models (Nakamura \\& Umemura 2001, 2002) have shown that the first stars formed out of the primordial composition of gas are typically very massive with masses more than 100 $M_\\odot$. Such metal-free stars as massive as 140--260 $M_\\odot$ are considered to explode as pair-instability SNe (PISNe, Fryer et al.\\ 2001; Umeda \\& Nomoto 2002; Heger et al.\\ 2002) and pollute the ISM by ejecting large amounts of metal and dust at the end of their lives.  Dust grains newly condensed in the SN ejecta are subsequently processed due to sputtering in the hot gas swept up by the reverse and forward shocks, and their size and mass can be greatly modified before being injected into the ISM. Thus, in order to clarify what kind of and how much dust are supplied to the ISM by SNe, it is necessary to investigate the formation of dust in the ejecta and its evolution in the shocked gas within the SN remnants (SNRs). In this proceedings, we describe the properties of dust in the early universe revealed by a series of our studies; the calculations of dust formation in the ejecta of primordial SNe II and PISNe are presented in \\S~2, and the evolutions of dust in those SNRs are described in \\S~3. In \\S~4 we discuss the role of dust in the early universe, focusing on the impacts of dust on the formation processes of the next generation of stars. In \\S~5 the extinction curve expected at high redshift is derived based on the results of calculations given in \\S~3 and is compared with the existing observational results. The concluding remarks are presented in \\S~6. ",
        "conclusions": "In this paper, we briefly outline recent advances in theoretical studies of the formation and evolution of dust in the early universe, with emphasis on the composition, size distribution, and amount of dust ejected from primordial SNe. The investigations of the evolution of newly formed dust in SNRs show that the transport and destruction of dust grains within SNRs heavily depend on their initial radii. Small grains are efficiently decelerated by the gas drag and are predominantly destroyed by sputtering in the shocked gas. A fraction of large grains surviving the destruction are finally trapped in the dense shell behind the forward shock, and the larger grains are injected into the ISM without undergoing significant deceleration and erosion.  Dust grains surviving the destruction in SNRs bring a few important conclusions. The dust grains accumulated in the dense SN shell can enrich the primordial gas gathered by the forward shock and cause the formation of low-mass Population II.5 stars in the dense shell. Those dust grains can also affect the elemental compositions of the Population II.5 stars and produce the metal abundance patterns resembling those seen in HMP and UMP stars. In addition, the large grains ejected from SNe prescribe the extinction property of the ISM and lead to a flat extinction curve at the high-redshift universe where dust grains are mainly supplied by primordial SNe. Similar situations appear in the circumnuclear regions of AGNs and the GRB host galaxies at low redshifts. Finally, the mass of dust surviving in primordial SNRs reaches up to 0.1--15 $M_\\odot$, which is consistent with the dust yield per SN required to explain a large quantity of dust observed in high-redshift systems. Thus, SNe play the dominant role in the enrichment history of dust as dust factories in the early universe.  Recently, Cherchneff \\& Lilly (2008) have suggested that primordial SNe are the first molecule providers to the early universe. They show that the mixing of hydrogen from the progenitor envelope facilitates the formation of molecules in the He core, which acts as a bottleneck process against the formation of dust in the SN ejecta. However, it has been still debated how large the mixing of elements extends into the ejecta and whether the mixing is at the knotty level or at the atomic level. We also note that although the mass of dust formed in nearby SNe has been considered to be less than $10^{-3}$ $M_\\odot$, the estimated mass is derived for a limited number of observations by assuming that the ejecta is optically thin (see Kozasa et al.\\ in this volume, see also Sugerman et al. 2006; Miekle et al.\\ 2007; Nozawa et al.\\ 2008).  In order to reveal the evolution of dust grains in the early universe, we need further investigations of the fundamental processes of formation and destruction of dust expected at high redshift. Elvis et al.\\ (2002) have proposed the possible condensation of a large amount of dust in quasar winds, while it is suggested that the formation of dust in AGB stars can dominate over the contributions from SNe even at $z=5$ (Valiante, R. et al.\\ in preparation). Hence, the formation of dust and its survival in these possible formation sites in the early universe should be examined. In addition, modellings of the spectral energy distribution of high-$z$ galaxies (e.g., Takeuchi \\& Ishii 2004) as well as nearby young blue compact dwarfs (e.g., Takeuchi et al.\\ 2003, 2005; Galliano et al.\\ 2005) are crucial to specify the properties of dust in the primeval universe. Such sophisticated studies that treat the chemical evolution and the radiation transport calculation self-consistently can be powerful probes to elucidate the origin and nature of the cosmic dust in the early universe, in particular when their results are compared with the future observations with {\\it Herschel}, {\\it JWST}, and {\\it ALMA}."
    },
    "0812/0812.0352_arXiv.txt": {
        "abstract": "We describe results from the first astronomical adaptive optics system to use multiple laser guide stars, located at the 6.5-m MMT telescope in Arizona. Its initial operational mode, ground-layer adaptive optics (GLAO), provides uniform stellar wavefront correction within the 2 arc minute diameter laser beacon constellation, reducing the stellar image widths by as much as 53\\%, from 0.70 to 0.33 arc seconds at $\\lambda = 2.14$ $\\mu$m. GLAO is achieved by applying a correction to the telescope's adaptive secondary mirror that is an average of wavefront measurements from five laser beacons supplemented with image motion from a faint stellar source. Optimization of the adaptive optics system in subsequent commissioning runs will further improve correction performance where it is predicted to deliver 0.1 to 0.2 arc second resolution in the near-infrared during a majority of seeing conditions. ",
        "introduction": "Until very recently, adaptive optics (AO) systems supporting astronomical observing have used single guide stars to measure atmospheric turbulence. This has limited the best optical correction to a single target, with correction degrading as a function of increasing field angle. By using multiple guide stars, many types of wide field adaptive optics correction can be implemented. The simplest of these, ground-layer adaptive optics (GLAO), was suggested by Rigaut \\cite{rigaut} as a way to improve wide field imaging for large telescopes. Wavefront measurements from guide stars located far from each other (2 -- 10+ arc minutes) can be averaged to estimate the turbulence close to the telescope aperture. A partially corrected field over the guide star constellation is produced when only this low-lying turbulence is compensated. Alternatively, single Rayleigh laser guide star (LGS) systems at smaller telescopes such as the 4.2-m William Herschel Telescope\\cite{WHT} and planned for the 4.1-m Southern Astrophysical Research\\cite{SAM} telescope can produce wavefront compensation over fields of 1.5 and 3 arc minutes respectively. However, a single beacon GLAO implementation will not be effective for larger apertures which suffer from stronger focal anisoplanatism effects that limit the accuracy of the ground-layer turbulence measurements\\cite{andersen}. Multiple laser guide stars will also be required for GLAO implementations that achieve full sky coverage on the next generation of extremely large (25+ m) class telescopes.  Simulations have predicted that multiple-beacon GLAO will effectively and consistently improve the atmospheric seeing \\citep{rigaut, andersen, louarn, tokovinin} and open-loop studies at the 6.5-m MMT and 1.5-m Kuiper telescopes predict that GLAO can reduce wavefront aberration by up to $\\sim 45\\%$ \\cite{apj, spie, ngsglao}. Using simultaneous wavefront measurements of 5 laser guide stars and a stellar source at the MMT, synthetic point-spread-functions (PSFs) have been calculated from the residual wavefront errors of open-loop ground layer correction \\cite{opex, spie}. During roughly median seeing conditions ($r_0 = 14.8$ cm at $\\lambda = 500$ nm), the full-width at half-maximum (FWHM) of the corrected PSF was calculated to be 0.12 and 0.16 arc seconds in the $K$ and $H$ bands respectively. Azimuthally averaged PSF profiles of both the uncorrected and open-loop ground-layer corrected stellar PSFs are presented in figure 1. During more favorable seeing conditions ($r_0 = 22.6$ cm), the calculated FWHM size of the open-loop ground-layer corrected PSFs were 0.10 and 0.14 arc seconds in the $K$ and $H$ bands. Although the performance of GLAO is affected by both the changing seeing conditions and the relative strength of ground-layer turbulence, it has been found at many sites, including at the nearby Mt. Graham and Mt. Bigelow, that typically one-half to two-thirds of the atmospheric turbulence lies in this ground layer \\cite{andersen, avila, egner, tok_cn2a, tok_cn2b, ngsglao, velur}. Further statistics on the strength of the ground-layer at the MMT need to be measured; however it is expected that the GLAO system will be capable of delivering 0.1 to 0.2 arc second resolution in the near-infrared during a majority of seeing conditions.  GLAO was first demonstrated on-sky in closed-loop using three bright natural guide stars on a 1.5 arc minute diameter with the Multi-conjugate Adaptive Optics Demonstrator fielded at the Very Large Telescope (VLT) in early 2007 \\citep{mad1,mad2}. Strehl ratios from 5 to 22.5$\\%$ in $K$ band were observed within the guide star constellation during seeing of $\\sim$ 0.7 arc seconds. However, sky coverage and the number of accessible science targets will be limited because there are very few suitably bright natural guide star constellations.  The MMT multiple laser adaptive optics system is being commissioned to support wide-field science with GLAO and narrow-field science with laser tomography adaptive optics (where Strehl is maximized at the cost of corrected field-of-view). In February 2008, the system demonstrated an initial GLAO correction where stellar image widths were reduced by as much as 53\\%, from 0.70 to 0.33 arc seconds at $\\lambda$ = 2.14 $\\mu$m. Further commissioning in May 2008 demonstrated wide field aberration compensation of a constellation of stars, and a reduction of the stellar wavefront errors in the control space of the laser wavefront sensor by $38\\%$. The observations reported here demonstrate significant image improvement across a 110 arc second field, but not yet at the anticipated resolution of 0.1 to 0.2 arc seconds. At the time of these observations, the GLAO system was affected by several issues in the tip-tilt loop which compromised the overall closed-loop performance. The primary problem were network delays and dropouts between the tip-tilt sensor and real-time reconstructor, responsible for unwanted random control loop delays, causing the system to be unstable at anything other than at very low control loop gain values. The tip-tilt issues have since been diagnosed and corrected; thus future observations should not be affected. ",
        "conclusions": "The multiple LGS AO program at the MMT is exploring the practical techniques required for wavefront sensing for ground-layer and tomographic correction. Closed-loop GLAO operation was successfully demonstrated in February 2008 where image correction reduced stellar image widths by as much as $53\\%$ at $\\lambda$ = 2.14 $\\mu$m. Wide field, $\\sim2$ arc minute diameter, correction was demonstrated in May 2008 upon successful testing of static aberration correction and additional system characterization. Optimization of the system in subsequent commissioning runs in the fall of 2008 will further improve correction performance and is expected to deliver $0.1$ to $0.2$ arc second resolution during a majority of seeing conditions. Additional narrow field PSF characterization will be done in the thermal infrared (3.5 -- 4.8 $\\mu$m) where Strehl ratios of $30\\%$ to $40\\%$ are expected using Clio with a plate scale of 48 milli arc seconds per pixels. Imaging with the Clio and PISCES cameras will allow for the exploration of parameters which affect the ground-layer AO correction. In particular, variables such as control gain, reconstruction basis set and the number of controlled modes will be optimized. The effect of observing conditions, such as the brightness and field location of the tilt star, will also be explored to support science programs by anticipating the improved resolution and sensitivity of ground-layer AO correction.  Early shared risk scientific programs will focus on seeing improvement with GLAO, taking advantage of existing near infrared instrumentation \\citep{mmtsci}, with plans for developing further capability over larger fields of view \\cite{loki}. The exploitation of routine near infrared seeing of $0.2$ arc seconds or better over a field of several arc minutes is likely to be very productive, both for imaging and high resolution multi-object spectroscopy where the many-fold improvement in encircled energy within $0.2$ arc seconds will be of particular value.  Development of the MMT laser adaptive optics system will lead to a greater understanding of the important factors in designing and operating future multiple guide star adaptive optics systems such as the ones being planned for the 8.1-m Gemini North \\citep{gemini1} and Gemini South \\citep{gemini2} telescopes, the 8.2-m VLT\\citep{ galacsi, vlt, hawki}, the $2 \\times 8.4$-m Large Binocular Telescope\\\\\\citep{lbtao, nirvana}, the 10-m Keck\\\\\\citep{keck, keck2} telescopes, the 25.4-m Giant Magellan Telescope \\\\\\citep{gmt}, 30-m Thirty Meter Telescope\\citep{tmt}, and 42-m European Extremely Large Telescope\\\\\\citep{eelt}."
    },
    "0812/0812.2161_arXiv.txt": {
        "abstract": "We present a new multi--fluid, grid MHD code PIERNIK, which is based on the Relaxing TVD scheme. The original scheme has been extended by an addition of dynamically independent, but interacting fluids: dust and a diffusive cosmic ray gas, described within the fluid approximation, with an option to add other fluids in an easy way. The code has been equipped with shearing--box boundary conditions, and a selfgravity module, Ohmic resistivity module, as well as other facilities which are useful in astrophysical fluid--dynamical simulations. The code is parallelized by means of the MPI library. In this paper we shortly introduce basic elements of the Relaxing TVD MHD algorithm, following Trac \\& Pen~(\\cite{trac}) and Pen \\etal~(\\cite{pen}), and then focus on the conservative implementation of the shearing box model, constructed with the aid of the Masset's~(\\cite{masset}) method. We present results of a test example of a formation  of a gravitationally bounded object (planet) in a self--gravitating and differentially rotating fluid. ",
        "introduction": "The \\textbf{Relaxing--TVD} conservative scheme (Jin \\& Xin~\\cite{Jin95}) presented by \\linebreak Trac~\\&~Pen~(\\cite{trac}) and Pen \\etal~(\\cite{pen}),  who provided short  codes with a basic implementation of the method, is a second order algorithm in  space and time. The scheme efficiently deals with shocks without artificial viscosity. The code is very flexible and can be extended with new modules representing additional physical processes.  The simplicity and robustness of the code is reflected in general performance of $10^5$ zone--cycles/s (on 2 GHz AMD Opteron processors). \\par The conservative form of MHD equations serves as a starting point \\begin{equation}\\label{basiceq} \\partial _t \\mathbf{u} + \\partial_x \\mathbf{F}\\mathbf{(u,B)} + \\partial_y \\mathbf{G}\\mathbf{(u,B)} + \\partial_z \\mathbf{H}\\mathbf{(u,B)} = 0. \\end{equation} where $\\mathbf{u} = \\left( \\rho, m_x, m_y, m_z, e \\right)$ and $\\mathbf{F}\\mathbf{(u,B)}$, $\\mathbf{G}\\mathbf{(u,B)}$, $\\mathbf{H}\\mathbf{(u,B)}$ are fluxes of the conserved fluid variables in $x$, $y$ and $z$ directions, respectively. Other symbols have their common meaning i.e. $\\rho$ denotes density of the fluid, whereas $m_x$, $m_y$, $m_z$ are components of momentum density and $e$ stands for total energy density. \\par Firstly, a dimensional splitting of equation (\\ref{basiceq}) is made by constructing numerical solution with timestep $\\Delta t$ to the equation \\begin{equation}\\label{tv_eq1} \\partial_t \\mathbf{u} + \\partial_x \\mathbf{F}\\mathbf{(u,B)} = 0, \\end{equation} separately for each dimension. Then $\\mathbf{u}$ and $\\mathbf{F}$ are split into waves that propagate leftwards and rightwards \\begin{equation} \\mathbf{u}^L \\equiv \\frac{1}{2}\\left(\\mathbf{u} -\\frac{\\mathbf{F}}{c}\\right),\\; \\mathbf{u}^R \\equiv \\frac{1}{2}\\left(\\mathbf{u} +\\frac{\\mathbf{F}}{c}\\right),\\; \\mathbf{u}=\\mathbf{u}^L+\\mathbf{u}^R, \\end{equation} \\begin{equation} \\mathbf{F}^L = c \\mathbf{u}^L, \\; \\mathbf{F}^R = c \\mathbf{u}^R, \\;  \\mathbf{F}= \\mathbf{F}^R + \\mathbf{F}^L \\end{equation} where $c$, called the \\textit{freezing speed}, is a function that satisfies $c\\ge \\max\\left(|v\\pm c_f|\\right)$, $v$~is the fluid speed and $c_f$ is the fast magnetosonic speed. Now, equation (\\ref{tv_eq1}) is equivalent to two independent equations: \\begin{equation} \\label{tv_eq2} \\partial_t \\mathbf{u}^L - \\partial_x \\mathbf{F}^L =0, \\quad \\partial_t \\mathbf{u}^R + \\partial_x \\mathbf{F}^R =0. \\end{equation} The above pair of equations  (\\ref{tv_eq2})  is solved by means of an upwind conservative scheme, separately for right-- and left--going waves, using cell--centered fluxes. To~achieve 2nd order spatial accuracy, a monotone upwind interpolation of fluxes onto cell boundaries is made,  with the aid of a \\textit{flux limiter}. Second order accuracy of time integration is achieved through the Runge--Kutta scheme (details see Trac~\\&~Pen~\\cite{trac}, Pen \\etal~\\cite{pen}) ",
        "conclusions": ""
    },
    "0812/0812.0813_arXiv.txt": {
        "abstract": "We present observations of the European Large-Area {\\it ISO} Survey-North 1 (ELAIS-N1) at 325 MHz using the Giant Metrewave Radio Telescope (GMRT), with the ultimate objective of identifying active galactic nuclei and starburst galaxies and examining their evolution with cosmic epoch. After combining the data from two different days we have achieved a median rms noise of  $\\approx40 \\mu$Jy beam$^{-1}$, which is the lowest that has been achieved at this frequency.  We detect 1286 sources with a total flux density above $\\approx270 \\mu$Jy. In this paper, we use our deep radio image to examine the spectral indices of these sources by comparing our flux density estimates with those of Garn et al. at 610 MHz with the GMRT, and surveys with the Very Large Array at 1400 MHz. We attempt to identify very steep spectrum sources which are likely to be either relic sources or high-redshift objects as well as inverted-spectra objects which could be Giga-Hertz Peaked Spectrum objects. We present the source counts, and report the possibility of a flattening in the normalized differential counts at low flux densities which has so far been reported at higher radio frequencies. ",
        "introduction": "One of the key questions in extragalactic studies is to understand the evolution of galaxies, namely their star formation history, formation of active galactic nuclei (AGN), building-up of large disks and bulges and the formation and evolution of supermassive black holes over the redshift range from $z\\approx5$ to the present epoch. This is the range where the global star-formation rate has passed through a maximum between a redshift of 1 and 2 (\\citeauthor{1996MNRAS.283.1388M}~\\citeyear{1996MNRAS.283.1388M}; \\citeauthor*{1998ApJ...498..106M}~\\citeyear{1998ApJ...498..106M}; see the recent compilation by \\citeauthor{2006ApJ...651..142H}~\\citeyear{2006ApJ...651..142H}).  It is conjectured that merging of galaxies has occurred significantly at these epochs, triggering collapse of molecular clouds and star formation, often in dusty environments. Many of these galaxies may also harbour an AGN and are copious emitters in the infrared region of the spectrum. It has become clear over the last decade that a population of galaxies which radiate most of their power in the infrared constitutes an important and significant component of the Universe (\\citeauthor*{2005ARA+A..43..727L}~\\citeyear{2005ARA+A..43..727L} for a review).  Multiwavelength observations of high-redshift infrared galaxies could provide valuable insights into issues of galaxy formation and evolution. The spectral energy distribution (SED) of the cosmic infrared (IR) background (CIB, which is the emission at wavelengths larger than a few microns) is due to the formation and evolution of both AGN and starburst galaxies (e.g. \\citeauthor{1996A+A...308L...5P}~\\citeyear{1996A+A...308L...5P}; \\citeauthor{1998ApJ...508...25H}~\\citeyear{1998ApJ...508...25H}; \\citeauthor{1999A+A...344..322L}~\\citeyear{1999A+A...344..322L}; \\citeauthor{2000A+A...360....1G}~\\citeyear{2000A+A...360....1G}; \\citeauthor{2005PhR...409..361K}~\\citeyear{2005PhR...409..361K}; \\citeauthor{2006A+A...451..417D}~\\citeyear{2006A+A...451..417D}). The SED of the CIB peaks around 150 $\\mu$m, accounting for about half the total energy in the optical/infrared extragalactic background light (cf. \\citeauthor{2001ARA+A..39..249H}~\\citeyear{2001ARA+A..39..249H}; \\citeauthor{2004NewAR..48..465W}~\\citeyear{2004NewAR..48..465W}; \\citeauthor{2006Natur.440.1018A}~\\citeyear{2006Natur.440.1018A}). It has been shown recently that the mid-infrared (MIR) 24~$\\mu$m selected sources contribute more than 70 per cent of the CIB at 70~and~160~$\\mu$m. Galaxies contributing the most to the total CIB are z$\\sim$1 luminous infrared galaxies, which have intermediate stellar masses \\citep{2006A+A...451..417D}.  To explore these issues related to galaxy formation and evolution and possible relationships between AGN and starburst activity require large samples over wide enough areas, to avoid biases due to large-scale structure variations, as well as multi-wavelength coverages to have a global appraisal of the energy output and of its origin. This means, in particular, a far-IR coverage to access the ``dusty'' objects to get the bolometric energy output; an optical coverage as this is the range where the usual spectroscopic diagnostics are available; the ultraviolet (UV) range as this is where the massive stars have their impact and where the energy is seen  in the absence of dust; the radio range to distinguish thermal from non-thermal sources and obtain the best spatial resolution; and possibly a high-energy coverage in X-rays to detect the most powerful AGNs. Amongst the different fields which are being studied extensively, a well-known one selected at infrared wavelengths is the European Large Area ISO (Infrared Space Observatory) Survey (ELAIS), from which we have selected one of the northern fields, the ELAIS-N1, observable with GMRT.  The ELAIS-N1 field has been chosen in a region of the sky with low-IR foreground emission, to allow detection of fainter (and presumably more distant) galaxies; also, because of this advantage, it  has already  been covered at far-IR wavelengths by the ISO-FIRBACK (Far-Infrared Background) survey at 170~$\\mu$m, to find dusty, high-z galaxies expected to significantly contribute to the far-IR cosmic background and is covered by Spitzer as well, as part of the SWIRE (Spitzer Wide-area Infrared Extragalactic) legacy survey \\citep{2004ApJS..154...54L}. The details on the first far-IR data can be found in \\cite{2001A+A...372..364D}. Our team has been studying this field in detail, and has particularly made some of the optical spectroscopic observations. Most of the  sources identified up to now, essentially the IR-brighter ones, are local, cold, moderately star-forming galaxies rather than the violent starbursters expected at larger distances \\citep{2005A+A...440....5D}. It is now understood that this is partly due to the sensitivity limit of ISO, so that the higher z sources, needed to reproduce the number counts, have to be found among the fainter IR sources detected by Spitzer. Due to its interest in studies of galaxy evolution, it has also been included in the Galex UV survey. It has already optical imaging coverage from LaPalma as part of the INT (Isaac Newton Telescope) Wide Field Survey \\citep{2001NewAR..45...97M} and a band-merged catalogue is available \\citep{2004yCat..73511290R}. Redshift surveys of the region are underway \\citep[e.g.][]{2008MNRAS.386..697R} and it is also partly covered by the Sloan Digital Sky Survey \\citep[SDSS,][]{2008ApJS..175..297A}.  Radio surveys provide important constraints in our understanding of the evolution of the Universe. Traditionally counts of the number of sources as a function of the radio flux density have provided information on the evolution of radio source populations with cosmic epoch (e.g. \\citeauthor{1966MNRAS.133..421L}~\\citeyear{1966MNRAS.133..421L}; \\citeauthor{1993MNRAS.263..123R}~\\citeyear{1993MNRAS.263..123R}). More recently, deep radio surveys have shown a flattening of the source counts at about a mJy at both 610 and 1400 MHz (e.g. \\citeauthor*{1990ASPC...10..389W}~\\citeyear{1990ASPC...10..389W}; \\citeauthor{2006A+A...457..517P}~\\citeyear{2006A+A...457..517P} and references therein; \\citeauthor{2007MNRAS.378..995M}~\\citeyear{2007MNRAS.378..995M} and references therein). The change of slope is believed to be due to a new population of radio sources, the so-called sub-mJy population, consisting largely of low-luminosity AGN and starburst galaxies \\citep[e.g.][]{2007ASPC..380..205P, 2008ApJS..177...14S}. The source counts at higher flux densities are dominated by the classical double radio galaxies and quasars \\citep[e.g.][]{1989ApJ...338...13C}.  Concerning the ELAIS-N1 field, a Very Large Array (VLA) radio survey has been made at $\\lambda$20 cm by \\citet{1999MNRAS.302..222C}, reaching a brightness limit of $\\approx$1 mJy beam$^{-1}$  (5 sigma) over the 1.54 square degrees of coverage, and deeper in smaller areas. This field has been studied at $\\lambda$50 cm using the Giant Metrewave Radio Telescope (GMRT) by \\citet{2008MNRAS.383...75G}, hereafter referred to as G2008. They have covered a total area of $\\approx$~9~deg$^2$ with a resolution of 6~$\\times$~5~arcsec$^2$ with a total of 19 pointings. In 4 of their pointings they reach an rms of $\\approx$~40~$\\mu$Jy~beam$^{-1}$ and an rms of $\\approx$~70~$\\mu$Jy~beam$^{-1}$ in the rest of the pointings. They have catalogued 2500 sources and present a mosaic of maps from their survey. Preliminary results from a study of polarised compact sources \\citep{2007ApJ...666..201T} at 1420 MHz using the Dominion Radio Astrophysical Observatory Synthesis Telescope (DRAO ST) are also available. They currently present 30\\% of observations from a survey of 7.4~deg$^2$ centered on $16^h11^m,~+55^\\circ00^\\prime$. They present maps in Stokes I, Q and U with a maximum sensitivity of 78~$\\mu$Jy~beam$^{-1}$, with a resolution of $\\approx$~1~arcmin$^2$.  A GMRT coverage of this field, at $\\lambda=$~90 and 200~cm, provides an ideal complement to get accurate spectral shapes and indexes over a large frequency range and help establish the nature of the objects. In particular it will help to separate the respective contribution of AGNs and starbursts, a long standing problem for the interpretation of the far-IR emission. It is also relevant to note here that there is a tight correlation between the far-IR and radio luminosities of galaxies (e.g. \\citeauthor*{1985ApJ...298L...7H}~\\citeyear{1985ApJ...298L...7H}; \\citeauthor{1992ARA+A..30..575C}~\\citeyear{1992ARA+A..30..575C}) which applies to both local and global properties of galaxies in the nearby and distant Universe (\\citeauthor{2002A+A...384L..19G}~\\citeyear{2002A+A...384L..19G}; \\citeauthor{2005ChJAA...5..448L}~\\citeyear{2005ChJAA...5..448L}; \\citeauthor{2006MNRAS.370..363H}~\\citeyear{2006MNRAS.370..363H}; \\citeauthor{2006ApJ...638..157M}~\\citeyear{2006ApJ...638..157M}; \\citeauthor{2005ApJ...622..772C}~\\citeyear{2005ApJ...622..772C}).  As part of our program to explore the issues outlined earlier, we are making deep observations of the ELAIS-N1 field with the GMRT at low frequencies. In this paper we present the results of our GMRT observations at 325 MHz pointed towards the centre of the field ($16^h10^m, +54^\\circ36^\\prime$). We have made the observations on two different days and have made images of the first day separately to check for consistency.  The final image made by combining the data of both the days has a median rms noise figure of $\\approx$40 $\\mu$Jy beam$^{-1}$ towards the centre of the field, and is amongst the deepest images made at this frequency. It is more sensitive than the low-frequency survey of the XMM-LSS (Large Scale Structure) field at 325 MHz with the VLA which has a 5$\\sigma$ brightness limit of $\\sim$4 mJy beam$^{-1}$ \\citep{2006A+A...456..791T}. For comparison, the 5$\\sigma$ limiting brightness of the Westerbork Northern Sky Survey (WENSS; \\citeauthor{1997A+AS..124..259R}~\\citeyear{1997A+AS..124..259R}) is $\\sim$18 mJy beam$^{-1}$, while for the deeper surveys of selected fields with the Westerbork telescope the corresponding values range from 2.4 to 3.5 mJy beam$^{-1}$ \\citep{1992A+A...256..331W}. The details of our observations and analysis are described in Section~\\ref{ef0:sec:obs}, while the results are presented in Section~\\ref{ef0:sec:result}. The discussion and conclusions are summarised in Section~\\ref{ef0:sec:discussions}.  \\begin{table} \\begin{center} \\caption{Observation summary} \\label{ef0:table:obs_sum} \\begin{tabular}{| l | l |} \\hline Date & 2005 June 26, 2005 June 27  \\\\ Working antennas & 25 \\\\ Centre Frequency & 325 MHz \\\\ Bandwidth & 32 MHz \\\\ Visibility integration time & 4.19 s \\\\ Total observation time & 10.3 hrs $\\times$ 2 \\\\ \\hline \\multicolumn{2}{l}{~~~~~~~~ Flux Calibrator}\\\\ \\hline Source & 3C286\\\\ Time & 0.3 hrs $\\times$ 2\\\\ Flux density & 25.97 Jy\\\\ Scale & Perley-Taylor 99\\\\ \\hline \\multicolumn{2}{l}{~~~~~~~~ Phase Calibrator}\\\\ \\hline Source & J1459+7140 \\\\ Time & 1.8 hrs $\\times$ 2\\\\ Flux density & 20.17 Jy\\\\ \\hline \\multicolumn{2}{l}{~~~~~~~~ Target Field}\\\\ \\hline Phase centre & J1610+5440 \\\\ Time & 8.2 hrs $\\times$ 2\\\\ Fraction ($f_d$) of data used & $\\approx$ 40\\% \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{table} \\begin{center} \\caption{Imaging summary} \\label{ef0:table:imag_sum} \\begin{tabular}{| l | l |} \\hline Package & AIPS++ (version: 1.9, build \\#1556)\\\\ Image size & $5000\\times5000$\\\\ Pixel size & $2^{\\prime\\prime}\\times2^{\\prime\\prime}$ \\\\ Clean Algorithm & Wide-field Clark clean\\\\ Imaging Weight & Briggs (rmode=`norm', robust=0)\\\\ \\hline \\multicolumn{2}{l}{Widefield Imaging Parameters}\\\\ \\hline Number of facets & 9 with 1.5 overlap\\\\ Wproject planes & 256\\\\ Primary Beam correction & Gaussian, FWHM=84.4$^{\\prime}$\\\\ Rms noise & $\\approx 40 \\mu$Jy beam$^{-1}$\\\\ Synthesised beam & $9.35^{\\prime\\prime}\\times7.38^{\\prime\\prime}$,\\ PA=73.25$^\\circ$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "We have presented deep observations of the ELAIS-N1 field at 325 MHz using the GMRT. This is part of an ongoing program to examine the source counts and the nature of radio sources identified from sensitive observations at low radio frequencies with the objective of identifying AGN and starburst galaxies and examining their evolution with cosmic epoch. For this data set we have achieved a median rms noise of $\\approx 40~\\mu$Jy beam$^{-1}$ towards the phase centre by combining data from two different days, which is about the lowest that has been achieved at this frequency.  We detect 1286 sources with a total flux density above $\\approx270 \\mu$Jy within a radius of 1.1$^\\circ$ of the phase centre. By comparing our results with those of G2008 and the FIRST survey, whose resolutions are within a factor of two we have identified candidate very steep spectrum sources with a $\\alpha\\geq$1.3 and candidate GPS objects for further investigations. Considering only those sources with a flux density $>$315~$\\mu$Jy, so that any effects of incompleteness are negligible, our data show evidence of a flattening of the source counts at low flux densities at 325 MHz, which has been reported earlier at higher frequencies and attributed to both starburst galaxies and low luminosity AGN. This needs to be investigated further by combining it with sensitive observations from other fields as well at this frequency."
    },
    "0812/0812.0807_arXiv.txt": {
        "abstract": "We study the generation of non-Gaussianity in models of hybrid inflation with two inflaton fields, (2--brid inflation). We analyse the region in the parameter and the initial condition space where a large non-Gaussianity may be generated during slow-roll inflation which is generally characterised by a large $\\fnl$, $\\tau_{NL}$ and a small $g_{NL}$. For certain parameter values we can satisfy $\\tau_{NL}\\gg\\fnl^2$. The bispectrum  is of the local type but may have a significant scale dependence. We show that the loop corrections to the power spectrum and bispectrum are suppressed during inflation, if one assume that the fields follow a classical background trajectory. We also include the effect of the waterfall field, which can lead to a significant change in the observables after the waterfall field is destabilised, depending on the couplings between the waterfall and inflaton fields. ",
        "introduction": "The increasingly accurate observations of the cosmic microwave background~\\cite{Komatsu:2008hk} motivate the study of inflation beyond linear order in perturbation theory. Since there are many models of inflation which give similar predictions for the spectral index of the scalar perturbations and the tensor-to-scalar ratio~\\cite{Liddle:2000cg,non-gaussian1} it is important to consider higher order observables. Non-Gaussianity has emerged as a powerful discriminant between different models of inflation. Any detection of primordial non-Gaussianity would rule out the simplest models of single field inflation.  Many ways to generate a large non-Gaussianity have been proposed in the literature~\\cite{Alabidi:2006hg,Bartolo:2001cw,Enqvist:2004ey,cline,Dutta:2008if,Misra:2008tx}. If the inflaton has a non-canonical kinetic term then the inflaton field perturbations may already be non-Gaussian at Hubble exit during inflation~\\cite{DBI}. Otherwise the perturbations are Gaussian at Hubble exit and any non-Gaussianity must then be generated on super Hubble scales \\cite{Maldacena:2002vr,Seery:2005gb}. Popular methods to generate a large non-Gaussianity after inflation include having a feature in the inflaton potential~\\cite{Chen:2006xjb}, the curvaton scenario~\\cite{curvaton}, modulated reheating/preheating~\\cite{modulatedreheating,Suyama:2007bg}, an inhomogeneous end of inflation~\\cite{Bernardeau:2002jy,endofinflation,Alabidi:endinf,Sasaki:2008uc,Naruko:2008sq} or by loop domination of the bispectrum or trispectrum~\\cite{Cogollo:2008bi,Rodriguez:2008hy}. Work has also been done on calculating the non-Gaussianity during inflation with a separable potential \\cite{Vernizzi:2006ve,Choi:2007su,Battefeld:2006sz,Seery:2006js} and with more general potentials~\\cite{Rigopoulos:2005us,Yokoyama:2007uu}. Recently the authors of this paper have shown that it is possible to generate a large non-Gaussianity during multiple field inflation while keeping all of the slow-roll parameters much less than unity~\\cite{Byrnes:2008wi}.  In this paper we study this possibility of generating a large non-Gaussianity during slow-roll inflation in detail for the specific model of hybrid inflation. We consider carefully the parameter constraints and initial conditions required to do this. Unlike in the previous work~\\cite{Byrnes:2008wi} we here also consider the effect of the waterfall field required to end inflation. We can therefore also compare and contrast our work to a recent paper on generating a large non-Gaussianity at the end of inflation~\\cite{Naruko:2008sq}. Depending on the values of the couplings between the two inflaton fields and the waterfall field observable quantities may change when the waterfall field is destabilised. This change in the primordial curvature perturbation is possible because there are isocurvature perturbations present at the end of inflation.  We also calculate further observables for this model, such as the trispectrum and the scale dependence of the bispectrum. In most cases the trispectrum is large through a large $\\tau_{NL}$ whenever the bispectrum is large but the relation between them also depends on the initial conditions.  Recently it has also been shown that for two-field hybrid inflation it is possible to generate a large loop correction to the bispectrum and trispectrum~\\cite{Cogollo:2008bi,Rodriguez:2008hy}. For a special trajectory it was shown that the bispectrum can be observably large even when the tree level value of the bispectrum is slow-roll suppressed. Here we show that this is not possible for a classical trajectory, since we require that the classical motion of the field dominates over its quantum fluctuations in order that a slow-roll calculation is valid.  The plan of our paper is as follows: In Section~\\ref{sechybrid} we review the generation of a large bispectrum during hybrid inflation and calculate the trispectrum. In Section~\\ref{sec:waterfall} we include the waterfall field and consider its effect at the end of inflation for different couplings between the waterfall and inflaton fields and further evolution after inflation. In Section~\\ref{sec:scale} we calculate the scale dependence of the non-Gaussianity parameter. We conclude in Section~\\ref{sec:conclusion}. In the Appendix, we show that the loop correction is subdominant to the tree level term for the power spectrum, bispectrum and trispectrum. ",
        "conclusions": "\\label{sec:conclusion}  We have made an in depth study of a model of hybrid inflation with two inflaton fields (two-brid inflation). We have studied the parameter space where non-Gaussianity may be large during slow-roll inflation. In particular we have calculated the observable parameters $\\Pz$, $n_{\\zeta}-1$, $r$, $\\fnl$, $g_{NL}$ and $\\tau_{NL}$. The spectral index depends on the weighted sum of the slow-roll parameters $\\etap$ and $\\etac$ while the bispectrum depends exponentially on the difference of the same slow-roll parameters, so it is possible to have a very large bispectrum and a scale invariant spectrum. We have shown that the trispectrum is generally large through $\\tau_{NL}$ whenever the bispectrum is large, but that the other parameter which parameterises the trispectrum, $g_{NL}$ is always smaller than $\\fnl$ and strongly suppressed compared to $\\tau_{NL}$. We have also shown that for certain initial conditions it is possible that $\\tau_{NL}$ is the only large non-linearity parameter, so the first observational signature of non-Gaussianity from this model could be the trispectrum through $\\tau_{NL}$.  Furthermore we have shown that during slow-roll inflation the loop corrections to the power spectrum, $\\fnl$ and $\\tau_{NL}$ are always suppressed compared to the tree level terms. The suppression follows from the constraint on the background trajectory that the classical background value of the field should dominate over its quantum fluctuations.  We then investigated how the large non-Gaussianity which was generated during slow-roll may be changed by the effects from the end of inflation. We included the waterfall field which ends inflation when its effective mass becomes negative and it is destabilised. The effect of the waterfall field depends on the values of the coupling constants between the waterfall field and the two inflaton fields. If we choose the coupling constants so that surface where the waterfall field is destabilised corresponds to a hypersurface of uniform energy density then there is no change to the observables we calculated during slow roll. However if we choose the two coupling constants to be equal then observables at both linear and higher order are changed by an amount that depends on the ratio $\\etap/\\etac$."
    },
    "0812/0812.2033_arXiv.txt": {
        "abstract": "We use the recently completed one billion particle {\\it Via Lactea II} $\\Lambda$CDM simulation to investigate local properties like density, mean velocity, velocity dispersion, anisotropy, orientation and shape of the velocity dispersion ellipsoid, as well as  structure in velocity space of dark matter haloes. We show that at the same radial distance from the halo centre, these properties can deviate by orders of magnitude from the canonical, spherically averaged values, a variation that can only be partly explained by triaxiality and the presence of subhaloes. The mass density appears smooth in the central relaxed regions but spans four orders of magnitude in the outskirts, both because of the presence of subhaloes as well as of underdense regions and holes in the matter distribution. In the inner regions the local velocity dispersion ellipsoid is aligned with the shape ellipsoid of the halo. This is not true in the outer parts where the orientation becomes more isotropic. The clumpy structure in local velocity space of the outer halo can not be well described by a smooth multivariate normal distribution. {\\it Via Lactea II} also shows the presence of cold streams made visible by their high 6D phase space density. Generally, the structure of dark matter haloes shows a high degree of graininess in phase space that cannot be described by a smooth distribution function. ",
        "introduction": "Spherical averaging is a commonly used method to describe the characteristics of dark matter haloes that form by gravitational collapse in a cosmological environment. For example one of the basic characteristics of a dark matter halo is its spherically-averaged density profile which can be well described by a smooth function within resolved scales \\citep[e.g.][]{1996ApJ...462..563N,1998ApJ...499L...5M,2005MNRAS.364..665D,2007ApJ...657..262D,2008Natur.454..735D,2008arXiv0808.2981S}. Other examples include the velocity dispersion and anisotropy profiles.  Consider a set of observers at a given distance from the halo centre and capable of measuring local properties of the halo (e.g., density, velocity dispersion, velocity anisotropy). In a smooth, spherically symmetric halo, these measurements would yield identical values, up to statistical fluctuations. However, it is well known that dark matter haloes that form in pure dark matter simulations, which neglect possible baryonic effects in the centre, are in general triaxial: close to prolate in the central part and becoming rounder in the outskirts of the halo \\citep[e.g.][]{1991ApJ...368..325K, 1991ApJ...378..496D, 2006MNRAS.367.1781A, 2007MNRAS.376..215B, 2007ApJ...671.1135K}. Such a shape variation obviously leads to a significant variance in local properties compared to the spherically averaged value at a given radius \\citep{2006PASA...23..125K}. In addition, haloes have a high level of subhaloes which also effect results of local measurements. But local deviations from a smooth model go beyond these shape and subhalo driven variations, as we demonstrate in this paper. For example, overdensities are expected due to the presence of subhaloes, but we also find underdense regions, which are unexpected in a smooth triaxial background halo with subhaloes (see section \\ref{sec:den}).  One aim of this paper is to quantify the degree of variation between local and spherically averaged values for key quantities such as the density and velocity dispersion. A better understanding of the local variations of these properties is essential in order to obtain a better description of the phase space structure of dark matter haloes and their formation process.  Such local variations can have profound consequences. For example, in dynamical models of dark matter haloes and galaxies it is often assumed that most properties are just a function of radius \\citep[e.g.][]{2005MNRAS.363.1057D,2007A&A...471..419B,2008arXiv0809.0901V}. From the work presented here it becomes clear that this is not a valid assumption for the anisotropy (see section \\ref{sec:ani}) and a lot of information is smeared out by spherically averaging. Models based on simple analytical forms of the distribution function, which are also used for generating $N$-body realisations of dark matter haloes and galaxies \\citep[e.g.][]{2004ApJ...601...37K,2008MNRAS.386.1543Z,2008MNRAS.387.1719Z}, often have the property that the local bulk motion is zero and the velocity distribution function is similar to a smooth multivariate normal distribution (if it is not even deliberately set to a multivariate normal distribution as in \\cite{1993ApJS...86..389H}). Both of these characteristics are not valid in the outskirts of dark matter haloes that form in a hierarchical cosmological context (see sections \\ref{sec:meanvel} and \\ref{sec:velspace}). As we see, most of these effects are not taken into account in current dynamical models of dark matter haloes and galaxies and they will be a challenge to be modelled with more realistic distribution functions.  In addition, it is also important to quantify these variations locally at the Earth's position within the Galaxy in order to interpret results from present and future direct dark matter detection experiments on Earth \\citep[e.g.][]{2001PhRvD..64f3508M,2001PhRvD..64h3516S,2002PhRvD..66f3502H,2003PhRvL..90u1301S,2008PhRvD..77j3509K,2008MNRAS.385..236V,2008arXiv0808.0704F}. While small subhalos contribute significantly to the indirect detection signal \\citep[e.g.][]{2006ApJ...649....1D,2008ApJ...686..262K} and infall caustics have no significant effect \\citep{2008ApJ...680L..25D}, it is worthwhile also to investigate whether additional lumpiness from tidal debris might enhance the annihilation rate (see section \\ref{sec:den})\\footnote{After submitting this work an independent, analytic study of this question \\citep{2008arXiv0811.1582A} has appeared. Our simulation results agree that such an enhancement is rather small (see section \\ref{sec:den}).}.  Also recent results from surveys like SDSS and RAVE need a better understanding of the structure of the outer halo. There is a lot of known structure in the Galactic stellar halo like for example the Sagittarius stream, the Monoceros stream or the Virgo overdensity \\citep[e.g.][]{2007ApJ...660.1264M,2008ApJ...683..750C,2008AJ....135.2013C}. Obviously, these features are baryonic, but stars are collisionless too, and it's probable that these features originate from infalling dark matter dominated objects.  In general, a more detailed picture of the phase space structure of dark matter haloes is needed in order to understand and model their properties that are set by the hierarchical formation process. This work is a further important step towards that aim.  Here we present a study of the distribution of these properties as a function of galactocentric distance of a Milky Way size dark matter halo with data from the Via Lactea II project \\citep{2008Natur.454..735D}. In section \\ref{sec:data} we give a summary of the simulation data used in this study and present the results in section \\ref{sec:properties}. We then discuss our results and present the conclusions in section \\ref{sec:conclusions}. In appendix \\ref{sec:comparison} we give a comparison of the results with lower resolution simulations and different definitions of locality. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we present a study of local properties of the Via Lactea II dark matter halo. Commonly, characteristics of dark matter haloes are described by spherically symmetric profiles. Here, we show that locally, at a fixed galactocentric distance, these properties can vary by orders of magnitude from their canonical, spherically averaged values. For example, while the density is smooth in the centre, in the outskirts the density range spans four orders of magnitude. This is due to the presence of both subhaloes and underdense regions or holes in the matter distribution. The widespread assumption that there are no bulk flows is only warranted in the central region, where we also find the local velocity dispersion ellipsoid to be aligned with the shape ellipsoid of the halo. Neither of these findings hold in the outer parts of the halo, where particles exhibit bulk motions and a more isotropic distribution of the velocity dispersion ellipsoid's orientation is found. The local velocity space structure can, in general, not be well described by a smooth multivariate normal distribution, at least not in the outer parts of the halo. A qualitatively new feature in Via Lactea II is the detection of streams that are made visible only through their true 6D phase space density.  Such variations of local properties is due to both the triaxial shape of the dark matter halo and its generally clumpy structure in phase space. We find that the phase space structure of a dark matter halo shows a significant departure from the canonical picture of a smooth background density profile with subhaloes in position space and multivariate normal distributions in velocity space. Spherically averaged quantities do not adequately describe properties of the dark matter halo at a given galactocentric distance - especially not the wealth of structure we find locally - since a lot of information gets smeared out by spherically averaging. In general, dark matter haloes show a high degree of graininess - clumps in phase space - which will probably show up at even smaller scales that now seem to be smooth in future simulations with higher resolution. Therefore, a smooth and featureless distribution function does not accurately describe dark matter haloes that form in a cosmological structure formation simulation.  We find several correlations between the shape ellipsoid and the local velocity dispersion ellipsoid. It is unclear at the moment to what degree our findings are universal or how these correlations depend on the hierarchical build up history.  \\citet{2006PASA...23..125K} earlier estimated analytically the effect of a triaxial halo shape on the variance of the local density. Their values of the dispersion normalised to the spherically averaged value for a halo with a similar triaxiality correspond quite good with our values in the centre of the halo but we get much higher dispersion values (larger than the mean) in the outskirts of the halo. This difference is due to their assumption of a smooth triaxial density profile and we get a higher scatter due to the presence of subhaloes and underdense regions.  For dark matter direct detection experiments it is essential to know the local dark matter properties at 8 kpc. The spherically averaged value for the density in Via Lactea II is $\\tilde{\\rho}(8~\\kpc) = 1.056 \\times 10^{-2} ~\\Mo~\\pc^{-3} = 0.4008~\\GeV~c^{-2}~\\cm^{-3}$, close to $0.3~\\GeV~c^{-2}~\\cm^{-3}$ which is often used as a canonical value in the literature \\citep{1998PhRvD..57.3256K,2008PhLB..667..212P}. But as we have shown here, this value can vary locally. One of the large uncertainty factors is the missing information about the orientation of the disk with respect to the dark matter halo. There are different claims from observations and theory about a possible alignment of the angular momentum axis of the disk with the shape axes of the halo \\citep[see e.g.][and references therein]{2004ApJ...613L..41N,2005ApJ...627L..17B,2005ApJ...628...21S} so that a clear answer is not possible at the moment. But often it is claimed that the angular momentum axis of the disk and the short axis ($z$) tend to be aligned \\citep{2005ApJ...627L..17B} which would mean that the disk would lie preferentially in the $xy$-plane in our coordinate system. An additional problem is that we measure these local properties within a sphere of $r_{\\mathrm{sph}}(8~\\kpc) = 500~\\pc$ radius but for dark matter detection experiments more a scale of $1~\\AU = 4.848 \\times 10^{-6}~\\pc = 9.696 \\times 10^{-9}~r_{\\mathrm{sph}}(8~\\kpc)$ is relevant and it is not clear what the local properties of the dark matter distribution on a 1 AU scale are or how one could reasonably extrapolate that over 8 orders of magnitude - especially considering the highly non-linear numerical effects mentioned above and in the appendix that affect the local phase space structure. Also the missing baryonic physics in Via Lactea II like adiabatic contraction, stellar disk and bulge, inspiralling compact objects like black holes etc. can modify the central dark matter structure in either way. Therefore, it is still not clear what the detailed structure of the dark matter locally is."
    },
    "0812/0812.0200_arXiv.txt": {
        "abstract": "Modifications to the gravitational potential affect the nonlinear gravitational evolution of large scale structures in the Universe.  To illustrate some generic features of such changes, we study the evolution of spherically symmetric perturbations when the modification is of Yukawa type; this is non-trivial, because we should not and do not assume that Birkhoff's theorem applies.  We then show how to estimate the abundance of virialized objects in such models.  Comparison with numerical simulations shows reasonable agreement:  When normalized to have the same fluctuations at early times, weaker large scale gravity produces fewer massive halos.  However, the opposite can be true for models that are normalized to have the same linear theory power spectrum today, so the abundance of rich clusters potentially places interesting constraints on such models.  Our analysis also indicates that the formation histories and abundances of sufficiently low mass objects are unchanged from standard gravity.  This explains why simulations have found that the nonlinear power-spectrum at large $k$ is unaffected by such modifications to the gravitational potential.  In addition, the most massive objects in CMB-normalized models with weaker gravity are expected to be similar to the high-redshift progenitors of the most massive objects in models with stronger gravity.  Thus, the difference between the cluster and field galaxy populations is expected to be larger in models with stronger large-scale gravity. ",
        "introduction": "When we consider the cosmological data from WMAP, supernovae Ia, galaxy clustering on large scales, and cross-correlations between galaxies and the CMB, we are faced with several possible conclusions. One is that we live in a spatially flat Friedmann-Robertson-Walker universe currently dominated by either a cosmological constant or repulsive dark energy. The best fit for the dimensionless energy density parameters are $\\Omega_m=0.28$ and $\\Omega_\\Lambda=0.72$, within the concordance $\\Lambda$CDM Model \\cite{wmap5}. The advantage of this model is that nearly all available observational data supports it; the disadvantage is that it requires the vast majority of the energy density in the universe to be in two unknown substances, dark matter and dark energy.  This conclusion rests on the accuracy of our current gravity model, General Relativity.  The key equation of General Relativity, Einstein's equation, relates the curvature and the expansion rate of the Universe to its matter and energy content. The current paradigm is to modify the matter content of the universe, by including dark matter and dark energy, to account for observations. Instead, however, we might modify how the universe curves in response to matter, which would mean modifying our theory of gravity.  There are many possible ways to modify gravity, depending on what one wishes to ``fix''. For example, MOND \\cite{milgrom, bekenstein} removes the need for dark matter to account for galactic rotation curves and has several other interesting results, but seems to fail at the scale of galaxy clusters, with even optimistic accountings needing roughly as much dark matter as baryonic matter. At the other end is something like DGP or conformal gravity, which hopes to account for the acceleration of the universe without invoking a cosmological constant \\cite{dgp, lss, mannheim, diaferio}.  In addition, there are other models which seek to unify dark matter and dark energy \\cite{chaplygin, makler, scherrer05, scherrer08, bertacca}. What we seek to do in this paper is to study the problem more phenomologically. For example, regardless of how the force law for gravity is modified, it will often be stronger or weaker, relative to the standard model, at larger or shorter scales.  One way to parameterize this is to introduce a modified Yukawa-like potential for a point mass: \\begin{equation} \\label{initialpotential} \\phi(r)=G m \\frac{1+\\alpha (1-e^{-r/r_s})}{r}, \\end{equation} \\cite{sealfon,shirata1,shirata2,fritz} where $\\alpha$ indicates the strength and $r_s$ the scale (constant in physical rather than comoving coordinates) on which this modification is most relevant.  On scales smaller than $r_s$, $\\phi(r)$ reduces to the standard Newtonian potential; on larger scales it transitions to the Newtonian potential multiplied by a factor of $(1+\\alpha)$. This is similar to the interaction considered in \\cite{ngp}, in which a long range dark matter interaction is introduced yielding a different Yukawa-like potential that instead modifies gravity on short length scales.  Note that this is not a cosmological model:  there is no prescription for determining things like the expansion factor and the resulting Hubble factor.  Like previous workers, we will assume that these are the same as in the standard cosmology.  The main goal of studying such a model is to gain intuition for some generic effects of modifications to standard gravity.  For example, in the linear theory of such a model, the growth of fluctuations is $k$-dependent \\cite{shirata1} -- it is not in standard gravity.  As a result, a smooth spherical region within which the density is the same as the background universe will evolve.  This qualitatively different behavior from standard gravity has not been emphasized -- so it is worth showing the argument explicitly.  Consider the density field smoothed on scale $R$ at some early time $t_i$.  We can write this field in terms of its Fourier modes, and the (Fourier Transform of the) smoothing kernel as \\begin{equation} \\delta_R(\\mathbf{x},t_i) = \\int d\\mathbf{k}\\, \\exp(i\\mathbf{k}\\cdot\\mathbf{x})\\,\\delta(\\mathbf{k})\\, W(kR). \\end{equation} The linearly evolved field is \\begin{equation} \\delta_R(\\mathbf{x},t) = \\int d\\mathbf{k}\\, \\exp(i\\mathbf{k}\\cdot\\mathbf{x})\\, \\frac{D(k,t)}{D(k,t_i)}\\,\\delta(\\mathbf{k})\\, W(kR), \\end{equation} where $D$ is the linear theory growth factor.  In standard gravity, $D$ is independent of $k$, so if $\\delta_R(\\mathbf{x},t_i)=0$ then $\\delta_R(\\mathbf{x},t)=0$ also.  But if $D$ depends on $k$, then if $\\delta_R(\\mathbf{x})=0$ at some time $t$, it will, in general, be non-zero at other times (the exception being if the $k$-dependence of $W$ happens to exactly cancels that of $D$).  For the Yukawa-like modification considered here, the $k$ dependence of the spherical top hat filter does not cancel that of $D$.  Thus, we are led to the rather remarkable conclusion that, when the gravitational potential has been modified, then linear theory predicts that a spherical tophat patch within which the density is the same as the background will evolve!  The reason why can be traced to the fact that Birkhoff's Theorem no longer applies once the Newtonian potential has been modified.  Without this Theorem, the spherical top hat filter is no longer special, and our common sense prejudice from standard gravity -- that initially overdense regions become denser, underdense regions less dense, but regions within which the density is the same as the background do not evolve -- must be treated with caution.   The evolution of nonlinear clustering in this model has been studied using numerical simulations by \\cite{fritz,shirata2}.  Our goal in what follows is to provide a framework for understanding this nonlinear evolution more completely.  To this end, we will use the spherical evolution model, which has found extensive use in the standard model -- it is used to motivate estimates of the abundance of nonlinear objects \\cite{pressSchechter}, a crucial ingredient in methods which describe the growth of nonlinear gravitational clustering \\cite{cooraySheth}.  It also provides a framework for discussing the environmental dependence of clustering \\cite{moWhite,st02}.  Section~\\ref{model} summarizes what is known from the linear theory of this Yukawa-like model, and then shows our estimate of the key, and sometimes subtle, changes to the spherical evolution model.  Section~\\ref{sims}  compares this work to the simulations of \\cite{fritz}. A final section summarizes. ",
        "conclusions": "We presented a study of nonlinear effects in a model with a modified gravitational potential (equation~\\ref{potentialeq}).  In particular, we showed how the spherical evolution model is modified, and the effect this has on the abundance of virialized objects.  Halos are more massive in models where gravity is stronger on large scales (Figure~\\ref{fig:nmod}), although this effect is only important for sufficiently massive objects whose evolution brings them close to the scale $r_s$ on which gravity was modified.  The effect this has on the abundance of massive objects depends on how the models are normalized.  If normalized so that the fluctuation field is the same at early-times (CMB-normalized), then the models with $\\alpha>0$ have more massive halos (solid curves in Fig.~\\ref{fig:comparisonlinear} differ from unity).  This remains true, although the dependence on $\\alpha$ is reduced, if the models are normalized to have the same ($\\alpha$-dependent) linear theory rms fluctuations today (compare solid and dashed curves in Fig.~\\ref{fig:comparisonlinear}).  If normalized to have the same (linear theory) rms fluctuations today, whatever the value of $\\alpha$, then the trends can be reversed (compare dotted curves with unity in Fig.~\\ref{fig:comparisonlinear}).  This last normalization convention is sometimes known as `cluster-normalized':  our work suggests that, in the context of modified gravity models, this jargon is misleading!  We showed that our analysis captures the essence of the trends seen in the simulations (Figure~\\ref{fig:compareall}), so, in principle, the abundance of rich clusters should place interesting constraints on modifications to the gravitational potential.  In particular, the modification to cluster abundances cannot be reproduced by standard gravity with initial conditions modified to match the change in the linear theory power spectrum; the differences in abundance can be larger than ten percent for sufficiently massive halos (Figure~\\ref{fig:comparisonlinear}).  However, to use cluster counts quantitatively in this way, our analysis should be extended to models in which objects form from an ellipsoidal collapse, as was necessary for standard gravity \\cite{st99,smt01}.  Our analysis also helps to understand an interesting fact about the shape of the nonlinear power spectrum in modified gravity theories.  Figure~7 in \\cite{fritz} shows that whereas the large scale power spectrum in modified theories may be rather different than in standard gravity (because the linear growth factor is modified), the power on small scales ($k>1$) is unchanged.  Our analysis shows that, because small mass objects were never larger than the scale $r_s$ on which gravity was modified, they are not affected by the modification, so their abundance is the same as in the standard case. This is not affected by the initial conditions that we choose so that in both when $\\alpha$ is non-zero, and when the power spectrum is changed so that we end up with a final power spectrum the same as in the case of modified gravity, the abundances of small halos is unaffected.  In addition, their formation histories will also be unchanged, so their internal structural parameters (shapes, density profiles) are also unchanged.  In the halo model description of large scale structure \\cite{cooraySheth}, the power at $k>1$ is dominated by small mass halos.  Since these are the ones for which gravity is essentially unmodified, the small-scale power spectrum is also unchanged.  \\begin{figure*}[tp] \\centering \\includegraphics[width=0.45\\hsize]{fig8a.eps} \\includegraphics[width=0.45\\hsize]{fig8b.eps} \\caption{(Color online) Spatial distribution of a massive halo chosen from the $\\alpha=0.5$ (left) and $-1$ (right) simulations. The black, green, blue, and red contours show the positions of the particles that make up the chosen halo in the $\\alpha = 0.5, 0, -0.5$, and $-1$ simulations, respectively. In the left panel, note that the central dense halo spreads out for weakening gravity, and in the right panel the upper left halo merely moves.  Axes show the $(x,y)$ coordinates of the halos in grid coordinates, $100/256 h^{-1}$Mpc.} \\label{fig:contours} \\end{figure*}  Figure~\\ref{fig:contours} shows the result of a complementary study.  In this case, we selected a massive halo from one of the simulations (say $\\alpha=0.5$), and then looked at where its particles were in the other runs (with different $\\alpha$).  The figure illustrates clearly that when $\\alpha=0.5$, then the particle distribution is more compact.  E.g., in the top panel, the large halo in the $\\alpha=0.5$ run is broken up into three smaller haloes in the $\\alpha = -1$ run.  The bottom panel shows another effect:  that the particles which made up a halo in the $\\alpha=-1$ run are in a different location in the $\\alpha=0.5$ run, suggesting that the peculiar velocities of the most massive halos are also sensitive to $\\alpha$.  The sequence of contours associated with the different $\\alpha$ runs look rather similar to the time evolution of an object in, say, an $\\alpha=0$ model.  Thus, the most massive halos in models with $\\alpha<0$ may be like high-redshift versions of the most massive halos in models with $\\alpha>0$.  Therefore, these figures suggest that the galaxy populations in massive clusters may be rather different in models with large $\\alpha$ than when $\\alpha$ is small.  In particular, it is likely that the difference between the cluster and field galaxy populations increases as the strength of gravity on large scales increases.  This is largely a consequence of the different $\\sigma_8$ values in these runs -- so cluster $M/L$ ratios, currently used to to constrain $\\sigma_8$ \\cite[e.g.][]{tinkerML,flatLsat}, may one day be used to constrain modifications to gravity.  Furthermore, in standard gravity models, there is a strong correlation between evolution and environment \\cite{moWhite, st02, effectiveCosmo}.  There are two reasons to suspect that this will be different if the gravitational potential is modified.  First, in the standard model, the correlation between local environment and evolution is a consequence of Birkhoff's Theorem.  Birkhoff's Theorem is lost when the force law is modified (it is this which modified our spherical evolution model from the standard case).  And second, Figure~\\ref{fig:contours} shows that the time scale for the assembly of objects is modified.  The environmental dependence of galaxy properties is in rather good agreement with the standard model \\cite{drexel, abbasSheth, blantonBerlind}, so it may be that this will one day provide interesting constraints on $\\alpha$.  This is the subject of work in progress.  In standard gravity, the formation and abundance of voids can be estimated using similar methods to those used for clusters \\cite{voids} -- extending our analysis of the modified potential to voids is also work in progress.  \\bigskip We would like to thank Francis Bernardeau, Antonaldo Diaferio, Lam Hui, Kyle Parfrey and Roman Scoccimarro for helpful discussions and encouragement."
    },
    "0812/0812.2496_arXiv.txt": {
        "abstract": "The disk instability mechanism for giant planet formation is based on the formation of clumps in a marginally-gravitationally unstable protoplanetary disk, which must lose thermal energy through a combination of convection and radiative cooling if they are to survive and contract to become giant protoplanets. While there is good observational support for forming at least some giant planets by disk instability, the mechanism has become theoretically contentious, with different three dimensional radiative hydrodynamics codes often yielding different results. Rigorous code testing is required to make further progress. Here we present two new analytical solutions for radiative transfer in spherical coordinates, suitable for testing the code employed in all of the Boss disk instability calculations. The testing shows that the Boss code radiative transfer routines do an excellent job of relaxing to and maintaining the analytical results for the radial temperature and radiative flux profiles for a spherical cloud with high or moderate optical depths, including the transition from optically thick to optically thin regions. These radial test results are independent of whether the Eddington approximation, diffusion approximation, or flux-limited diffusion approximation routines are employed. The Boss code does an equally excellent job of relaxing to and maintaining the analytical results for the vertical ($\\theta$) temperature and radiative flux profiles for a disk with a height proportional to the radial distance. These tests strongly support the disk instability mechanism for forming giant planets. ",
        "introduction": "High precision Doppler surveys provide the best estimates to date of the frequency of gas giant planets around solar-type stars in the sun's neighborhood of the galaxy. Cumming et al. (2008) analyzed 8 years of HIRES spectrometer data from the Keck Planet Search effort, begun in 1996, taking into account the detection threshold for each star. Cumming et al. (2008) estimated that 17\\% to 20\\% of solar-type stars have planets with masses of 0.3 $M_J$ (i.e., Saturn-mass) and above, orbiting within 20 AU. Considering that this estimate is based on an extrapolation to longer period orbits from detections of planets with orbital periods of 2000 days or less, this frequency estimate may well be a lower bound, if the frequency of gas giant planets continues to increase with orbital period: the frequency of gas giant planets with orbital periods less than 300 days is only 1/5 that of planets with orbital periods between 300 days and 2000 days. Gas giant planets are thought to form primarily with orbital periods greater than 2000 days, so there may well be an as yet unseen extensive population of gas giant planets with orbital periods greater than 2000 days that have not suffered inward migration. Such planets would have formed in regions with relatively short-lived outer disks (e.g., Orion, Carina), while planets formed in regions with relatively long-lived outer disks (e.g., Taurus, Ophiuchus) would have migrated inward (Boss 2005; Eisner et al. 2008).  Given this high frequency of gas giant planets, it is clear that there must be at least one efficient mechanism for gas giant planet formation. Two such mechanisms have been proposed, core accretion (e.g., Mizuno 1980) and disk instability (e.g., Boss 1997), and both appear to be necessary to explain the entire range of planets detected to date. Core accretion is the preferred mechanism for giant planets with very large inferred core masses, such as HD 149026b, which appears to have a core mass of $\\sim$ 70 $M_\\oplus$ and a gaseous envelope of $\\sim$ 40 $M_\\oplus$ (Sato et al. 2005). On the other hand, disk instability would seem to be the preferred mechanism for forming gas giants in very low metallicity systems, such as the M4 pulsar planet, where the metallicity [Fe/H] = -1.5 (Sigurdsson et al. 2003). Disk instability is also able to form gas giants around M dwarf stars (Boss 2006b), unlike core accretion (Laughlin et al. 2004; Ida \\& Lin 2005). Both Doppler and microlensing searches have detected gas giants in orbit around M dwarfs, most recently a 3 $M_J$ planet in the OGLE-2005-BLH-071L microlensing event (Dong et al. 2008). Microlensing has detected a Jupiter/Saturn analog system (Gaudi et al. 2008) around an M dwarf that is consistent with photoevaporative gas loss from an outer gas giant formed by disk instability in a region of massive star formation (Boss 2006b,c). Doppler searches have detected numerous hot super-Earths, accompanied by outer gas giant planets (e.g., HD 181433b and HD 47186b [Bouchy et al. 2008] and HD 40307bcd [Mayor et al. 2008]), a situation similar to that of our solar system and consistent with the ``best of both worlds'' scenario (Boss 1997), where rocky planets form by collisional accumulation inside the orbits of gas giants formed by disk instability.  In spite of this observational evidence, theoretical work on disk instability is split between those whose studies are either supportive (e.g., Boss 1997, 2000, 2001, 2005, 2006a,b,c, 2007, 2008; Mayer et al. 2002, 2004, 2007) or dismissive (e.g., Pickett et al. 2000; Cai et al. 2006, 2008; Boley et al. 2006, 2007a,b; Rafikov 2005, 2007; Stamatellos \\& Whitworth 2008) of disk instability being able to form giant planets. This divergence appears to be the result of a variety of modeling issues (Nelson 2006; Boss 2007), such as numerical spatial resolution, gravitational potential solver accuracy, artificial viscosity, stellar irradiation, radiative transfer, and spurious numerical heating. A combination of joint code test cases (e.g., Boss 2007) and rigorous testing is needed to try to resolve these differences.  Boley et al. (2006, 2007b) presented the results of a series of tests of their radiative hydrodynamics code on a ``toy problem'' (a plane-parallel grey atmosphere) with sufficient assumptions to permit an analytical solution for the vertical temperature and radiative flux profiles. This toy problem is well-suited to their cylindrical coordinate code, as their vertical cylindrical coordinate can be used to simulate a plane-parallel atmosphere. However, the Boley et al. (2006, 2007b) test cases are unsuitable for a code based on spherical coordinates (Boss 2008), as is the situation for the Boss \\& Myhill (1992) code used in Boss (1997, 2000, 2001, 2005, 2006a,b,c, 2007, 2008).  Boley et al. (2006) ``... challenge all researchers who publish radiative hydrodynamics simulations to perform similar tests or to develop tests of their own and publish the results.'' This purpose of this paper is to address this challenge by the latter route, as the former route is precluded. Two analytical solutions for radiative transfer in spherical coordinates are first derived and then used to test extensively the spherically symmetric (one dimensional, $r$) and azimuthally symmetric (two dimensional, $r$, $\\theta$) versions of the Boss \\& Myhill (1992) three dimensional code ($r$, $\\theta$, $\\phi$), the same code that has been used in all of the author's radiative hydrodynamics models of disk instability. ",
        "conclusions": "The primary motivation for this paper was to answer the challenge issued by Boley et al. (2006) for other workers with radiative hydrodynamics codes ``... to develop tests of their own and publish the results.'' The new test results show that the Boss \\& Myhill (1992) radiative hydrodynamics code does a superb job of relaxing to and maintaining  the analytical solutions for a spherical cloud or an axisymmetric disk derived here. This excellent performance is independent of whether the Eddington approximation, diffusion approximation, or flux-limited diffusion approximation is employed for the cloud, as well as of the optical depth assumed. Considering that typical disk instability calculations with the diffusion approximation last for only $\\sim 10^3$ yr, the fact that these models show that the radiative transfer scheme is highly accurate over time scales of at least $\\sim 10^6$ yr is reassuring for the mechanism of giant planet formation by disk instability.  These models presented here test only the radiative transfer routines and other thermodynamical aspects of the Boss codes, not the coupling between these processes and the hydrodynamics that occurs in full disk instability models. Ideally, one would test the full radiative hydrodynamics codes against analytical solutions. In the absence of such solutions, one can test the codes with respect to their ability to represent convective motions, which do involve a coupling of hydrodynamics and thermodynamics. Boss (2004) analyzed in detail the convective stability of his disk instability models, and found a good agreement between where transient, convective-like upwellings and downwellings occurred and where the Schwarzschild criterion for convection was met. Boley et al. (2007b) presented the results of several tests for convection in their codes, finding that convection occurred when it should have and did not occur when it should not have. Convection and convective-like motions thus appear to be appropriately modeled by both the Boss and Boley et al. codes.  Boss \\& Myhill (1992) described a variety of other tests to which the code has been subjected, including the standard nonisothermal test case for protostellar collapse (tested on two different codes by Myhill \\& Boss 1993), whose results have since been confirmed by Whitehouse \\& Bate (2006). Further tests of the Boss \\& Myhill (1992) code have been presented as follows: spatial resolution (Boss 2000, 2005); gravitational potential solver (Boss 2000, 2001, 2005), artificial viscosity (Boss 2006a); and radiative transfer (Boss 2001, 2007, 2008). Given the ongoing theoretical debate over the viability of disk instability for giant planet formation, it will continue to be important for other workers to conduct their own tests of these key numerical issues."
    },
    "0812/0812.0346_arXiv.txt": {
        "abstract": "{Nearby edge-on galaxies showing a synchrotron halo are nearly ideal objects for studying the transport of cosmic rays (CRs) in galaxies. Among them, the nearby starburst galaxy NGC\\,253 hosts a galactic wind indicated by various ISM phases in its halo.} {The diffusive and convective CR transport from the disk into the halo is investigated using the local CR bulk speed. The connection between the CR transport and the galactic wind is outlined.} {We observed NGC\\,253 with the VLA at $\\lambda 6.2\\,{\\rm cm}$ in a mosaic with 15 pointings. The missing zero-spacing flux density of the VLA mosaic was filled in using observations with the 100-m Effelsberg telescope. We also obtained a new $\\lambda 3.6\\,{\\rm cm}$ map from Effelsberg observations and reproduced VLA maps at $\\lambda 20\\,{\\rm cm}$ and $\\lambda90\\,{\\rm cm}$. The high dynamic range needed due to the strong nuclear point-like source was addressed with a special data calibration scheme for both the single-dish and the interferometric observations.} {We find a thin and a thick radio disk with exponential scaleheights of $0.3\\,{\\rm kpc}$ and $1.7\\,{\\rm kpc}$ at $\\lambda 6.2\\,{\\rm cm}$. The equipartition total magnetic field strength between $7\\,{\\rm\\mu G}$ and $18\\,{\\rm\\mu G}$ in the disk is remarkably high. We use the spectral aging of the cosmic ray electrons (CREs) seen in the vertical profiles of the spectral index to determine a lower limit for the global CR bulk speed as $170\\pm70\\,{\\rm km\\, s^{-1}}$. The linear correlation between the scaleheights and the CRE lifetimes, as evident from the dumbbell shaped halo, requires a vertical CR transport with a bulk speed of $300\\pm 30\\,{\\rm km\\, s^{-1}}$ in the northeastern halo, similar to the escape velocity of $280\\,\\rm km\\,s^{-1}$. This shows the presence of a ``disk wind'' in NGC\\,253. In the southwestern halo, the transport is mainly diffusive with a diffusion coefficient of $2.0\\pm0.2\\times10^{29}\\,\\rm cm^2\\,s^{-1}$.} {Measuring the radio synchrotron scaleheight and estimating the CRE lifetime allow us to determine the bulk speed of the CR transport into the halo. The transport is convective and more efficient in the northeastern halo, while it is diffusive in the southwestern halo. The luminous material is transported by the disk wind, which can explain the different amounts of extra-planar \\ion{H}{I}, H$\\alpha$, and soft X-ray emission in the two halo parts. Future low-frequency radio observations will provide the data to analyze the vertical velocity profile of galactic winds.} ",
        "introduction": "NGC\\,253 is a prototypical starburst galaxy at a distance of 3.94\\,Mpc \\citep{karachentsev_03a}. It is a late-type spiral galaxy with a high inclination angle of $78.5^\\circ\\pm0.5^\\circ$ \\citep{pence_80a}. This allows observations of extra-planar components of the interstellar medium (ISM) at different wavelengths: ROSAT Position Sensitive Proportional Counter (PSPC) observations have revealed a very extended halo of soft diffuse X-ray emission extending into the halo to a distance of $8\\,{\\rm kpc}$ from the disk. The prominent lobes of soft X-ray emission show a distinct ``X''-shaped pattern centered on the nucleus \\citep{pietsch_00a}. Moreover, H$\\alpha$ observations by \\citet{hoopes_96a} show plumes emerging from the disk into the halo, which spatially coincide with the lobes of the soft X-ray emission. Observations of \\citet{boomsma_05a} in \\ion{H}{I} show large plumes of extra-planar \\ion{H}{I} emission enclosing the soft X-ray emission.  Given the different phases of the ISM in the halo of NGC\\,253, the question arises of how the observed halo structure was formed. The distribution of the different ISM phases in shells has led to the suggestion that the halo can be considered as one huge bubble that is fueled by an active star-forming region around the nucleus of the galaxy \\citep{strickland_02a}. This scenario is supported by the finding of a small plume of hot X-ray gas in CHANDRA maps associated with a conical outflow from the nuclear starburst \\citep{strickland_00a}. Spectroscopic measurements of H$\\alpha$-emitting gas in the southern nuclear outflow cone by \\citet{schulz_92a} give an outflow velocity of $390\\,{\\rm km\\, s^{-1}}$. Numerical models by \\citet{suchkov_94a} follow the evolution of a multi-phase halo powered by a compact source located in the center of the disk. A close resemblance of the ``X''-shaped pattern of the soft X-ray emission was indeed found, which could be interpreted as the termination shock of the hot gas where the gas cools effectively.  A possible explanation for the amount of X-ray and H$\\alpha$-emitting gas in the halo is the dissipation of kinetic energy via shock heating \\citep{lehnert_96a} if the thermalization fraction is high enough \\citep{strickland_02a}. A precondition for this model is the existence of cool gas in the halo that condenses to clouds from which the shock can form. ISOPHOT observations discussed by \\citet{radovich_01a} show ``X''-shaped far-infrared emission overlapping with the H$\\alpha$ and X-ray emission, indicating the existence of such a cool gas. However, the origin of the gas is not clear. It could be a preexisting population of cool gas in the halo, possibly accreted from a merger with a satellite. Another possibility is that the gas of the disk can be transported into the halo along the walls of the cavity as Kelvin-Helmholtz instabilities rip off clumps of cool gas in the disk \\citep[see e.g.][]{suchkov_94a, heckman_00a}. Summarizing, it is important to understand how the gas is transported from the disk into the halo.\\\\  Cosmic rays (CRs) and magnetic fields are important for the evolution of the ISM, because they may possess an energy density comparable to that of the thermal gas \\citep{beck_96a}. They are even more important for the study of galactic winds, because the thermal gas itself can drive a galactic wind only when the heating processes are sufficiently strong and the gravitational potential is comparably weak. A good example of this situation is the edge-on galaxy M\\,82, which possesses a huge halo of hot X-ray, H$\\alpha$, and radio continuum emission. However, the hot gas cools radiatively preventing the escape from the gravitational potential and falls back onto the disk in condensed clouds. This {\\it galactic fountain} scenario proposed by \\citet{field_69a} might explain the infall of the so-called High Velocity Clouds that could be interpreted as cooled remnants of the hot gas.  If we incorporate the CR pressure the galactic wind can be sustained against the radiative losses \\citep{ipavich_75a}, because the CRs suffer only little from radiative losses (mainly via synchrotron losses of the electrons). As the protons carry most of the energy this does not influence the CR energy density. Thus the pressure support from below can drive a galactic wind against the gravity of the underlying galactic disk. \\citet{breitschwerdt_91a,breitschwerdt_93a} started a series of papers treating the CR transport along ``flux density tubes'' in a one-dimensional model. This model was extended to a full three-dimensional treatment (with rotational symmetry) by \\citet{zirakashvili_96a}, including the influence of galactic rotation. It was found that CRs can drive a galactic wind under a broad variety of conditions in the disk and that the loss of mass and momentum over the lifetime of a galaxy is significant \\citep{ptuskin_97a}. A recent study by \\citet{everett_08a} showed that the extra-planar X-ray emission in the Milky Way can be explained by a CR driven galactic wind.  The theory of a CR driven galactic wind can be constrained with radio continuum studies: as the relativistic CRs are the most pervasive component of the ISM, they are expected to be associated with the other phases in the halo of NGC\\,253. Indeed, a radio halo around NGC\\,253 was detected by \\citet{beck_79a} using the 100-m Effelsberg telescope at $\\lambda 3.6\\,{\\rm cm}$. However, no correction for the sidelobes of the strong nuclear point-like source was applied and thus it was not possible to disentangle the intrinsic halo emission from the contribution by the sidelobes. The study was extended later with VLA interferometric observations by \\citet{carilli_92a}, which show an extended radio halo over the length of the optical disk very different in shape from the halo seen in X-rays. The finding of a large ``radio spur'' southeast of the nucleus already led to speculations of an outflow connected with the spur.  Extending the observations to other wavelengths and including polarization, a prediction of the magnetic field structure in the disk was presented by \\citet{beck_94a}. The sensitivity of the observations, however, was not sufficient to get information for the extra-planar polarized emission so that the magnetic field structure in the halo could not be determined. With the new VLA mosaic at $\\lambda 6.2\\,{\\rm cm}$ we are able to increase the sensitivity at short wavelengths, allowing us to study in detail the extra-planar emission and possible relations to the outflow of hot gas. The advantage is that the high synchrotron losses at this short wavelength impose tighter constraints on the CR transport from the disk into the halo. Moreover, for the first time sensitive polarization measurements were carried out to study the large-scale magnetic field structure in the halo of NGC\\,253, which will be addressed in a subsequent paper (thereafter Paper~II).  This paper is organized as follows: in Sect.~\\ref{sec:cr_observations} we introduce our observations and present the data reduction. We present our radio continuum maps along with a discussion of the morphology in Sect.~\\ref{sec:cr_morphology}. The distribution of the radio spectral index is presented in Sect.~\\ref{sec:cr_si}. The finding of an extended radio halo with a thick radio disk leads to the determination of the vertical CR bulk speed (Sect.~\\ref{sec:cr_transport}). In Sect.~\\ref{sec:cr_discussion} we discuss the impact of our findings on the feeding of the radio halo with CRs from the star-forming disk. Finally, we summarize our results and present our conclusions in Sect.~\\ref{sec:cr_conclusions}. ",
        "conclusions": " \\begin{enumerate} \\item The scaleheight of the thick radio disk in the northeastern halo is a linear function of the CRE lifetime in the underlying disk at $\\lambda90\\,\\rm cm$, $\\lambda 20\\,\\rm cm$, and $\\lambda 6.2\\,\\rm cm$. This can be understood if the CR transport proceeds preferentially in the vertical direction. \\item The steepening of the spectral index, which is a linear function of $z$, can be explained by strong synchrotron losses of the CREs. The contribution of ionization losses and bremsstrahlung losses can be neglected. The adiabatic losses are of the same order of magnitude as the synchrotron losses at $\\lambda 20\\,\\rm cm$ but dominate at $\\lambda 90\\,\\rm cm$. \\item We computed the CR bulk speed as the ratio of the total power scaleheight and the CRE lifetime. The CR bulk speed in the northeastern halo of $300\\pm 30\\,{\\rm km\\, s^{-1}}$ is similar to the escape velocity of $280\\,{\\rm km\\,s^{-1}}$, indicating the existence of a disk wind which escapes from the gravitational potential. The local CR bulk speed is independent of the CR energy, the SFR, and the magnetic field strength. \\item The CR transport is mainly convective in the northeastern halo while in the southwestern halo the diffusion of CRs is significant. We measure an average diffusion coefficient of $2.0\\pm0.3\\times10^{29}\\,\\rm cm^2\\,s^{-1}$ \\item The disk wind is the driving source for the luminous gas in the halo. In the convective northeastern halo the gas is transported more efficiently from the disk into the halo than in the diffusive southwestern halo. This can explain the different amounts of extra-planar gas in these two halo parts as observed in \\ion{H}{I}, H$\\alpha$, and soft X-rays. \\end{enumerate} The CR transport is always the superposition of convection and diffusion. The data presented here provide a basis for more detailed calculations using the combined diffusion and convection equation. These calculations should include vertical profiles of the magnetic field strength, the CR number density, and the velocity of the galactic wind."
    },
    "0812/0812.1495_arXiv.txt": {
        "abstract": "{} {By pushing the magnitude limit of high proper motion surveys beyond the limit of photographic Schmidt plates we aim at the discovery of nearby and very fast low-luminosity objects of different classes: cool white dwarfs (CWDs), cool subdwarfs (sd), and very low-mass stars and brown dwarfs at the very faint end of the main sequence (MS). } {The deep multi-epoch Sloan Digital Sky Survey data in a 275 square degrees area along the celestial equator (SDSS stripe 82) allow us to search for extremely faint ($i>21$) objects with proper motions larger than 0.14~arcsec/yr. A reduced proper motion diagram $H_z/(i-z)$ clearly reveals three sequences (MS, sd, CWD) where our faintest candidates are representative of the still poorly known bottom of each sequence. We classify 38 newly detected objects with low-resolution optical spectroscopy using FORS1 @ ESO VLT. Together with our targets we observe six known L dwarfs in stripe 82, four (ultra)cool sd and one CWD as comparison objects. Distances and tangential velocities are estimated using known spectral type/absolute magnitude relations. } {All 22 previously known L dwarfs (and a few of the T dwarfs) in stripe 82 have been detected in our high proper motion survey. However, 11 of the known L dwarfs have smaller proper motions (0.01$<$$\\mu$$<$0.14~arcsec/yr). Although stripe 82 was already one of the best investigated sky regions with respect to L and T dwarfs, we are able to classify 13 new L dwarfs. Two previously known L dwarfs have been reclassified by us. We have also found eight new M7.5-M9.5 dwarfs. The four new CWDs discovered by us are about 1-2 mag fainter than those previously detected in SDSS data. All new L-type, late-M and CWD objects show thick disk and halo kinematics. {Since our high-velocity late-M and L dwarfs do not show indications of low metallicity in their spectra, we conclude that there may be a population of ultracool halo objects with normal metallicities. There are 13 objects, mostly with uncertain proper motions, which we initially classified as mid-M dwarfs. Among them we have found 9 with an alternative subdwarf classification (sdM7 or earlier types), whereas we have not found any new spectra resembling the known ultracool ($>$sdM7) subdwarfs. Some M subdwarf candidates have been classified based on spectral indices with large uncertainties.} We failed to detect new nearby ($d<50$~pc) L dwarfs, probably because the SDSS stripe 82 area was already well-investigated before. With our survey we have demonstrated a higher efficiency in finding Galactic halo CWDs than previous searches. The space density of halo CWDs is according to our results about 1.5-3.0 $\\times$ 10$^{-5}$~pc$^{-3}$. } {} ",
        "introduction": "Deep sky surveys are not only important for the study of the most distant galaxies in the Universe but also for near-field stellar statistics and physics. It has been realised for a long time that the immediate Solar neighbourhood is the only region of a galaxy where there is a real chance for a complete registration of its stellar and sub-stellar content reaching to the lowest-luminosity objects. A complete volume-limited sample provides the groundwork for the knowledge of the local luminosity function, the mass function of field stars and brown dwarfs, and the Galactic star formation history (Reid et al.~\\cite{reid02}; Gizis et al.~\\cite{gizis02}; Cruz et al.~\\cite{cruz03}). Moreover, the nearest representatives of a given class of objects can be considered as benchmark sources allowing detailed follow-up studies, including the search for possible planetary companions.  However, our knowledge of the Solar neighbourhood is still remarkably incomplete: Within a 25~pc horizon more than 60\\% of the stars (Henry et al.~\\cite{henry02}) and probably more than 90\\% of the brown dwarfs remain to be discovered. The Research Consortium on Nearby Stars (RECONS) has documented the recent progress achieved for the 10~pc sample\\footnote{http://www.chara.gsu.edu/RECONS/census.posted.htm}.  It should be noted that this sample is also not complete since it neglects recently discovered sources likely within 10~pc but without accurately measured parallaxes. It is defined as all known objects with trigonometric parallaxes of 100~mas or more, and an error in that parallax of less than 10~mas. An increase of about 20\\% has been achieved since 2000, so that the census has reached about 350 objects. About 240 of them are M dwarfs, which also contributed to the increase with the largest number (+40). The only other object classes with improved statistics are the L/T dwarfs and the planetary companions around nearby stars. The number of white dwarfs (18) is stagnating, but there are ongoing efforts to find missing nearby white dwarfs (Subasavage et al.~\\cite{subasavage04}; \\cite{subasavage08}). Subdwarfs represent the smallest group with only 4 objects in the 10~pc sample (Henry et al.~\\cite{henry06}), which all have been known for a long time, and which all have extremely large proper motions (3.7$<$$\\mu$$<$8.7~arcsec/yr).  Before time-consuming trigonometric parallax measurements can be started, one needs to know good nearby candidates of low-luminosity stars and substellar objects. Most nearby stars have been first detected in high proper motion (HPM) surveys (e.g. Luyten~\\cite{luyten79a,luyten79b}; Luyten \\& Hughes~\\cite{luyten80}). HPM objects represent a mixture of {\\it very nearby} neighbours of the Sun in the local population of the Galactic thin disk, and {\\it very fast} representatives of the Galactic thick disk and halo, just passing through the neighbourhood. Although halo stars are relatively rare when compared to the number of disk stars in a local volume-limited sample, they are over-represented in HPM samples. The main observational material used in Luyten's HPM surveys were photographic Schmidt plates. Due to their full sky coverage in different colours and over long time intervals these archival data continue to play an important role in the search for nearby candidates. In particular, the deeper and partly redder ($I$-band) second epoch Schmidt survey plates have by far not yet been fully exploited to their limiting magnitudes.  The SuperCOSMOS Sky Surveys (SSS) represent one of the best available photographic archives suitable for the identification of hitherto unknown faint HPM stars, since they provide multi-colour catalogue data with cross-matching between the different passbands and epochs as well as images from the individual digitised photographic plates taken with three different Schmidt telescopes (Hambly et al.~\\cite{hambly01a,hambly01b,hambly01c}). The comparison of optical SSS data with near-infrared Two Micron All Sky Survey (2MASS; Skrutskie et al.~\\cite{skrutskie06}) led to the discovery of new faint HPM objects with $\\mu>$ 1 arcsec/yr including some of the nearest brown dwarfs (Scholz et al.~\\cite{scholz03}; Kendall et al.~\\cite{kendall07}; Folkes et al.~\\cite{folkes07}), cool subdwarfs (Scholz et al.~\\cite{scholz04a,scholz04b}), late-type M dwarfs (Pokorny et al.~\\cite{pokorny03}; Hambly et al.~\\cite{hambly04}), and cool white dwarfs (Scholz et al.~\\cite{scholz02}; Hambly et al.~\\cite{hambly04}).  New near-infrared sky surveys like the 2MASS and the DEep Near-Infrared Survey (DENIS; Epchtein et al.~\\cite{epchtein97}) and deep multi-colour optical surveys like the Sloan Digital Sky Survey (SDSS; York et al.~\\cite{york00}) uncovered large numbers of objects cooler than the latest-type M dwarfs. New spectral classes, L (2500~K$>$$T_{eff}$$>$1500~K) and T (1500~K$>$$T_{eff}$$>$700~K), describing these ultracool objects, had to be defined less than ten years ago (Mart\\'{\\i}n et al.~\\cite{martin99}; Kirkpatrick et al.~\\cite{kirkpatrick99}; Burgasser et al.~\\cite{burgasser02} Geballe et al.~\\cite{geballe02}). Since then almost 650 L and T dwarfs have been discovered, mostly in 2MASS ($\\approx$50\\%) and SDSS ($\\approx$30\\%) (see Gelino et al.~\\cite{gelino08}, and references therein). About 50\\% of the L dwarfs and all T dwarfs in the Solar neigbourhood are brown dwarfs, whereas among late-M dwarfs only a few are sub-stellar objects (see review by Kirkpatrick~\\cite{kirkpatrick05}). Note that one of the latter is LP\\,944-20 (Tinney~\\cite{tinney98}), the only isolated brown dwarf catalogued already by Luyten~(\\cite{luyten79b}; Luyten \\& Hughes~\\cite{luyten80}) as an HPM object, long before the first free-floating brown dwarfs were discovered from a new HPM survey by Ruiz et al.~(\\cite{ruiz97}) and from DENIS (Delfosse et al.~\\cite{delfosse97}).  Cool, low-mass subdwarfs are interesting as tracers of Galactic structure and chemical evolution. They are metal-poor and typically exhibit halo kinematics (hence they show up preferentially in HPM surveys). The current classification scheme of cool subdwarfs ends at sdM7 (Gizis~\\cite{gizis97}). However, more and more cooler subdwarfs are being discovered in recent years (e.g. L{\\'e}pine et al.~\\cite{lepine03}; Burgasser et al.~\\cite{burgasser03}; Scholz et al.~\\cite{scholz04a,scholz04b}; Burgasser \\& Kirkpatrick~\\cite{burgasser06b}). A systematic search in the SDSS spectroscopic database has uncovered 23 new subdwarfs with spectral sub-types of M7 or later and more than doubled the number of known ultracool subdwarfs (L{\\'e}pine \\& Scholz~\\cite{lepine08}). There is some evidence that these so-called ultracool subdwarfs are somewhat hotter compared to normal dwarfs of the same spectral sub-type. Burgasser et al.~(\\cite{burgasser07}) derived 3100~K$>$$T_{eff}$$>$2400~K for subdwarfs with spectral sub-types between M7.5 and L4 from comparison with model spectra. An extension of the classification scheme for the ultracool subdwarfs, including sub-stellar subdwarfs, is now developing, where the spectral features indicating subtle changes in their  physical properties (metallicity, surface temperature/clouds, gravity) are still under debate (Kirkpatrick~\\cite{kirkpatrick05}; Gizis \\& Harvin~\\cite{gizis06}; L{\\'e}pine et al.~\\cite{lepine07}; Burgasser et al.~\\cite{burgasser07}).  A third class of cool objects, completely different in their evolutionary stage and physical parameters from low-mass stars/brown dwarfs and cool subdwarfs, but also hard to detect due to their faintness, are the cool white dwarfs (CWDs). CWDs have long been considered as important objects for a number of reasons, such as determining the age of the Galactic disk (Leggett et al.~\\cite{leggett98}) and as possible components of the Galactic dark matter halo (Ibata et al.~\\cite{ibata00}; Oppenheimer et al.~\\cite{oppenheimer01}; M{\\'e}ndez~\\cite{mendez02}; Flynn et al.~\\cite{flynn03}; Salim et al.~\\cite{salim04}; Reid~\\cite{reid05}). Among the coolest known CWDs ($T_{eff}$$<$4500~K) in the Solar neighbourhood there are many objects already catalogued in the LHS (Luyten~\\cite{luyten79a}) or discovered as HPM objects from photographic plates (Hambly et al.~\\cite{hambly97}; Harris et al.~\\cite{harris99}; Ibata et al~\\cite{ibata00}; Scholz et al.~\\cite{scholz00}; Monet et al.~\\cite{monet00}; Scholz et al.~\\cite{scholz02}; L{\\'e}pine et al.~\\cite{lepine05}; Carollo et al.~\\cite{carollo06}; Rowell et al.~\\cite{rowell08}). Other CWDs have been found in the SDSS (Harris et al.~\\cite{harris01}; Gates et al.~\\cite{gates04}; Kilic et al.~\\cite{kilic06}; Vidrih et al.~\\cite{vidrih07}; Harris et al.~\\cite{harris08}). Hall et al.~(\\cite{hall08}) have recently found a nearby halo CWD with a large proper motion detected by combining the photographic USNO-B1.0 catalogue (Monet et al.~\\cite{monet03}) with the SDSS (Munn et al.~\\cite{munn04}). The space density of CWDs in the Galactic halo, possibly responsible for observed microlensing events, is still under debate (see e.g. Torres et al.~\\cite{torres08}). Reid et al.~(\\cite{reid01} had shown that the halo CWDs detected by Oppenheimer et al.~(\\cite{oppenheimer01}) in a HPM survey based on photographic Schmidt plates are in fact members of the thick disk. We conclude that a deeper HPM survey could either detect cooler objects or uncover the CWDs with even higher velocities, i.e. the real halo members.  As we have shown above, extending HPM surveys to fainter magnitude limits has led to the discovery of many benchmark sources of new classes of cool objects. Several new large-area HPM surveys relying {\\it only} on modern deep optical (SDSS) and/or near-infrared (2MASS and SIMP) data have been successfully started (Artigau et al.~\\cite{artigau06}; Looper et al.~\\cite{looper07}; Vidrih et al.~\\cite{vidrih07}; Bramich et al.~\\cite{bramich08}; Metchev et al.~\\cite{metchev08}). The UKIRT Infrared Deep Sky Survey (UKIDSS; Lawrence et al.~\\cite{lawrence07}) will eventually provide its own second epoch data suitable for HPM searches for the coolest brown dwarfs. The large-area survey of UKIDSS overlaps with SDSS, hence these different epoch data enable already deep HPM searches.  In this paper we present the first results of one of the deepest optical HPM searches done so far over a relatively large sky area.  It is based on multi-epoch SDSS data in an equatorial stripe (called stripe 82) of about 275 square degrees. In Sect.~\\ref{sec2} we describe our HPM survey and the selection of targets for follow-up observations. Sect.~\\ref{sec3} deals with the spectroscopic classification of our extremely faint HPM objects. In Sect.~\\ref{sec4} we discuss their distance estimates and the resulting kinematics. In Sect.~\\ref{sec5} we draw some conclusions and give an outlook on further work needed.  \\begin{figure*}[!] \\centering \\includegraphics[angle=270,width=19cm]{s82hpmfig1n.ps} \\caption{Reduced proper motion diagram for objects with proper motions exceeding 0.14~arcsec/yr: all initially found objects after step 1) of our HPM survey (dots) with three separate sequences (MS, sd, CWD) marked, known L dwarfs (filled lozenges) and T dwarfs (filled squares) falling in the S82 survey area, M9 dwarfs (stars), L-type subdwarfs (filled triangles), a late-M extreme (esdM6.5) subdwarf (open triangle), CWDs (open circles and squares), and the confirmed faintest HPM objects after step 2) and 3) of our survey (crosses). Symbols overplotted by large open circles mark objects with uncertain proper motions (see Tab.~\\ref{tabnewhpm} and Tab.~\\ref{tabLTpm}). } \\label{fig1}% \\end{figure*} ",
        "conclusions": "\\label{sec5}  \\begin{enumerate} \\item We have carried out a HPM survey of faint objects in S82 using their individual epoch positions from 1998 to 2004. The 38 faintest HPM objects ($i$$>$21, $\\mu$$>$0.14~arcsec/yr) have been selected (together with some comparison objects) for our spectroscopic follow-up observations. Very recently, in an independent attempt to find moving sources in S82, Lang et al.~(\\cite{lang08}) used a method of fitting a model of a moving point source to all S82 imaging data. Their initial candidate list was selected from the S82 ``Co-add catalog'' (J. Annis et al., in prep.) aiming at the detection of HPM objects that are too faint to detect at any individual epoch. Among their 19 HPM brown dwarf candidates (with proper motions between 0.07 and 0.65~arcsec/yr) there are 10 previously known objects (nine L dwarfs and one T dwarf, which were all detected in our survey, too. Among the remaining objects of Lang et al. there are only four with proper motions $\\mu>0.14$~arcsec/yr. Two of the latter (SDSS J011014.30$+$010619.1 and SDSS J234841.38$-$004022.1) were also on our target list and classified as L4 and L1.5, respectively. We conclude that most of their objects are still detectable at individual epochs as we used in our HPM survey. \\item The low-resolution (and relatively low signal-to-noise) spectra of our faint targets observed with FORS1@VLT allowed us to classify them as early-L (13 objects) , late-M (8), mid-M (13) dwarfs, and CWDs (4). In addition, two previously known L dwarfs were reclassified by us. None of the 38 spectra looked like those of the three ultracool ($>$sdM7) subdwarf spectra observed by us for comparison, but 9 out of 13 mid-M dwarfs were alternatively classified as subdwarfs (sdM7 and earlier types) from a comparison with additional subdwarf templates. Relatively uncertain spectral indices measurements do also hint at a possible subdwarf nature of seven M dwarfs. One mid-M dwarf is likely a very low-metallicity (usdM) object or may have an unresolved white dwarf companion. \\item The number of L dwarfs in the S82 area, which was already before our study one of the best investigated fields in the sky, has been further increased (from 22 to 35). The newly added L dwarfs are preferentially thick disk and halo objects as are the new late-M dwarfs and CWDs found in S82. The nearest new object is an L4 dwarf (SDSS J011014.30$+$010619.1) at 64~pc. All other new M and L dwarfs have spectroscopic distance estimates above 100~pc. \\item The re-discovery of all formerly known L dwarfs and even of some T dwarfs in S82, including the L8 dwarf SDSSp J010752.33$+$004156.1 and the T4.5 dwarf SDSS J020742.48$+$000056.2 with trigonometric parallaxes (Vrba et al.~\\cite{vrba04}) placing them at 16 and 29~pc, respectively, shows that our HPM survey was sensitive to discover very nearby objects, too. {\\it However, we failed to detect additional very nearby objects.} For two previously known objects reclassified by us, the L4.5 dwarf 2MASS J00521232$+$0012172 and the L9 dwarf SDSSp J023617.93$+$004855.0, we have determined new spectroscopic distances of 53~pc (formerly 81~pc) and 18~pc (formerly 21~pc), respectively. \\item On the one hand, we may have been just unlucky in finding new L-type neighbours of the Sun since our survey covered only 1/150 part of the sky and the total number of missing L dwarfs in the Solar neighbourhood (within about 50~pc) may be of the order of several hundreds. On the other hand, the S82 area had already been well-investigated for the apparently brighter and consequently more nearby L dwarfs. Using the L and T dwarf compendium by Gelino et al.~(\\cite{gelino08}) and the $M_J$/spectral type relation given in Dahn et al.~(\\cite{dahn02}), we have estimated that about 400 of the known L dwarfs have spectroscopic distances of less than 50~pc. 11 of them (seven L0-L4.5 and four L5.5-L9.5 dwarfs) lie in the S82 area, which corresponds to a 4 times higher surface density than the average. Therefore, we conclude that the census of nearby L dwarfs in S82 was probably already near-complete before our study. The space density of L dwarfs determined from the 11 objects within 50~pc falling in the S82 area is about 0.003~pc$^{-3}$. This is comparable with the value we get from the probably more complete list of about 200 L dwarfs with spectroscopic distances within 25~pc in Gelino et al.~(\\cite{gelino08}), again using the Dahn et al.~(\\cite{dahn02}) $M_J$/spectral type relation. \\item With our deep HPM survey we have discovered preferentially high-velocity late-M and L dwarfs. However, the optical spectra of these objects do not show indications of a lower metallicity typical of ultracool subdwarfs. As long as we have no infrared data on the new fast late-M and L dwarfs, we can not say whether they are unusually blue (according to specific near-infrared spectral features) or blue photometric outliers (acording to their $J-K_s$ colours) (cf. Faherty et al.~\\cite{faherty08}). Our findings suggest that there may exist a population of normal-metallicity ultracool halo objects. \\item 10 of the 38 HPM objects have uncertain proper motions. Most of these uncertain proper motion objects turned out to be mid-M dwarfs at large distances. Most of them have been alternatively classified as subdwarfs or can be considered as subdwarf candidates, which may in fact have large velocities. However, the extremely large tangential velocities (700...2500~km/s) of nine objects suggest an overestimation of their proper motions. A more careful investigation of the kinematics of these sources makes sense only after a verification of their proper motions with additional epoch measurements. When the large proper motions will be confirmed we would also need higher resolution and signal-to-noise spectra to investigate their metallicities and radial velocities. \\item Follow-up observations (better spectroscopy and near-infrared photometry allowing a fit of the observed spectral energy distribution with CWD atmosphere models for determining the effective temperatures) are needed to better characterise the four new CWDs. Trigonometric parallax measurements are challenging in view of their faint magnitudes and possible large distances, but would provide the ultimate values of their absolute magnitudes. Their estimated distances of more than 100~pc are based on the assumption that their absolute magnitudes compare with other CWDs detected in the SDSS. \\item If we consider the optimistic case that all four new CWDs are Galactic halo members within 170~pc, then we determine a space density of 2.9 $\\times$ 10$^{-5}$~pc$^{-3}$ {\\it for halo CWDs alone}, which is almost exactly the same number as the density of 3.0 $\\times$ 10$^{-5}$~pc$^{-3}$ derived by Gates et al.~(\\cite{gates04}) for a mix of {\\it all} CWDs in various Galactic components. Even with the conservative assumtion that only two of our CWDs are representatives of the halo population the estimated space density of halo CWDs reaches already 50\\% of the above cited value. Nevertheless, our derived space density of halo CWDs is consistent with other estimates indicating that ancient CWDs do not contribute significantly to the baryonic fraction of the dark halo of the Galaxy (see Carollo et al.~\\cite{carollo07} and references therein). Our estimate of the density of halo CWDs corresponds to less than 0.4\\% of the local dark matter density (0.0079~$M_{\\odot}$pc$^{-3}$) adopted by Torres et al.~(\\cite{torres08}). \\end{enumerate}"
    },
    "0812/0812.3729_arXiv.txt": {
        "abstract": "The standard lower limit for the mass of white dwarfs (WDs) with a C/O core is roughly 0.5 M$_{\\odot}$. In the present work we investigated the possibility to form C/O WDs with mass as low as 0.33 M$_{\\odot}$. Both the pre-WD and the cooling evolution of such nonstandard models will be described. ",
        "introduction": "It is by now a well-known and commonly accepted result of the stellar evolution theory that stars of low and intermediate mass will end their lives as white dwarfs (WDs). Since the pioneering papers by Chandrasekhar (1931) and Stoner (1932), it has been known that there must be an upper limit of the mass of WDs. The current value for such a limit, now universally known as Chandrasekhar's mass, is  M$_{Ch}$= 1.456 (2/$\\mu_e$)$^2$ M$_{\\odot}$(where $\\mu_e$ is the electronic mean molecular weight, Hansen, Kawaler \\& Trimble 2004). Below this critical mass there are essentially three different families of WDs concerning the chemical composition of their cores: the O/Ne/Mg, C/O and He. The upper mass limit for C/O WDs, which is still subject to some debate, since it is the aftermath of the evolution of a star with initial mass slightly lower than M$_{up}$, the lowest mass for star that succeeds in igniting carbon burning before the asymptotic giant branch (AGB), it is around 1.05 M$_{\\odot}$(Dominguez et al. 1999, Weideman 2000, Prada Moroni \\& Straniero 2007, Catalan et al. 2008, Meng, Cheng \\& Han 2008).  On the other hand, the value for the minimum mass of WDs with a C/O core is commonly believed to be around 0.5 M$_{\\odot}$ (Weideman 2000, Meng et al. 2008). The theory of stellar evolution predicts that an isolated single star with mass lower than about 0.5 M$_{\\odot}$, the exact value depending on the chemical composition (0.46 and 0.50 for Z=0.04 and Z=0.0001, respectively, Dominguez et al. 1999, Girardi 1999), do not ignite the helium-burning and die as WDs with an helium rich core. For these stars, during the red giant phase the electronic thermal conduction and the emission of neutrinos, which cool the core, prevail on the heating due to the contraction caused by the accretion of fresh helium processed by the hydrogen-burning shell. This is the main reason why WDs of mass less than about 0.5 M$_{\\odot}$ are commonly believed to have a He-core.  On the other hand, as early as 1985 Iben \\& Tutukov, in a very beautiful and instructive paper, showed that a star with an initial mass of 3 M$_{\\odot}$ evolving in a close binary can produce a remnant of 0.4 M$_{\\odot}$ with a sizeable C/O core. A similar result was obtained 15 years later by Han, Tout \\& Eggleton (2000), who computed the evolution of close binary systems with several different parameters. They showed that in a binary system with the primary star of mass M$_1$= 2.51 M$_{\\odot}$, the mass ratio q=M$_1$/M$_2$=2 and the initial period P=2.559 d, the primary star succeeded to ignite the helium-burning. They did not followed the evolution of the remnant C/O WD, since their code met numerical problems when the primary star was as low as 0.33 M$_{\\odot}$ with a C/O core of 0.11 M$_{\\odot}$.  In this paper we will investigate both a possible evolutionary scenario able to produce C/O WDs of low mass, that is M $< 0.5$ M$_{\\odot}$, and their cooling evolution. We will show the results of detailed evolutionary computations performed with a full Henyey code able to follow consistently the evolution of stars from the initial pre-main sequence to the final cooling phase of WDs. ",
        "conclusions": "In conclusion, as already anticipated by Iben \\& Tutukov (1985), we proved, by means of fully and consistent evolutionary computations, that the minimum mass for a C/O WD is about 0.33 M$_{\\odot}$, near the minimum possible mass of the helium-core at the onset of He-burning, which occurs at the RGB phase transition (M$\\approx$ 2.3 M$_{\\odot}$).  As a consequence, in the mass range 0.33-0.5 M$_{\\odot}$ both He and C/O core WDs can exist. As expected and already shown by Panei et al. (2007), the cooling times of these two classes of WDs are quite different, being the He-core remnants significantly slower than the C/O ones. In addition, we showed also that the He WDs are more expanded that their C/O counterparts at a given effective temperature. \\subsection"
    },
    "0812/0812.4615_arXiv.txt": {
        "abstract": "{We report on an ongoing effort to image active galactic nuclei simultaneously observed at 2.3 and 8.6~GHz in the framework of a long-term VLBI project RDV (Research \\& Development -- VLBA) started in 1994 aiming to observe compact extragalactic radio sources in the astrometric/geodetic mode. Observations of bright extragalactic sources are carried out bi-monthly making up to six sessions per year with participation of all ten VLBA antennas and up to nine additional (geodetic and EVN) radio telescopes. Analysis of single-epoch results for 370 quasars, BL Lacs and radio galaxies is presented. We discuss VLBI core properties (flux densities, sizes, brightness temperatures), spectral characteristics of the cores and jets, evolution of brightness temperatures in the jets. }  \\FullConference{The 9th European VLBI Network Symposium on The role of VLBI in the Golden Age for Radio Astronomy and EVN Users Meeting\\\\ September 23-26, 2008\\\\ Bologna, Italy}  \\begin{document} ",
        "introduction": "Long-term VLBI project RDV (Research \\& Development -- VLBA) aimed at observations of bright compact extragalactic radio sources was started in 1994 under coordination of NASA and NRAO \\cite{Petrov08}. The simultaneous observations at 2.3 and 8.6~GHz are carried out bi-monthly making up to five-six sessions per year with participation of all ten 25-m VLBA antennas and up to nine geodetic and EVN stations. The participation of the southern antennas such as HartRAO (South Africa) and TIGO (Chile) allowed to successfully observe sources with declination up to $-47^\\circ$. Sample of observing objects consist of $\\sim$500 sources. It is important to note that the sample is not flux density complete. In each experiment 80-90 active galactic nuclei are observed and $\\sim$50 objects form a core of the sample and are scheduled continuously. The sample is dominated by quasars, with the weak-lined BL Lacs and radio galaxies making up 8.3\\% and 7.8\\% of the sample, respectively. ",
        "conclusions": "We discuss first-epoch results for 370 active galactic nuclei (Fig.~\\ref{fig1}, left) on the basis of high dynamical range ($\\sim$1000) images obtained at 2.3 and 8.6~GHz. The sources are bright (in 92\\% the correlated flux density was more than 200~mJy). One half of the sources are compact and core-dominated (VLBI compactness greater than 0.51, core dominance greater than 0.75). The median values of the core size are 0.28 and 1.04~mas at 8.6 and 2.3~GHz, respectively.  Dual-frequency VLBI observations provide a possibility to study spectral properties on parsec scales. On Fig.~\\ref{fig1} (right) we plot spectral index distribution map for J1800+7828 as a typical example. Most of VLBI cores have flat spectra ($\\alpha_{\\rm core}\\sim0$, $S\\propto\\nu^\\alpha$) since the radiation from these regions is dominated by optically thin emission from the jet base. We have cross-identified 48 jet components in 38 sources, the median value of obtained spectral indices is $\\alpha_{\\rm jet}=-0.75$ indicating optically thin radiation. This results in the median value of power index energy distribution of radiative particles to be $2.5$. Determination of spectral index of the jet components has been done taking into account a correction for the respective core shifts found in these sources~\\cite{Kovalev08}.   \\begin{figure} \\centering \\includegraphics[height=.47\\textwidth,angle=-90]{figure2a.eps} \\includegraphics[height=.47\\textwidth,angle=-90]{figure2b.eps} \\caption{Brightness temperature histograms for the core components at 2.3~GHz (left) and 8.6~GHz (right).} \\label{fig2} \\end{figure}  We have also measured the core brightness temperatures in the source rest frame. The respective distributions at 2 and 8~GHz shown in Fig.~\\ref{fig2} have close median values of $2.5\\times10^{11}$K. The empty bins are the lower limits and represent the cases when either the source has unknown redshift or an upper resolution limit has been used for the size of the component.  The sources with the prominent and well modeled jets having at least three jet components at both frequencies were the cases of our particular interest, since they provided us with information about brightness temperature evolution along the jets as a function of a distance to the core, $r$. The brightness temperature gradients can be well fitted with power law $T_{\\rm b}\\propto r^{-k}$ for 12 selected sources. The power index $k$ varies with values typically between $1.2$ and $3.6$ with the median value of $k=2$. Applying synchrotron radiation theory for conical jet model~\\cite{Lobanov98,Kadler05} and taking into account the median values for jet spectral index $\\alpha=-0.75$ and power law index $k=2$ we obtained the dependencies of electron density $n_e\\propto r^{-1.5}$ and magnetic field $B\\propto r^{-0.9}$ along the jet."
    },
    "0812/0812.1801_arXiv.txt": {
        "abstract": "We present a novel numerical implementation of radiative transfer in the cosmological smoothed particle hydrodynamics (SPH) simulation code {\\small GADGET}. It is based on a fast, robust and photon-conserving integration scheme where the radiation transport problem is approximated in terms of moments of the transfer equation and by using a variable Eddington tensor as a closure relation, following the `OTVET'-suggestion of Gnedin \\& Abel.  We derive a suitable anisotropic diffusion operator for use in the SPH discretization of the local photon transport, and we combine this with an implicit solver that guarantees robustness and photon conservation.  This entails a matrix inversion problem of a huge, sparsely populated matrix that is distributed in memory in our parallel code. We solve this task iteratively with a conjugate gradient scheme.  Finally, to model photon sink processes we consider ionisation and recombination processes of hydrogen, which is represented with a chemical network that is evolved with an implicit time integration scheme.  We present several tests of our implementation, including single and multiple sources in static uniform density fields with and without temperature evolution, shadowing by a dense clump, and multiple sources in a static cosmological density field. All tests agree quite well with analytical computations or with predictions from other radiative transfer codes, except for shadowing. However, unlike most other radiative transfer codes presently in use for studying reionisation, our new method can be used on-the-fly during dynamical cosmological simulation, allowing simultaneous treatments of galaxy formation and the reionisation process of the Universe. ",
        "introduction": "The absence of Gunn-Peterson troughs in the spectra of high redshift quasars up to $z \\leq 6$ \\citep{Fan2006, White2003} suggests that hydrogen is highly ionised at low redshift, with a volume averaged neutral hydrogen fraction of $x_{\\rm HI} \\leq 10^{-4}$ \\citep{Fan2007}. The current generation of simulation codes for cosmological structure formation calculates the self-gravity of dark matter and cosmic gas, and the fluid dynamics of the cosmic gas, but radiation processes are typically not taken into account, or only at the level of a spatially uniform, externally imposed background field. However, we know that the radiation field has been highly inhomogeneous during certain phases of the growth of structure, and may have in fact provided important feedback effects for galaxy formation \\citep[e.g.][]{Iliev2005, Yoshida2007, Croft2008}. Therefore, it would be ideal to be able to calculate self-consistent simulations that {\\em simultaneously} follow cosmic structure growth and the ionisation history of the Universe.  However, the high calculation cost and complexity of radiative transfer has so far prevented such simulations, apart from very few exceptions \\citep[e.g.][]{GO1997, Kohler2007, Shin2008, Wise2008}. Instead, the cosmic reionisation process has been most often treated separately in postprocessing, based on static simulated density fields \\citep[e.g.][]{Ciardi2003, Sokasian2004, Iliev2006, Zahn2007, Croft2008, Li2008}. It is the goal of this study to develop a new numerical scheme that makes self-consistent radiation-hydrodynamic simulations possible based on the high-resolution Lagrangian cosmological code {\\small GADGET-3} \\citep{gadget1, gadget2}.  A large number of different approximate methods have been developed to make problems of radiative transfer numerically tractable, which are in their full generality only solvable analytically for very few special cases.  The known numerical methods include long characteristics ray-tracing schemes \\citep{MM, ANM1999, Sokasian2001, Cen2002, Abel2002b, Razoumov2005}, short characteristic ray-tracing schemes \\citep{KA1988, NUS2001, MIAS2006, Whalen2006, Alvarez2006a, Susa2006, Altay2008, CFMR2001, MFC2003}, hybrid ray-tracing schemes \\citep{Rijkhorst2006, Trac2007},  moment methods and direct solvers \\citep{Mellema1998, GA2001, Shapiro2004, Whitehouse2005, AT2007, Finlator2008} and other particle-based transport methods \\citep{Ritzerveld2003, Pawlik2008}.  In the long characteristics method each test cell in the computational volume is connected to all other relevant cells. Then the equation of radiative transfer is integrated individually from that test cell to each of the selected cells. While this method is relatively simple and straight-forward, it is also very time consuming, since it requires $N^2$ interactions between the cells. Moreover, parallelization of this approach is cumbersome and requires large amounts of data exchange between the different processors. Short characteristic methods integrate the equation of radiative transfer along lines that connect nearby cells. Here a cell is connected only to its neighboring cells, and not to all other cells in the computational domain. This reduces the redundancy of the computations and makes the scheme easier to parallelize.  Stochastic integration methods, specifically Monte Carlo methods, employ a ray-tracing strategy where the rays are discretized into photon packets. For each photon packet, the emission, its frequency and its direction of propagation are determined by sampling the appropriate distribution functions that have been assigned in the initial conditions. A particular advantage of this approach is that comparatively few approximations to the radiative transfer equations need to be made, so that the quality of the results is primarily a function of the number of photon packets employed, which can be made larger in proportion to the CPU time spent. A disadvantage of these scheme is the comparatively high computational cost and the sizable level of noise in the discretized radiation field.  Using moments of the radiative transfer equations instead of the full radiative transfer equation can lead to very substantial simplifications that can drastically speed up the calculations. In this approach, the radiation is represented by its mean intensity field throughout the computational domain. Instead of following rays, the moment equations are solved directly on the grid or for the particles. Due to its local nature, the moment approach is comparatively easy to parallelize, but it of course depends on the nature of the investigated problem whether the simplifications made still provide sufficient accuracy.   In this paper we develop a moments method that is closely related in spirit to the Optically Thin Variable Eddington Tensor (OTVET) scheme proposed by \\citet{GA2001}. However, we try to implement it directly on top of the irregular set of positions sampled by the particles of Smoothed Particle Hydrodynamics (SPH) simulations, and we use quite different numerical techniques to solve the resulting transport equations.  In OTVET, the system of moment equations is closed by estimating the local Eddington tensor with a simple optical thin approximation, i.e.~one pretends that all sources of light are `visible' at a given location. Once the Eddington tensors are found, the local radiation transfer reduces to an anisotropic diffusion problem. The particular attraction of this moment-based formulation is that it is potentially very fast, allowing a direct coupling with cosmological hydrodynamic simulations. In particular, if a rapid method for calculating the Eddington tensors can be found, the scheme should be able to easily deal with an arbitrary number of sources. Also, the radiation intensity field does not suffer from the Poisson shot noise inherent in Monte Carlo approaches. Together with the local nature of the diffusion problem, this makes this approach particularly attractive for trying to address the cosmological reionisation problem with self-consistent simulations of galaxy formation, since it is likely that low-mass star-forming galaxies of high number density play an important role for the reionisation process.  We therefore adopt in this paper the suggestion of \\citet{GA2001} and work out an implementation of the OTVET scheme in SPH. As we shall see, this entails a number of numerical challenges in practice. We will describe our solutions for these problems, and carry out a number of tests to evaluate the accuracy of the resulting implementation.   The near complete lack of analytical results for non-trivial radiative transfer problems makes it actually hard to validate different numerical techniques and to compare their performance with each other. A very useful help in this respect is provided by the cosmological radiative transfer code comparison project, carried out by \\citet{Iliev2006b}. In their paper, they present a comparison of 11 independent cosmological radiative transfer codes when applied to a variety of different test problems. A number of our tests are based on this study, which hence allows a comparison with results of these other codes.  We start this paper with a brief introduction to the radiative transfer equations in Section~\\ref{SecIntro}. We then describe in Section~\\ref{SecMoments} the moment-based method that is the basis for our approximate treatment of the radiation transfer problem.  In Section~\\ref{SecNumerics} we elaborate in detail the numerical implementation of this scheme in SPH. This is followed by a presentation of results for various test problems in Section~\\ref{SecResults}.  Finally, we conclude with a summary and an outlook in Section~\\ref{SecConclusions}. In an appendix, we give for reference further details on our rate equations, and on the conjugate gradient approach. ",
        "conclusions": "\\label{SecConclusions}  We have presented a novel method for solving the radiative transfer equations within SPH, which is based on moments of the radiative transfer equation that are closed with a variable Eddington tensor. The radiation transport effectively becomes an anisotropic diffusion problem in this formulation. We have developed a new discretization scheme for anisotropic diffusion in SPH together with an implicit time integration method which for the first time allows a calculation of such anisotropic diffusion in SPH. Together with a scheme to estimate Eddington tensors based on the optically thin approximation, this yields a very fast approximate treatment of radiative transfer that can be used in dynamical SPH calculations.   We have implemented our method into the cosmological simulation code {\\small GADGET-3} and presented several test problems where we varied the initial conditions and different numerical parameters to investigate the accuracy of the method. Our test results agree in general very well with analytical predictions and data from other simulations, except that the long-term evolution of sharp geometric shadows is clearly not followed accurately. While this clearly limits the range of applicability of the method, we expect that the method can still provide reasonably accurate results for problems where shadowing is comparatively unimportant, such as cosmological reionisation, where the SPH-based variable Eddington tensor approach can be competitive with other techniques.  However, our method has two important strengths not shared by most other techniques: It can easily cope with an arbitrary number of sources since its speed is essentially independent of the number of sources, and furthermore, it is fast enough to be included into a cosmological simulation code where radiative transfer is calculated on-the-fly together with the ordinary dynamics. This is especially promising for future calculations of galaxy formation and reionisation that we want to carry out with our new code.   The epoch of reionisation is an important and extremely interesting process in the evolution of the Universe. There is already very important observational evidence that constrains reionisation, but few precise statements about the onset, duration, and end of reionisation can be made at this point. The situation will likely change in the near future through upcoming observations, ranging from new CMB observations with PLANCK, to 21cm tomography at high redshift with radio telescopes such as LOFAR. In the meantime, we heavily rely on cosmological simulations to advance our understanding of the reionisation process.  However, the simulations run so far leave many questions still unanswered, and their lack of self-consistency with the actual dynamics and the rapidly evolving source population may have introduced sizable inaccuracies. It is therefore highly desirable to have new numerical techniques that do not rely on solving the radiative transfer equation in a post-processing mode by treating only individual snapshot files from simulations.  Rather, accurate calculations should account for radiative transfer `on-the-fly', so that the gas properties are affected simultaneously by gravity and radiation. Our new method promises to be an important step in this direction. In future work, one of our most important goals is therefore to carry out high-resolution simulations of structure formation where the build-up of stellar and quasar sources in galaxies is followed self-consistently with the radiation field, allowing us to make more accurate predictions for the temporal evolution of reionisation, and the topology of reionised regions as a function of time.  All tests problems we have presented in this thesis agree well with theoretical predictions or results obtained with other radiative transfer codes. We should be able to obtain a realistic and accurate temperature evolution of the Universe during reionisation, which is important for setting the `cosmic equation of state' that regulates the absorption seen as Lyman-$\\alpha$ forest in the spectra of distant quasars. Finally, we plan to include the photoionisation of other elements besides hydrogen, most importantly of helium. Helium reionisation probably happened sometime at redshift $z\\sim 2-4$, where it may have left a sizable imprint in the temperature evolution of the intergalactic medium. Surprisingly, recent observations suggest that the temperature-density relation of the IGM may be inverted \\citep{Bolton2007}, which could be caused by radiative transfer effects related to helium reionisation. Whether this is indeed possible can only be clarified with simulations. Studying this question with our new methods would therefore be particularly timely."
    },
    "0812/0812.2878_arXiv.txt": {
        "abstract": "A cloud of gas collapsing under gravity will fragment. We present a new theory for this process, in which layers shocked gas fragment due to their gravitational instability. Our model explains why angular momentum does not inhibit the collapse process. The theory predicts that the fragmentation process produces objects which are significantly smaller than most stars, implying that accretion onto the fragments plays an essential role in determining the initial masses of stars. This prediction is also consistent with the hypothesis that planets can be produced by gravitational collapse. ",
        "introduction": "\\label{sec: 1}  Understanding the formation of stars by gravitational collapse is a key problem in astrophysics. Sophisticated simulations (see, for example, \\cite{Bat+05,Goo+04}) are used to model star formation within cold molecular clouds in the interstellar medium, but there are significant limitations in the resolution and range of lengthscales achievable with existing computers. Furthermore, it is desirable to have simple models which can be used as a framework for understanding the results of simulations as well as fundamental questions, some of which have not yet been given satisfactory answers.  A theory of star formation should explain why stars have a rather narrow distribution of masses, despite the very broad distribution of physical properties of the molecular clouds. It should also estimate the typical mass of a star and the mass of the smallest objects which are formed by gravitational collapse. When a cloud of gas undergoes gravitational collapse, conservation of angular momentum implies an increase in its angular velocity.  One might expect that this increase in angular speed limits the extent of the collapse, and one would like to estimate the importance of this effect in limiting the collapse process. Also it is known that, on the high-mass side of its peak, the mass distribution is well approximated by a power-law \\citep{Sal55}, an observation which theory should explain.  In this paper we propose a model which gives a simplified but physically well-grounded account of gravitational collapse. We show that fragmentation arises because of the gravitational instability of sheets of relatively dense gas produced by shocks. Our theory does not give precise quantitative predictions, but it does give a framework which could be used for producing a quantitative understanding of simulations. We find that angular momentum does not create a barrier to gravitational collapse. We predict that fragmentation ceases when the fragment size is considerably smaller than the mass of a typical star, implying that that mass of a protostar increases by a large factor due to accretion of residual gas after the initial formation of a gravitationally collapsed core (and coalescence of cores could also play a role). There is support from simulations for the hypothesis that accretion is important for determining the initial mass of stars \\citep{Bat+05}. A theory for the initial mass distribution must therefore go beyond the distribution of masses of the gravitationally condensed cores which we discuss in this paper. \\cite{Bat+05} have described a simple model for this accretion process which gives results consistent with their simulations, but a physically well-motivated model will require substantial additional work. ",
        "conclusions": "\\label{sec: 8}  We have described a mechanism whereby a gas cloud fragments due to the formation of shocks. When these occur, only a finite fraction of the material in the cloud becomes shocked and undergoes fragmentation. The remainder of the gas remains unaffected, and continues to evolve at a slower rate than the shocked material. This residual material may itself form further shocks as it collapses, or it may become stabilised against further collapse by centrifugal effects. Whatever is the detailed sequence of the collapse process, when non-fragmenting dense cores form, they will do so in an environment of low density gas, which was not able to take the fastest route to gravitational collapse. This reservoir of low density gas can end up feeding the growth of protostars by accretion processes. A theoretical description of this late stage is outside the scope of the theory presented here: we cannot readily estimate how much the mass of a nascent protostar will be increased by accretion processes, nor the effect of the accretion on the form of the initial mass function of stars.  It is instructive to compare our approach with the turbulent fragmentation theory advanced by \\cite{Pad+02,Pad+04}. Shock waves play a central role in both theories, but in other respects the theories are very different. The turbulent fragmentation theory suggests an origin for the Salpeter mass distribution, a power-law distribution for the mass of heavy stars. The power-law in the mass distribution is proposed to result from a power-law relation between relative velocity and separation in the gas cloud, analogous to that which characterises fully-developed turbulence (discussed by \\cite{Fri97}). This implies that the shock strength is predicted to increase with the size of the structures in the turbulent gas cloud \\citep{Pad+04}. By way of contrast, in our model the shock strengths approach a universal distribution which is independent of the size scale. Also, we predict that the initial mass of fragments is much smaller than the typical stellar mass, with the protostellar mass determined by subsequent accretion. In the turbulent fragmentation theory the peak of the mass distribution is determined by some adjustable parameters, whereas in our model the fragment mass is determined by radiation transfer considerations.  Our model is consistent with the numerical simulation by \\cite{Bat+05} in that the gravitationally bound cores which form protostars are initially much smaller than typical stellar masses. One point of difference is that our model predicts that sheets of dense material formed by shocks will fragment into clouds of material with a size similar to the thickness of the sheet, whereas it is usually reported that simulations show sheets tending to break up into filamentary structures. In our discussion of the gravitational instability we assumed that the local motion of the shocked material is purely rotational. We hypothesise that filamentary structures arise when the local velocity gradient in the shocked layer is anisotropic.  It has been suggested that some of the puzzling features of extrasolar planetary systems, such as the high probability of finding highly eccentric orbits, can be explained by the hypothesis that planets formed by gravitational collapse \\citep{Rib+07,Wil+08} rather than by the standard model \\citep{Saf69} involving aggregation of dust particles in a circumstellar disc. According to our theory, the size of the smallest fragments produced by gravitational collapse are comparable to the sizes of gas-giant planets, which lends support to the picture of planet formation described by \\cite{Wil+08}. In this picture the planets would arise from dense cores which were not able to grow by accretion because of competition from a larger core nearby, which grows to become a protostar."
    },
    "0812/0812.1756_arXiv.txt": {
        "abstract": "We present a new method based on the N-point probability distribution (pdf) to study non-Gaussianity in cosmic microwave background (CMB) maps. Likelihood and Bayesian estimation are applied to a local non-linear perturbed model up to third order, characterized by a linear term which is described by a Gaussian N-pdf, and a second and third order terms which are proportional to the square and the cube of the linear one. We also explore a set of model selection techniques (the Akaike and the Bayesian Information Criteria, the minimum description length, the Bayesian Evidence and the Generalized Likelihood Ratio Test) and their application to decide whether a given data set is better described by the proposed local non-Gaussian model, rather than by the standard Gaussian temperature distribution. As an application, we consider the analysis of the WMAP 5-year data at a resolution of $\\approx 2^\\circ$. At this angular scale (the Sachs-Wolfe regime), the non-Gaussian description proposed in this work defaults (under certain conditions) to an approximative local form of the weak non-linear coupling inflationary model~\\citep[e.g.][]{komatsu01} previously addressed in the literature. For this particular case, we obtain an estimation for the non-linear coupling parameter of $ -94 < {\\rm f}_{\\nl} < 154$ at 95\\% CL. Equally, model selection criteria also indicate that the Gaussian hypothesis is favored against the particular local non-Gaussian model proposed in this work. This result is in agreement with previous findings obtained for equivalent non-Gaussian models and with different non-Gaussian estimators. However, our estimator based on the N-pdf is more efficient than previous estimators and, therefore, provides tighter constraints on the coupling parameter at degree angular resolution. ",
        "introduction": "\\label{intro}  The Cosmic Microwave Background (CMB) fluctuations, a relic radiation originated around 400,000 years after the Big-Bang, is one of the most outstanding sources for understanding the evolution and energy/matter content of the Universe. The large amount of high quality data provided by recent CMB experiments and other complementary astronomical observations, have provided with a consistent picture for a flat Universe filled with cold dark matter (CDM) and dark energy in the form of a cosmological constant ($\\Lambda$), plus the standard baryonic and electromagnetic components: the \\emph{concordance model} \\citep[e.g.][]{komatsu08}. However, beyond the strength of the CMB measurements to put constraints on the cosmological parameters, like the ones already provided by the NASA \\emph{Wilkinson Microwave Anisotropy Probe} \\citep[WMAP,][]{hinshaw08} and the ones expected from the incoming ESA \\emph{Planck satellite}, the CMB is a unique tool to probe fundamental principles and assumptions of the so-called \\emph{standard model}. In particular, the application of sophisticated statistical analysis to CMB data might help us to understand whether the temperature fluctuations of the primordial radiation are compatible with the fundamental isotropic and Gaussian predictions from the \\emph{inflationary phase}. The basic inflationary scenario relates the CMB fluctuations, as well as the large-scale structure of the Universe, to the Gaussian quantum energy density perturbations present during the early Universe~\\citep[see for instance][]{liddle00}. The present homogeneity and isotropy of the Universe is compatible with an inflationary era in the early universe and this idea is the only way, nowadays, to explain efficiently current observations ranging from galaxies to the CMB. In particular, this fundamental hypothesis predicts that the CMB fluctuations follow an isotropic and Gaussian random. In fact, the estimation of the cosmological parameters defining the concordance model is done by assuming these statistical properties.  The quality of current CMB data (in particular the ones provided by the WMAP satellite) has allowed for a systematic probe of the statistical properties of the relic radiation. Indeed, the interest of the scientific community in this field has experimented a significant growth, since the analysis of the WMAP data has reported several hints for departure from isotropy and Gaussianity of the CMB temperature distribution. Some of these works are the following. \\cite{park04} detected a Gaussianity deviation with a genus-based statistic; \\cite{vielva04} found a significant non-Gaussian signature on the 1-year WMAP data on the kurtosis of the Spherical Mexican Hat (SMHW) wavelet coefficients at scales of around 10 degrees, pointing out a very large \\emph{cold spot} (CS) on the southern hemisphere as a possible source for this non-Gaussianity. A posterior analysis \\citep{cruz05} confirmed the anomalous nature of the CS by performing an area-based statistical analysis of the wavelet coefficients. This detection, as well as new ones, were confirmed by analyzing WMAP with various wavelet bases and several statistical estimators \\citep{mukherjee04,cayon05,mcewen05,cruz06,martinez06,cruz07a,mcewen06,pietrobon08,wiaux08}. Isotropy deviations were also reported in different manners: an anomalous alignment of the low multipoles of the CMB \\citep{copi04,deOliveira04,katz04,schwarz04,bielewicz05,land05b,abramo06,freeman06,land07}; north-south asymmetries of the CMB fluctuations \\citep{eriksen04a,eriksen04b,hansen04a,hansen04b,donoghue05,eriksen05,land05a,bernui06,bernui07,eriksen07,rath07,gordon07}; an anomalous variance value \\citep{monteserin08}; unexpected correlation among the CMB phases \\citep{chiang03,coles04,chiang06}; unbalanced distribution of the temperature extrema \\citep{tojeiro06}; and anomalous alignment of CMB structures \\citep{wiaux06,vielva06,vielva07}.  All the previous analyses can be considered as \\emph{blind approaches}, since null tests were performed to probe the CMB compatibility with the isotropic and Gaussian hypotheses. Complementary to these works, the reader can also find in the literature several studies where \\emph{targeted} departures from Gaussianity are explored, based on non-standard physical models. In particular, several analyses have studied the WMAP data compatibility with anisotropic universes: the Bianchi VII$_h$ model \\citep{jaffe06a,jaffe06b,jaffe06c,bridges07a,bridges08}. Recently, some works have proposed non-rotational invariant models like \\cite{bohmer08} and \\cite{ackerman07} \\citep[exprored by][in the context of WMAP data]{groeneboom08}. Although these non-rotational invariant models are promising and provide us with anisotropic templates that could help to fix some anomalies in WMAP data, more work is still needed to connect these anisotropic patterns of the CMB fluctuations to a satisfactory physical model describing the evolution of the anisotropic field \\citep[][]{himmetoglu08a,himmetoglu08b}.  In addition to the previous analyses, several works have studied different hypotheses to explore the anomalous nature of the CS; for instance \\cite{cruz08} explored the CS compatibility with different non-standard models, pointing out that an explanation in terms of a cosmic texture \\citep[as proposed by][]{cruz07b} is much more favored than other alternatives already discussed in the literature, like a very large void in the large scale structure of the Universe or contamination, in the form of the Sunyaev-Zeldovich emission, due to a large and nearby galaxy cluster. However, the study of non-standard inflationary models is the problem that has attracted a larger attention. For instance, and for the WMAP case, the non-linear coupling parameter f$_{\\nl}$ that describes the non-linear evolution of the inflationary potential \\citep[see e.g.][and references therein]{bartolo04} has been constrained by several groups: using the angular bispectrum \\citep{komatsu03,creminelli06,spergel07,komatsu08}; applying the Minkowski functionals \\citep{komatsu03,spergel07,gott07,hikage08,komatsu08}; and using different statistics based on wavelets \\citep{mukherjee04,cabella05,curto08}. Besides a claim for f$_{\\nl} >$ 0 with a probability greater than 95\\% \\citep{yadav08}, there is a general consensus on the WMAP compatibility with the predictions made by the standard inflationary scenario at least at 95\\% confidence level. The current best limits \\citep{curto08} are: $-8 < {\\rm f}_{\\nl} < 111$ at 95\\% CL.   The present paper is related to \\emph{non-blind} or \\emph{targeted} probes for Gaussianity deviations, and, more specifically, it addresses two common topics arose in these kind of studies: the definition of an optimal estimator for the non-standard model (and, therefore, the optimal estimation of the parameters that define such a model) and the complementary issue of model selection. The former aspect is, usually, quite complex, since, for many physical models and realistic observational limitations (e.g. incomplete sky coverage) there is not a trivial solution and, therefore, non-optimal estimators are usually proposed (which are posteriorly characterized by simulations). The latter issue on model selection is equally complex, since, generally, relies on heuristic principles. Some authors adopt a fully Bayesian view, whereas others adopt asymptotic measurements for the distance between two hypothesis, and others prefer to rely on the information that can be obtained from the likelihood itself. In this work we propose to study a non-Gaussian model for the CMB that is a local non-linear expansion of the temperature fluctuations. For this model ---that, at large scales, is considered by some authors as an approximation to the non-linear coupling parameter f$_{\\nl}$--- we are able to build the exact likelihood on pixel space. To work in pixel space allows one to include easily non-ideal observational conditions, like incomplete sky coverage and anisotropic noise. We show how, from this likelihood, it is straightforward to obtain an analytical expression for the optimal estimator for the non-Gaussian term. In addition, we also explore different model selection criteria, like the Akaike information criterion \\citep{akaike73}, the Bayesian information criterion \\citep{schwarz78}, the minimum description length \\citep{rissanen01}, the Bayesian evidence, and the generalized likelihood ratio test.   The paper is organized as follows. In Section~\\ref{sec:model} we describe the physical model based on the local expansion of the CMB fluctuations and derive the full posterior probability. In Section~\\ref{sec:statistics} we address the issue of the parameter estimation (\\ref{subsec:estimation}) and model selection (\\ref{subsec:selection}). The methodology is explored on WMAP-like simulations in Section~\\ref{sec:simulations} and it is applied to WMAP 5-year data in Section~\\ref{sec:wmap}. Finally, conclusions are given in Section~\\ref{sec:final}. ",
        "conclusions": "\\label{sec:final} We have presented a parametric non-Gaussian model for the CMB temperature fluctuations. The non-Gaussian model is a local perturbation (up to third order) of the standard CMB Gaussian field which recovers (for the case of $b \\equiv 0$ in equation~\\ref{eq:physical_model}) an approximative form of the weak non-linear coupling inflationary model \\citep[e.g.][]{komatsu01,liguori03} at scales larger than the horizon scale at the recombination time (i.e. above the degree scale). For this model, we are able to build the posterior probability of the data given the non-linear parameter $\\epsilon$ (see equation~\\ref{eq:model}), from which, in principle, an \\emph{optimal estimator} (i.e., unbiased and with minimum variance) can be derived. Analytical expressions for the maximum-likelihood estimation of the non-linear parameter ($\\hat{\\epsilon}$) and its associated error ($\\sigma_{\\hat{\\epsilon}}$) are derived. In addition, we also discuss an alternative Bayesian estimation (in terms of the posterior probability of the non-linear parameter), for the hypothetical case in which we might have some prior information for $\\epsilon$. As an example, two cases are addressed: a non-informative (i.e, uniform) and a Gaussian priors. We also investigate an issue very much linked to the parameter estimation: the model selection. Indeed, we discuss several well known techniques to perform hypotheses test, like the Akaike information criterion~\\citep[AIC,][]{akaike73}, the Bayesian information criterion~\\citep[BIC,][]{schwarz78}, the minimum description length~\\citep[MDL,][]{rissanen01}, the generalized likelihood ratio test (GLRT) and the Bayesian evidence (BE). We derive analytical expressions, for the particular local non-Gaussian model proposed in this work, for all these model selection techniques.  The performance of both, parameter estimators and model selection criteria, are investigated by analyzing non-Gaussian simulations, as they could be observed by WMAP. We check that the maximum-likelihood estimation provides an unbiased and efficient estimation of the non-linear parameter defining the deviations from Gaussianity. We find that, for the HEALPix resolution considered in this work (\\nside=32), results are consistent up to a value of $\\epsilon = 0.025$, which approximately corresponds (at the Sachs-Wolfe regime) to a value of the the non-linear coupling parameter f$_{\\nl} \\approx 350$. This parameter is the one commonly used to described the weak non-linear coupling inflationary model \\citep[e..g][]{komatsu01}. We also find that, among the model selection criteria, AIC is the asymptotic method that provides the less restrictive decision rule, whereas, on the other hand, MDL is the most strict one. We also find that BE, for a uniform prior given by $\\epsilon \\in \\left[-0.025, 0.025\\right]$, is even more restrictive than MDL.   The proposed methodology is applied to WMAP 5-year data. We obtain a value for the non-linear coupling parameter of $\\hat{{\\rm f}}_\\nl = 30 \\pm 124$ at 95\\% CL. This result provides a more efficient estimation than previous works in the literature, at the same angular scales. For instance, comparing with the work by \\cite{curto07} using Minkowski functionals, we can infer that the maximum-likelihood error bar is $\\approx 40\\%$ smaller than the one obtained with those geometrical estimators. Application of model selection criteria to WMAP data confirms that standard hypothesis of Gaussianity is favored against the alternative hypothesis of non-Gaussianity, for the specific local model proposed in this work, and for the adopted resolution  of $\\approx 2^\\circ$.  Finally, we would like to comment that, currently, we are extending the technique based on the N-pdf presented in this work, to deal with a more realistic non-local non-Gaussian model, where higher resolution CMB data are considered, including as well the effect of anisotropic noise."
    },
    "0812/0812.3279_arXiv.txt": {
        "abstract": "In Paper I we presented a detailed formulation of the relativistic shocks and synchrotron emission in the context of Gamma-Ray Burst (GRB) physics. To see how well this model reproduces the observed characteristics of the GRBs and their afterglows, here we present the results of some simulations based on this model. They are meant to reproduce the prompt and afterglow emission in some intervals of time during a burst. We show that this goal is achieved for both short and long GRBs and their afterglows, at least for part of the parameter space. Moreover, these results are the evidence of the physical relevance of the two phenomenological models we have suggested in Paper I for the evolution of the {\\it active region}, the synchrotron emitting region in a shock. The dynamical active region model seems to reproduce the observed characteristics of prompt emissions and late afterglow better than the quasi-steady model which is more suitable for the onset of afterglows. Therefore these simulations confirm the arguments presented in Paper I about the behaviour of these models based on their physical properties. ",
        "introduction": "\\label{sec:spectbehave} In the GRB data analysis usually all the photons detected during a given interval of time, e.g. the duration of the observation by the gamma-ray telescope - $T_{100}$ - is used for the determination of an averaged spectrum. Therefore for comparing the theoretical spectrum, equation (60-Paper I), with data we must integrate it with respect to $r$ for a given radius/time interval and divide it by the duration length. However, here we do not consider a lower limit for the detectable flux in the simulated models, and therefore $T_{100}$ is the same as the simulation time. For this reason we simply determine the total spectrum for the duration of the simulations without dividing it by time. The results are shown in Fig. \\ref{fig:specdar} for dynamical and in Fig. \\ref{fig:specqsar} for quasi steady active region models explained in Table \\ref{tab:param}.  \\subsection{Spectrum of simulated models with dynamical active region} \\label{sec:specdar} From variety of behaviours we have seen in the light curve of dynamical models we must expect the same level of variation in the spectrum, and it is exactly what we find in Fig. \\ref{fig:specdar}. We should remind that the spectrum of the observed bursts are usually time averaged on the total duration of the burst, or belongs to the total duration of a peak. Thus, their behaviour can be very different from the spectrum in a shorter interval of time. Each of The simulations presented in this work correspond to a fraction of a burst, and therefore we should not expect that it has a spectrum similar to time averaged observed spectrum of complete bursts.  The spectra of models 1, 2, and 6 are much harder than observed time averaged spectra. However, considering their light curves, they correspond to a short time at the beginning of the burst, and therefor their impact on the integrated spectrum would be small. In fact some bright hard bursts such as GRB 060105~\\citep{grb060105}, GRB 080319B~\\citep{grb080319b}, and GRB 080706~\\citep{grb080607} (long bursts) and GRB 051221~\\citep{grb051221, grb051221-1}, GRB 060313~\\citep{grb060313}, and GRB 080123~\\citep{grb080123} (short bursts) seem to have a very flat or even rising spectrum at the onset of the burst.  At highest energies, $E \\gtrsim$ few times $10^4$ eV, models 3, 4, and 5 which present the end of a peak, see Fig. \\ref{fig:lcgmzero} have a power-law with an exponential cutoff spectrum at $E \\gtrsim 10^4$ eV and at lower energies the slope is positive. This is consistent with magnetohydrodynamics simulations~\\citep{simulspec,simulspec0} and shock emission models~\\citep{spectheory}. At $E \\sim 10^4 - 10^5$ eV where their spectra are roughly flat, with an error at the level or higher than what is shown in the plots ($10\\%$ of the simulated values), they can be fit with a power-law consistent with observation of long \\swift GRBs~\\citep{batcat}. Although these models present only the end of a peak, the general behaviour of the broad band spectra is consistent the broad band spectrum of GRB 080916C which has been observed simultaneously by a large number of space and ground based instruments such as {\\it Fermi}~\\citep{grb080916c}, \\swift~\\citep{grb080916c-1,grb080916c-3}, {\\it Konus-Wind}~\\citep{grb080916c-2}, etc. This means that the effect of hard initial rise and soft tail emission after the decay of emission is small. Moreover, from similarity of these spectra with the spectrum of the only burst for which we have broad spectrum up to GeV bands we can conclude that these models present {\\it typical prompt emissions}.  A common feature between all these spectra is their low luminosity at optical energies. This is consistent with the lack of very bright peak in optical frequencies when multiple-band observations of late gamma-ray peaks were possible. Examples are GRB 060526~\\citep{grb060526,grb060526-1} and GRB 061121~\\citep{grb061121,grb061121-1}. In GRB 060526 the X-ray emission increased about $100$ times during a relatively soft but long peak observed also in gamma-ray. The optical emission during this peak increased only by a factor of $\\lesssim 2$ in V band from the continuum. In GRB 061121 also the X-ray increased about $100$ times but V band peak was only $\\sim 5$ times larger than the continuum and in B band only $\\sim 3$ times. Most authors consider the null assumption of a single power-law spectrum at low energies and try to explain the low luminosity of bursts at low energies by dust absorption. In the case of GRB 061121 apriori a dust to gas model similar to the value for Milky Way can explain observations~\\citep{grb061121}. However, we ignore the exact amount of dust in the host galaxy and we can not rule out a break at low energies consistent with magnetohydrodynamics simulations~\\citep{simulspec,simulspec0}.  Models 7 to 10 simulate external shocks and have a much steeper spectrum at high energies. The spectrum of model 7 which can correspond to the rise of the external shock is hard with a break at few keV. The spectrum of model 8 and the soft X-ray spectra of models 9 and 10 can be fit by a single power-law with negative index. At very soft X and UV bands the spectrum index of models 9 and 10 becomes steeper but at optical bands it becomes again flatter. This should correspond to the transition region from $E > \\bar{\\omega}_m$ to $E < \\bar{\\omega}_m$ where $\\bar{\\omega}_m$ is the average minimum synchrotron characteristic energy in the observer frame. The spectrum indexes in X-ray band varies from $\\sim -0.4$ in model 8 to $\\sim -1.3$ for model 9 and roughly flat for $10^3 < E < 10^4$ eV for model 7. The large scatter of the X-ray spectrum index in these simulations is consistent with the observations of the \\swift-XRT~\\citep{xrt}.  \\subsection{Spectrum of simulated models with quasi-steady active region} \\label{sec:specqsar} Fig. \\ref{fig:specqsar} shows that the spectrum of external shock simulations according to this model is very similar to what have been obtained for the dynamical models. Except model 11 which is hard and corresponds to the rising of the emission, others can be fit by a simple power-law or power-law with exponential cutoff at high and/or low energies. Like dynamical models, their spectrum index in X-ray band are mostly shallow with $\\alpha > -1$ and get even shallower in optical bands. But when the averaged minimum characteristic energy is in the plotted bands, as in model 12, there is a transient steep region similar to what we saw in dynamical models. Models 11 and 15 are hard with a positive index at low energies. As these models can only be relevant for the rising of the external shock emission which is usually overlapped with the tail emission of the prompt, it is very difficult to compare this results directly with observations.  Fig. \\ref{fig:specqsardr} shows the spectrum of the same models but with an energy dependent $\\Delta r$ as explained in Sec.\\ref{sec:lcqsar}. As expected, they are softer. Models 11 and 15 have negative slopes at low energies and a cutoff at high energies. Other models have the same form of the spectrum as in Fig. \\ref{fig:specqsar} but softer.  Models 17 to 19 in Fig. \\ref{fig:specqsar} which present internal shocks with quasi-steady state active region are hard and has spectra similar to prompt models with dynamical active region.  For various fundamental and observational reasons it is very difficult to compare these models with observations and with previous simulations or theoretical estimations. As we mentioned at the beginning of this section, these simulations does not correspond to a complete burst but just to part of a burst that can be simulated with a single and simple evolution model. Moreover, the emission after the prompt gamma-ray peak is most probably a combination of emission from the prompt internal shock and the afterglow / external shocks~\\citep{xrtafterglow}. It is not always easy to distinguish these components and in any case the separation would be model dependent. Nonetheless, if we assume that the missing sectors do not change the spectrum significantly, we can compare these simulations with observations and previous theoretical estimations. For instance models 7, 11, and 15 which present the early emission of the afterglow with what is called {\\it fast cooling regime} in~\\citep{emission1}, the general form of the spectrum is similar. Notably, their slope at high energies is $\\lesssim -1$ when $p = 2.5$ for all models except model 16. This value is consistent with the value of $-p/2$ predicted in~\\cite{emission1}. In model 16 the index $p = 3$ and the value of slope at high energies is $\\sim -1.5$ consistent with estimation of~\\cite{emission1}.  As for comparison with observations, it strongly depends on the time interval of observations. For instant, the available spectrum fit for the \\swift bursts belong mainly to the first couple of thousands of seconds because at later times the emission is faint and it is very difficult to determine the spectrum. However, as we mentioned before the early X-ray and optical emission are most probably due to the tail emission from the remnant of the prompt shock and therefore they can not be compared with what is expected from external shock. This mixing can be the reason for the large scatter in the spectral index of the X-ray spectrum~\\citep{xrtafterglow}. For instance GRB 070724A~\\citep{grb070724a}, GRB 070721B~\\citep{grb070721b}, GRB 061004~\\citep{grb061004} have an average spectrum index of $\\sim -1.4$ when others such as GRB 050128~\\citep{grb050128}, GRB 070208~\\citep{grb070208}, and GRB 060804~\\citep{grb060804} have an index of $\\sim -2.3$.  On the other hand, one of the principle differences between the prediction of these models and what is usually considered in the literature for the broad X-ray and optical spectrum is a single power law with or without cutoff at high or low energies. The simulated models here show the possibility of variant slopes and presence of cutoffs even in the energy range accessible to one instrument. This can be the reason for the absence of an optical afterglow specially when the burst is hard. Nonetheless, the under-luminous optical emission can be also explained by the presence of dust in the host galaxy or even in the circumburst environment~\\citep{dust}. One argument against the presence of a large amount of dust in the host galaxy of GRBs is the fact that most of them have relatively low metalicity~\\citep{grbhostmetal}. On the other hand, if the absorbers are limited to the star-forming clouds and their surroundings, then the local quantity of dust can be significant although the total fraction of dust in the host galaxy can be low. Present observations and reduction techniques are not yet sufficient to clarify this issue. In conclusion, these simulations show that the spectrum of emission itself can be in part responsible for the lack of observation of optical and lower energy afterglows in the observed GRBs. ",
        "conclusions": "We discussed the results of a few numerical simulations of GRBs based on the model explained in Paper I. We showed that the behaviour of these models is in general consistent with some regimes in the observed GRBs. It is difficult to verify the reality of some of the features seen in these simulations. For instance, a long lag between the peak of the X-ray and optical afterglows is predicted by some of these simulations. We need long multi-wavelength observations to be able to see such a feature in real bursts, and usually such data do not exist. Nonetheless, there are exceptions. Long follow-up of some bursts such as GRB 081029~\\citep{grb080129} and GRB 081007~\\citep{grb081007} show a wide late peak in optical wavelengths undetected in the X-ray. It can be from the afterglow, although other possibility such as the appearance of a supernova can not be ruled out. These observations increase the confidence on this model and analytical results presented here. Moreover, they encourage long term monitoring of GRBs to see if similar behaviour can be found in other bursts.  To complete these simulations we need to put together the simulation of various regimes and make the light curves and spectra corresponding to the whole burst. This should help to better understand the contribution of internal and external shocks in the light curves and spectrum as well as features such as the shallow regime, breaks, etc. We leave these investigations for future works."
    },
    "0812/0812.4173_arXiv.txt": {
        "abstract": "{} {The aim of this work was to search for double mode pulsators among RR~Lyr variables of globular cluster $\\omega$~Cen.} {We conducted a systematic frequency analysis of CASE photometry of $\\omega$~Cen RR~Lyr stars. We searched for periodicities using Fourier and ANOVA periodograms, combined with consecutive prewhitening technique.} {We discovered six double mode pulsators, with the first overtone and a secondary mode of higher frequency simultaneously excited. These are the first double mode RR~Lyr stars identified in $\\omega$~Cen. In variable V10 period ratio of the two modes is 0.80, which corresponds to pulsations in the first and second radial overtones. In V19 and V105 we found unexpected period ratio of 0.61. Three other stars display period ratios of either $\\sim\\! 0.80$ or $\\sim\\! 0.61$, depending on the choice of aliases.} {While the period ratio of $\\sim\\! 0.80$ is easy to interpret in terms of two lowest radial overtones, the value of $\\sim\\! 0.61$ cannot be explained by any two radial modes. Thus, V19 and V105 are the first members of a new class of double mode RR~Lyr pulsators.} ",
        "introduction": "The RR~Lyr type pulsating stars are one of the most interesting and most important objects in astrophysics, currently representing one of the basic rungs in the \"distance ladder\". Already from the moment of their discovery it was realized that their light curves were not uniform and could be divided into three classes: a, b and c (Bailey 1902). Later these classes were simplified to two types: RR$ab$ (radial pulsations in the fundamental mode) and RR$c$ (radial pulsations in the first overtone).  For the next 75 years it appeared that these were the only forms of pulsations possible in this type of stars. However, in the late 1970s and early 1980s it was found that RR~Lyr stars can pulsate with both fundamental mode and the first overtone simultaneously excited (Jerzykiewicz \\& Wenzel 1977; Cox et al. 1983). Later it was discovered that in addition to radial pulsations, RR~Lyr stars can also display non-radial modes of similar periods (Olech et al. 1999). Possibility of pulsations in the second radial overtone was discussed in the literature as well. For example, Alcock et al. (1996) found three maxima in the period distribution of RR~Lyr stars in the LMC and proposed that the shortest period maximum corresponds to the second overtone. Cluster RR~Lyr stars with nearly sinusoidal, low amplitude light curves and very short periods were interpreted as second overtone pulsators, too ({\\em e.g.} Walker \\& Nemec 1996, Kaluzny et al. 2000). However, presented evidence was weak, and all these stars could be pulsating in the first overtone after all. Theoretical models actually predict existence of such short period, low amplitude RR$c$ stars ({\\em e.g.} Bono et al. 1997).  While in RR~Lyr stars there is no conclusive evidence for excitation of the second overtone, such a mode is observed in many Cepheids. In the the latter case, unambiguous evidence for single mode second overtone pulsations have been found only recently (Soszy\\'nski et al. 2008), but double mode Cepheids with the first and second overtones simultaneously excited have been known for years (Mantegazza 1983; Alcock et al. 1995).   \\begin{table*}[!t] \\caption[]{Pulsation frequencies of V10, V350, V81, V87, V105 and V19.} \\centering \\begin{tabular}{|l|c|c|c|c|c|c|} \\hline \\hline Star                 & $P_1$\\thinspace [day] & $P_2$\\thinspace [day] & $P_2/P_1$ & $A_1$\\thinspace [mag] & $A_2$\\thinspace [mag] & $\\sigma$\\thinspace [mag] \\\\ \\hline V10 ~~~~1st\\,\\, sol. & ~0.3749759(5)~ & ~0.299176(9)~ & ~0.79785~ & ~0.1831(4)~ & ~0.0067(4)~ & ~0.0071~ \\\\ V10 ~~~~2nd     sol. & ~0.3749758(5)~ & ~0.230126(6)~ & ~0.61371~ & ~0.1825(5)~ & ~0.0064(5)~ & ~0.0074~ \\\\ \\hline V350 ~~1st\\,\\, sol.  & ~0.3791074(3)~ & ~0.230632(5)~ & ~0.60836~ & ~0.2126(6)~ & ~0.0062(6)~ & ~0.011~ \\\\ V350 ~~2nd     sol.  & ~0.3791077(4)~ & ~0.303810(9)~ & ~0.80138~ & ~0.2117(6)~ & ~0.0058(6)~ & ~0.011~ \\\\ V350 ~~3rd\\,   sol.  & ~0.3749758(5)~ & ~0.230126(6)~ & ~0.61371~ & ~0.1825(5)~ & ~0.0064(5)~ & ~0.011~ \\\\ \\hline V81 ~~~~1st\\,\\, sol. & ~0.3893907(3)~ & ~0.238990(4)~ & ~0.61375~ & ~0.2214(5)~ & ~0.0072(5)~ & ~0.010~ \\\\ V81 ~~~~2nd     sol. & ~0.3893907(3)~ & ~0.314313(7)~ & ~0.80719~ & ~0.2217(5)~ & ~0.0065(6)~ & ~0.010~ \\\\ \\hline V87 ~~~~1st\\,\\, sol. & ~0.3964889(3)~ & ~0.246584(4)~ & ~0.62192~ & ~0.2314(6)~ & ~0.0063(5)~ & ~0.010~ \\\\ V87 ~~~~2nd     sol. & ~0.3964884(3)~ & ~0.320542(8)~ & ~0.80845~ & ~0.2322(6)~ & ~0.0055(6)~ & ~0.010~ \\\\ \\hline V105                 & ~0.3353278(2)~ & ~0.205811(2)~ & ~0.61376~ & ~0.2465(6)~ & ~0.0125(6)~ & ~0.010~ \\\\ \\hline V19                  & ~0.2995510(2)~ & ~0.183302(2)~ & ~0.61192~ & ~0.2279(5)~ & ~0.0071(5)~ & ~0.010~ \\\\ \\hline \\hline \\end{tabular} \\end{table*}   \\begin{figure*} \\vspace{4cm} \\caption{Double mode RR~Lyr variable V10. (a) light curve phased with the primary period $P_1$, (b) power spectrum of original light curve, (c) power spectrum of prewhitened light curve.} \\special{psfile=Rysunek1.ps hoffset=30 voffset=-115 vscale=85 hscale=85 angle=0} \\end{figure*}  \\begin{figure*} \\vspace{11.6cm} \\caption{Same as Fig.\\thinspace 1, but for double mode RR~Lyr variables V350, V81, V87, V105 and V19.} \\special{psfile=Rysunek2.ps hoffset=72 voffset=-95 vscale=70 hscale=70 angle=0} \\end{figure*} ",
        "conclusions": "Systematic frequency analysis of RR~Lyr stars of globular cluster $\\omega$~Cen resulted in discovery of six RR$c$ variables, which in addition to the dominant first radial overtone, display also a weak secondary mode of higher frequency. These are the first double mode RR~Lyr stars identified in this cluster. In variable V10 the most probable period ratio of the two modes is 0.80. This value points towards pulsations in the first and second radial overtones. In three other stars, the period ratios are either $\\sim\\! 0.80$ or $\\sim\\! 0.61$, depending on the choice of aliases. Finally, in the last two stars (V19 and V105) an unambiguous period ratio of 0.61 was found. Such a period ratio cannot be explained by two radial modes, which implies that the secondary mode must be nonradial ({\\em cf.} Moskalik \\& Ko{\\l}aczkowski 2009 for discussion of this point). Thus, V19 and V105 belong to a new class of double mode RR~Lyr pulsators. We recall, that a similar period ratio of $\\sim\\! 0.61$ was also discovered in the LMC first overtone Cepheids (Moskalik \\& Ko{\\l}aczkowski 2008; Soszy\\'nski et al. 2008) and in the field double mode RR~Lyr star AQ~Leo (Gruberbauer et al. 2007).  Discovery of RR~Lyr stars pulsating in the first and second overtones has been claimed previously in the LMC (Alcock et~al. 2000). However, Soszy\\'nski et~al. (2003) have shown, that all these objects are about 1\\thinspace mag brighter than typical RR~Lyr stars in the LMC, and suggested that they might be short-period Cepheids instead. In contrast, all six double mode variables discovered in $\\omega$~Cen belong to the RR~Lyr population of the cluster without any doubt. Thus, V10 is the first solid candidate for the FO/SO double mode RR~Lyr pulsator."
    },
    "0812/0812.0029_arXiv.txt": {
        "abstract": "We present photometry of WASP-10 during a transit of its short--period Jovian planet. We employed the novel PSF--shaping capabilities of the OPTIC camera mounted on the UH 2.2\\,m telescope to achieve a photometric precision of $4.7\\times10^{-4}$ per 1.3~min sample. With this new light curve, in conjunction with stellar evolutionary models, we improve on existing measurements of the planetary, stellar and orbital parameters. We find a stellar radius $R_\\star = 0.698 \\pm 0.012$~\\rsun\\ and a planetary radius $R_P = 1.080 \\pm 0.020$~\\rjup. The quoted errors do not include any possible systematic errors in the stellar evolutionary models. Our measurement improves the precision of the planet's radius by a factor of 4, and revises the previous estimate downward by 16\\% (2.5$\\sigma$, where $\\sigma$ is the quadrature sum of the respective confidence limits). Our measured radius of WASP-10b is consistent with previously published theoretical radii for irradiated Jovian planets. ",
        "introduction": "When an exoplanet transits its star it provides a valuable opportunity to study the physical characteristics of planets outside our solar system. Photometric survey teams are discovering a rapidly growing number of short--period planets transiting relatively bright ($9 \\lesssim V \\lesssim 13$) Sun--like stars by monitoring the brightnesses of hundreds of thousands of stars for weeks or months, and following up on transit candidates with spectroscopy to rule out false positives \\citep{tres1,xo1,bakos07,wasp1,corot1}. With this windfall of new planet discoveries comes a need for efficient and precise photometric follow--up. Measuring planet properties precisely enough for meaningful tests of theoretical models of planetary interiors requires transit light curves with good photometric precision ($<$10$^{-3}$) and high cadence ($<$1~min). In the past these requirements have often been met by combining photometry from many transit events \\citep{holman06, winn07c}, or observing from space using the \\emph{Hubble}\\, and \\emph{Spitzer}\\, space telescopes \\citep{brown01, gillon07b}. These approaches work but they are costly in time and/or resources.  Our approach is to capitalize on the strengths of a relatively new type of detector: the orthogonal transfer array \\citep[OTA; ][]{tonry97}. OTA detectors are charge-coupled devices (CCDs) employing a novel four--phase gate structure to allow accumulated charge to be shifted in two dimensions during an exposure. Usually, this capability is used to compensate for image motion, providing on--chip tip--tilt correction and a narrower point spread function (PSF). However, as demonstrated by \\citet{howell03}, OTAs can also be utilized to deliberately broaden PSFs in a manner that is well suited for high--precision photometry. The method of ``PSF shaping'' uses a raster scan to distribute the accumulated signal from each star over many pixels within a well-defined region of the detector, thereby collecting more light per exposure and increasing the observing duty cycle compared to traditional CCDs. This OTA observing mode provides benefits similar to defocusing, but without the downside of a spatially and chromatically varying PSF \\citep{howell03, tonry05}.  In this Letter, we present new OTA-based photometry for a transiting planetary system recently announced by the Super Wide-Angle Search for Planets (SuperWASP) consortium. The planet, WASP-10b, is a 3.3~\\mjup\\ planet in an eccentric, 3.09~day orbit around a $V=12.7$ K--type dwarf \\citep[][hereafter C08]{wasp10}. C08 measured a radius for WASP-10b that is much larger than predicted by models of planetary interiors. In \\S~\\ref{data} we present our observations. We describe our data reduction procedures in \\S~\\ref{reduction} and our light curve modeling procedure and error analysis in \\S~\\ref{model}. We present our results in \\S~\\ref{results} and conclude in \\S~\\ref{discussion} with a brief discussion of our findings. ",
        "conclusions": "\\label{discussion}  We have presented high--precision photometry of the K--type star WASP-10 during the transit of its close--in Jovian planet. Our analysis has improved the precision with which the planetary, orbital and stellar properties are known. From the photometry, we improved the uncertainty in the orbit inclination angle by a factor of 2.7 and the transit depth by a factor of 3.8, and we measured the transit mid-time to 7 seconds. We also measured the stellar density ($\\rho_*$) from the transit light curve, which we used together with stellar evolution models to measure the stellar mass to a precision of 4.6\\%. We used the refined  stellar mass to improve the precision in the planetary radius by a factor of 4. Notably, the 1.8\\% uncertainty in $R_P$ is dominated by the uncertainty in the stellar mass, rather than the photometry.  Our determination of the planetary radius represents a 16\\% downward revision of the value reported by C08. Those authors found that the radius of WASP-10b was much larger than predicted by even a coreless planet model. Following C08, we compared our revised radius of WASP-10b to the theoretical predictions of \\citet{fortney07}, who provide tabulated radii for planets with a range of semimajor axes, masses, ages and heavy-element core masses. For a hydrogen-helium planet with properties appropriate to WASP-10, Fortney et al.\\ predict a radius of 1.13~\\rjup\\ for the case of a coreless planet with an age of 4.5 Gyr. Adding 25, 50, and 100~$M_\\oplus$ of heavy elements (nominally in the form of a solid core) decreases the radius to 1.11, 1.09 and 1.06~\\rjup, respectively.  The models are able to reproduce our measured radius ($R_P = 1.08 \\pm 0.02$~\\rjup) by including a heavy-element core with a mass ranging from 35~$M_\\oplus$ to 95~$M_\\oplus$.  We also compared our WASP-10b mass and radius measurements to the theoretical models of \\citet{bodenheimer03}, who predicted the radii of gas giant planets as a function of mass, age and equilibrium temperature. Assuming a planet temperature $T_{\\rm eq} = T_{\\rm eff}(R_*/a)^{1/2} = 1370$ (Table~2), \\citet{bodenheimer03} predict $R_P = 1.11$~\\rjup\\ for a 4.5 Gyr, solar-composition planet. The predicted radius is essentially the same whether the composition is taken to be purely solar or if a 40~$M_\\oplus$ core of heavy elements is included. It is also fairly insensitive to the value of $T_{\\rm eq}$. Thus our measured radius is only 3\\% ($1.5\\sigma$) smaller than predicted by \\citet{bodenheimer03}.  We conclude that WASP-10b is not significantly ``inflated'' beyond the expectations of standard models of gas giant planets. Instead, we find that WASP-10b is among the densest exoplanets so far discovered. Its mass and radius similar to another transiting planet, HD\\,17156b (Winn et al.\\ 2008c). However, WASP-10 and HD\\,17156 are very different host stars. Whereas HD\\,17156 is metal-rich ([Fe/H]$ = +0.24$) and has a mass of 1.26~\\msun, WASP-10 is a 0.75~\\msun\\ K dwarf with a much lower atmospheric abundance ([M/H]$ = +0.03 \\pm 0.2$; C08). WASP-10 stands out as one of only nine known low-mass ($M_* < 0.8$~\\msun) planet host stars within 200~pc, and its planet is among the most massive so far detected in this sample of low-mass stars\\footnote{http://exoplanets.org}.  We have found the OPTIC instrument to be capable of delivering single-transit photometry that is precise enough for meaningful comparisons of observed planetary properties and theoretical models. Our WASP-10 light curve has a precision of $4.7\\times10^{-4}$ per 1.3~min sample, or $6.8\\times10^{-4}$ per 1~min bin, for a $V=12.7$ star (Figure~\\ref{fig:lc1}). This is comparable to the precision that has been reported for composite light curves, based on observations of many transits, with traditional CCD cameras. For example, \\citet{johnson08b} presented a composite $z$-band light curve for HAT-P-1 ($V=10.4$) based on observations of 8 transits, with a precision of $5.7\\times10^{-4}$ in 1~min bins. OTAs may also prove useful for the future detection and characterization of transiting Neptune--mass and ``super--Earth'' planets."
    },
    "0812/0812.2908.txt": {
        "abstract": "\\label{abstract} % One of the scientific objectives of NASA's Fermi Gamma-ray Space Telescope is the study of Gamma-Ray Bursts (GRBs). The Fermi Gamma-Ray Burst Monitor (GBM) was designed to detect and localize bursts for the Fermi mission. By means of an array of 12 NaI(Tl) (8~keV to 1~MeV) and two BGO (0.2 to 40~MeV) scintillation detectors, GBM extends the energy range (20~MeV to >~300 GeV) of Fermi's main instrument, the Large Area Telescope, into the traditional range of current GRB databases. The physical detector response of the GBM instrument to GRBs is determined with the help of Monte Carlo simulations, which are supported and verified by on-ground individual detector calibration measurements. We present the principal instrument properties, which have been determined as a function of energy and angle, including the channel-energy relation, the energy resolution, the effective area and the spatial homogeneity. % %Include keywords, PACS and mathematical subject classification numbers as needed. ",
        "introduction": "% The Fermi Gamma-ray Space Telescope (formerly known as GLAST), which was successfully launched on June 11, 2008, is an international and multi-agency space observatory \\cite{ATW94,MIC96} that studies the cosmos in the photon energy range of 8 keV to greater than 300 GeV. The scientific motivations for the Fermi mission comprise a wide range of non-thermal processes and phenomena that can best be studied in high-energy gamma rays, from solar flares to pulsars and cosmic rays in our Galaxy, to blazars and Gamma-Ray Bursts (GRBs) at cosmological distances \\cite{GLAST01}. Particularly in GRB science, the detection of energy emission beyond 50 MeV \\cite{DIN01,KAN08} still represents a puzzling topic, mainly because only a few observations by the Energetic Gamma-Ray Experiment Telescope (EGRET) \\cite{THO93} on-board the Compton Gamma-Ray Observatory (CGRO) \\cite{HUR94,GON03} and more recently by AGILE \\cite{GIU08} are presently available above this energy . Fermi's detection range, extending approximately an order of magnitude beyond EGRET's upper energy limit of 30 GeV, will hopefully expand the catalogue of high-energy burst detections. A greater number of detailed observations of burst emission at MeV and GeV energies should provide a better understanding of bursts, thus testing GRB high-energy emission models \\cite{MES94,WAX97,DER00,SAR01}. Fermi was specifically designed to avoid some of the limitations of EGRET, and it incorporates new technology and advanced on-board software that will allow it to achieve scientific goals greater than previous space experiments.  The main instrument on board the Fermi observatory is the Large Area Telescope (LAT), a pair conversion telescope, like EGRET, operating in the energy range between 20 MeV and 300 GeV. This detector is based on solid-state technology, obviating the need for consumables (as was the case for EGRET's spark chambers, whose detector gas needed to be periodically replenished) and greatly decreasing (< 10 $\\mu$s) dead time (EGRET's high dead time was due to the length of time required to re-charge the HV power supplies after event detection). These features, combined with the large effective area and excellent background rejection, allow the LAT to detect both faint sources and transient signals in the gamma-ray sky. Aside from the main instrument, the Fermi Gamma-Ray Burst Monitor (GBM) extends the Fermi energy range to lower energies (from 8 keV to 40 MeV). The GBM helps the LAT with the discovery of transient events within a larger FoV and performs time-resolved spectroscopy of the measured burst emission. In case of very strong and hard bursts, the GRB position, which is usually communicated by the GBM to the LAT, allows a repointing of the main instrument, in order to search for higher energy prompt or delayed emission.  The GBM is composed of unshielded and uncollimated scintillation detectors (12 NaI (Tl) and two BGO) which are distributed around the Fermi spacecraft with different viewing angles, as shown in Fig. \\ref{Schema_Spacecraft}, in order to determine the direction to a burst by comparing the count rates of different detectors. The reconstruction of source locations and the determination of spectral and temporal properties from GBM data requires very detailed knowledge of the full GBM detectors' response. This is mainly derived from computer modeling and Monte Carlo simulations \\cite{KIP04,HOO08b}, which are supported and verified by experimental calibration measurements. % %*********************************** Fermi SPACECRAFT ********************************************************** \\begin{figure}[t!] \\centering \\begin{tabular}{cc} \\includegraphics[height=65mm,bb=54 63 486 569,clip]{Fig1a.eps} & \\includegraphics[height=65mm,bb=9 0 603 792,clip]{Fig1b.eps} \\end{tabular} \\caption{{\\it On the left}: Schematic representation of the Fermi spacecraft, showing the placement of the 14 GBM detectors: 12 NaI detectors (from {\\it n0} to {\\it nb}) are located in groups of three on the spacecraft edges, while two BGOs ({\\it b0} and {\\it b1}) are positioned on opposite sides of the spacecraft. {\\it On the right}: Picture of Fermi taken at Cape Canaveral few days before the launch. Here, six NaIs and one BGO are visible on the spacecraft's side. Photo credit: NASA/Kim Shiflett (\\texttt{http://mediaarchive.ksc.nasa.gov})} \\label{Schema_Spacecraft} \\end{figure} %************************************************************************************************************** %  In order to perform the above validations, several calibration campaigns were carried out in the years 2005 to 2008. The calibration of each individual detector (or detector-level calibration) comprises three distinct campaigns: a main campaign with radioactive sources (from 14.4 keV to 4.4 MeV), which was performed in the laboratory of the Max-Planck-Institut f\\\"ur extraterrestrische Physik (MPE, Munich, Germany), and two additional campaigns focusing on the low energy calibration of the NaI detectors (from 10 to 60 keV) and on the high energy calibration of the BGO detectors (from 4.4 to 17.6 MeV), respectively. The first one was performed at the synchrotron radiation facility of the Berliner Elektronenspeicherring-Gesellschaft f$\\rm{\\ddot{u}}$r Synchrotronstrahlung (BESSY, Berlin, Germany), with the support and collaboration of the German Physikalisch-Technische Bundesanstalt (PTB), while the second was carried out at the SLAC National Accelerator Laboratory (Stanford, CA, USA).  Subsequent calibration campaigns of the GBM instrument were performed at system-level, that comprises all flight detectors, the flight Data Processing Unit (DPU) and the Power Supply Box (PSB). These were carried out in the laboratories of the National Space Science and Technology Center (NSSTC) and of the Marshall Space Flight Center (MSFC) at Huntsville (AL, USA) and include measurements for the determination of the channel-energy relation of the flight DPU and checking of the detectors' performance before and after environmental tests. After the integration of GBM onto the spacecraft, a radioactive source survey was performed in order to verify the spacecraft backscattering in the modeling of the instrument response. These later measurements are summarized in internal NASA reports and will be not further discussed.  This paper focuses on the detector-level calibration campaigns of the GBM instrument, and in particular on the analysis methods and results, which crucially support the development of a consistent GBM instrument response. It is organized as follows: Section \\ref{GBM Det} outlines the technical characteristics of the GBM detectors; Section \\ref{Calib_camp} describes the various calibration campaigns which have been done, highlighting simulations of the calibration in the laboratory environment performed at MPE (see Section \\ref{Sec_Lab_Sim}); Section \\ref{Line_Ana_Proc} discusses the analysis system for the calibration data and shows the calibration results. In Section \\ref{Concl}, final comments about the scientific capabilities of GBM are given and the synergy of GBM with present space missions is outlined. % % %******************************NaI DETECTORS******************************************************************* \\begin{figure}[b!] \\centering \\begin{tabular}{c} \\includegraphics[width=84mm,bb=126 83 769 423,clip]{Fig2a.eps}  \\\\ \\includegraphics[width=84mm,bb=94 0 886 425,clip]{Fig2b.eps} \\end{tabular} \\caption{On the {\\it left}: Schematic cross-section of a GBM NaI detector showing the main components. A picture of a detector flight unit mounted on the calibration stand was taken in the laboratory during detector level calibration measurements and is shown in the {\\it right panel} } \\label{Fig_NaI_detectors} \\end{figure} %************************************************************************************************************** % % % %=============================== THE GBM DETECTORS ============================================================ % ",
        "conclusions": "% After the successful launch of the Fermi mission and the proper activation of its instruments, the 14 GBM detectors started to collect scientific data. The spectral overlap of the two BGO detectors (0.2 to 40 MeV) with the LAT lower limit of $\\sim$~20 MeV opens a promising epoch of investigation of the high-energy prompt and afterglow GRB emission in the yet poorly explored MeV-GeV energy region.  On ground, the angular and energy response of each GBM detector was calibrated using various radioactive sources between 14.4 keV and 4.4 MeV. The channel-energy relations, energy resolutions, on- and off-axis effective areas of the single detectors were determined. Additional calibration measurements were performed for NaI detectors at PTB/BESSY below 60 keV and for BGO detectors at SLAC above 5 MeV, thus covering the whole GBM energy domain. As already mentioned in the introduction, further calibration measurements at system level and after integration onto the spacecraft were carried out. All those measurements crucially contribute to the validation of Monte Carlo simulations of the direct GBM detector response. These incorporate detailed models of the Fermi observatory, including the GBM detectors, instruments, and all in-flight spacecraft components, plus the scattering of gamma-rays from a burst in the spacecraft and in the Earth's atmosphere \\cite{HOO08b}. The response as a function of photon energy and direction is finally captured in a Direct Response Matrix (DRM) database, allowing the determination of the true gamma-ray spectrum from the measured data.  The results reported in this paper directly contribute to the DRM final determination, and they fully follow physical expectations \\cite{KIE04}. Measurements and fit results for two sample detectors (NaI FM 04 and BGO FM 02) are given in Tables \\ref{Tab_Results_NaI_FM04} and \\ref{Tab_Results_BGO_FM02}, respectively (see Appendix). Fit parameters for the energy-channel relations and energy resolutions of those detectors are always reported in the captions of the corresponding plots (see Figures \\ref{figNaI_FM04_low-high}, \\ref{Fig_BGO_CE} and \\ref{fig_R}). It is worth noting that these parameters reflect the characteristics of all other NaI and BGO detectors not reported in this paper, thus showing that all detectors behave the same within statistics.  The channel-energy relation parameters obtained throughout this analysis are not directly used for in-flight calculations because of the different electronic setup which was used during detector level calibration and in-flight (see Section \\ref{MPE_calib_setup}). The same analysis method described in Section \\ref{Ana_Photopeak} and the same fitting procedures of Section \\ref{CE Relation} were adopted to analyze data collected during system level calibration. These results confirm that the systematic uncertainties in the channel-energy conversions arise from the discussed sources, i.e. fitting procedure, limited statistics in the case of high-energy lines of BGOs, electronics and non-uniform responses of detectors, and are fully consistent with measurements presented in this paper.  The GBM detectors will play an important role in the GRB field in the next decade. The unprecedented synergy between the GBM and the LAT will allow to observe burst spectra covering $\\sim$~7 decades in energy. Moreover, simultaneous observations by the large number of gamma-ray burst detectors operating in the Fermi era will complement each other. A nice overview of the currently operating space missions can be found in Table 1 of \\cite{BAN08}. Here, instrument characteristics such as FOVs, effective areas, localization uncertainties and energy bands are compared. The GBM detectors fit in this overall picture by providing a higher trigger energy range (50-300 keV) than e.g. Swift-BAT \\cite{GEH04} (15-150 keV) and a spectral coverage up to 30 MeV, an energy limit which can only be investigated with the LAT and the Mini-Calorimeter on-board AGILE \\cite{TAV06}. New insights into the GRB properties are therefore expected from GBM, thus advancing the study of GRB physics. % %"
    },
    "0812/0812.1169_arXiv.txt": {
        "abstract": "{The sunspot solar cycle has been usually explained as the result of a dynamo process operating in the sun. This is a classical problem in Astrophysics that until the present is not fully solved. Here we discuss  current problems and limitations with the solar dynamo modeling and their possible solutions using the kinematic dynamo model with the Babcock-Leighton approximation as a tool. In particular, we discuss the importance of the turbulent magnetic pumping versus the meridional flow circulation in the dynamo operation.}    \\addkeyword{Sun: magnetic fields}  \\begin{document} ",
        "introduction": "\\label{sec:intro}  The sunspot cycle is one of the most interesting magnetic phenomenon in the Universe. It was discovered more than 150 years ago by \\citet{schwabe}, but until now it remains  an open problem in astrophysics. There are several large scale observed phenomena that evidence that the solar cycle corresponds to a dynamo process operating inside the sun. These can be summarized as follows:  \\begin{inparaenum}[(1)] \\item The sunspots usually appear in pairs at both sides of the solar equator; the leading spot of a pair  (i.e. the one that points to the E-W direction) has opposite polarity to the other one. Besides, leading spots in the northern hemisphere have the opposite polarity to that of the leading spots in the southern hemisphere. Sunspots invert their polarity every 11 years; the total period of the cycle is then 22 years. This is known as the Hale's law; \\item A straight line connecting the leading and the companion spots of a pair has always an inclination of $10^{\\circ}$ to $30^{\\circ}$ with respect to the equatorial line. This is known as the Joy's law; \\item When the toroidal field reaches its maximum, i.e., when the number of sunspots is maximum, the global poloidal field inverts its polarity, so that there is a phase lag of $\\pi/2$ between the toroidal and poloidal components of the magnetic field; \\item The sunspots appear only in a belt of latitudes between $\\pm 30^{\\circ}$ (at both sides of the solar equator), these are known as the latitudes of activity; \\item The strength of the magnetic fields in the sunspots is around $10^3$ G. The magnitude of the diffuse poloidal  field is of tens of G. \\end{inparaenum}  \\citet{parker55} was the first to try to explain the solar cycle as a hydromagnetic phenomenon, since then although there has been important improvements in the observations, theory and simulations, a definitive model for the solar dynamo is still missing. Helioseismology has mapped the solar internal rotation showing a detailed profile of the latitudinal and radial shear layers, which seems to confirm the usually accepted idea that the first part of the dynamo process is the transformation of an initial poloidal field into a toroidal one. This stage is known as the $\\Omega$ effect. The second stage of the process, i.e., the transformation of the toroidal field into a new poloidal field of opposite polarity is a less understood process, and has been the subject of intense debate and research. Two main hypotheses have been formulated in order to explain the nature of this effect, usually denominated the $\\alpha$ effect: the first one is based on the Parker's idea of a turbulent mechanism where the poloidal field results from cyclonic convective motions operating at small scales in the toroidal field. These small loops should reconnect to form a large scale dipolar field. However, these models face an important problem: in the non-linear regime, i.e. when the back reaction of the toroidal field on the motions becomes important, the $\\alpha$ effect can be catastrophically quenched \\citep{vainsh92} leading to an ineffective dynamo \\citep{cattaneo96}\\footnote{Nevertheless, it is noteworthy that when the magnetic helicity is included in dynamical computations of $\\alpha$, it does not become catastrophycally quenched as long as the flux of the magnetic helicity remains non null \\citep[see][for a complete review of this subject]{bs05}. The shear in the fluid could be the way through which the helicity flows to outside of the domain \\citep{vishniac}.}.  The second one is based on the formulations of \\citet{babcock} and \\citet{leighton} (BL). They proposed that the inclination observed in the bipolar magnetic regions (BMR's) contains a net dipole moment. The supergranular diffusion causes the drift of half of each of these active regions to the equator and the drift of the other half in direction to the poles. so that this large scale poloidal structure annihilates the previous dipolar field. The new dipolar field is transported by the meridional circulations to the higher latitudes in order to form the observed polar field. This second mechanism has the advantage of being directly observed at the surface \\citep{wangetal89,wangetal91}, but it does not discard the existence of other $\\alpha$ sources underneath.  Following the BL idea, the physical model for the solar dynamo begins with a dipolar field. The differential rotation stretches the poloidal lines and form a belt of toroidal field at some place within the solar interior, in the convective layer. This toroidal field is somehow pushed through the turbulent convective eddies and forms strong and well organized magnetic flux tubes. When the magnetic field is intense enough and the density inside a tube is lower than the density of the surrounding plasma, it becomes unstable and begins to emerge towards the surface where it will form the BMRs. By diffusive decay, a BMR will form a net dipolar component with the opposite orientation of the original one, this new dipolar field is amplified during the cycle evolution until it reverses the previous dipolar field. In order to complete the cycle, it is necessary to transport this new poloidal flux first to the poles and then to the internal layers where the toroidal field will be created again, and so on. Most of the models in the BL mechanism use the meridional circulation  flow as the main agent of transport.  Recent works have invoked the turbulent pumping as an additional mechanism to advect the magnetic flux (see below).  In the absence of direct observations to confirm  the model above, several numerical studies have been performed in order to simulate the solar dynamo. These can be divided in two main classes: global dynamical models \\citep{brun04} and mean field kinematic models \\citep{dik99,ccn04,kuketal01,bon02,gue04,kapia06}. The first class integrates the full set of MHD equations in the solar convection zone and employ  the inelastic approximation in order to overcome the numerical constrain imposed by fully compressible convection on the time-step. These models are able to reproduce the observed differential rotation pattern, but they do not generate a cyclic, and well organized pattern of toroidal magnetic field. The second class of simulations solves the induction equation only and uses observed and/or estimated profiles for the velocity field and the diffusion terms. These models are relatively successful in reproducing the large scale features of the solar cycle, but the lack of the dynamical part of the problem has led to uncertainties in  the dynamo mechanism.  We have carried out several numerical tests with a mean field dynamo model in the Babcock-Leighton approximation in order to search for answers to four main issues:  \\begin{inparaenum}[(1)] \\item where is the solar dynamo located?; \\item what is the dominant flux transport mechanism?; \\item how to explain the observed latitudes of the solar activity?;  and \\item why is the solar parity  anti-symmetric? \\end{inparaenum}  In the following paragraphs, we briefly summarize our model assumptions and, step by step, draw our approach to the questions above in the light of the numerical simulations. ",
        "conclusions": "We have used a mean field dynamo model in the BL approach in order to look for answers to four current problems widely reported in the literature of solar dynamo modeling.  Our results confirm the idea that there should be a magnetic layer below the convection zone where magnetic fields which are mainly produced in the convection zone are stored. We have found that it is possible to have a flux-transport dynamo without a well defined meridional flow pattern,  dominated by the pumping advection. However, efforts in order to obtain a more realistic profile for the meridional velocity profile are still required, as well as to obtain a better comprehension of the real contribution of the pumping velocities.  Nevertheless, numerical simulations including near-surface shear as observed, provide also support for a near-surface magnetic layer, since toroidal fields with intensities between $10^3$ - $10^4$ G are formed there and since the radial differential rotation is negative at lower latitudes, the direction of migration of the butterfly wings constructed with these fields reproduces the one observed (e.g., Fig. 6).  The latitude of activity established by the observations (between $\\pm 30^{\\circ}$) can be explained by both scenarios above for the magnetic layer. For magnetic fields being stored in the overshoot region, a thin ($\\lesssim 0.2 \\sr$) tachocline could be the solution to avoid strong toroidal fields at the polar regions. On the other hand, if the near-surface shear is considered, then the magnetic pumping provides the required downwards flux of the weak polar fields, letting only the toroidal fields to survive at the active latitudes.  The parity in a dynamo solution is a problem that requires especial attention. We have investigated this problem looking for the role that the pumping could play in the solutions. Our simulations support the idea that the quadrupolar solution in most of the models is due to the strong quadrupolar imprint due to the one-cell meridional flow pattern. This imprint is larger at the surface than at the bottom of the convection zone, therefore the results of \\citet{dikgil01,bon02} which suggest that an $\\alpha$ effect operating at the tachocline results an anti-symmetric solution, could be explained by this fact. We suggest that models with an $\\alpha$ effect operating close to the surface are also able to generate anti-symmetric solutions, as observed, if a mechanism such as pumping cleans the quadrupolar imprint.  More observational and theoretical efforts are necessary in order to determine where the feet of the magnetic flux tubes responsible for the sunspots are located. This is an issue that needs to be solved before a more realistic coherent dynamo model can be constructed."
    },
    "0812/0812.3446_arXiv.txt": {
        "abstract": "{We present the results of a morphological study performed to a sample of Ultracompact (UC) H~II regions with Extended Emission (EE) using \\emph{Spitzer}--IRAC imagery and 3.6~cm VLA conf. D radio-continuum (RC) maps. Some examples of the comparison between maps and images are presented. Usually there is an IR point source counterpart to the peak(s) of RC emission, at the position of the \\uchii source. We find that the predominant EE morphology is the cometary, and in most cases is coincident with IR emission at 8.0~$\\mu$m. Preliminary results of \\emph{Spitzer}--IRAC photometry of a sub-sample of 13 \\uchii regions with EE (\\uchiie) based on GLIMPSE legacy data are also presented. Besides, individual IRAC photometry was performed to 19 \\uchii sources within these 13 regions. We show that \\uchii sources lie on specific locus, both in IRAC color-color and AM-product diagnostic diagrams. Counts of young stellar sources are presented for each region, and we conclude that a proportion of $\\sim$2\\%, $\\sim$10\\%, and $\\sim$88\\% of sources in \\uchiie are, in average, Class~I, II, and III, respectively.}  \\resumen{Presentamos los resultados de un estudio morfol\\'ogico realizado a una muestra de regiones H~II ultracompactas (UC) con emisi\\'on extendida (EE) usando im\\'agenes \\emph{Spitzer}--IRAC y mapas de radio-continnum (RC) del VLA a 3.6~cm en su configuraci\\'on D. Presentamos algunos ejemplos de la comparaci\\'on entre mapas e im\\'agenes. Usualmente hay una fuente puntual infrarroja que es contraparte de el(los) pico(s) de emisi\\'on en RC, en la posici\\'on de la fuente \\uchii. Encontramos que la morfolog\\'{\\i}a predominante de la EE es cometaria, y en la mayor\\'{\\i}a de los casos es coincidente con emisi\\'on IR a 8.0~$\\mu$m. Tambi\\'en presentamos resultados preliminares de la fotometr\\'{\\i}a \\emph{Spitzer}--IRAC de una sub-muestra de 13 regiones \\uchii con EE (\\uchiie) basada en datos del legado de GLIMPSE. Adem\\'as, realizamos fotometr\\'{\\i}a IRAC individual a 19 fuentes \\uchii dentro de estas 13 regiones. Mostramos que las fuentes \\uchii caen en lugares espec\\'{\\i}ficos, en ambos diagramas de diagn\\'ostico, color-color de IRAC y del producto AM. Presentamos conteos de fuentes estelares j\\'ovenes para cada regi\\'on, y concluimos que una proporci\\'on de $\\sim$2\\%, $\\sim$10\\%, y $\\sim$88\\% de las fuentes en \\uchiie son, en promedio, Clase~I, II, y III, respectivamente.}  \\addkeyword{H~II regions: Ultracompact} \\addkeyword{Stars: Pre-main sequence} \\addkeyword{Stars: High mass}   \\begin{document} ",
        "introduction": "\\label{sec:intro}  The ultracompact \\hii (\\uchii) regions (term coined by Israel, Habing \\& de Jong 1973), are small (size $\\leq$ 0.1 pc), dense ($\\gtrsim$ 10$^4$ cm$^{-3}$), photoionized Hydrogen regions with high emission measures ($\\geq$ 10$^7$ {${\\rm pc\\ cm}^{-6}$}), surrounding a recently formed ionizing (OB type) star(s). They are considered as good tracers of recent massive star formation, and generally they are surrounded by a natal dust `cocoon'.  Recent works show that several \\uchii regions have extended emission (EE) at arcmin scales (de la Fuente 2007; F07 hereafter, de la Fuente et al. 2009; F09 hereafter, Ellingsen et al., 2005; Kim \\& Koo, 2001; Kurtz et al. 1999). It is not clear if this EE is physically associated with the ultracompact (UC) emission at arcsec scales, but if it is, there will be strong implications and changes in our understanding of these objects, such as definition, environment, energetics and excitation mechanism (F07; F09; Kurtz 2002; Kurtz et al. 1999). This association can be performed comparing VLA with IR observations (F07; F09). The former traces ionized gas and the latter dust and environment at large scales, where previous VLA radio-observations only detect sources smaller than 20$''$--30$''$ (e.g. Kurtz, et al. 1994; Wood \\& Churchwell 1989; using the conf. A and B), and extended emission at scales up to $\\sim$~3$'$ (Kurtz et al. 1999) and to $\\sim$~7--15$'$ (Condon et al. 1998; Kim \\& Koo 2001) using the conf. D.  The \\emph{Spitzer} legacy project GLIMPSE (Benjamin et al. 2003\\footnote{see also \\emph{http://www.astro.wisc.edu/glimpse/}}), via photometry and imaging, has provided excellent tools to study star formation regions, discover star clusters, and to characterize the YSO population (e.g. F07; F09; Chavarria et al. 2008; Kumar \\& Grave 2007, KG07 hereafter; Hartmann et al. 2005 and references therein).  A summary of the results reported in de la Fuente Ph. D. thesis (2007), regarding the comparison between IRAC imagery and VLA low resolution maps is presented in \\S~2 and \\S~3, while the IRAC photometry analysis and young stellar classification also for \\uchii sources, are presented in \\S~4. Preliminary conclusions are given in \\S~5. ",
        "conclusions": "\\label{sec:fin} Because this is a work in progress, we list our preliminary conclusions.  \\begin{enumerate}  \\item Based on a comparison between RC maps at 3.6~cm (VLA conf.~D) and GLIMPS--IRAC images of a sample of \\uchiie, a coincidence of morphologies is observed. This indicate that heated dust coexists with ionized gas in the EE areas. \\item The peaks of RC emission are in most cases coincident with luminous IR counterparts. We identify broad (FWHM~$\\sim$8\\arcsec) IR point sources in the \\uchii position in several cases. \\item We perform \\emph{Spitzer}--IRAC photometry of 19 \\uchii sources in 13 \\uchiie regions (a subsample of F07), that are not available at GLIMPSE photometry catalogs. \\item This 19 \\uchiis lie in a region located (locus) at [5.8]--[8.0]$\\simeq$1.7 and 0.5$\\lesssim$[3.6]--[4.5]$\\lesssim$2.0, in the ``classic'' IRAC color-color diagram. \\item UC~\\hii sources also lie at the extreme border of Class~I objects in the AM product diagram (KG07), suggesting that at least some of them might be Class~0 candidates. \\item Following an $\\alpha_{IRAC}$ classification, we find that in average, the population of sources are: $\\sim$2\\% of Class~I (proto-stars), $\\sim$10\\% of Class~II (stars with disks), and $\\sim$88\\% of Class~III (stellar photospheres), in \\uchiie regions. \\end{enumerate}"
    },
    "0812/0812.1263_arXiv.txt": {
        "abstract": "Casimir energy is calculated for the 5D electromagnetism and 5D scalar theory in the {\\it warped} geometry. It is compared with the flat case(arXiv:0801.3064). A new regularization, called {\\it sphere lattice regularization}, is taken. In the integration over the 5D space, we introduce two boundary curves (IR-surface and UV-surface) based on the {\\it minimal area principle}. It is a {\\it direct} realization of the geometrical approach to the {\\it renormalization group}. The regularized configuration is {\\it closed-string like}. We do {\\it not} take the KK-expansion approach. Instead, the position/momentum propagator is exploited, combined with the {\\it heat-kernel method}. All expressions are closed-form (not KK-expanded form). The {\\it generalized} P/M propagators are introduced. Rigorous quantities are only treated (non-perturbative treatment). The properly regularized form of Casimir energy, is expressed in a closed form. We numerically evaluate $\\La$(4D UV-cutoff), $\\om$(5D bulk curvature, warp parameter) and $T$(extra space IR parameter) dependence of the Casimir energy. We present two {\\it new ideas} in order to define the 5D QFT:\\ 1) the summation (integral) region over the 5D space is {\\it restricted} by two minimal surfaces (IR-surface, UV-surface) ; or 2) we introduce a {\\it weight function} and require the dominant contribution, in the summation, is given by the {\\it minimal surface}. Based on these, 5D Casimir energy is {\\it finitely} obtained after the {\\it proper renormalization procedure.} The {\\it warp parameter} $\\om$ suffers from the {\\it renormalization effect}. The IR parameter $T$ does not. In relation to characterizing the dominant path, we {\\it classify} all paths (minimal surface curves) in AdS$_5$ space. We examine the meaning of the weight function and finally reach a {\\it new definition} of the Casimir energy where {\\it the 4D momenta( or coordinates) are quantized} with the extra coordinate as the Euclidean time (inverse temperature). We comment on the cosmological constant term and present an answer to the problem at the end. Dirac's large number naturally appears. ",
        "introduction": "}  \\begin{figure} \\caption{ The configuration of the Casimir energy measurement. The radiation cavity bounded by two parallel perfectly-conducting plates separated by $2l$. The plate size is $2L\\times 2L$. } \\begin{center} \\includegraphics[height=8cm]{2Plates} \\end{center} \\label{2Plates} \\end{figure} In the dawn of the quantum theory, the {\\it divergence} problem of the specific heat of the radiation cavity was the biggest one (the problem of the blackbody radiation). It is historically so famous that the difficulty was solved by Planck's idea that the energy is quantized. In other words, the phase space of the photon field dynamics is not continuous but has the \"cell\" or \"lattice\" structure with the unit area ($\\Del x\\cdot\\Del p$) of the size $2\\pi\\hbar$ (Planck constant). The radiation energy is composed of two parts, $E_{Cas}$ and $E_\\be$: \\bea E_{4dEM}=E_{Cas}+E_\\be\\com\\nn E_{Cas}=\\sum_{m_x,m_y,n\\in\\bfZ}\\omtil_{m_xm_yn}\\com\\q E_{\\be}=2\\sum_{m_x,m_y,n\\in\\bfZ}\\frac{\\omtil_{m_xm_yn}}{\\e^{\\be\\omtil_{m_xm_yn}}-1}\\com\\nn {\\omtil_{m_xm_yn}}^2=(m_x\\frac{\\pi}{L})^2+(m_y\\frac{\\pi}{L})^2+(n\\frac{\\pi}{l})^2 \\com\\q\\q l\\ll L\\com \\label{4dEM18X} \\eea where the parameter $\\be$ is the inverse temperature, $l$ is the separation length between two perfectly-conducting plates, and $L$ is the IR regularization parameter of the plate-size. See Fig.\\ref{2Plates}. The second part $E_\\be$ is, essentially, Planck's radiation formula. The first one $E_{Cas}$ is the vaccuum energy of the radiation field, that is, the Casimir energy. It is a very delicate quantity. The quantity is formally {\\it divergent}, hence it must be defined with careful {\\it regularization}. (See App.B). $E_{Cas}/(2L)^2$ does depend only on the {\\it boudary} parameter $l$. The quantity is a quantum effect and , at the same time, depends on the {\\it global} (macro) parameter $l$. \\footnote{ See the recent reviews on Casimir energy: Ref.\\cite{Cas01t06}. } \\bea \\frac{E_{Cas}}{(2L)^2}=\\frac{\\pi^2}{(2l)^3}\\frac{B_4}{4!}=-\\frac{\\pi^2}{720}\\frac{1}{(2l)^3}\\com\\q B_4\\mbox{(the fourth Bernoulli number)}=-\\frac{1}{30} \\com \\label{4dEM24x} \\eea   In Fig.\\ref{PlanckDistB}, Planck's radiation spectrum distribution is shown. \\begin{figure} \\caption{ Graph of Planck's radiation formula. $ \\Pcal (\\be,k)=\\frac{1}{(c\\hbar)^3}\\frac{1}{\\pi^2}k^3/(\\e^{\\be k}-1)\\ \\ (1\\leq\\be\\leq 2,\\ 0.01\\leq k\\leq 10),\\ \\ k\\mbox{(photon energy)}=\\hbar \\om=h\\n=\\hbar c\\frac{2\\pi}{\\la}, \\ \\be=1/k_BT$ where $c\\hbar=2000~\\mbox{eV~\\AA}, \\ k_B\\mbox{(Boltzmann's constant)}=8.6\\times 10^{-5}~\\mbox{eV/K}$. $k$=(1,10)~eV correspond to $\\frac{\\la}{2\\pi}$=(2000,200)\\AA, $\\be$=(1,2)~eV$^{-1}$ correspond to T=($10^5/8.6,~10^5/17.2$) K. $\\Pcal=0.1$ corresponds to $\\Pcal\\Del k=(2000)^{-3}\\times 0.1\\times \\Del k$[eV/\\AA $^3$]. } \\begin{center} \\includegraphics[height=8cm]{PlanckDistB} \\end{center} \\label{PlanckDistB} \\end{figure} Introducing the axis of the inverse temperature($\\be$), besides the photon energy or frequency ($k$), it is shown stereographically. Although we will examine the 5D version of the zero-point part (the Casimir energy), the calculated quantities in this paper are much more related to this Planck's formula. \\footnote{ Planck's formula depends only on the temperature $1/\\be$, not on $l$. (Comparatively the Casimir energy part does not depend on $\\be$, but on the separation $l$. ) It is known that $\\be$ can be regarded as the periodicity for the axis of the inverse temperature (Euclidean time). The axis corresponds to the {\\it extra axis} in the following text. } We see, near the $\\be$-axis, a sharply-rising surface, which is the Rayleigh-Jeans region (the energy density is proportional to the {\\it square} of the photon frequency). The {\\it damping} region in high $k$ is the Wien's region. \\footnote{ We recall that the old problem of the {\\it divergent} specific heat was solved by the Wien's formula. This fact strongly supports the present idea of introducing the {\\it weight function} (see Sec.\\ref{uncert}). } The ridge (the line of peaks at each $\\be$) forms the {\\it hyperbolic} curve (Wien's displacement law). When we will, in this paper, deal with the energy distribution over the 4D momentum and the extra-coordinate, we will see the similar behavior (although top and bottom appear in the opposite way). In order to compare this 4D case with the 5D case of the present paper, we do some preparation for the Casimir energy in App.B.  In the quest for the unified theory, the higher dimensional (HD) approach is a fascinating one from the geometrical point. Historically the initial successful one is the Kaluza-Klein model\\cite{Kal21,Klein26}, which unifies the photon, graviton and dilaton from the 5D space-time approach. The HD theories \\footnote{ The HD theories we consider here, are the HD generalization of the familiar renormalizable (in 4D) theories such as 5D free scalar theory, 5D QED and 5D Yang-Mills theory. } , however, generally have the serious defect as the quantum field theory(QFT) : un-renormalizability. \\footnote{ Note that the ordinary power counting criterion is about the divergence degree for the coupling expansion. In this sense, there are some renormalizable HD field theories such as 6D $\\Phi^3$-theory. In the present paper, however, we have focused on the coupling-independent part, the Casimir part. This part is generally divergent in the higher dimensions. } The HD quantum field theories, at present, are not defined within the QFT. One can take the standpoint that the more fundamental formulation, such as the string theory and D-brane theory, can solve the problem. In the present paper, we have the {\\it new standpoint} that the HD theories should be defined by themselves within the QFT.  In order to escape the dimension requirement D=10 or 26 from the quantum consistency (anomaly cancellation)\\cite{StringText}, we treat the gravitational (metric) field only as the background one. This does {\\it not} mean the space-time is not quantized. See later discussions (Sec.\\ref{weight}). We present a way to define 5D quantum field theory through the analysis of the Casimir energy of 5D electromagnetism.  In 1983, the Casimir energy in the Kaluza-Klein theory was calculated by Appelquist and Chodos\\cite{AC83}. They took the cut-off ($\\La$) regularization and found the quintic ($\\La^5$) divergence and the finite term. The divergent term shows the {\\it unrenormalizability} of the 5D theory, but the finite term looks meaningful \\footnote{ The gauge independence was confirmed in Ref.\\cite{SI85PLB}. } and, in fact, is widely regarded as the right vacuum energy which shows {\\it contraction} of the extra axis. In this decade, triggered by the development of the string and D-brane theories, new treatments or new ideas were introduced to calculate the vacuum energy or the effective potential in HD. (The motivation is to settle the stability problem of the moduli parameters\\cite{GoldWise9907,GoldWise9911}. ) One is to regard the system as the bulk and boundary, and do renormalization in both parts\\cite{Toms00,GoldWise01}. Various regularization methods were carefully re-examined for the bulk-boundary theory\\cite{GoldRoth00}. They are applied to various theories including realistic models\\cite{FlachiToms01,GPT01,PonPop01,FGPT03,FlaPujo03}. From the regularization viewpoint, the zeta-function (or dimensional) regularization combined with some summation formula is most commonly taken. The renormalization procedure, however, does {\\it not} seem satisfactory. They succeed in calculating the properly regularized quantity and in separating the divergent terms. They found the finite part, but its physical meaning is obscure because the treatment of the divergent part is {\\it not established}. They simply say, based on the analogy to the case of the ordinary (4D) renormalizable thoeries, the {\\it local counterterms} can cancel divergences. \\footnote{ See the first footnote of Sec.8 about the present treatment of the local counterterms. } They try to absorb divergences by the renormalization of parameters such as the brane tension (cosmological constant) and the gravitational constant. But it is fair to say that the $\\La^5$-divergence problem, posed by Appelquist and Chodos, is {\\it not} yet solved. All these come from the unsatisfactory situation of the quantum treatment of the brane dynamics and the HD quantum field theories.  In the development of the string and D-brane theories, a new approach to the renormalization group was found. It is called {\\it holographic renormalization} \\cite{SusWit9805,HenSken9806,SkenTown9909,DFGK9909,FGPW9904,dBVV9912}. We regard the renormalization flow as a curve in the bulk (HD space). The flow goes along the extra axis. The curve is derived as a dynamical equation such as Hamilton-Jacobi equation. It originated from the AdS/CFT correspondence\\cite{Malda9711,GKP9802,Witten9802}. Spiritually the present basic idea overlaps with this approach. The characteristic points of this paper are:\\ a) We do {\\it not} rely on the 5D supergravity;\\ b) We do {\\it not} quantize the gravitational(metric) field;\\ c) The divergence problem is solved by reducing the degree of freedom of the system, where we require, not higher symmetries, but some restriction based on the {\\it minimal area principle};\\ d) No local counterterms are necessary.  In the previous paper\\cite{SI0801}, we investigated the 5D electromagnetism in the {\\it flat} geometry. For the later use of comparison (with the present {\\it warped} case), we list the main results here. \\footnote{ We do not require the reader to read ref.\\cite{SI0801}. The necessary key procedures are explained in the text. } The extra space is {\\it periodic} (periodicity $2l$) and $Z_2$-parity is taken into account: \\bea ds^2=\\eta_\\mn dx^\\m dx^\\n+dy^2\\com\\q -\\infty < x^\\m, y < \\infty\\com\\q y\\ra y+2l,\\ y\\change -y\\com\\nn (\\eta_\\mn)=\\mbox{diag}(-1,1,1,1)\\ , (X^M)=(X^\\m=x^\\m,X^5=y)\\equiv (x,y)\\ ,\\nn M,N=0,1,2,3,5;\\ \\m,\\n=0,1,2,3. \\label{5dEM1} \\eea The IR-regularized geometry of this 5D flat space(-time) is depicted in Fig.\\ref{RegGeomFlat}. \\begin{figure} \\caption{ IR-regularized geometry of 5D flat space (\\ref{5dEM1}). The 4D world (3-brane) is Euclideanized and is shown as 4D ball (shaded disk region) surrounded by $S^3$ sphere with radius $\\mu^{-1}$. $\\mu$ is the 4D IR regularization parameter, and is taken to be $\\mu=1/l$. UV-regularization is introduced, in Sec.\\ref{UIreg}, by replacing the 4D ball with the \"sphere lattice\" composed of many small (size: $1/\\La$) 4D balls. See Sec.\\ref{surf} for detail. } \\begin{center} \\includegraphics[height=8cm]{RegGeomFlat} \\end{center} \\label{RegGeomFlat} \\end{figure} The Casimir energy is {\\it rigorously} (all KK-modes are taken into account) expressed as \\bea E_{Cas}(\\La,l)=\\frac{2\\pi^2}{(2\\pi)^4}\\int_{1/l}^{\\La}d\\ptil\\int_{1/\\La}^ldy~\\ptil^3 W(\\ptil,y)F(\\ptil,y)\\com\\nn F(\\ptil,y)\\equiv F^-(\\ptil,y)+4F^+(\\ptil,y) =\\int_\\ptil^\\La d\\ktil\\frac{-3\\cosh\\ktil(2y-l)-5\\cosh\\ktil l}{2\\sinh(\\ktil l)} \\pr \\label{UIreg2} \\eea where $\\La$ is the 4D-momentum {\\it cutoff}, and $W(\\ptil,y)$ is the {\\it weight function} to suppress the IR and UV divergences. \\footnote{ Z$_2$-odd part $F^-$ comes from the quantum fluctuation of the extra component $A_5(X)$, while Z$_2$-even part $F^+$ from the 4D components $A_\\m(X)$. } \\footnote{ The expression (\\ref{UIreg2}) of $E_{Cas}$ is negative definite. The same thing can be said about the warped case (\\ref{HKA8}) in the later description. } We obtained the following $\\La$ and $l$ dependence by the numerical analysis.  1) Un-weighted case: $W=1$ \\bea \\mbox{Un-restricted integral region}:\\q\\qqq\\qqq\\nn E_{Cas}(\\La,l)=\\frac{1}{8\\pi^2}\\left[ -0.1249 l\\La^5 -(1.41, 0.706, 0.353)\\times 10^{-5}~l\\La^5\\ln (l\\La)\\right] \\com\\nn \\mbox{Randall-Schwartz integral region}:\\q E^{RS}_{Cas}=\\frac{1}{8\\pi^2}[-0.0894~\\La^4] \\pr \\label{UIreg5X} \\eea The {\\it quintic} divergence of the upper one of (\\ref{UIreg5X}) shows the {\\it unrenormalizability} of the 5D theory in the ordinary treatment. The triplet data show the unstable situation of numerical results. \\footnote{ The results of (\\ref{UIreg5X}) are based on the numerical integral of (\\ref{UIreg2}) for $l=(10,20,40), \\La=10\\sim 10^3$. The triplet coefficients correspond to the three values of $l$. This unstable situation does {\\it not} appear in the present case of warped geometry. See (\\ref{UIreg5}) and (\\ref{AppD3}). The same thing can be said about the weighted case 2) in the following. } As for the lower case, the $(\\ptil,y)$-integral region is restricted to below the {\\it hyperbolic curve} $\\ptil y =1$.\\footnote{ This restriction was taken in Ref.\\cite{RS01} to suppress the UV-divergence. }  2) Weighted case   \\bea E^W_{Cas}/\\La l=\\mbox{\\hspace{10cm}}\\nn \\nn \\left\\{ \\begin{array}{cc} -2.50\\frac{1}{l^4}+(-0.142,1.09,1.13)\\cdot 10^{-4}\\frac{\\ln l\\La}{l^4} & \\mbox                                          {for}\\q W=\\frac{1}{N_1}\\e^{-(1/2) l^2\\ptil^2-(1/2) y^2/l^2}\\equiv W_1\\\\ -6.04~10^{-2}\\frac{1}{l^4}-(24.7,2.79,1.60)\\cdot 10^{-8}\\frac{\\ln l\\La}{l^4} & \\mbox{for}\\q W=\\frac{1}{N_2}\\e^{-\\ptil y}\\equiv W_2\\\\ -2.51\\frac{1}{l^4}+(19.5,11.6,6.68)\\cdot 10^{-4}\\frac{\\ln l\\La}{l^4}  & \\mbox{for}\\q W=\\frac{1}{N_8}\\e^{-(l^2/2) (\\ptil^2+1/y^2)} \\equiv W_8\\end{array} \\right. \\label{uncert1b} \\eea where some representative cases ($W_1$:\\ elliptic, $N_1=1.557/8\\pi^2$ ; $W_2$:\\ hyperbolic, $N_2=2(l\\La)^3/8\\pi^2$; $W_8$:\\ reciprocal, $N_8=0.3800/8\\pi^2$) are shown. (See ref.\\cite{SI0801} for other cases. ) The quantity $\\La l$ is the normalization factor in the numerical analysis. The Casimir energy behavior of the case $W_2$ is consistent with the Randall-Schwartz's one of 1). These results imply the {\\it renormalization of the compactification size} $l$. \\bea E^W_{Cas}/\\La l =-\\frac{\\al}{l^4}\\left( 1-4c\\ln (l\\La) \\right) =-\\frac{\\al}{{l'}^4}\\com\\q \\be=\\frac{\\pl}{\\pl (\\ln \\La)}\\ln\\frac{l'}{l}=c\\com \\label{uncert1cc} \\eea where $\\al$ and $c$ should be uniquely fixed by clarifying the meaning of the weight function $W$ and the unstable situation of the triplet data.  The aim of this paper is to examine how the above results change for the 5D warped geometry case. The {\\it IR-regularized} geometry of the 5D warped space(-time) is depicted in Fig.\\ref{RegGeomWarp}. One additional massive parameter, that is, the warp (bulk curvature) parameter $\\om$ appears. The limit $\\om\\ra 0$ leads to the flat case.  This introduction of the \"thickness\" $1/\\om$ comes from the expectation that it softens the UV-singularity, which is the same situation as in the string theory. See Ref.\\cite{SI0909} and \\cite{SI0912} besides this work. \\begin{figure} \\caption{ IR-regularized geometry of 5D warped space (\\ref{KKexp2}). The 4D world (3-brane) is Euclideanized and is shown as a 4D ball (shaded disk region) surrounded by $S^3$ sphere with radius $\\mu^{-1}$. $\\mu$ is the 4D IR regularization parameter. UV-regularization is introduced  by replacing the 4D ball with the \"sphere lattice\" composed of many small (size: $1/\\La$) 4D balls. See Sec.\\ref{surf} for detail. } \\begin{center} \\includegraphics[height=8cm]{RegGeomWarp} \\end{center} \\label{RegGeomWarp} \\end{figure}      The content is organized as follows. We start with the familiar approach to the 5D warped system: the Kaluza-Klein expansion, in Sec.\\ref{KKexp}. In Sec.\\ref{HKA}, the same content of Sec.\\ref{KKexp} is dealt in the {\\it heat-kernel} method and the Casimir energy is expressed in a {\\it closed} form in terms of the P/M propagator. The closed expression of Casimir energy enables us {\\it numerically} evaluate the quantity in Sec.\\ref{UIreg}. Here we introduce UV and IR regularization parameters in (4D momentum, extra coordinate)-space. A new idea about the UV and IR regularization is presented in Sec.\\ref{surf}. The {\\it minimal area principle} is introduced. The {\\it sphere lattice} and the {\\it renormalization flow} are explained. In Sec.\\ref{uncert} an improved regularization procedure is presented where a {\\it weight function} is introduced. Here again the {\\it minimal surface principle} is taken. The meaning of the weight function is given in Sec.\\ref{weight}. In Sec.\\ref{conc} we make the concluding remarks. {\\it Renormalization of the warp parameter $\\om$} is explicitly shown. We argue the (4D) coordinates or momenta look quantized in the present treatment. The cosmological constant is addressed. We prepare five appendices to supplement the text. App.A  deals with the {\\it classification} of all minimal surface curves in the 5D warped space. App.B reviews 4D Casimir energy (the ordinary radiation cavity problem) where the features of the cut-off and zeta-function regularizations are examined. App.C explains the numerical confirmation of the (approximate) equality of the minimal surface curve and the dominant path in the Casimir energy calculation. The results, appearing in this paper, heavily relies on some {\\it numerical calculations}. We explain them in App.D. Normalization constants of various weight functions are explained in App.E. ",
        "conclusions": "}  The log-divergence in (\\ref{uncert1c}) is the familiar one in the ordinary QFT. It can be {\\it renormalized} in the following way. \\bea E^W_{Cas}/\\La T^{-1} =-\\al\\om^4\\left( 1-4c\\ln (\\La/\\om)-4c'\\ln (\\La/T) \\right) =-\\al (\\om_r)^4\\com\\nn \\om_r=\\om\\sqrt[4]{1-4c\\ln (\\La/\\om)-4c'\\ln (\\La/T) }\\com \\label{conc1} \\eea where $\\om_r$ is the {\\it renormalized} warp factor and $\\om$ is the {\\it bare} one. No local counterterms are necessary. Note that this renormalization relation is {\\it exact} (not a perturbative result). In the familiar case of the 4D renormalizable theories, the coefficients $c$ and $c'$ depend on the coupling, but, in the present case, they are {\\it pure numbers}. \\footnote{ In the usual case, the log-terms (divergent terms) are separated and are canceled by the local counter-terms. It is difficult, in the present case, to take such a renormalization procedure because the theory is free and has no interaction terms (no couplings). The only choice, if we stick to the usual procedure, is the renormalization of the wave function and the mass parameter. It does not seem work well. Here we take a new approach. We regard the starting boundary parameter $\\om$ is a {\\it bare} quantity and the $\\om_r$ defined in (\\ref{conc1}) is a renormalized one. The boundary parameters flow {\\it by themselves}. {\\it No} local counter-terms are necessary. } It reflects the {\\it interaction between (EM) fields and the boundaries}. When $c$ and $c'$ are sufficiently small \\footnote{ In the list (\\ref{uncert1bX}), all data show $|4c|\\sim 10^{-1}$ and $|4c'|\\sim 10^{-1}$ . } we find the renormalization group function for the warp factor $\\om$ as \\bea |c|\\ll 1\\com\\q |c'|\\ll 1\\com\\q \\om_r=\\om (1-c\\ln (\\La/\\om)-c'\\ln (\\La/T) )\\com\\nn \\be (\\mbox{$\\be$-function})\\equiv \\frac{\\pl}{\\pl(\\ln \\La)}\\ln \\frac{\\om_r}{\\om}=-c-c' \\pr \\label{conc2} \\eea We should notice that, in the flat geometry case, the IR parameter (extra-space size) $l$ is renormalized (see Sec.\\ref{intro}). In the present warped case, however, the corresponding parameter $T$ is {\\it not renormalized}, but the {\\it warp parameter} $\\om$ {\\it is renormalized}. Depending on the sign of $c+c'$, the 5D bulk curvature $\\om$ {\\it flows} as follows. \\footnote{ Both IR and UV boundary interactions ($c'$ and $c$ ) contribute to the scaling behavior of the system. This should be compared with the usual perturbative (w.r.t. the coupling) case where $\\be$ is  calculated from one region, for example, UV-region. } When $c+c'>0$, the bulk curvature $\\om$ decreases (increases) as the the measurement energy scale $\\La$ increases (decreases). \\footnote{ This is the case of \"asymptotic free\" in the usual renormalization of the gauge coupling of 4D YM. } When $c+c'<0$, the flow goes in the opposite way. When $c+c'=0$, $\\om$ does not flow ($\\be=0$) and is given by $\\om_r=\\om (1+c\\ln(\\om/T))$.  The final result (\\ref{conc1}) is the {\\it new} type Casimir energy, $-\\om^4$. $\\om$ appears as a boundary parameter like $T$. The familiar one is $-T^4$ in the present context (See (\\ref{uncert1cc}). Note $T$ is the IR parameter and is related to $l$ as $T=\\om \\e^{-\\om l}$.). In ref.\\cite{IM05NPB}, another type $T^2\\om^2$ was predicted using a \"quasi\" Warped model (bulk-boundary theory).  Through the Casimir energy calculation, in the higher dimension, we find a way to quantize the higher dimensional theories within the QFT framework. The quantization {\\it with respect to the fields} (except the gravitational fields $G_{AB}(X)$) is done in the standard way. After this step, the expression has the summation {\\it over the 5D space(-time) coordinates or momenta} $\\int dz\\prod_adp^a$. We have proposed that this summation should be replaced by the {\\it path-integral} $\\int \\prod_{a,z}\\Dcal p^a(z)$ with the {\\it area} action (Hamiltonian) $A=\\int\\sqrt{\\det g_{ab}}d^4x$ where $g_{ab}$ is the {\\it induced} metric on the 4D surface. This procedure says the 4D momenta $p^a$ (or coordinates $x^a$) are {\\it quantum statistical} operators and the extra-coordinate $z$ is the inverse temperature (Euclidean time). We recall the similar situation occurs in the standard string approach. The space-time coordinates obey some uncertainty principle\\cite{Yoneya87}.  Recently the dark energy ( as well as the dark matter ) in the universe is a hot subject. It is well-known that the dominant candidate is the cosmological term. We also know the proto-type higher-dimensional theory, that is, the 5D KK theory, has predicted so far the {\\it divergent} cosmological constant\\cite{AC83}. This unpleasant situation has been annoying us for a long time. If we apply the present result, the situation drastically improve. The cosmological constant $\\la$ appears as \\bea R_\\mn-\\half g_\\mn R-\\la g_\\mn =T_\\mn^{matter}\\com\\nn S=\\int d^4x \\sqrt{-g}\\{\\frac{1}{G_N}(R+\\la) \\} +\\int d^4x \\sqrt{-g}\\{\\Lcal_{matter}\\}\\ \\com\\q g=\\mbox{det}~g_{\\mn}\\com \\label{conc3} \\eea where $G_N$ is the Newton's gravitational constant, $R$ is the Riemann scalar curvature. We consider here the 3+1 dim Lorentzian space-time ($\\mu,\\nu=0,1,2,3$). The constant $\\la$ observationally takes the value. \\bea \\frac{1}{G_N}\\la_{obs}\\sim \\frac{1}{G_N{R_{cos}}^2}\\sim m_\\n^4\\sim (10^{-3} eV)^4 \\com\\q \\la_{obs}\\sim \\frac{1}{R_{cos}^{~2}}\\sim 4\\times 10^{-66}(eV)^2\\com \\label{conc4} \\eea where $R_{cos}\\sim 5\\times 10^{32}\\mbox{eV}^{-1}$ is the cosmological size (Hubble length), $m_\\n$ is the neutrino mass. \\footnote{ The relation $m_\\n\\sim \\sqrt{M_{pl}/R_{cos}}=\\sqrt{1/R_{cos}\\sqrt{G_N}}$, which appears in some extra dimension model\\cite{SI0012,SI01Tohwa}, is used. The neutrino mass is, at least empirically, located at the {\\it geometrical average} of two extreme ends of the mass scales in the universe. } On the other hand, we have theoretically so far \\bea \\frac{1}{G_N}\\la_{th}\\sim \\frac{1}{{G_N}^2}={M_{pl}}^4\\sim (10^{28} eV)^4 \\pr \\label{conc5} \\eea This is because the mass scale usually comes from the quantum gravity. (See ref.\\cite{SI83NPB} for the derivation using the Coleman-Weinberg mechanism.) We have the famous huge discrepancy factor: \\bea \\frac{\\la_{th}}{\\la_{obs}}\\sim N_{DL}^{~2}\\com\\q N_{DL}\\equiv M_{pl}R_{cos}\\sim 6\\times 10^{60} \\com \\label{conc5b} \\eea where $N_{DL}$ is the Dirac's large number\\cite{PD78}. If we use the present result (\\ref{conc1}), we can obtain a natural choice of $T, \\om$ and $\\La$ as follows. By identifying $T^{-4}E_{Cas}=-\\al_1\\La T^{-1}\\om^4/T^4$ with $\\int d^4x\\sqrt{-g}(1/G_N)\\la_{ob}=R_{cos}^{~2}(1/G_N)$, we obtain the following relation. \\bea N_{DL}^{~2}=R_{cos}^{~2}\\frac{1}{G_N}=-\\al_1 \\frac{\\om^4\\La}{T^5}\\com\\q \\al_1\\ :\\ \\mbox{some coefficient} \\pr \\label{conc6} \\eea The warped (AdS$_5$) model predicts the cosmological constant {\\it negative}, hence we have interest only in its absolute value. \\footnote{ This fact strongly suggests de Sitter (dS$_5$) version of the present work could solve this sign problem. Details are under way. } We take the following choice for $\\La$ and $\\om$. \\bea \\La=M_{pl}\\sim 10^{19}GeV\\com\\q \\om\\sim \\frac{1}{\\sqrt[4]{G_N{R_{cos}}^2}}=\\sqrt{\\frac{M_{pl}}{R_{cos}}}\\sim m_\\n\\sim 10^{-3}\\mbox{eV} \\pr \\label{conc6b} \\eea The choice for $\\La$ is accepted in that the largest known energy scale is the Planck energy. The choice for $\\om$ comes from the experimental bound for the Newton's gravitational force.  As shown above, we have the standpoint that the cosmological constant is mainly made from the Casimir energy. We do not yet succeed in obtaining the value $\\al_1$ negatively, but succeed in obtaining the finiteness of the cosmological constant and its gross absolute value. The smallness of the value is naturally explained by the renormalization flow as follows. Because we already know the warp parameter $\\om$ {\\it flows} (\\ref{conc2}), the $\\la_{obs}\\sim 1/R_{cos}^2$ expression (\\ref{conc6b}), $\\la_{obs}\\propto \\om^4$, says that the {\\it smallness of the cosmological constant comes from the renormalization flow} for the non asymptotic-free case ($c+c'<0$ in (\\ref{conc2})). \\footnote{ A.M. Polyakov presented, in an early stage, the idea that the cosmological constant may be screened by the IR fluctuation of the metric\\cite{Pol82}. It was clearly shown, using the 2 dim R$^2$-gravity, such thing really occurs\\cite{SI95NPB}. } \\footnote{ We claim here the smallness of the cosmological constant is {\\it dynamically} explained (without fine tuning). }   The IR parameter $T$, the normalization factor $\\La/T$ in (\\ref{uncert1c}) and the IR cutoff $\\mu=\\La\\frac{T}{\\om}$ are given by \\bea T=R_{cos}^{~-1}(N_{DL})^{1/5}\\sim 10^{-20}eV\\com\\q \\frac{\\La}{T}=(N_{DL})^{4/5}\\sim 10^{50}\\com\\q \\mu=M_{pl}N_{DL}^{-3/10}\\sim 1GeV\\sim m_N \\com \\label{conc8} \\eea where $m_N$ is the nucleon mass. \\footnote{ Note that the present model predicts the nucleon mass scale (1 Gev) from the 3 data:\\ Planck mass $M_{pl}$, the neutrino mass $m_\\nu$, and the cosmological size $R_{cos}$ (or the cosmological constant $\\la_{obs}\\sim R_{cos}^{~-2}$). } The Fig.\\ref{IRUVRegSurfW} strongly suggests that the degree of freedom of the universe (space-time) is given by \\bea \\frac{\\La^4}{\\m^4}=\\frac{\\om^4}{T^4}=N_{DL}^{~6/5}\\sim 10^{74}\\sim (\\frac{M_{pl}}{m_N})^4 \\pr \\label{conc9} \\eea  In the recent exciting work by Ho\\u{r}ava\\cite{Hora0812,Hora0901}, the vastness of the string theory is stated as \"the string theory represents a logical completion of quantum field theory, not a single theory\". He tackles the renormalizability problem of the quantum gravity not from the string theory but from a \"small\" one, that is, a 3+1 dim local field theory with spacially higher-derivative interactions. The idea comes from the success of Lifshitz theory\\cite{Lifshitz41}, a spatially-higher-derivative scalar theory, in the condensed matter physics. The present approach has some points which can be compared with \\Hora 's. Basically both are the gravity-matter local field theory inspired by the string, brane and membrane theories. \\Hora 's one is basically 3+1 dimensional, while the present one 4+1 or 5 dimensional. In both ones, the renormalization flow plays a key role in the renormalizability (UV-completion). In the \\Hora 's, the Lorentz symmetry is abandoned as the starting principle, but is regarded as the dynamically emergent one (in IR region of RG flow). At the cost of Lorentz symmetry, he introduces spatially higher-derivative terms in order to suppress divergences in UV-region. This anisotropy between space and time (non-relativistic aspect) does not appear in the present approach because we treat the 4D world isotropically. Instead we have anisotropy between the 4D world and the extra space. In the present case, the renormalizability is realized by the warped configuration (thickness) and the appropriate suppression by the weight function. We need {\\it not} higher-derivative terms. The origin of the suppression is, at present, not established. Also in \\Hora 's, an uncertain procedure \"detailed balance condition\" is introduced in order to reduce the number of independent coupling constants. It looks important to take a {\\it new} standpoint or view, which has been overlooked so far, about some basic things, such as the starting symmetries, the quantum treatment of gravitational and matter fields, the regularization method and the meaning of the extra axis(es), to solve the {\\it divergence problem} of gravity."
    },
    "0812/0812.1041_arXiv.txt": {
        "abstract": "A study of the precision of the semiempirical methods used in the determination of the chemical abundances in gas-rich galaxies is carried out. In order to do this  the oxygen abundances of a total of 438 galaxies were determined using the electronic temperature, the $R_{23}$ and the P methods. The new calibration of the P method gives the smaller dispersion for the low and high metallicity regions, while the best numbers in the turnaround region are given by the $R_{23}$ method. We also found that the dispersion correlates with the metallicity. Finally, it can be said that all the semiempirical methods studied here are quite insensitive to metallicity with a value of $8.0\\pm0.2$ dex for more than $50\\%$ of the total sample. ",
        "introduction": "The determination of nebular abundances in H\\,{\\sc ii} regions is not a simple matter. It is well known that the ionic abundances depend on the intensity of the lines involved, and on the electronic density and temperature (Peimbert \\& Torres-Peimbert 1977; Aller 1984; Osterbrock 1989). For nearby bright galaxies, the weak auroral lines needed for the electronic temperature ($T_e$) determination can be detected when the signal-to-noise (S/N) is high enough. When the $T_e$ is determined, the ionic abundances can be easily obtained through the emissivities of the specific ions, when a certain ionization structure for the nebulae is assumed (e.g., Aller 1984;  Izotov et al. 2006 for a more recent review). Such a procedure is normally called the ``standard method'' for the abundance determination. The most important caveat concerning the standard method is that the auroral lines are always very weak. The typical uncertainty associated with the standard method is about $0.1$ dex (Kewley \\& Ellison 2008) but this increase significantly for low S/N spectra.  The situation gets worse when the forbidden auroral lines are absent. The most common reason for such absence is the low S/N of the spectra but it is not the only one (see Hoyos \\& D\\'{\\i}az 2006). In such situations the so-called semiempirical, or bright-line, methods need to be used in order to determine the metal content of the H\\,{\\sc ii} region. The most common and popular of these methods is the $R_{23}$ calibrator proposed by Pagel et al. (1979) and improved by McGaugh (1991), among others. Other calibrations of the $R_{23}$ method have been proposed using larger samples and more complete stellar evolutionary grids (e.g. Zaristky et al. 1994; Kewley \\& Dopita (2002); Kobulnicky \\& Kewley 2004). We prefer to use the calibration by McGaugh (1991). The main reason is that, although this calibration is old, it takes into account the influence of the ionization parameter on the chemical abundance determination.  In recent years, many other  semiempirical methods have been proposed, based on the same lines as the $P$ method (Pilyugin 2000,2001; Pilyugin \\& Thuan 2005), or on the nitrogen (the $N2$, Denicol\\'o et al. 2002) or the sulfur lines (the  S$_{23}$, D\\'{\\i}az \\& P\\'erez-Montero 2000; the $S_{234}$, Oey \\& Shields 2000)  All of the semiempirical methods present problems: bivaluation, large dispersion, dependence on other parameters (such as the ionization parameter or the nitrogen abundance), very large wavelength range between the lines involved, etc. Moreover, the large dispersion in the metallicity values obtained with the semiempirical methods can be due to the different H\\,{\\sc ii} regions geometries (which amount to different ionization parameters) and differences in age of the H\\,{\\sc ii} regions used for the calibration of the semiempirical methods. In addition, there are other influences such as the different apertures for different spectra which might mask the real calibration of the semiempirical methods (A.M. Hidalgo-G\\'amez 2009, in preparation).   The main interest of this investigation is to determine whether any of the semiempirical methods available or any of their different calibrations give better oxygen abundances than the others. In order to make such a comparison we think that the standard method abundances are good touchstones. The idea for this investigation comes out from to the neccesity of realiable abundance determinations with any of the semiempirical methods for a sample of dwarf spiral galaxies where the [OIII] $\\lambda$4363 is not detected (A.M. Hidalgo-G\\'amez et al 2009, in preparation). It is well known that galaxies with and without the [O\\,{\\sc iii}]$\\lambda$4363\\AA~ line have different properties on the $T_e$ range, ionization parameter, and luminosity (Hoyos \\& D\\'{\\i}az 2006). Therefore, the metal contents might significantly differ. Although we are aware of this, studies such as the one presented here might give some clues about the goodness of each of the semiempirical methods.     In the next section, the description of the sample of galaxies used in this investigation is carried out, while the comparison of the standard method with the semiempirical methods used is presented in Section 3. A comparison with other studies already published is presented in Section 4, and a brief discussion is given in Section 5. Finally, conclusions are presented in Section 6. ",
        "conclusions": "We have made a detailed comparison of the abundances obtained with the so-called standard method (SM) and some of the semiempirical methods used: $R_{23}$ (McGaugh's calibration) and $P$ (Pilyugin \\&  Thuan 2005). Our main interest is to obtain the closest abundances to the SM abundances when the standard method cannot be applied. In order to do that we used a large sample of late-type galaxies observed by SDSS and reduced by Kniazev et al. (2004). For all of the galaxies in the sample the oxygen forbidden line [O\\,{\\sc iii}]$\\lambda$4363 is detected, and therefore the SM abundances can be determined. The main advantage of this sample is that it is large enough for obtaining conclusive results and very homogeneous (being observed by only one facility and reduced and analyzed in the same way for all the galaxies); therefore, the systematic errors will affect all the data in the same way. This is not the usual procedure in the literature, where data, gathered from different sources, are obtained from different facilities, and analyzed and reduced by different researchers. A.M. Hidalgo-G\\'amez (2009, in preparation) found that this procedure might be a large source of uncertainties in the abundances determination. One of the disadvantages of the sample is that the use of only those galaxies where the auroral oxygen line was detected restrict the metallicity range, especially at high metallicity. Therefore, and keeping this last statement in mind, the conclusions obtained in this investigation are restricted to sample with similar metallicity range and characteristics.   The main conclusions could be listed as follows.  1- The distribution of the metallicity of the galaxies is a single gaussian for both the SM and the $R_{23}$ method. The main differences between then is the smaller width of the gaussian for the $R_{23}$ abundances. This clustering of the $R_{23}$ abundances around a single value ($\\approx 8.05$) is clearly visible at figures ~\\ref{fig2}b and ~\\ref{fig3}. Such behaviour might be real or due to the sample itself, which is restricted to an ``special'' type of galaxies (those with the auroral oxygen line detected). Moreover, as the number of this type of galaxies is a tiny fraction of the total SDSS catalog, SDSS imposes important bias to any sample. In order to find out if the clustering is real, a similar study has to be done with other large sample obtained from choosing the galaxies observed without any ``a priori'' assumption.   2- The $P$ calibration gives the lowest dispersion in both the high- and the low-metallicity regions, while the $R_{23}$ gives the best results in the ``turnaround'' region.  3- For both methods, there is a dependence of the dispersion with the metallicity for the low- and high-metallicity regions, while they are very insensitive to the metallicity in the ``turnaround'' region.  4- There is a dependence of the residuals with the abundances and the $T_e$ for both methods. For galaxies with log[O\\,{\\sc ii}]/O\\,{\\sc iii}]$>$ 1, the residuals are larger.  5- Finally, those galaxies with larger values of the equivalent width of H$\\beta$ seem to have very different abundances than the SM abundances, especially at low metallicity, but there is no clear trend between the residuals and the EW(H$\\beta$).   Finally, we might note that the results obtained here are similar to those found in other investigations, but due to the large and homogeneous sample used here, the results are more robust and the dispersion lower than in previous investigations."
    },
    "0812/0812.1794_arXiv.txt": {
        "abstract": "Principal component analysis is considered as an addition to the well-tested parametrization $w(a)=w_0+w_a(1-a)$ for the dark energy equation of state.  This brief note cautions against some unjustified assumptions in interpretation of PCA calculations, giving quantified examples. ",
        "introduction": "}  Dark energy is a premier mystery of cosmology and high energy physics. To address this, NASA and the Department of Energy (DOE) are funding the Joint Dark Energy Mission (JDEM).  Because the physical nature of dark energy is so unknown, it is challenging to quantify simply and accurately the science requirements for learning about the physical origin.  To assess approaches to a figure of merit for the dark energy science reach of different JDEM architectures, NASA and DOE formed the Figure of Merit Science Working Group (FOMSWG).  FOMSWG found in a preliminary report \\cite{fomswg} that the figure of merit used by the Dark Energy Task Force \\cite{detf}, given by the inverse area of the likelihood contour in the dark energy equation of state plane $w_0$-$w_a$ was reasonable for the task.  Here the equation of state (EOS) as a function of scale factor is given by $w(a)=w_0+w_a(1-a)$. This form has been tested for physical accuracy and against bias, and has been shown to faithfully reproduce relations for observables such as distances and Hubble parameters to an accuracy of $10^{-3}$ \\cite{calib}, better than needed for JDEM.  These results, and especially the calibration relations of \\cite{calib}, should substantially allay concerns about using a particular functional form.  Recall that the current issue is how to {\\it project\\/} simulated constraints of various JDEM scenarios; once the data are in hand one will carry out the analysis through a diversity of methods, and test specific models directly against the data without an intermediate form. Nevertheless, one can reasonably consider an {\\it additional\\/} (not replacement) approach to build confidence that the conclusions on science design were robust.  An example is principal component analysis (see, e.g., \\cite{hutstar,albern,deplpca,mort} for application to dark energy), which has the capability of covering a wide variety of functional forms.  The assessment of whether principal components (PCs) add new insights to JDEM design projections depends on appropriate scientific criteria. Here we present brief cautions about possible oversimplifications of interpretation, titling the following section headings with some of the possible misunderstandings. ",
        "conclusions": "Principal component analysis is a valid technique, used appropriately.  Oversimplifying PCA interpretation or inadequately appreciating the effect of assumptions employed can lead to misunderstandings and false beliefs.  We present cautionary examples of three apparently plausible but unjustified extrapolations.  While data should be analyzed in every reasonable manner, for understanding the generic cosmology reach the more complicated PCA approach demonstrates no extraordinary advantage over the well-tested and highly calibrated phase space dynamics approach of $w_0$-$w_a$."
    },
    "0812/0812.3272_arXiv.txt": {
        "abstract": "Understanding the origin and evolution of dwarf early-type galaxies remains an important open issue in modern astrophysics. Internal kinematics of a galaxy contains signatures of violent phenomena which may have occurred, e.g. mergers or tidal interactions, while stellar population keeps a fossil record of the star formation history, therefore studying connection between them becomes crucial for understanding galaxy evolution.  Here, in the first paper of the series, we present the data on spatially resolved stellar populations and internal kinematics for a large sample of dwarf elliptical (dE) and lenticular (dS0) galaxies in the Virgo cluster. We obtained radial velocities, velocity dispersions, stellar ages and metallicities out to 1--2 half-light radii by re-analysing already published long-slit and integral-field spectroscopic datasets using the {\\sc NBursts} full spectral fitting technique. Surprisingly, bright representatives of the dE/dS0 class ($M_B = -18.0 \\ldots -16.0$~mag) look very similar to intermediate-mass and giant lenticulars and ellipticals: (1) their nuclear regions often harbour young metal-rich stellar populations always associated with the drops in the velocity dispersion profiles; (2) metallicity gradients in the main discs/spheroids vary significantly from nearly flat profiles to $-0.9$~dex~$r_e^{-1}$, i.e. somewhat 3 times steeper than for typical bulges; (3) kinematically decoupled cores were discovered in 4 galaxies, including two with very little, if any, large scale rotation. These results suggest similarities in the evolutionary paths of dwarf and giant early-type galaxies and call for reconsidering the role of major mergers in the dE/dS0 evolution. ",
        "introduction": "Dwarf elliptical (dE) and lenticular (dS0) galaxies, low-luminosity ($M_B > -18.0$~mag) stellar systems with early-type morphology, represent the numerically dominant galaxy population in nearby galaxy clusters and groups \\citep{FB94}. In the cold dark matter models of hierarchical galaxy formation they are considered as the building blocks of presently observed stellar systems (e.g. \\citealp{WF91}). Thus, understanding their formation and evolution becomes one of the important questions of modern astrophysics. Several dE/dS0 formation scenarii were proposed (see review and discussion in \\citealp{deRijcke+05}), but none of them is able to fully explain all observational properties of these galaxies simultaneously. A possible diversity of evolutionary scenarii has been suggested by \\cite{vZBS04} and later investigated by \\cite{LGB08} to explain properties of different dE subclasses, although \\cite{BBCG08_1} insist on the ram pressure stripping to be the unique channel of dE formation.  Dwarf and giant early-type galaxies form two distinct sequences in the absolute magnitude ($M_B$) vs effective surface brightness ($\\langle \\mu \\rangle_e$) diagram (see e.g. \\citealp{FB94}) joining around $M_B = -18.0$~mag, which is often referred as a luminosity separating these two classes (see \\citealp{BBCG08} and \\citealp{KFCB08} for the recent discussions). Although, some arguments are provided (e.g. \\citealp{GG03}) for the continuity of the dwarf--giant sequence on the $M_B$ vs $\\langle \\mu_e \\rangle$ plot, it is clear that the effective surface brightness of dE galaxies correlates with their luminosity, or, in other words, less luminous galaxies are fainter in terms of surface brightness, making their observations a challenging task. One has to keep in mind that dE galaxies usually lack of the interstellar medium (ISM), therefore their spectra normally contain no emission lines. This explains why the first kinematical study of dEs \\citep{BN90} appeared more than a decade after the publication of the absorption-line kinematics of giant elliptical galaxies (e.g. \\citealp{BC75}).  New debates addressing the dE formation appeared recently: \\cite{BBCG08_1,BBCG08} argue for the ram pressure stripping to be the only scenario of dE formation, hence making these galaxies evolutionary different from giant early-type systems and reproducing their locus on the $M_B$ vs $\\langle \\mu \\rangle_e$ plot; while \\cite{JL08} succeeded to explain the observed discontinuity of giant and dwarf ellipticals in the size--luminosity relation by comparing a large homogeneous imaging dataset to the predictions of semi-analytical models of elliptical galaxy formation, thus suggesting their common origin. Comparing the structural properties of dwarf galaxies to their internal kinematics and stellar populations therefore becomes a question of major importance for understanding dE/dS0 formation and evolution.  The significant progress in the astronomical instrumentation and new detectors having low readout noise became the cornerstones to the accomplishments of several projects assessing spatially-resolved kinematics of diffuse elliptical galaxies in different environmental conditions. Kinematical profiles presented in \\cite{DRDZH01,DRDZH03,GGvdM02,GGvdM03,Pedraz+02,SP02,vZSH04} revealed a diversity of the degree of rotational support and presence of kinematically-decoupled components in some objects \\citep{DRDZH04,GGvdM05,PCSA05,Thomas+06}. Detailed studies of the morphological properties of dEs revealed embedded structures such as stellar discs, sometimes harbouring low-contrast spiral arms, or bars in many galaxies \\citep{JKB00,BBJ02,LGB06}. These discoveries strengthen the evolutionary connection between dwarf ellipticals and discy late-type dwarfs.  At the same time, studies of the stellar population properties of dwarf elliptical galaxies still remain a challenge both for observations and for the data analysis, because absorption line strengths, a traditional spectroscopic indicator of stellar populations (see e.g. \\citealp{WFGB94} for the description of the Lick system) require high signal-to-noise ratios in order to reach reasonable quality of age and metallicity estimates. Some attempts were made to derive dE stellar populations from the narrow-band photometry \\citep{Rakos+01}. However, the calibrations used by the authors were based on Galactic globular clusters, therefore an assumption about very old ages was done, resulting in low metallicity estimates ($-1.4$ to $-0.6$~dex). The integrated measurements of dE stellar populations from the Lick indices were made later \\citep{GGvdM03,vZBS04} and revealed populations having intermediate ages of 3--5~Gyr, and considerably higher metallicities around $-0.5$~dex. These results were confirmed and extended to a larger sample of Virgo cluster dEs recently \\citep{Michielsen+08}.  A new family of techniques for recovering stellar population properties from integrated light spectra appeared recently \\citep{CidFernandes+05,OPLT06,CPSA07}. These methods are based on the direct fitting of the the stellar population models against the observed spectra in the pixel space. Since every pixel of an absorption-line galactic spectrum bears information about the stellar content of a galaxy, these techniques are potentially much more sensitive than the line-strength indices dealing with the individual spectral features. Another advantage of the pixel space fitting is an easy way of getting rid of residual cosmic ray hits and regions of the spectra contaminated by emission lines by excluding these regions from the fitting procedure.  The full spectral fitting using the {\\sc NBursts} code (see preliminary description in \\citealp{Chilingarian06,CPSK07}) allowing to extract internal kinematics simultaneously with the parametrized star formation history, was successfully used in the studies of integrated kinematical and stellar population properties of dE galaxies in the Fornax \\citep{Michielsen+07} and Abell~496 \\citep{Chilingarian+08} clusters. The method was shown to be stable and working well even for the data having low signal-to-noise ratios.  Up-to now, a published sample of dwarf and low-luminosity early-type galaxies with the spatially-resolved stellar population information (apart from Local Group dwarfs studied by means of colour-magnitude diagrams) contains only seven objects, all of them observed with the Multi-Pupil Field Spectrograph \\citep{ADM01} at the 6-m Bolshoi Teleskop Al'tazimutal'nij (BTA) of the Special Astrophysical Observatory of the Russian Academy of Sciences (SAO RAS). The sample includes four dE/dS0s in the Virgo cluster \\citep{CPSA07,CSAP07} and three low-luminosity E/S0s in groups: NGC~770 \\citep{PCSA05}, NGC~126 and NGC~130 \\citep{CSAP08}. In all cases the two-dimensional maps of kinematics and stellar population properties were obtained using the {\\sc NBursts} technique. Chemically and evolutionary decoupled structures which are quite common in giant elliptical and lenticular galaxies \\citep{Sil06,Kuntschner+06} were detected in the central regions of six low-luminosity galaxies.  Presently, the growing amount of high-quality optical spectroscopic data on dwarf galaxies becomes available in the archives of major observatories. Therefore, we decided to re-analyse the existing and available spectroscopic data for dE/dS0 galaxies using the full spectral fitting technique. The project aims at studying the connection between internal kinematics and stellar populations of dE galaxies and understanding the mechanisms of their formation and evolution, primarily in the cluster environment.  In the first paper of the series we provide the kinematical and stellar population profiles of Virgo cluster dE/dS0 galaxies obtained from the new analysis of several published datasets carried out using the {\\sc NBursts} full spectral fitting technique \\citep{CPSK07,CPSA07}. The thorough discussion of the obtained results will be given in the next papers. The paper is organized as follows: in Section 2 we describe the spectroscopic data, data reduction procedures, and define a sample of galaxies, Section 3 provides details about the techniques used to analyse the data, Section 4 presents the results of this analysis, the conclusions are given in Section 5. ",
        "conclusions": "We have presented the first large dataset of spectroscopically derived radial profiles of stellar population parameters for dwarf early-type galaxies made by re-analysing some of the published spectroscopic data. Thanks to the usage of the {\\sc NBursts} full spectral fitting technique we have improved the stellar kinematics compared to \\cite{vZSH04} reaching comparable quality of measurements to those presented in \\cite{GGvdM02,GGvdM03} but generally going out to 1~$r_e$ or further from the centre. We have performed the stellar population analysis of the spectra where only kinematical data were available \\citep{SP02}. Our measurements of the nuclear stellar populations are in a good agreement with the results obtained from the fitting of SDSS DR6 spectra.  Significant metallicity gradients are often observed in dE/dS0 galaxies, although in some of the objects the metallicity distribution is completely constant along the radius. There is a tendency to the dichotomy in the distribution of the metallicity gradients -- metallicity profiles are either nearly flat or steeply decreasing along the radius, reaching $-0.9$~dex~$r_e^{-1}$ in VCC~437. In none of the cases we see a statistically significant positive gradient. Age distribution along the radius usually exhibits flat behaviour. If any gradients in age exist, they cannot be clearly detected using the existing datasets.  Levels of $\\alpha$-enhancement as derived from the Lick indices exhibit no statistically significant changes along the radius and between the nuclei and main discs/spheroids in all objects, but VCC~1036, where the value in the nuclear region is nearly solar ($[$Mg/Fe$]$=$-0.05$~dex), whereas it stays super-solar (+0.17~dex) along the radius out to 1.2~$r_e$.  Stellar populations in the nuclear regions are usually more metal-rich than in the main galactic bodies, and the central values of the luminosity-weighted metallicity are higher than what one would expect from the simple interpolation of the metallicity gradients toward the centres. A remarkable fraction of galaxies (11 or 27) demonstrate luminosity-weighted ages in the nuclear regions to be younger than those of the main galactic bodies, with the difference exceeding 7~Gyr in VCC~917 and VCC~490.  These chemically and evolutionary decoupled nuclear regions are always associated with the drops in the velocity dispersion profiles. Only in those cases, where the velocity dispersions are too low to be measured with the present data, the $\\sigma$-drops are not observed. In these cases, the mass-to-light ratios change significantly along the radius, pointing out to unrighteous use of constant $M/L$ dynamical models.  In certain objects (VCC~990, 1122, 1250, 1261, 2019, and 2050) we detected faint emission lines (H$\\beta$, [OIII]) in the nuclear regions, associated with the chemically and evolutionary decoupled cores.  We discovered four kinematically decoupled rotating cores in bright members of our sample (VCC~917, 1036, 1261, and 1545), two of them (VCC~917 and 1261) being hosted by the galaxies with very low if any major axis rotation. In all cases KDCs are associated with young metal-rich stellar populations and velocity dispersion drops.  Finally, we see remarkable resemblance between internal structure of dE/dS0s and of intermediate-mass and giant early-type galaxies. Some low-luminosity objects after detailed analysis look like ``scaled down'' versions of more luminous counterparts (e.g. VCC~917/1261 and NGC~5813, VCC~1545 and NGC~4365), suggesting possible similarities between their evolutionary paths. Some physical processes, e.g. major mergers and secular evolution, are usually considered to rule the life of giant early-type galaxies. We have to revise their applicability to explain the dE/dS0 evolution."
    },
    "0812/0812.2618_arXiv.txt": {
        "abstract": "{We are engaged in a multi-wavelength study of several Galactic \\HII\\ regions that exhibit signposts of triggered star formation on their borders, where the collect and collapse process could be at work.} {When addressing the question of triggered star formation, it is critical to ensure the {\\emph{real}} association between the ionized gas and the neutral material observed nearby. In this paper we stress this point and present CO observations of the RCW~82 star forming region.} {The velocity distribution of the molecular gas is combined with the study of young stellar objects (YSOs) detected in the direction of RCW~82. We discuss the YSO's evolutionary status using near- and mid-IR data. The spatial and velocity distributions of the molecular gas are used to discuss the possible scenarios for the star formation around RCW~82.} {Several massive molecular condensations, together with star formation sites, are observed on the borders of RCW~82.  The shapes of the three brightest condensations suggest that they were pre-existent, i.e. not formed through the collect and collapse process. A thin layer of molecular material is observed surrounding the ionized gas, adjacent to the ionization front. This results from the sweeping up of neutral material around the expanding region. Several Class~I YSOs are detected in the direction of this layer.} {The numerous YSOs observed towards the bright molecular condensations bordering (and velocity-associated with) the ionized gas reveal the intense star formation activity in RCW~82. But this region is probably too young to have triggered star formation via the collect and collapse process.} ",
        "introduction": "}  The formation of massive stars can be triggered by \\HII\\ regions. Massive molecular condensations on the borders of Galactic \\HII\\ regions are the most likely sites for star formation and are consequently the most likely locations for observing the earliest stages of stellar births.  Several recent \\HII\\ region studies (Deharveng et al.~\\cite{deh03}, Zavagno et al.~\\cite{zav06}, Watson et al.~\\cite{Wat08}, Cappa et al.~\\cite{cap08}, Kirsanova et al.~\\cite{kir08}) have taken up the challenge to understand the triggered star formation processes, even if some points still remain not clear.  In the collect and collapse process, proposed by Elmegreen \\& Lada (\\cite{elm77}), the expansion of the ionized region leads to the formation of a dense layer of material, collected between the ionization front (IF) and the shock front, during the supersonic expansion of the \\HII\\ region. This layer eventually becomes gravitationally unstable along its length and fragments, leading to the formation of massive condensations that are potential sites for subsequent star formation. In 2005, Deharveng et al.\\ proposed a list of candidate Galactic \\HII\\ regions where the collect and collapse process may be the main triggering agent for the star formation observed on their borders. We have shown this process to be at work in a number of regions, among them Sh~104 (Deharveng et al. \\cite{deh03}) and RCW~79 (Zavagno et al.~\\cite{zav06}). We have also shown that this process works in inhomogeneous turbulent clouds as in the cases of Sh2-212 (Deharveng et al.~\\cite{deh08}), Sh2-217 (Brand et al. in preparation) and Sh2-241 (Pomar\\`es et al. in preparation).  The observations of an increasing number of sources show that the layer of collected neutral material is always present around the ionized gas and that various triggering processes are at work there.  Information about the distribution of associated molecular material is frequently lacking for these regions. We therefore carried out line observations of the molecular material, at velocities bracketing those of the surrounding ionized gas. Without such radial velocities we are never sure that a given condensation is indeed associated with the \\HII\\ region we are studying. The same is true of the YSOs that are observed {\\emph{in the direction of}} the region, with no more than a reasonable probability of being associated with it, based on the spatial distribution of these types of sources. We therefore carried out a series of molecular observations around \\HII\\ regions using the Mopra radio-telescope. The first results are presented here for RCW~82, along with a discussion of the YSO population using new near-IR, Spitzer-GLIMPSE and MIPSGAL mid-IR data.  This paper is organized as follows: Sect.~2 presents the Galactic \\HII\\ region RCW~82. New observations towards this region are presented in Sect.~3. Results are given in Sect.~4. Sect.~5 presents a general discussion about the morphology of the region and the star formation processes at work there. Conclusions are given in Sect.~6. ",
        "conclusions": "  $\\bullet$ RCW~82 is seen on the border of a large shell observed at 8~$\\mu$m, which probably surrounds an \\HII\\ region. Unfortunately the spatial coverage of our molecular observations is not sufficient to prove the link between the two regions;  $\\bullet$ We have identified four candidate exciting stars of RCW~82. Among these sources, stars $\\#$e2 and $\\#$e3 are probably O-type stars responsible for the ionization of RCW~82;  $\\bullet$ Molecular material is associated with RCW~82. We observe a molecular layer around the \\HII\\ region, corresponding to interstellar material accumulated around the ionized gas during the expansion of RCW~82. We also observe elongated structures, oriented almost perpendicular to the ionization front. We suggest that these structures correspond to pre-existing condensations, possibly shaped by UV radiation. Some condensations have masses greater than 2000~$\\msol$;  $\\bullet$ Many YSOs are observed at the periphery of RCW~82, among which is a non-negligible number of potential Class~I sources. Their distribution is not uniform and most of them are observed on the borders of RCW~82, close to the IF. It indicates that triggered star formation is at work around this \\HII\\ region;  $\\bullet$ The presence of a collected layer of material around RCW~82 shows that the collect part of the collect and collapse process has occurred. However, our age estimate of RCW~82 shows that the \\HII\\ region is probably not evolved enough for the collected shell to be in the fragmentation phase. This means that the star formation activity observed around RCW~82 is due to other processes;  $\\bullet$ Among the YSOs we have identified a candidate isolated massive young object, source $\\#$106; its evolutionary stage remains to be more precisely defined. The formation of this object may have been triggered by RCW~82, as it is observed in the direction of the shell of material collected around the \\HII\\ region. However this point remains to be confirmed as the estimated age of the \\HII\\ region does not allow enough time for the collapse process to occur. On the other hand, the young cluster, observed around source $\\#$51, could be associated with a massive molecular component observed around $-55$~km~s$^{-1}$, for which the association with RCW~82 is uncertain."
    },
    "0812/0812.0337_arXiv.txt": {
        "abstract": "The stability of magnetized strange quark matter (MSQM) is investigated within the phenomenological MIT bag model, taking into account the variation of the relevant input parameters, namely, the strange quark mass, baryon density, magnetic field and bag parameter. We obtain that the energy per baryon decreases as the magnetic field increases, and its minimum value at vanishing pressure is lower than the value found for SQM. This implies that MSQM is more stable than non-magnetized SQM. Furthermore, the stability window of MSQM is found to be wider than the corresponding one of SQM. The mass-radius relation for magnetized strange quark stars is also derived in this framework. ",
        "introduction": "\\label{sec1}  Strange quark matter (SQM), in contrast to ordinary nuclei (where quarks are confined into colorless nucleons), is thought to be a form of macroscopic matter which contains a large quantity of deconfined quarks ($u$, $d$ and $s$) in $\\beta$-equilibrium, with electric and color charge neutrality. Theoretically, Bodmer~\\cite{Bodmer:1971we} and Witten~\\cite{Witten:1984rs} proved that SQM could be the more stable phase of nuclear matter. The reason behind this crucial conclusion is the presence of the quark $s$ with a mass lower than the Fermi energy, which leads to a decreasing of the total energy per baryon $E/A$, if some of the $u$ and $d$ quarks on the Fermi surface are converted into $s$ quarks via weak interactions. This is, in other words, a consequence of the Pauli exclusion principle.  If the hypothesis of the existence of SQM is true, this form of matter could have interesting astrophysical consequences. For instance, SQM could be realized in the inner core of compact objects, such as neutron stars, or even form the hypothetical strange stars. It may be possible as well that small quarks droplets could be formed in relativistic heavy ion collider (RHIC) experiments~\\cite{RHIC}. As SQM is more stable and bound at finite density, one expects configurations of quark stars with macroscopic properties quite different from neutron stars~\\cite{itoh,Hansel}. These objects would be self-bound and their masses would scale with the radius as $M \\sim R^{3}$, in contrast to neutron stars which have masses that decrease with increasing radius ($M \\sim R^{-3}$) and are bound by gravity. Thus, they could explain a set of peculiar astrophysical observations which are awaiting for theoretical explanations~\\cite{Lugones:2002va, Lugones:2002zd,Quan,Nice,Xu}.  SQM has been widely studied using the MIT bag model, where the non-interacting quark gas at zero temperature is considered to be confined into a phenomenological bag, $B_{\\rm bag}$, which mimics the strong interaction. Such studies have revealed that even heavy strangelets might be stable, with the corresponding $E/A$ lower than that of the iron nuclei, $E/A\\,(^{56}{\\rm Fe}) \\simeq 930$~MeV~\\cite{Farhi:1984qu}. Being a phenomenological model, the MIT bag model has some limitations that are well known. The confinement is put by hand, and the model is unable to reproduce the chiral symmetry breaking at zero density. Nevertheless, the model is quite satisfactory in describing quark matter at finite density.  Focusing on a real astrophysical scenario, one should also note the important role played by the strong magnetic fields expected in this context. Pulsars, magnetars, neutron stars, the emission of intense sources of X-rays could be associated to sources with intense magnetic fields around $10^{13}-10^{15}$~G or even higher fields~\\cite{duncan,kouve}. We remark that the magnetic field intensity may vary significantly from the surface to the center of the source. Although there is no direct observational evidence for the magnetic field strength in the inner core of stars, theoretical estimates indicate that fields as high as $10^{18}$~G could be allowed. Therefore, SQM in the presence of strong magnetic fields is a subject that it is worth addressing. In a recent work~\\cite{Felipe:2007vb}, we investigated the SQM properties in the presence of a strong magnetic field, taking into account $\\beta$- equilibrium and the anomalous magnetic moments (AMM) of the quarks. It was found that the stability of the system requires a magnetic field value $B \\lesssim 10^{18}$~G. This is in contrast to the bound $B \\lesssim 10^{19}$~G, obtained when AMM are not included~\\cite{Chakrabarty:1996te}.  The scope of the present paper is to establish a bridge between the microscopic properties of magnetized SQM (MSQM), given by the equation of state (EoS) of the system, and the macroscopic observables of strange quark stars (SQS), derived from this kind of magnetized matter. In the framework of the MIT bag model, we study the behavior of the system with the variation of its parameters, namely, the bag parameter $B_{\\rm bag}$, the strange quark mass $m_s$, the baryon density $n_B$ and the magnetic field $B$. The possible mass-radius configurations of magnetized strange quark stars (MSQS) are then obtained by solving the Tolman-Oppenheimer-Volkoff (TOV) equation. As it turns out, stable stars with smaller radii are allowed due to the compactness of matter and the presence of a strong magnetic field, since the energy per baryon at vanishing pressure is lower in this case.  The paper is organized as follows. In section \\ref{sec2} we briefly review the thermodynamic properties of SQM in the presence of a strong magnetic field. In section \\ref{sec3} the stability windows of SQM in the presence of a magnetic field are obtained varying the relevant input parameters of the model. Section \\ref{sec4} is devoted to the study of the mass-radius relation for MSQS by numerically solving the TOV equations. A comparison with some of the available observational data~\\cite{Psaltis} is also presented. Finally, our conclusions are given in section \\ref{sec5}. ",
        "conclusions": "\\label{sec5}  In the present work, we have investigated the stability of magnetized strange quark matter within the phenomenological MIT bag model. We have studied the stability windows of MSQM taking into account the variation of the strange quark mass, the baryon density, the magnetic field and the bag parameter. We found that MSQM is indeed more stable than non-magnetized SQM. While for SQM the stability range for the baryon density is $1.8 \\lesssim n_B/n_0 \\lesssim 2.4$ for $50 \\leq m_s \\leq 300$~MeV, MSQM allows densities in the range $1.85 \\lesssim n_B/n_0 \\lesssim 2.6$ for $50 \\leq m_s \\leq 240$~MeV and a magnetic field value of $5\\times 10^{18}$~G. Moreover, the allowed range for the bag parameter is $57 \\lesssim B_{\\rm bag} \\lesssim 90$~MeV/fm$^3$.  At vanishing pressure and finite density, the derived EoS for MSQM exhibits a minimum of the energy per baryon lower than that of $^{56}$Fe. This gives the possibility of existence of stable configurations for magnetized strange quark stars. In this simple framework for MSQM, we have then derived the mass-radius relation for such compact stars and compared the theoretically predicted relations with some of the recent observational data. We have concluded that in the presence of a strong magnetic field, stable strange stars with smaller masses and radii could exist.   \\ack{ The authors acknowledge the fruitful discussion with J. Horvath and his important suggestions and comments. A.P.M. thanks CFTP-IST~(Lisbon,Portugal) for their hospitality and financial support. This work was partially supported by \\emph{Funda\\c c\\~ ao para a Ci\\^ encia e a  Tecnologia} (FCT, Portugal) through the projects POCI/FP/81919/2007, and CFTP-FCT UNIT 777,  which are partially funded through POCTI (FEDER). The work of R.G.F. has been partially supported by the Marie Curie Research Training Network MRTN-CT-2006-035505. The work of A.P.M. has been supported by \\emph{Ministerio de Ciencia, Tecnolog\\'{\\i}a y Medio Ambiente} under the grant CB0407 and the ICTP Office of External Activities through NET-35.  A.P.M. also acknowledges TWAS-UNESCO for financial support at CBPF-Brazil.\\\\}"
    },
    "0812/0812.2568_arXiv.txt": {
        "abstract": "We present two-dimensional hydrodynamic models of thermally driven winds from highly irradiated, close-in extra-solar planets.  We adopt a very simple treatment of the radiative heating processes at the base of the wind, and instead focus on the differences between the properties of outflows in multidimensions in comparison to spherically symmetric models computed with the same methods.  For hot ($T \\gapprox 2\\times10^{4}$~K) or highly ionized gas, we find strong (supersonic) polar flows are formed above the planet surface which produce weak shocks and outflow on the night-side.  In comparison to a spherically symmetric wind with the same parameters, the sonic surface on the day-side is much closer to the planet surface in multidimensions, and the total mass loss rate is reduced by almost a factor of four.  We also compute the steady-state structure of interacting planetary and stellar winds.  Both winds end in a termination shock, with a parabolic contact discontinuity which is draped over the planet separating the two shocked winds.  The planetary wind termination shock and the sonic surface in the wind are well separated, so that the mass loss rate from the planet is essentially unaffected. However, the confinement of the planetary wind to the small volume bounded by the contact discontinuity greatly enhances the column density close to the planet, which might be important for the interpretation of observations of absorption lines formed by gas surrounding transiting planets. ",
        "introduction": "A significant fraction (about 20\\%) of extra-solar giant planets (EGPs) discovered to date have an orbit with a semi-major axis of less than 0.1 AU (Schneider 2008).  For such close-in EGPs, heating of the upper layers of the atmosphere by irradiation from the central star, especially in the extreme ultraviolet (EUV), can produce an extended envelope of gas, and perhaps even drive a wind (e.g.  Moutou et al. 2000; see Ehrenreich 2008 for a review).  Direct detection of an extended envelope surrounding HD209458b has been reported (Vidal-Madjar et al. 2003; 2008; Ehrenreich et al. 2008), based on the absorption of stellar Ly-$\\alpha$ during transits. The fraction of the stellar Ly-$\\alpha$ flux absorbed by HD209458b is so large that it must be surrounded by a cloud of neutral hydrogen that extends beyond the Roche lobe of the planet, and therefore is unbound (although see Ben-Jaffel 2007 for an alternative analysis of the same observations).  The observations also place a lower limit on the mass loss rate from the planet of $\\dot{M} \\gapprox 10^{10}$~g~s$^{-1}$, although the actual rate could be much higher since the absorption is saturated.  Detailed theoretical models of winds from highly irradiated EGPs are of interest not only to interpret the observations, but also to place firmer constraints on the mass loss rates.  If these rates are a factor of few hundred times higher than the lower limit observed in HD209458b, then over its lifetime the planet will lose a significant enough fraction of its total mass to alter its structure and evolution (Hubbard et al. 2007; Lecavelier des Etangs 2007).  A useful measure of the strength of the wind expected from a highly irradiated EGP is the ratio of gravitational potential to thermal energy at the top of the atmosphere, usually termed the hydrodynamic escape parameter, that is \\begin{equation} \\lambda = \\frac{GM_{p} \\mu}{R_{p}kT} \\end{equation} where $M_{p}$ and $R_{p}$ are the mass and radius of the planet, and $T$ and $\\mu$ are the temperature and mean mass per particle in the atmosphere. For $\\lambda \\gg 1$, the atmosphere is tightly bound and a hydrodynamic wind is not expected, although a weak outflow may still be produced by a variety of non-thermal processes, e.g. Hunten (1982).  For $\\lambda \\lapprox 10$, a thermally-driven hydrodynamic (Parker) wind will be produced (for reference, $\\lambda \\approx 15$ for $10^{6}$~K plasma in the solar corona at 2$R_{\\odot}$).  When evaluated at the effective temperature $T_{\\rm eff} \\approx 10^{3}$~K of HD209458b, $\\lambda \\approx 140$, indicating a hydrodynamic wind is unlikely.  However, it has been pointed out the upper layers of close-in EGPs will be heated to $T \\sim 10^{4}$~K by the intense EUV radiation from the central star (Lammer et al. 2003; Yelle 2004; see Ballester et al. 2007 for possible observational detection of this hot gas).  At this temperature, and if the gas is mostly neutral, $\\lambda \\approx 14$ (a value similar to that in the solar corona) so that a thermally-driven wind is possible.  On the other hand, if the gas is ionized (reducing $\\mu$), or if the temperature is slightly higher (both of which are relevant for EGPs around YSOs or higher mass stars), then $\\lambda$ can be even smaller.  Calculating the structure and mass loss rate from a thermally-driven hydrodynamic wind from close-in EGPs requires solving the time-dependent equations of hydrodynamics, with a proper accounting of the radiative heating at the base of the wind.  The latter is probably the most challenging aspect of the problem, and different approaches have led to different estimates of the mass loss rates.  To simplify the calculation, all of the models presented to date are one-dimensional (spherically symmetric).  Watson et al. (1981) presented models for the escape of hydrogen from the terrestrial planets in the early solar system, using a simplified treatment of the heating that deposits all of energy in a single zone at the base of the wind.  Scaled-up versions of these models have been applied to EGPs.  Tian et al (2005) improved these solutions by using a multi-dimensional radiative transfer calculation to estimate the radial distribution of the heating rate (the underlying hydrodynamic models are still 1D, however).  Lecavelier des Etangs et al. (2004) pointed out that tidal forces could distort the gravitational potential isosurfaces around the planet, and this could affect the mass loss rate (Erkaev et al. 2007), but only if the shape and location of the sonic surface are significantly altered.  Yelle (2004) has presented more realistic models of the wind from close-in EGPs including chemistry, photo-ionization and recombination, and thermal and molecular diffusion in a hydrodynamic model (again assuming spherical symmetry).  New models with a similar level of sophistication have recently been computed for HD209458b by Garc\\'{i}a Mu\\~{n}oz (2007), and for EGPs in general by Murray-Clay et al. (2008).  As has been pointed out by many authors, the fact that close-in EGPs are irradiated on only one side calls into question the assumption of spherical symmetry.  The problem is exasberated by the fact that close-in EGPs will be tidally locked, so that the only mechanisms that can transport heat to the night-side are bulk flows in the atmosphere and thermal conduction. Although global circulation models are beginning to be investigated for close-in EGPs (e.g. Dobbs-Dixon \\& Lin 2008; Showman et al. 2008), these models do not apply to the weakly bound hot ($T \\sim 10^{4}$~K) gas in the upper atmosphere. For HD209458b, the advection time for gas moving at the sound speed to move half way around the planet is $\\approx 11$ hours, much longer than the radiative cooling time for the hot gas.  Thus, there is likely to be a large contrast in the temperature between the day and night sides in the upper layers of the atmosphere which serve as the base of the wind, even if the temperature at the infrared photosphere is more uniform. The amplitude of this contrast is, however, uncertain.  The goal of this paper is to investigate the effect that anisotropic heating has on the winds from close-in EGPs using two-dimensional numerical hydrodynamic calculations.  Given the uncertainties in the mass loss rate produced by different treatments of the radiative heating and microphysics in the wind, we adopt the simplest possible approach, and focus on the {\\rm relative} difference between our multidimensional calculations and spherically symmetric models computed with identical techniques.  We find there are important differences in multidimensions. For example, the sonic surface in the wind is moved much closer to the planet, and the overall mass loss rate is reduced by almost a factor of four for a nearly isothermal wind.  Our numerical approach also allows us to consider the interaction between planetary and stellar winds. Although a stellar wind does not affect the mass loss rate, it confines the planetary wind to a small volume and greatly enhances the column density close to the planet.  These results suggest that incorporating multidimensional effects may be as important as improved treatments of the radiative transfer and microphysics in order to interpret the observations of transiting planets such as HD209458b.  In addition to a simplified treatment of the thermodynamics, we have made a number of other simplifying assumptions in this first investigation. For example, our calculations are two-dimensional (axisymmetric), which precludes us from studying the effect of Coriolis forces (orbital motion) on the outflow.  Therefore, our models are appropriate only for the flow close to the planet.  Fully three-dimensional hydrodynamic models of an {\\em isotropic} wind from orbiting planet have recently been presented by Schneiter et al. (2007).  Such calculations are challenging, even with an adaptive mesh these authors could only afford 3 grid points per planetary radius in the model.  In additions, our models are hydrodynamic rather than MHD, even though at large radii the wind may be significantly ionized, and interact with a primarily MHD stellar wind.  Finally, there are a number of kinetic plasma effects (e.g. charge-exchange reactions between the planetary and stellar wind particles) that might become important as the wind becomes very diffuse at large radii.  These effects are known to be important in the MHD of the heliopause (e.g. Borovikov et al. 2008).  Fully 3D MHD models of anisotropic winds from planets including multidimensional radiative transfer, microphysics, and the appropriate kinetic plasma effects are an interesting and important direction for future work.  The organization of this paper is as follows.  In the next section, we describe our numerical methods.  In \\S3, we present results for anisotropic winds from isolated EGPs.  In \\S4, we consider the interaction of anisotropic winds from an EGP with a stellar wind.  Finally, in \\S5 we discuss our results, and in \\S6 conclude. ",
        "conclusions": "We have presented two-dimensional, hydrodynamic calculations of the steady-state structure of thermally driven winds from highly irradiated, close-in EGPs, including the interaction with a high-velocity wind from the central star.  Our primary conclusions are that \\begin{itemize} \\item the mass loss rate in an anisotropic wind is reduced by about a factor of four in a nearly isothermal wind ($\\gamma=1.01$) compared to a spherical wind with the same parameters, \\item the sonic point in an anisotropic wind on the day-side is located 2-3 times closer to the planet surface compared to a spherically symmetric wind with the same parameters, \\item a supersonic polar flow from the day-side to the night-side is generated just above the planet surface (within $r/R_{p}\\approx 2$).  This flow generates weak shocks due to geometrical compression on the night side. At large radial distances, the outflow is nearly spherically symmetric. \\item the interaction with a stellar wind strongly compresses the planet wind and greatly enhances the column density of the gas in the outflow surrounding the planet.  However, the termination shock in the planetary wind is located well above the sonic surface, so the overall mass loss rate is affected very little. \\end{itemize}  There are a number of limitations to our results.  Probably the most important is the very simplistic treatment of the thermodynamics in the wind that we have adopted in our computations.  Rather than attempting to model the heating, cooling, and thermal conduction processes in the wind directly, we have simply fixed the temperature at the base of the wind consistent with previous more detailed models, and computed polytropic models with different $\\gamma$.  However, by comparing our models to spherically symmetric solutions computed using the same techniques and parameters, we are able to isolate the effects that multidimensional dynamics have on the wind.  We conclude these effects can be important, and therefore it is desirable to include a more realistic treatment of the radiative transfer and microphysics (e.g. Murray-Clay et al. 2008) in multidimensional models in the future.  Our conclusions are mostly applicable to winds in which the hydrodynamic escape parameter $\\lambda_{0}=5$, a value which is three times smaller than that inferred for observed systems such as HD209458b (assuming the temperature at the base of the wind is $T=10^{4}$~K, and the wind is mostly neutral).  A slightly higher temperature, or an ionized wind would produce a value closer to that adopted in most of our models. For larger values of $\\lambda_{0}$, we find the sonic point is located much farther from the planet, and polar flows produce a spherically symmetric, extended atmosphere above the planet which serves as the base of a nearly spherically symmetric wind.  However, the assumptions we have adopted for the thermodynamics in the wind are likely not applicable in this case.  Thus, the study if the multidimensional structure of winds in this parameter regime will require more sopisticated treatments of the radiative heating and cooling.  Additional limitations to our calculations are that we have not included the orbital motion of the planet, or the gravitational field of the central star.  We also have not included the effects of magnetic fields that are likely to be important in the dynamics of the stellar wind. Finally, non-thermal plasma processes (such as charge exchange reactions) might be important in the extremely low density wind material far from the planet, and since the terminal velocity of the planet winds studied here is less than the orbital velocity of the planet, these processes might even dominate the escape of gas from the system.  These limitations suggest that fully 3D MHD models of anisotropic winds from EGPs are warranted, and that ideally these models would span the entire orbital plane of the planet, include possibly important kinetic plasma effects, and would be based on an improved treatment of the thermodynamics in the wind."
    },
    "0812/0812.3664_arXiv.txt": {
        "abstract": "{} {Polycyclic Aromatic Hydrocarbons (PAHs) have been detected toward molecular clouds and some young stars with disks, but have not yet been associated with embedded young stars. We present a sensitive mid-infrared spectroscopic survey of PAH features toward a sample of low-mass embedded young stellar objects (YSOs). The aim is to put constraints on the PAH abundance in the embedded phase of star formation using radiative transfer modeling.} {VLT-ISAAC L-band spectra for 39 sources and Spitzer IRS spectra for 53 sources are presented. Line intensities are compared to recent surveys of Herbig Ae/Be and T Tauri stars. The radiative transfer codes RADMC and RADICAL are used to model the PAH emission from embedded YSOs consisting of a pre-main-sequence star with a circumstellar disk embedded in an envelope. The dependence of the PAH feature on PAH abundance, stellar radiation field, inclination and the extinction by the surrounding envelope is studied.} {The 3.3 \\micron\\ PAH feature is undetected for the majority of the sample (97\\%), with typical upper limits of 5 $\\times 10^{-16}$ W m$^{-2}$. One source originally classified as class I, IRS~48, shows a strong 3.3 \\micron\\ feature from a disk. Compact 11.2 \\micron\\ PAH emission is seen directly towards 1 out of the 53 Spitzer Short-High spectra, for a source that is borderline embedded. For all 12 sources with both VLT and Spitzer spectra, no PAH features are detected in either. In total, PAH features are detected toward at most 1 out of 63 (candidate) embedded protostars ($\\lesssim$ 2\\%), even lower than observed for class II T Tauri stars with disks (11--14\\%). Models predict the 7.7 \\micron\\ feature as the best tracer of PAH emission, while the 3.3 \\micron\\ feature is relatively weak. Assuming typical class I stellar and envelope parameters, the absence of PAHs emission is most likely explained by the absence of emitting carriers through a PAH abundance at least an order of magnitude lower than in molecular clouds but similar to that found in disks. Thus, most PAHs likely enter the protoplanetary disks frozen out in icy layers on dust grains and/or in coagulated form. } {} ",
        "introduction": "Polycyclic Aromatic Hydrocarbons (PAHs) have been observed toward a wide range of astrophysical environments \\citep{all89,pee02}, including the interstellar medium (ISM) and star-forming regions, first hinted at by the discovery of widespread broad emission features in the mid-infrared. Since PAHs are a good tracer of UV radiation, they are an indirect probe of star formation in the high opacity environments of molecular clouds as well as circumstellar disks. Recent spectroscopic studies have detected PAHs toward a significant fraction, 11--14\\% and 54\\% of low-mass and intermediate mass pre-main-sequence (PMS) stars respectively (\\citealt{ack04}; \\citealt[][chapter 2]{gee07c}), and recent ground-based high spatial resolution observations show evidence for the PAH emission to originate from the circumstellar disks \\citep[]{res03,hab06,gee07,gee07b}. PAHs are also prominently seen toward photon dominated regions (PDRs) in dense clouds exposed to massive young stars \\citep[e.g.,\\ ][]{ver96}. However, to this date, no significant PAH emission has been reported to be directly associated with the earlier class 0--I phase, when a (proto)star with circumstellar disk is still embedded in an envelope of gas and dust. Deeply embedded high-mass young stellar objects (YSO's) also lack prominent PAH emission \\citep{dis00}.  PAHs are believed to form in the outflows of carbon-rich Asymptotic Giant Branch (AGB) stars, which deposit them in the interstellar medium \\citep[e.g.,\\ ][]{hel96}. Their ubiquitous presence in the ISM shows that at least larger PAH molecules (of 50--100 carbon atoms or more) will survive this phase, until being incorporated in molecular clouds. The amount of carbon locked up in PAHs is [C/H]$_{\\mathrm{PAH}} \\simeq 5 \\times 10^{-5}$ \\citep{hab04}, corresponding PAH abundances of $\\sim$~$5 \\times 10^{-7}$ relative to H, for PAH molecules of 100 carbon atoms, making them among the most abundant molecules after H$_2$ and CO.  PAHs can play an important role in the star-forming environments. As large molecules, they provide efficient heating of the gas in both the ISM and circumstellar disks. Small dust particles like PAHs have also been suggested as the main formation site of molecules such as H$_2$ and water in the later evolutionary phases of circumstellar disks when classical grains have grown to large sizes \\citep{jon06}.  The shape and relative strength of the PAH features in the 6--9 \\micron\\ region has been shown to vary between various astrophysical environments from AGBs, the ISM to the class II disk sources. This has provided evidence that the emission characteristics of PAHs are sensitive to local physical conditions and  that interstellar PAHs undergo processing in space environments \\citep{pee02}. Further support for this comes from recent values for the PAH abundance in disk surface layers, which is a factor of  10-100 lower than in the ISM for sources with sufficient UV radiation to excite them \\citep{gee06}. PAHs frozen out on icy grains can chemically react with other molecules to form a large variety of complex species, including pre-biotic molecules \\citep{ber99,ehr06}. A study of PAHs toward embedded low-mass young stars is important to find out what happens to PAHs and how large a role they may play in these embedded objects.  In this paper we present the first mid-infrared spectroscopic survey for PAH emission from embedded low-mass young stellar objects using the Infrared Spectrometer And Array Camera (ISAAC) mounted on the ESO Very Large Telescope (VLT) and the Infrared Spectrograph (IRS) on board the NASA Spitzer Telescope, in a wide range of nearby star-forming regions. We subsequently use a radiative transfer code to model PAH emission from young stars with disks embedded in an envelope, and discuss the observed detection rate in the context of PAH abundance, the strength of UV emission from the protostar  and the extinction from the envelope. ",
        "conclusions": "No PAHs are detected toward any of the 53 confirmed embedded sources in neither our Spitzer nor ISAAC surveys. PAHs are detected at 3.3 \\micron\\ towards 1 of the 17 confirmed disk sources, IRS~48, and at 11.2 \\micron\\ for 1 of the 10 sources with an uncertain or borderline embedded classification, GY~23. For all 12 sources with ISAAC and Spitzer spectra, no PAH features are detected in either.  In total, PAH features are detected toward at most 1 out of 63 (candidate) embedded protostars (1.6\\%), even lower than observed for class II T Tauri stars with disks (11--14\\%).  The low PAH detection rate is compared with radiative transfer model calculations. It is shown that the effects of envelope mass on the absorption of PAH features is not enough to obscure PAH features. The 7.7 \\micron\\ feature is predicted by models as the best tracer of PAH emission, while the 3.3 \\micron\\ feature is relatively weak.  Assuming typical class I stellar and envelope parameters, the absence of PAH emission is most likely explained by the absence of emitting carriers through a much lower PAH abundance in the gas, e.g., due to freeze-out of PAHs on icy layers on dust grains or agglomeration. Our inferred low abundances are similar to those found for disks around T Tauri stars. Thus PAHs are expected to be removed from the gas already at earlier stages of star- and planet formation and enter the disks frozen out on grains. Further searches for PAH absorption features are needed to constrain the presence of PAHs in icy grains."
    },
    "0812/0812.4791_arXiv.txt": {
        "abstract": "Observations of the southern Cepheid $\\ell$~Car to yield the mean angular diameter and angular pulsation amplitude have been made with the Sydney University Stellar Interferometer (SUSI) at a wavelength of 696\\,nm.  The resulting mean limb-darkened angular diameter is 2.990$\\pm$0.017\\,mas (i.e. $\\pm$0.6\\,per cent) with a maximum-to-minimum amplitude of 0.560$\\pm$0.018\\,mas corresponding to 18.7$\\pm$0.6\\,per cent in the mean stellar diameter.  Careful attention has been paid to uncertainties, including those in measurements, in the adopted calibrator angular diameters, in the projected values of visibility squared at zero baseline, and to systematic effects.  No evidence was found for a circumstellar envelope at 696\\,nm.  The interferometric results have been combined with radial displacements of the stellar atmosphere derived from selected radial velocity data taken from the literature to determine the distance and mean diameter of $\\ell$~Car. The distance is determined to be 525$\\pm$26\\,pc and the mean radius 169$\\pm$8\\,$R_{\\sun}$.  Comparison with published values for the distance and mean radius show excellent agreement, particularly when a common scaling factor from observed radial velocity to pulsation velocity of the stellar atmosphere (the $p$-factor) is used. ",
        "introduction": "\\label{sec:intro}  The accurate determination of the zero point of the Cepheid Period-Luminosity relation is an important step in the refinement of the extragalactic distance scale.  The combination of interferometric and spectroscopic data, as discussed by \\citet{99jd}, is an approach to this task that is now feasible. The first demonstration of the combination of interferometric and spectroscopic data to determine both the mean diameter and distance of a Cepheid was made by \\citet{00lane} based on interferometric observations with the Palomar Testbed Interferometer (PTI) of $\\zeta$~Gem.  \\citet*{02lane} subsequently applied the technique to $\\eta$~Aql and, with additional data, to $\\zeta$~Gem, again based on PTI observations.  \\citet{04K1} have reported the measurement of the angular pulsations of seven Galactic Cepheids made with the European Southern Observatory's Very Large Telescope Interferometer (VLTI) and, in combination with spectroscopic data, have determined the distances to the Cepheids.  One of the prime programmes for which the Sydney University Stellar Interferometer (SUSI) \\citep{99susi1} was developed was the measurement of the mean diameters and angular pulsation amplitudes of Cepheids but, in its initial configuration, it lacked the required sensitivity. The recent commissioning of a red beam-combination system in SUSI \\citep{07susi} has resulted in a significant increase in sensitivity.  In this paper we report SUSI measurements of the mean angular diameter and angular pulsation amplitude of $\\ell$~Car, one of the seven Cepheids measured with the VLTI by \\citet{04K1}.  A preliminary analysis was presented at a European Southern Observatory symposium in 2005 \\citep{07eso} and \\citet{08apj} has discussed the SUSI Cepheid programme including an analysis of a subset of the $\\ell$~Car observations reported here.  Spectroscopic radial velocity measurements from the literature have been combined to produce a radial velocity versus pulsation phase curve and, after scaling with a fixed $p$-factor as discussed in Section~\\ref{sec:p-factor}, this has been integrated to give the radial displacement of the Cepheid surface as a function of pulsation phase.  The radial displacement and limb-darkened angular diameter data have been combined to determine the distance and mean radius of $\\ell$~Car.  The resulting distance and mean radius are compared with values in the literature determined by the same technique and by the infrared surface brightness method and excellent agreement is found, particularly when all values are scaled to a common $p$-factor. ",
        "conclusions": "\\label{sec:summary}  The measurement of the angular pulsations of Cepheid variables and the combination with spectroscopically determined radial displacements of the stellar atmospheres was one of the key programmes for which SUSI was developed. Only recently has adequate sensitivity been achieved for this programme to be undertaken and the first results are those for $\\ell$~Car presented here.  The angular pulsation of $\\ell$~Car has been measured with good cover in pulsation phase and careful attention has been paid to uncertainties in the measurements, in the adopted calibrator angular diameters, in the projected value of $V^{2}$ at zero baseline, and to systematic effects.  A phase-dependent limb-darkening factor, to convert uniform-disk angular diameters to limb-darkened angular diameters, was established based on Kurucz model atmospheres and its use, rather than a fixed mean value for the limb-darkening factor, was justified for 696\\,nm. The resulting mean limb-darkened angular diameter is 2.990$\\pm$0.017\\,mas (i.e. $\\pm$0.6\\,per cent) with a maximum-to-minimum amplitude of 0.560$\\pm$0.018\\,mas corresponding to 18.7$\\pm$0.6\\,per cent in the mean stellar diameter.  The projected value of $V^{2}$ at zero baseline shows no evidence at 696\\,nm of the circumstellar envelope observed in the $N$ and $K$ bands by \\citet{06ketal}.  A radial displacement curve has been computed from carefully selected radial velocity measurements from the literature and this has been combined with the interferometric data to determine the mean radius and distance to $\\ell$~Car.  The value for the distance is 525$\\pm$26\\,pc and the mean radius is 169$\\pm$8\\,$R_{\\sun}$.  These values have been compared with values in the literature and excellent agreement is found, particularly when all values are reduced to a common $p$-factor.  The uncertainty in the $p$-factor is a major contributor to the uncertainties in the distance and mean radius. The SUSI data may be useful in the future if $\\ell$~Car specific hydrodynamic spherical models are generated that consider the possibility of a phase-dependence of limb-darkening and the $p$-factor as discussed, for example, by \\citet{94sass} and \\citet{02marengo}.  To facilitate such possible applications the calibrated $V^{2}$ values may be obtained from the lead author (JD)."
    },
    "0812/0812.1727_arXiv.txt": {
        "abstract": "{} {We investigate the photometric signature of magnetic flux tubes in the solar photosphere.} {We developed two dimensional, static numerical models of isolated and clustered magnetic flux tubes. We investigated the emergent intensity profiles at different lines-of-sight for various spatial resolutions and opacity models. } {We found that both geometric and photometric properties of bright magnetic features are determined not only by the physical properties of the tube and its surroundings, but also by the particularities of the observations, including the line/continuum formation height, the spatial resolution and the image analyses techniques applied. We show that some observational results presented in the literature can be interpreted by considering bright magnetic features to be clusters of smaller elements, rather than a monolithic flux tube.} {} ",
        "introduction": "Recent observations taken in the G-band with sub-arcsecond  resolution have allowed the study of photometric \\citep{berger2004,okunev2004,ishikawa2007,berger2007} and dynamic  \\citep{bovelet2003,nisenson2003,rouppe2005,ishikawa2007} properties of small magnetic flux concentrations in the solar photosphere. Observations away from disk center have revealed that bright magnetic flux concentrations are usually accompanied by a dark lane on their disk center side and by a bright tail on the limb-ward side \\citep{lites2004, hirzberger2005, berger2007}. However, some bright points show no associated tail or dark lane, and some dark lanes with enhanced magnetic flux density show no associated bright point \\citep{berger2007}. Somewhat contradictory measurements have been reported for the center-to-limb variation (CLV) of G-band bright point photometric contrast and size. For example \\citet{berger2007} reported maximum contrast at a position much closer to the limb than previous observers, while \\citet{hirzberger2005}, after analysis of extreme limb observations, reported no trend in contrast or size with position on the solar disk. Discrepant CLVs have also been obtained when comparing G-band with other continua measurements \\citep[e.g.][]{auffret1991,sutterlin1999,sanchez2002,okunev2004}. Together these observations have raised questions about the validity of viewing bright magnetic features as monolithic flux elements. In particular, it has been suggested that the observed CLVs of contrast and size are determined by properties of unresolved clusters of magnetic elements at the limb \\citep[e.g.][]{okunev2005,berger2007}.  The properties of small individual magnetic features have been intensively investigated using both two and three dimensional magneto-hydrodynamic numerical simulations \\citep[e.g.][]{keller2004, carlsson2004,steiner2005,hasan2005,uiten2006}. Some of the highlights of these modeling efforts include the following. Using both static and dynamic two-dimensional simulations, \\citet{steiner2005} showed that the dark lane observed disk-center-ward of bright magnetic elements is caused by the presence of cooler material adjacent to (both inside and outside) the disk-center-ward flank of the magnetic tube. Similar results were previously presented by \\citet{pizzo1993a,pizzo1993b,deinzer1984, knolker1988a, knolker1988b}, who also investigated the sensitivity of the contrast intensity profiles to the physical properties of the tube (such as the temperature boundary condition, the evacuation of the tube, the intensity of magnetic field, and the inhibition of convection) and its temporal evolution. Ultimately, the intensity contrast profiles are determined by the degree of evacuation of the tube and the properties (particularly temperature) of the atmosphere inside and surrounding it. Fully three-dimensional magneto-hydrodynamic simulations \\citep[e.g.][]{keller2004, carlsson2004} have demonstrated that the increase in magnetic element contrast with distance from disk center largely reflects the properties of the solar granulation behind the tube.  This becomes visible away from disk center because of the reduced opacity within the tube for highly oblique viewing angles. Moreover, \\citet{depontieu2006}, using both magneto-hydro-dynamic simulations and high spatial resolution observations, demonstrated that the temporal evolution of the photometric properties of magnetic flux tubes occurs on the characteristic granular time scale.  The literature addressing the properties of flux tube aggregations is less extensive. \\citet{caccin1979} showed that the shape of the contrast-CLV of clusters of flux tubes is independent of the shape of the single underlying  model tube (funnels or cylinders). More recently, \\citet{okunev2005}, following \\citet{karpinsky1998},  used three-dimensional static models  in pressure equilibrium, to show that the appearance of aggregations of flux tubes and their measured polarimetric signals are strongly influenced by their position on the solar disk, filling factor and spatial resolution.  In this paper we investigate how the photometric signatures of isolated and clusters of small magnetic elements  (such as the  center-to-limb variation of photometric and geometric properties) reflect, not only the physical properties of the tube and its surroundings (such as temperature stratification and magnetic flux density), but also the particularities of the observation itself (such as spatial resolution, formation height of the wavelength of observation and image analyses technique employed). We examine these sensitivities using two-dimensional models of static magnetic flux tubes in radiative equilibrium with their surroundings. We show that results obtained by the analyses of recent high resolution observations can be interpreted assuming that observed magnetic bright features are actually aggregations of unresolved magnetic structures.  In \\S2 we describe the model. In \\S3 we investigate thermal properties of aggregations of magnetic flux tubes. In \\S4 we  present the obtained contrast profiles and discuss their dependence on observational wavelength and spatial resolution. In \\S5 and \\S6 we  show how the observed CLVs of photometric and geometric properties of aggregations of magnetic features depend not only on peculiarities of observations, but also on the numerical techniques employed to detect features on images. In \\S7 we investigate the size dependence of contrast. Conclusions are presented in \\S8. ",
        "conclusions": "We have investigated thermal, photometric and geometric  properties of various configurations of clusters of magnetic flux tubes by two dimensional, static numerical simulations. The main results, obtained by the analyses of tubes with Wilson depression 60 km when in RE, are the following.  At depths for which radiative heating is efficient, the temperature stratification within and around flux tubes is function not only of the physical properties of the flux tubes, but also of the tube number. In fact we found that the temperature inside the tubes increases with the number of adjacent tubes in the optically thin part of the domain and decreases at $0<{\\rm log}(\\tau)<1$. At larger values of optical depth temperature increases again with the filling factor. The area around and between the tubes is cooled at $0.5<\\tau<2$, while is heated at greater depths.  Since the temperature stratification determines photometric properties,  contrast profiles of flux tubes also vary in clusters. At disk center contrast increases with the number of flux tubes. Moreover regions in between the tubes are characterized by negative contrast, which is a signature of the decrease of temperature, generated by the radiative cooling, in these areas. These dark features gradually disappear with the increase of the inclination of the line of sight. We found that for each cluster model there always exist an inclination beyond which clusters appear as compact monolithic features, as if they were isolated tubes. At these inclinations, the contrast of a cluster is larger than the contrast of a single flux tube element, and smaller than the contrast of an isolated tube with the same size and Wilson depression of the cluster.  As for isolated flux tubes, off-disk center contrast profiles of clusters are very asymmetric, with long tails toward the limb and dark lanes at disk center side. Nevertheless, projection onto the plane of the sky and the decrease of the spatial resolution (due for instance to both instrumental and atmospheric degradation), largely makes profiles more symmetric. This finding is in agreement with measurements off-disk center of G-band bright points contrast profiles by \\citet{berger2007}. These authors found  quite symmetric profiles at all disk positions investigated. In both our models and observations, profiles at $\\mu=0.6$ are more symmetric than profiles at $\\mu=0.4$.  This is true in the models for both isolated and clusters of flux tubes (Fig. \\ref{CLV_asym}). Reduction of spatial resolution also decreases the depth and size of dark rings and dark lanes, which disappear in the worst resolution cases investigated (0.3\"). Moreover, dark lanes are not ubiquitous features of flux tubes. Since these are signatures of cooler regions within and around the tubes, they are very faint or absent in tubes models with larger radiative heating (see also models presented in \\citet{steiner2005} or \\citet{pizzo1993b}), or can be observed in some wavelengths and not in others. This helps explain why \\citet{berger2007} did not find dark lanes in all the contrast profiles they observed off-disk center.   CLVs of properties of magnetic bright features are also affected by clustering and reduction of resolution. The contrast-CLV is for instance steeper for isolated flux tubes than it is for clusters of tubes. Since most of feature identification techniques on images are based on intensity threshold, isolated small (and therefore less brilliant) features are less likely to be detected. Measurements at the limb are therefore biased toward larger tubes and clusters and the resulting CLV is flatter than the one obtained for single structures (see Fig. \\ref{CLV_contrast_configs}). For the same reasons, the CLV of sizes is also flatter for tube clusters (see Fig. \\ref{CLV_aree_configs}). This would explain CLV shapes obtained by recent observations at the extreme limb \\citep{hirzberger2005}. Reduction of resolution also shifts the peak of the contrast-CLV toward lower values of $\\mu$, especially for tubes models with marked dark lanes, and makes the CLVs flatter (see Fig. \\ref{CLV_risoluzione}). In this context, we notice that Berger et al. 2007, who analyzed better resolution images with respect to other authors, found the peak of the CLV of contrast to occur closer to the disk center with respect to some previous works (see Berger et al. 2007 and references therein).  Finally, we have investigated the size-contrast relation for clusters of various numbers and spatial resolutions. We found a slight increase of contrast with size at disk center, which has to be partially ascribed to the increase of temperature with the number of adjacent tubes and partially to the larger effects of smoothing on smaller features. Increase of contrast with size (and therefore with the number of flux tubes) is found for inclined lines of sight due to the reduced opacity inside the tubes.  Since for a given inclination of the line of sight only a finite number of tubes is crossed by radiation (see \\S4), there is a size threshold beyond which contrast does not vary. This threshold increases with the inclination of the line of sight. With the decrease of spatial resolution, even for vertical lines of sight the contrast increases with size, the threshold size value increases, and the contrast-size curves obtained at various lines of sight can be fitted by straight parallel lines. These findings are in agreement with recent contrast-size measurements obtained on both high \\citep{berger2007} and medium \\citep{ermolli2007} resolution images.   Results presented in this work suggest that measurements of properties of magnetic bright features are significantly influenced by spatial resolution, image analysis techniques, and the wavelength of observation. Geometric properties of contrast profiles, such as size and asymmetries, are particularly affected. These findings can partially explain the discrepancies presented in the literature. Moreover, some observational results can be better interpreted by considering observed bright magnetic features as aggregation of smaller elements, rather than a monolithic entity.  The results presented in this work should be properly taken into account when trying to derive physical properties of magnetic flux tubes by observations."
    },
    "0812/0812.4473_arXiv.txt": {
        "abstract": "We study the possibility that minor mergers resolve the loss cone depletion problem, which is the difficulty occured in the coalescence process of the supermassive black hole (SMBH) binary, by performing numerical simulations with a highly accurate $N$-body code. We show that the minor merger of a dwarf galaxy disturbs stellar orbits in the galactic central region of the host galaxy where the loss cone depletion is already caused by the SMBH binary. The disturbed stars are supplied into the loss cone. Stars of the dwarf galaxy are also supplied into the loss cone. The gravitational interactions between the SMBH binary and these stars become very effective. The gravitational interaction decreases the binding energy of the SMBH binary effectively. As a result, the shrink of the separation of the SMBH binary is accelerated. Our numerical results strongly suggest that the minor mergers are one of the important processes to reduce the coalescence time of the SMBH binary much less than the Hubble time. ",
        "introduction": "It is well known that the merging of galaxies with central supermassive black holes (SMBHs) is an important process for growth of the SMBH mass. Recent studies, however, have shown the possibility that the SMBHs cannot coalesce within the Hubble time in the merger remnant \\citep{beg80,mak04}.  \\citet{beg80} have examined the merging process of galaxies with central SMBHs with masses of $10^8 M_{\\odot}$. They estimate the timescale of coalescence of SMBHs in the merger remnant and show that the timescale is more than the Hubble time. They describe the reason for this as follows: SMBHs sink into the center in the merger remnant because of the dynamical friction from field stars. During this process, the field stars near the center are scattered by SMBHs and the number of the field stars decreases. The loss cone depletion occurs because of the scattering. In this case, dynamical friction force on the SMBHs becomes very weak. Then, it is difficult for the SMBH binary to shrink to a small distance enough to emit significant gravitational waves in the final coalescence stage of SMBHs.  \\citet{mak04} have studied the dynamical evolution of a SMBH binary, that each mass of the SMBH is $10^{8} M_{\\odot}$, in the stellar system by performing high-resolution $N$-body simulations. Their results have shown that the hardening timescale of the binary strongly depends on the relaxation time of the host galaxy as predicted by \\citet{beg80}. These results have confirmed the prediction that the SMBH binary cannot coalesce within the Hubble time by only gravitational interactions between SMBHs and field stars of the host galaxy since the relaxation time in a galactic stellar system is larger than the Hubble time. This difficulty in the coalescence process of SMBHs is called the \"loss cone depletion problem\".  To resolve the loss cone depletion problem, several ideas to accelerate the orbital decay of the binary are proposed. Gaseous torque in a massive gas disk is proposed in the wet merger cases \\citep{esc04,esc05,dot06,dot07,hay08}. In the dry merger cases which are observed in neaby galaxies \\citep{whi08}, the effect of a galactic triaxial potential \\citep{ber06}, large mass ratio between a SMBH and Intermediate-Mass BH \\citep{mat07}, and a triple SMBH system \\citep{iwa06} have been proposed. In the triple SMBH system, two SMBHs coalesce through the three body instability and the Kozai mechanism and the coalescence possibility is roughly $50$ \\% \\citep{iwa06}. \\citet{per07} and \\citet{per08} have proposed the role of massive perturbers. In their analytical studies, it has been shown the possibility that the massive perturbers of giant molecular clouds or molecular gas clumps accelerate the relaxation of stars in the galactic central region and, as a result, trigger the rapid coalescence of the SMBH binary. They pointed out the importance of three body interactions between the SMBH binary and stars. These ideas have possibility to lead to the coalescence of two SMBHs within the Hubble time, if some suitable conditions are realized.  In this paper, in order to resolve the loss cone depletion problem, we propose a new scenario that a minor merger triggers the rapid shrink of the SMBH binary. Similar idea was studied by \\citet{per07} and \\citet{per08} Our scenario is as follows: If a dwarf galaxy is compact enough, it can come close to the galactic center, and then, stellar orbits of the host galaxy are highly disturbed by the dwarf galaxy. In this case, many stars will be supplied into the loss cone. Moreover, if the dwarf galaxy is not destroyed before it closes enough to the central region, stars of the dwarf galaxy will be also supplied into the loss cone. In this way, gravitational interactions of the SMBH binary with these stars become effective and the hardening rate of the SMBH binary will become high. As a result, the binary can close to the unstable separation at which significant gravitational wave is emitted.  To demonstrate our scenario, we perform highly accurate $N$-body simulations. We show that a merging of a dwarf galaxy is an effective process to decrease the binding energy of a SMBH binary and trigger the rapid shrink of the binary, if the core of the dwarf galaxy is not destroyed by the tidal force and comes very close to the galactic central region. We present our simulation model and method in \\S \\ref{simulation}, numerical results in \\S \\ref{results}, and discussions in \\S \\ref{discussion}. ",
        "conclusions": "\\label{discussion}  \\subsection{Minor Mergers of the Compact Dwarf Galaxies} \\label{minor_merger}  We demonstrate that separation of a SMBH binary shrinks rapidly after a compact dwarf galaxy comes close to the central region of a host galaxy. It is important for our scenario that the dwarf galaxies are compact and are able to come close to the galactic central region without their destruction by the tidal force of the host galaxy. In this section, we discuss the possibility that such minor mergers occur.  Dwarf galaxies formed in the early universe are compact, since the mean density of dark matter halos is higher in the earlier universe, $\\rho _{\\rm{DM}}(z)\\propto (1+z)^{-3}$. Therefore, many compact dwarf galaxies are expected to form in the early universe and merge to their host galaxy. This is confirmed by the cosmological numerical simulations of the galaxy formation (e.g., \\citet{sai06}).  To trigger the rapid shrink of a SMBH binary, such compact dwarf galaxies need to come close to the central region of a host galaxy. To investigate the possibility, we calculate the motions of the dwarf galaxies with various initial orbital parameters which are in the range expected from the cosmological numerical simulations. Here, the fourth order Runge-Kutta method is used for the time integration. We assume that the dark halo and stellar potentials of the host galaxy are fixed. For the dynamical friction force, we use the formula given by \\citet{fuk92}. The initial position of the dwarf galaxy is set at $50$ kpc from the center of the host galaxy which is vicinity of the virial radius of the dark halo. The initial parameter of the nondimensional orbital angular momentum of the dwarf galaxy is from $\\lambda=0.01$ to $0.04$ which is the range of the spin parameter distribution of sub halos in a host dark halo \\citep{sha05}.  The time which is needed for the dwarf galaxies to move to the galactic central region (within $100$ pc) is shown in the left panel of Fig.~\\ref{lambda_dwarf}. The dwarf galaxy with $10^9 M_{\\odot}$ can move to the galactic center for the spin parameter $\\lambda =0.01$-$0.04$ within $2\\times 10^{9}$ yr. Such dwarf galaxies can come close to the galactic center within much less than $10^{10}$ yr. For the dwarf galaxies with $10^{8} M_{\\odot}$, they can come close to the center within $10^{10}$ yr in the case of the spin parameter of $\\lambda = 0.01$ and $\\lambda = 0.02$.  The right panel of Fig.~\\ref{lambda_dwarf} shows the specific angular momentum of the dwarf galaxy when it passes through $r=1$ from the galactic center of the host galaxy. In the case of $\\lambda = 0.01$, the specific angular momentum is about $0.37$. For larger $\\lambda$, it ranges from $0.5$ to $0.6$. The models from $Run~4$ to $Run~6$ and $Run~10$ correspond to the case of $\\lambda = 0.01$. The models of $Run~7$ and $Run~8$ correspond to the case of $\\lambda = 0.02-0.04$. Then, it is needed for our scenario that dwarf galaxies have mass more than $10^8 M_{\\odot}$ and smaller $\\lambda$ than $\\lambda=0.03$.   \\subsection{Conclusions}  We have performed $N$-body simulations and have shown that a minor merger is an effective process to resolve the loss cone depletion problem. If the core of the dwarf galaxy is not destroyed by the tidal force and comes close to the galactic central region, disturbed stars of the host galaxy are supplied into the loss cone. If the dwarf galaxy passes through the SMBH binary directly, stars of the dwarf galaxy are also supplied into the loss cone. After that, SMBHs can interact gravitationally with these stars and the binary loses its binding energy. In this process, three body interactions of the SMBH binary with these stars should be important \\citep{per07,per08}. As a result, the hardening rate of the binary becomes high and the semi-major axis can shrink rapidly in the host galaxy.  We have also performed high resolution simulations of $200000$ $N$-body particles with for the host galaxy. We have confirmed that the minor merger triggers the high hardening rate and the rapid shrink of the SMBH binary in the high resolution simulations. We find that the difference of time evolution of the hardening rate and the semi-major axis between with and without a minor merger becomes clear in higher resolution simulations. This is because timescale of the two body relaxation of $N$-body particles becomes longer in the $200000$ $N$-dody simulations. This result indicates that effects of the minor merger appear clearly in the realistic stellar system in which the two body relaxation time is larger than the Hubble time.  It is important for our scenario that the dwarf galaxy is enough compact to come close to the galactic central region. In the hierarchical galaxy formation scenario, such minor mergers are expected to occur frequently. Therefore, we emphasize that our scenario is one of the effective processes to trigger the rapid orbital decay of the SMBH binary and helps its coalescence within the Hubble time together with other previous proposed mechanisms.    \\bigskip The authors appreciate the anonymous referee for helpful advices. We thank Junichro Makino for helpful advices and suggestions. We also thank Keiichi Wada, Koji Tomisaka, Takayuki R. Saitoh, Masayuki Y. Fujimoto, and Kazuo Sorai for their helpful discussions, and Masaki Iwasawa and Tamon Suwa for their helpful advices on simulation method. This work has been supported by Grant-in-Aid for the 21st Century COE Scientific Research Programme on ``Topological Science and Technology'' from the Ministry of Education, Culture, Sport, Science, and Technology of Japan (MECSST), in part by Grant-in-Aid for Scientific Research (14340058) of Japan Society for the Promotion of Science, and in part by Hokkaido University Grant Program for New Fusion of Extensive Research Fields. Numerical computations were carried out on GRAPE system at Center for Computational Astrophysics, CfCA, of National Astronomical Observatory of Japan."
    },
    "0812/0812.3912_arXiv.txt": {
        "abstract": "The negative pressure accompanying gravitationally-induced particle creation can lead to a cold dark matter (CDM) dominated, accelerating Universe (Lima et al.~1996~\\cite{lima96}) without requiring the presence of dark energy or a cosmological constant.  In a recent study Lima et al.~2008 \\cite{lss} (LSS) demonstrated that particle creation driven cosmological models are capable of accounting for the SNIa observations \\cite{sn1a} of the recent transition from a decelerating to an accelerating Universe.  Here we test the evolution of such models at high redshift using the constraint on $z_{eq}$, the redshift of the epoch of matter -- radiation equality, provided by the WMAP constraints on the  early Integrated Sachs-Wolfe effect (ISW)~\\cite{komatsu}.  Since the contribution of baryons and radiation was ignored in the work of LSS, we include them in our study of this class of models.  The parameters of these more realistic models with continuous creation of CDM is tested and constrained at widely-separated epochs ($z \\approx z_{eq}$ and $z \\approx 0$) in the evolution of the Universe.  This comparison reveals a tension between the high redshift CMB constraint on $z_{eq}$ and that which follows from the low redshift SNIa data, challenging the viability of this class of models. ",
        "introduction": "A large amount of complementary cosmological data have established the surprising result that the Universe experienced a recent transition from a decelerating to an accelerating expansion~\\cite{accel}. The current data are consistent with a flat, Friedman-Lemaitre cosmology whose dominant components at present consist of matter (cold dark matter plus baryons) and a cosmological constant ($\\Lambda$) or, equivalently, the energy density of the vacuum; the $\\Lambda$CDM concordance model~\\cite{darkenergy}.  The matter and radiation dominate earlier in the evolution of the Universe, leading to its decelerating expansion, while the cosmological constant or vacuum energy component has come to dominate more recently, driving the accelerating expansion.  While the concordance model provides the simplest, most economical explanation which is consistent with all extant data, provided that it is fine-tuned to fit that data, it is by no means the only candidate proposed to explain the accelerating expansion of the current Universe.  Dozens (if not hundreds!) of alternate possibilities have been explored in the literature (see Lima et al.~2008 \\cite{lss} (LSS) for an extensive list of references) and many of them remain viable candidate models.  Since the astronomical community is planning a variety of large observational projects intended to test and constrain the standard $\\Lambda$CDM concordance model, as well as many of the proposed alternative models, it is timely and important to identify and explore a variety of physical mechanisms (or substances) which could also be responsible for the late-time acceleration of the Universe.  In principle, any homogeneously distributed ``exotic\" fluid endowed with a sufficiently negative pressure, $p < -\\rho/3$, so-called dark energy, will suffice. The simplest candidates for the dark energy are a positive cosmological constant $\\Lambda$, or a non-zero vacuum energy density.  In both cases the equation of state of the exotic fluid is $w \\equiv p/\\rho = -1$.  All current observational data are in good agreement with the cosmic concordance ($\\Lambda$CDM) model, consisting of a cosmological constant or vacuum energy ($w_{\\Lambda} = -1$) whose current density parameter is $\\Omega_{\\Lambda} \\approx 0.74$, plus cold dark matter and baryons ($w_{\\rm M} = 0$) with $\\Omega_{\\rm M} = \\Omega_{\\rm CDM} + \\Omega_{\\rm B} = 1 - \\Omega_{\\Lambda} \\approx 0.26$. Nevertheless, the $\\Lambda$CDM model encounters several challenges.  For example, the value of the energy density of the vacuum required by the data ($\\Omega_{\\Lambda} \\approx 0.26$) needs to be fine-tuned, exceeding the estimates from quantum field theories by some 50 -- 120 orders of magnitude.  If the vacuum energy (or $\\Lambda$) differs from zero, why does it have the value needed for consistency with the data?  LSS noted recently \\cite{lss} that the key ingredient required to accelerate the expansion of the Universe is a sufficiently negative pressure, which occurs naturally as a consequence of cosmological particle creation driven by the gravitational field \\cite{partcreat}.  While completely different from it, there are some analogies between this class of models and the continuously accelerating, never decelerating, expansion of the steady-state cosmology driven, in the Hoyle-Narlikar \\cite{hoyle} version, by C-field induced particle creation.  The late-time evolution of the class of models explored by LSS is qualitatively very different from that of the standard, $\\Lambda$CDM concordance model in that they lack a cosmological constant and, when baryons are included, the ratio of the dark matter to baryon densities evolves with time or redshift.  Nonetheless, the LSS \\cite{lss} analysis of the late-time evolution of such a CDM-dominated model with no cosmological constant or dark energy, lacking baryons and radiation, showed that it is quantitatively possible to account for the SNIa data which reveal a recent transition from decelerated to accelerated expansion.  Indeed, while the flat $\\Lambda$CDM concordance model provides a good fit to the SNIa data, the best fit actually corresponds to a positively curved (closed) $\\Lambda$CDM model, so that for some choice of parameters the models considered by LSS may provide an even better fit to the supernovae data.  For a subset of the models explored by LSS to be discussed here, the early-time, high redshift (\\eg $z~\\ga 10$) evolution is indistinguishable from that of the $\\Lambda$CDM concordance model, so that the parameters controlling the early evolution of these models must agree with those of the concordance model. To set the stage for the constraints and tests discussed here, we note that for the $\\Lambda$CDM concordance model with $\\Omega_{\\Lambda} = 1 - \\Omega_{\\rm M} = 0.76$, the redshift of the transition from decelerated to accelerated expansion is $z_{t} = 0.79$, the redshift of the epoch of equal matter and radiation densities (related to the early, Integrated Sachs-Wolfe (ISW) effect) is $z_{eq} = 3223$ and, the present age of the Universe in units of the Hubble age ($H_{0}^{-1}$) is $H_{0}t_{0} = 1.00$.  Since at present baryons are subdominant to CDM and the contribution of radiation is negligible, the neglect of baryons and radiation by LSS was not a bad approximation for their study of the recent evolution of the Universe.  Here, we revisit the LSS model of dark matter creation driven acceleration in a more realistic model incorporating baryons and radiation.  While it will be seen that there is a subset of models with baryons and radiation whose late time evolution (\\ie $H = H(z)$ for $z~\\la 5$) is sufficiently close to those studied by LSS  and to the $\\Lambda$CDM concordance model (and, therefore, they too will be consistent with the SNIa data), these models differ from those studied by LSS in their early evolution, in particular at the epochs of recombination and of equal matter and radiation densities.  To further test these models we impose an additional, high-redshift constraint derived from the early Integrated Sachs-Wolfe (ISW) effect, provided by the WMAP~\\cite{komatsu} value of the redshift of matter -- radiation equality, $z_{eq} = 3141 \\pm 157$.  In \\S~\\ref{partcreat} we describe the class of models under consideration and solve for the evolution of the Hubble parameter as a function of redshift ($H = H(z)$). In \\S~III we compare the predictions of this class of models as a function of the two parameters related to the particle creation rate to the SNIa data~\\cite{sn1a} and to the constraint on the redshift of the epoch of equal matter and radiation densities provided by the WMAP observations of the CMB temperature anisotropy spectrum~\\cite{komatsu}.  Our results are summarized and our conclusions presented in \\S~\\ref{concl}. ",
        "conclusions": "\\label{concl}  As an alternative to the standard, $\\Lambda$CDM model, we have explored here a class of models whose late-time acceleration is driven by the creation of cold dark matter. In constrast to the $\\Lambda$CDM model, the models considered here have no cosmological constant or vacuum energy.  While the dark energy models have two or more adjustable parameters (the dark energy equation of state, $w$, and its evolution with redshift, in addition to the choice of $\\Omega_{\\rm DE}$), the flat $\\Lambda$CDM model has only one free parameter, $\\Omega_{\\Lambda} = 1 - \\Omega_{\\rm CDM}$.  Here we have considered a class of two-parameter (\\{$\\beta,\\gamma$\\}) models. We found that the supernovae data prefer $\\beta = 0$ and that the early ISW effect is consistent with this choice.  For this subset of creation-driven models, with only one free parameter ($\\gamma$), the high-redshift evolution of the Universe is qualitatively indistinguishable from that of the $\\Lambda$CDM model, while the recent evolution is sufficiently similar to it to allow consistency with the SNIa data.  From the high redshift constraint on the WMAP determined value of $z_{eq}$~\\cite{komatsu} provided by the early ISW effect, we determined $\\gamma$(ISW) = 0.52.  For $\\beta = 0$ and this value for $\\gamma$, $H_{0}t_{0} = 0.97$, consistent with the estimate $H_{0}t_{0} = 1.00$ from the $\\Lambda$CDM model.  This choice of $\\gamma$ corresponds to the WMAP determined value of $z_{eq} = 3142$, consistent, within the uncertainties, with the $\\Lambda$CDM value of $z_{eq} = 3223$.  However, for $\\beta = 0$, this choice of $\\gamma$ provides a poor fit to the SNIa data~\\cite{sn1a}.  In contrast, the SNIa data prefer $\\gamma = 0.66$ which, for $\\beta = 0$, corresponds to $H_{0}t_{0} = 1.13$, in good agreement with the $\\Lambda$CDM value and with estimates of the age of the Universe.  But, for $\\beta = 0$ and $\\gamma = 0.66$, the redshift of equal matter and radiation densities is $z_{eq} = 1798$, in conflict (at $\\sim 2.6\\sigma$) with the WMAP determined value.  While these models, which are consistent with the SNIa data, offer an intriguing alternative to the standard, $\\Lambda$CDM concordance model, they face challenges. For example, although the one-parameter ($\\beta \\neq 0$, $\\gamma = 0$) model with baryons provides a natural solution to the observed, late-time acceleration of the Universe, without any need for fine-tuning, its early evolution is baryon-dominated and inconsistent with the early ISW effect as constrained by the CMB data.  In general, there is a clear tension in this class of models between the SNIa data and the independent, high redshift constraint from the observed early ISW effect. It should also be noted that in these models the ratio of dark matter (CDM plus baryons) to baryons increases during the recent evolution of the Universe from $\\rho_{\\rm M}/\\rho_{\\rm B} \\approx 6.0$ at high redshifts ($z~\\ga 10$) to $\\rho_{\\rm M}/\\rho_{\\rm B} \\approx 24$ at present ($z = 0$).  Correspondingly, the baryon fraction, $f_{\\rm B} \\equiv \\rho_{\\rm B}/\\rho_{\\rm M}$ decreases from $f_{\\rm B} \\approx 0.17$ at high redshifts to $f_{\\rm B} \\approx 0.04$ at present. While this latter value appears to be in conflict with the x-ray cluster baryon fraction~\\cite{clusters}, if clusters were formed at sufficiently high redshifts, their baryon fraction may be representative of the earlier, higher value of $f_{\\rm B}$.  Due to the recent increase in the density of CDM, the late-time growth of structure and of the cluster baryon fraction in these models will differ from that of the $\\Lambda$CDM concordance model.  Therefore, before ruling out these models, it might be worthwhile to explore their late-time evolution more carefully, especially with regard to predictions for the CMB and for the growth of large scale structure."
    },
    "0812/0812.3409_arXiv.txt": {
        "abstract": "{Submillimeter galaxies (SMGs) are distant, dusty galaxies undergoing star formation at prodigious rates.  Recently there has been major progress in understanding the nature of the bright SMGs (i.e. $S_{\\rm 850\\mu m} > 5$ mJy).  The samples for the fainter SMGs are small and are currently in a phase of being built up through identification studies. } {We study the molecular gas content in two SMGs, SMMJ163555 and SMMJ163541, at redshifts $z=1.034$ and $z=3.187$ with un-lensed submillimeter fluxes of 0.4 mJy and 6.0 mJy.  Both SMGs are gravitationally lensed by the foreground cluster Abell~2218. } {IRAM Plateau de Bure Interferometry observations at 3mm were obtained for the lines CO(2-1) for SMMJ163555 and CO(3-2) for SMMJ163541.  Additionally we obtained CO(4-3) for the candidate $z=4.048$ SMMJ163556 with an unlensed submm flux of 2.7\\,mJy. } {CO(2-1) was detected for SMMJ163555 at redshift 1.0313 with an integrated line intensity of $1.2\\pm0.2$ Jy km/s and a line width of $410\\pm120$ km/s. From this a gas mass of $1.6\\times10^{9}$\\,M$_\\odot$ is derived and a star formation efficiency of 440\\,L$_\\odot$/M$_\\odot$ is estimated.  CO(3-2) was detected for SMMJ163541 at redshift 3.1824, possibly with a second component at redshift 3.1883, with an integrated line intensity of $1.0\\pm0.1$ Jy km/s and a line width of $280\\pm50$ km/s.  From this a gas mass of $2.2\\times10^{10}$\\,M$_\\odot$ is derived and a star formation efficiency of 1000\\,L$_\\odot$/M$_\\odot$ is estimated.  For SMMJ163556 the CO(4-3) is undetected within the redshift range $4.035-4.082$ down to a sensitivity of 0.15\\,Jy\\,km/s.  } {Our CO line observations confirm the optical redshifts for SMMJ163555 and SMMJ163541.  The CO line luminosity $L'_{\\rm CO}$ for both galaxies is consistent with the $L_{\\rm FIR}-L'_{\\rm CO}$ relation. SMMJ163555 has the lowest far-infrared luminosity of all SMGs with a known redshift and is one of the few high redshift LIRGs whose properties can be estimated prior to ALMA.} ",
        "introduction": "CO observations of high redshift galaxies are instrumental to assess the quantity of molecular gas available for star formation.  The submillimeter galaxies \\citep[SMGs; see][for review]{blain02} are dusty galaxies undergoing star formation at prodigious rates, $> 1000$\\,M$_\\odot$yr$^{-1}$ with a redshift distribution peaking at $z\\sim2.5$ \\citep{chapman05}.  Previous CO results of SMGs have shown that SMGs can harbour molecular gas reservoirs on the order of $10^{10}$\\,M$_\\odot$, enabling the massive starburst to build-up stars during 10-100 Myrs.  It is thus likely that SMGs are the progenitors of the nearby massive ellipticals \\citep[e.g.][]{chapman05,greve05,lapi06,tacconi06,nesvadba07}.  Hitherto the CO results of SMGs have primarily been obtained for SMGs with intrinsic fluxes $S_{\\rm 850\\mu m} > 2$\\,mJy \\citep{frayer98,frayer99,neri03,genzel03,downes03,tacconi06,tacconi08}.  Of the five known SMGs with lensing corrected $S_{\\rm 850\\mu m}< 1$\\,mJy \\citep{kneib04a,borys04,knudsen06,knudsen08a}, only the triply-imaged galaxy SMMJ16359+6612 has previously been detected in CO \\citep{sheth04,kneib05,weiss05}. Through statistical studies it is clear that other high redshift galaxies populations, such distant red galaxies, extremely red objects and BzK galaxies, have an average 850$\\mu$m flux $\\leq 1$\\,mJy \\citep[e.g.][]{webb04,knudsen05,daddi05}, and thus there will be an overlap between these and the faint SMGs.  However, none of these have individual detections at these low flux density levels due to the SCUBA 850$\\mu$m confusion limit making it difficult to establish the exact link between the faint SMGs and other galaxy populations.  In this paper we report the IRAM Plateau de Bure Interferometer (PdBI) observations of two strongly lensed SMGs, both present in the galaxy cluster field Abell~2218 \\citep{knudsen06}. SMM\\,J163555.2+661150 (henceforth SMMJ163555) has an optical counterpart at redshift $z=1.034$, which is also known as galaxy \\#289 in the notation of \\citet{pello92}.  It has an unlensed 850$\\mu$m flux of 0.4 mJy and is one of the lowest redshift SMGs \\citep[cf.\\ SMM\\,J02396-0134;][]{smail02}, making it the SMG with the lowest far-infrared luminosity known so far. SMM\\,J163541.2+661144 (henceforth SMMJ163541) has been identified with a galaxy at redshift $z=3.187$; optical spectroscopy is also presented in this paper.  With its unlensed 850$\\mu$m flux of 6.0\\,mJy it belongs to the bright SMGs as found typically in the blank field surveys. SMMJ163555.5+661300 (henceforth SMMJ163556) has an unlensed 850$\\mu$m flux of 2.7\\,mJy and has a candidate optical counterpart at redshift $z=4.048$ \\citep{knudsen08b} placing it possibly among the highest redshift SMGs known.  Throughout the paper we assume a cosmology with $H_0 = 70$\\,km/s/Mpc, $\\Omega_m = 0.3$ and $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "\\label{sect:conc}  We have presented IRAM 3mm spectroscopy for three SMGs in the field of A2218. CO detections were obtained for the sources SMMJ163555 at $z=1.0313$ (CO(2-1)) and SMMJ163541 $z=3.1824$ (CO(3-2)) confirming the optical redshifts.  Both lines are shifted by $-400$ km/s relative to the optical redshift.  The $z\\approx 4$ candidate counterpart SMMJ163556 is undetected in CO(4-3) in the redshift range $4.035-4.082$. Future observations at high angular resolution and also other J-transitions are required for a more detailed study of the SMGs.  Our results add to the yet limited number of successful CO detections of SMGs.  When correcting for gravitational lensing SMMJ163555 is so far the SMG known with the lowest far-infrared luminosity and for the first time allows us to probe the $L_{\\rm FIR}-L'_{\\rm CO}$ relation in the LIRG regime for high redshift galaxies.  Based on this one detection we speculate that the high redshift LIRGs have higher star formation efficiency than the nearby LIRGs. This can, however, only be clarified with a large sample using future sensitive instruments such as ALMA."
    },
    "0812/0812.2125_arXiv.txt": {
        "abstract": "Using turbulent MHD simulations (magnetic Reynolds numbers up to $\\approx8000$) and {{\\sl Hinode}} observations, we study effects of turbulence on measuring the solar magnetic field outside active regions.  Firstly, from synthetic Stokes~$V$ profiles for the {{\\tt FeI} }lines at 6301 and 6302 \\AA, we show that a peaked probability distribution function (PDF) for observationally-derived field estimates is consistent with a monotonic PDF for actual vertical field strengths.  Hence, the prevalence of weak fields is greater than would be naively inferred from observations.  Secondly, we employ the fractal self-similar geometry of the turbulent solar magnetic field to derive two estimates (numerical and observational) of the true mean vertical unsigned flux density.  We also find observational evidence that the scales of magnetic structuring in the photosphere extend at least down to \\AD{an order of magnitude smaller than $200\\,$km:} the self-similar power-law scaling in the signed measure from a {{\\sl Hinode}} magnetogram ranges (over two decades in length scales and including the granulation scale) down to the $\\approx200\\,$km resolution limit.  From the self-similar scaling, we determine a lower bound for the true quiet-Sun mean vertical unsigned flux density of $\\sim50\\,$G.  This is consistent with our numerically-based estimates that 80\\% or more of the vertical unsigned flux should be invisible to Stokes$-V$ observations at a resolution of $200\\,$km owing to the cancellation of signal from opposite magnetic polarities.  Our estimates significantly reduce the order-of-magnitude discrepancy between Zeeman- and Hanle-based estimates. ",
        "introduction": "Determining the strength of the magnetization of the ``quiet'' Sun is tied to the question of how much flux resides at small scales.  This is important, for example, in determining the energy budget available for chromospheric heating.  Observationally, two methods are employed to constrain the solar magnetic field: the Hanle and Zeeman effects.  The Hanle effect measures (in principle) the mean magnetic field strength, $\\left<|B|\\right>$, as there are no cancellation effects, but quantitative interpretation requires assumptions about the probability distribution function (\\pdf) of the turbulent magnetic field.  The Hanle de-polarization is measured in strong\\resp{er} lines formed in the \\resp{mid- to upper-}photosphere and estimates are made of $\\left<|B|\\right> \\sim 130\\,$G, the field residing primarily in the intergranular lanes \\citep{TrBuShAsRa2004}.  The Zeeman effect measures the longitudinal component (via Stokes~$V$) and transverse components (via Stokes~$Q$ and $U$) of the magnetic field but suffers from cancellation effects.  Hence, it is ``blind'' to any ``hidden'' mixed-polarity flux at scales smaller than the resolution limit of an instrument.  However, the benefit of Zeeman measurements is that their interpretation requires no assumption about the turbulent \\pdfa and measurements can actually be used to determine the \\pdfa on scales larger than the resolution limit.  Some attempts to incorporate the effects of cancellation into Zeeman-based Stokes inversions utilize the micro-structured field hypothesis \\citep{SaAlLadeMaPi+1996}. However, typical estimates based on the longitudinal Zeeman effect give $\\left<|B_z|\\right>\\sim10\\,$G for the mean unsigned vertical flux density (see Table 3 in \\citealt{BeGoYeChOk+2008} for a review of the spread of recent Zeeman-based estimates).  These values are significantly smaller than the Hanle-based estimates. The discrepancy between the results from Hanle ($\\sim100\\,$G) and Zeeman ($\\sim10\\,$G) measurements is not unexpected since the Zeeman observations see only the resolved flux while the Hanle interpretation depends on assumptions about an unknown \\pdf.  In this work, we ask the question ``How can the turbulent fractal geometry of the magnetic field in the solar photosphere be accounted for in interpretations of Zeeman-based observations?''  We attempt to derive the spectral/fractal properties from high resolution ($0\\arcsec.3$) Zeeman observations with the \\hinodea spectro-polarimeter (SP) \\citep{KoMaSa+2007,TsIcKa+2008,LiElSt2001} and high resolution (down to $4\\,$km$\\,\\approx 0\\arcsec.006$), turbulent dynamo simulations (up to magnetic Reynolds numbers, $Re_M \\approx 8000$) with the \\murama code \\citep{V03,VSS+05}.  First, we clarify the consistency of \\pdfsa derived from observations and simulations, respectively.  Zeeman observations typically show \\pdfsa of quantities derived from Stokes~$V$ which can be described as a peaked function \\citep{KhMaGoCo+2005,LiKuSoNa+2008}.  We demonstrate that the peaked \\pdfsa from Stokes~$V$ measurements and the monotonic \\pdfsa typically reported from numerical simulations are, in fact, compatible.  Failure to take into account this observational bias can lead to a gross underestimation of the prevalence of weak fields and, consequently, to incorrect estimates of the mean magnetic field strength and magnetic energy density in the solar photosphere.  Second, we extrapolate the results obtained with observations and simulations (based on their respective fractal properties) to scales below the resolution limits to estimate the amount of ``hidden'' flux. To do this, we exploit the fact that the fractal geometry is intimately connected to power-law scaling relations such as $N(l)\\propto~l^{-D_f}$.  Here, $N(l)$ is the number of boxes of edge length $l$ covering a fractal set (such as all pixels with apparent magnetic flux above some threshold) and $D_f$ is the fractal dimension.  It has long been known that the distribution of plage magnetic field is such a statistically self-similar fractal \\citep{ScZwBa+1992,BaScZw+1993}.  As the threshold can influence the fractal dimension inferred, however, the geometrical structure is more complicated than a simple fractal.  In this case, the fractal concept must be generalized by adding a measure defined by the absolute value of the net magnetic flux through each box of edge length $l$. This measure also displays self-similarity (a power-law scaling) for the solar quiet Sun network magnetic flux (down to $0\\arcsec.5$ resolution in \\citealt{LaRuCa1993} and \\citealt{CaLaRu+1994}; see also \\citealt{KrSo2004}) and also for numerical simulations of magnetoconvection \\citep{BrPrSe+1992}. Hence, the geometry of the magnetic field is said to be multifractal.  Cancellation effects do not play an important role in observations of the multifractal unipolar solar regions discussed above.  In contrast, if the magnetic field of the quiet Sun internetwork is multifractal (as we will show), cancellation can play a significant role for scales below $1\\arcsec$ \\AD{(see also \\citealt{SaAlEmCa2003}).}  For these mixed-polarity fields it is necessary to further generalize the fractal concept to a signed measure.  Here, the power-law scaling exponent is the cancellation exponent \\citep{ODS+92}.  If this self-similar power-law scaling extends below the observational resolution, small scale cancellation will occur and correct values of the mean field strength and energy cannot be established.  As was pointed out by \\citet{LaRuCa1993}, such a signed measure could also be employed to extrapolate moments of the magnetic field below resolvable limits. Until now, no attempt has been made to employ self-similar scaling to estimate the total cancellation, and, hence, the true mean quiet-Sun magnetic field strength. ",
        "conclusions": "On the basis of surface dynamo simulations, we have demonstrated that the \\pdfa generated from the Stokes~$V$ spectra are not necessarily equivalent in form to that of the \\pdfa of the vertical component of the underlying magnetic field.  The \\pdfa for Stokes~$V$ shows a reduction of likelihood for weak vertical magnetic field compared to the \\pdfa of the field itself.  This effect is not due to a reduction in horizontal resolution, but is caused by a combination of vertical radiative transfer through a turbulent fluid (via the Doppler effect) and noise.  % That is, any systematic sampling (by geometrical height, optical depth, or volume) of $B_z$ from the simulation yields a monotonic \\pdf, \\AD{but} due to Doppler shifts between different atmospheric heights, the Stokes $V$ signal is not such a systematic sampling.  Consequently, the \\pdfa of Stokes $V$ field estimates do not accurately represent the \\pdfa of the actual vertical magnetic field \\AD{even in the absence of noise.  Additionally, for two different levels of noise (``normal mode'' and ``deep mode'') we have demonstrated that the peak in the observational \\pdfa is dominated by the influence of noise. Because of these two effects,} a monotonic \\pdfa for the field can result in a peaked \\pdfa in observations and the assumption that \\pdf($B_z$) can be uniquely derived from Stokes~$V$ observations becomes dubious.   From the cancellation function for a \\hinodea observation of the apparent longitudinal flux density, we have demonstrated that the multi-fractal self-similar pattern of the quiet-Sun photospheric magnetic field covers two decades of length scales down to the resolution limit, $200\\,$km.  This constitutes observational evidence that the the smallest scale of magnetic structuring in the photosphere is \\AD{at least an order of magnitude smaller than $200\\,$km.}  The power law also allows us to constrain the quiet-Sun true mean unsigned vertical flux density.  We estimate the lower bound to be $\\approx46\\,$G.  Estimates based solely on our numerical simulations suggest that the vertical unsigned flux at \\hinode's resolution should be multiplied by 5 to obtain the true vertical unsigned flux (i.e., $\\sim50\\,$G). These two results are consistent and suggest that the order of magnitude disparity between Hanle and Zeeman-based estimates may be fully resolved by a proper consideration of the cancellation properties of the full vector field.   {\\bf Acknowledgments}  The authors would like to acknowledge fruitful discussions with S. Solanki, A. Pietarila, and R. Cameron.  \\hinodea is a Japanese mission developed and launched by ISAS/JAXA, collaborating with NAOJ as a domestic partner, NASA and STFC (UK) as international partners. Scientific operation of the Hinode mission is conducted by the Hinode science team organized at ISAS/JAXA. This team mainly consists of scientists from institutes in the partner countries. Support for the post-launch operation is provided by JAXA and NAOJ (Japan), STFC (U.K.), NASA, ESA, and NSC (Norway).  {\\bf Note added in proof}  \\resp{We would like to point out the correlations between our conclusions and the works of \\citet{SaAlLadeMaPi+1996} and \\citet{SaAlLi2000} who postulate structuring of the magnetic and velocity fields on scales much smaller than $100\\,$km. They find that synthetic profiles generated by 3-component Milne-Eddington atmospheres re-produce the observed Stokes-$V$ asymmetries found in $1\\arcsec$ resolution observations.  Though they did not estimate the undetected photospheric magnetic flux, the results indicated a significant fraction remaining undetected.  \\citet{DoCeSaAlKn2006} assumed the quiet-Sun \\pdfa can be approximated by a linear combination of the \\pdfa inferred from Zeeman observations and a log-normal distribution (accounting for the observed Hanle depolarization).  They determined that the Hanle and Zeeman signals are consistent with a single \\pdfa with $\\langle|B|\\rangle\\ga100\\,$G (see also \\citealt{SaAlEmCa2003}).  \\citet{SaAl2006} assuming that numerical simulations of magnetoconvection with no dynamo action ($20\\,$km horizontal resolution) had achieved the asymptotic rate of magnetic energy dissipation, derived an estimate for the unsigned magnetic flux contained in unresolved scales; in our notation their finding is $\\chi(100\\,$km$)\\sim0.5$ while our extrapolation estimates $\\sim0.36$. Finally, \\citet{SaAl2008} using observational data from various sources plot $\\langle|B_z|\\rangle_l$ versus $l$ (their Fig. 1).  The data are compared to a line, the slope of which corresponds to $\\kappa=1$, i.e., the result for white noise (\\citealt{VSP+94}, also set $D_f=0$ in Eq. (\\ref{eq:vsp94})).  There is, however, a large scatter about this line suggestive of either a large uncertainty in $\\kappa$ or an element of randomness in the calibration issues between the various data used.}"
    },
    "0812/0812.2063_arXiv.txt": {
        "abstract": "We present a detailed examination of the features of the Active Region (AR) NOAA 10798. This AR generated coronal mass ejections (CMEs) that caused a large geomagnetic storm on 24 August 2005 with the minimum Dst index of $-216$~nT. We examined the evolution of the AR and the features on/near the solar surface and in the interplanetary space. The AR emerged in the middle of a small coronal hole, and formed a {\\it sea anemone} like configuration. H$\\alpha$ filaments were formed in the AR, which have southward axial field. Three M-class flares were generated, and the first two that occurred on 22 August 2005 were followed by Halo-type CMEs. The speeds of the CMEs were fast, and recorded about 1200 and 2400~km~s$^{-1}$, respectively. The second CME was especially fast, and caught up and interacted with the first (slower) CME during their travelings toward Earth. These acted synergically to generate an interplanetary disturbance with strong southward magnetic field of about $-50$~nT, which was followed by the large geomagnetic storm. {\\it Accepted for publication in JGR. Copyright (2008) American Geophysical Union. Further reproduction or electronic distribution is not permitted.} ",
        "introduction": "Space weather has attracted a lot of attention in recent times. Space weather research involves various related fields, such as the solar surface, solar wind, interplanetary space, geomagnetosphere, ionosphere, and atmosphere, since a comprehensive understandings from active phenomena on the solar surface to the propagation of the disturbances toward Earth are crucially required for the studies.  Vast plasma ejected from the solar corona in the form of coronal mass ejections (CMEs) leads to a geomagnetic storm, and therefore, CMEs have been actively discussed. Moreover, large geomagnetic storms are associated with large flares (that are emitting strong X-ray), long duration events (LDEs), fast CMEs, and so on (e.g., Gopalswamy, Yashiro, \\& Akiyama 2007). Flare locations are another factor for a major geomagnetic storm, since the flare location close to the disk center indicates that the related CME is heading towards Earth and is likely to cause a large geomagnetic storm (Manoharan et al. 2004). Ejections that cause strong disturbances with southward magnetic field in the interplanetary space are also important. Coronal holes (CHs) are, on the other hand, related with fast solar wind because of their open magnetic field, and therefore, themselves have been another important factor for space weather studies. While large geomagnetic storms are caused by earth-directed CMEs (see e.g., Gosling et al. 1990), weaker storms are associated with high-speed streams from CHs (see e.g., Sheeley, Harvey, \\& Feldman 1976). However, storms related to high-speed streams from CHs cause larger flux enhancement of MeV electrons of the Earth's Van Allen belt than the CME-associated storms do on average (Kataoka \\& Miyoshi 2006). It has been, furthermore, reported that many fast Halo CMEs are associated with CHs (Verma 1998; Liu \\& Hayashi 2006). The recent work done by Liu (2007) showed that the speeds are faster even statistically than those of CMEs under the heliospheric current sheet. Therefore, in order to understand what kind of events can generate a large geomagnetic storm, it is necessary to study active phenomena on the solar surface and the propagation in the interplanetary space, in relation to the surrounding magnetic structure.  In this paper we examine in detail the evolution of an Active Region (AR) that emerged in a CH and the related flares, CMEs, and the geomagnetic storm to elucidate how such a magnetic configuration works to generate a geo-effective flares/CMEs. CMEs originating from AR NOAA 10798, generated a large geomagnetic storm on 24 August 2005, which was one of the 88 major geomagnetic storms reported by Zhang et al. (2007). This AR has been highly paid attention to, since it was one of the targets of the International CAWSES\\footnote{Climate And Weather of the Sun-Earth System} Campaign\\footnote{see, http://www.bu.edu/cawses/secondcampaign.html}, and at the related virtual conference\\footnote{http://workshops.jhuapl.edu/s1/index.html}, there were intensive discussions on the AR (e.g., Asai et al. 2006). Figure~\\ref{fig1} shows an overview of the geomagnetic storm and the related solar-terrestrial events. The Soft X-ray (SXR) flux in the {\\it GOES} 1.0 - 8.0~{\\AA} channel shows three M-class flares that occurred on 22 and 23 August 2005. The first two flares (marked with $\\times$) were associated with the CMEs responsible for the geomagnetic storm in question. In the second panel we can recognize the sufficient enhancements of the proton fluxes in $>$10~MeV (black line) and $>$50~MeV (gray line) channels obtained with {\\it GOES}. Both flares were followed by enhancements of solar energetic particles (SEPs), and the second flare's was larger. The bulk velocity of solar wind $V_{\\rm sw}$ in the third panel and the total $|B|$ (black line) and Z-component of the magnetic field $B_{z}$ (gray line) in the fourth panel were measured with the {\\it Advanced Composition Explorer} ({\\it ACE}). The first shock was recorded at 05:35~UT by {\\it ACE} as shown by the dashed line. The same shock was also recorded by the {\\it Geotail} satellite at 06:15~UT. The interplanetary magnetic field had a strong southward component of about $-50$~nT. The bottom panel shows the Dst index produced by the Kyoto University. The decrease of the Dst index was quite large, reaching $-216$~nT. In \\S 2 we describe the evolution of the AR, focusing on the photospheric magnetic configuration, and the H$\\alpha$ filament formed during the evolution of the AR and the coronal features are also presented. In \\S 3 we discuss the flares/CMEs that occurred on 22 August 2005. In \\S 4 we shortly review the associated interplanetary disturbances. In \\S 5 we summarize our results. ",
        "conclusions": "In order to make clear the importance of an AR that emerged in a CH to generate geo-effective flares/CMEs, we examined the evolution of the AR NOAA 10798, the solar events associated with a geomagnetic storm that occurred on 24 August 2005, and the related interplanetary disturbances. The summary of the features of the AR and the events is as follows: (1) Highly twisted and complex magnetic flux emerged within a small CH on 18 August 2005, which was named NOAA 10798, (2) An anemone type structure was generated, and H$\\alpha$ filaments that had southward axial fields were formed on 21 August 2005, (3) Two halo CMEs associated with M-class flares occurred on 22 August 2005. (4) The CME speeds were fast, especially the second one recorded 2400 km~s$^{-1}$, (5) The interplanetary disturbances with strong southward magnetic field of about $-50$~nT and strong compression of plasma were produced.  The CMEs were particularly geo-effective, and the minimum Dst index was $-216$ nT. The reasons for the CMEs to be so geo-effective were the high speeds of the two CMEs and their interaction as well as the CMEs traveled directly toward the earth. For the current case, the speed of CME2 was faster and pushed the slower CME1, which led to a unusual strong compression of the plasma at the front of CME2.  The high speeds of the CMEs are more notable. The AR was large and very complex, and violated the Hale-Nicholson's magnetic polarity law. These reverted polarity ARs are statistically favorable to produce large flares. However, it is suspicious whether just the violation of the Hale-Nicholson's law is responsible for high speed CMEs of about 2000~km s$^{-1}$, and it should be quantitatively and statistically clarified in the future. In this paper we suggest that the fast CMEs are probably a consequence of the eruption inside a CH from an anemone AR. This is consistent with the association between fast Halo CMEs and CHs as reported before (Verma 1998; Liu \\& Hayashi 2006; Liu 2007).  Eruptive activities of anemone ARs are usually low (Asai et al. 2008), and often confined to small-scale activities inside CHs that appears to be SXR bright points. In some cases, anemone ARs can produce large SXR coronal jets (Shibata et al. 1994b, Vourlidas et al. 1996, Kundu et al. 1999, Alexander \\& Fletcher 1999). This is because the situation of an emergence of a magnetic flux within a CH is suitable for magnetic reconnection with the surrounding field to generate SXR coronal jets and/or H$\\alpha$ surges (Yokoyama \\& Shibata 1995, 1996). Wang (1998) indicates the possibility that even polar plumes are associated with jets from anemone ARs at high latitudes. Anemone ARs are related to non-radial coronal streamers emanating from magnetically high latitudes (Saito et al. 2000). The relation between anemone ARs and fast solar winds have also been paid attentions to (Takahashi et al. 1994, Saito et al. 1994, Wang 1998). Saito et al. (2000) further discussed the rotational reversing model and the triple dipole model to explain the reversal of the solar surface magnetic field, and anemone ARs play an important role in this. This model implies that anemone ARs are more conspicuous in the decaying phase of a solar cycle as in the case of AR NOAA 10798.  On the other hand, the deflection of CMEs eastward by the interplanetary fields effectively worked in the current case as shown in the Figure~\\ref{fig6}d. As Wang et al. (2004) pointed, the faster CMEs are deflected more eastward, and therefore, the AR NOAA 10798 generated geo-effective CMEs, although it was quite close to the southwest limb. The azimuthal angle of the magnetic field measured from the x-axis $\\phi_{B}$ ($= \\arctan(B_{y}/B_{x})$) changed 90$^{\\circ}$ -- 180$^{\\circ}$ -- 270$^{\\circ}$ ($-$90$^{\\circ}$) during the passage of the FR, which is consistent with the guess that the deflection of the CMEs were so strong that the axis of the FR passed through the east of the earth. This is also consistent with the fact that the flares in the next rotation (and renamed as NOAA 10808) did not affect the magnetosphere so much (Wang et al. 2006, Nagashima et al. 2007). Furthermore, the extremely short duration of FR and the missing of CME1 (see Fig.~\\ref{fig9}) are possibly explained by the skimming encounter with the CMEs due to the strong deflection.  The nature of the interplanetary disturbances and their impact on the magnetosphere strongly depend on the features of emergence and evolution of an AR and the relation with the surrounding magnetic field. In this work we succeeded to follow in detail the evolution of the AR and the large geomagnetic storm resulting from eruptions in the AR. The reconstruction of the proposed scenario using numerical simulations will be attempted in the future."
    },
    "0812/0812.0585_arXiv.txt": {
        "abstract": "If the cosmological dark matter has a component made of small primordial black holes, they may have a significant impact on the physics of the first stars and on the subsequent formation of massive black holes. Primordial black holes would be adiabatically contracted into these stars and then would sink to the stellar center by dynamical friction, creating a larger black hole which may quickly swallow the whole star. If these primordial black holes are heavier than $\\sim 10^{22}\\;{\\rm g}$, the first stars would likely live only for a very short time and would not contribute much to the reionization of the universe. They would instead become $10 - 10^3\\; M_\\odot$ black holes which (depending on subsequent accretion) could serve as seeds for the  super--massive black holes seen at high redshifts as well as those inside galaxies today. ",
        "introduction": "The first stars in the Universe mark the end of the cosmic dark ages, reionize the Universe, and provide the enriched gas required for later stellar generations.  They may also be important as precursors to black holes (BHs) that coalesce and power bright early quasars. The first stars are thought to form inside Dark Matter (DM) halos of mass $ 10^5 \\, M_\\odot$--$ 10^6 \\, M_\\odot$ at redshifts $z=10-50$ ({Abel} {et~al.} (2002); {Bromm} {et~al.} (2002); {Yoshida} {et~al.} (2003)). These halos consist of 85\\% DM and 15\\% baryons in the form of metal-free gas made of H and He.  Theoretical calculations indicate that the baryonic matter cools and collapses via H$_2$ cooling ({Peebles} \\& {Dicke} (1968); {Matsuda} {et~al.} (1971); {Hollenbach} \\& {McKee} (1979); {Tegmark} {et~al.} (1997)) into a single small protostar ({Omukai} \\& {Nishi} (1998)) at the center of the halo (for reviews see {Ripamonti} \\& {Abel} (2005), {Barkana} \\& {Loeb} (2001) and {Bromm} \\& {Larson} (2004)).   It is interesting to study the effects on the evolution of the first stars due to the large reservoir of DM within which these stars form. The first protostars and stars are particularly good sites for this investigation because they form inside the highest density environment (compared to today's stars): they form at high redshifts (density scales as $(1+z)^3$) and in the high density centers of DM haloes. Previously, two of us ({Spolyar} {et~al.}(2008)) studied the effects of Weakly Interacting Massive Particles on the first stars: we found that the annihilation of these particles could provide a heat source for the star that stops the collapse of the protostar as well as potentially dominates over any fusion luminosity for a long time.   In this paper, we consider instead the effects on these first stars of a different candidate for the DM: Primordial Black Holes (PBHs). These are small black holes that may be formed in the very early universe (see the next section for more detail) and may exist in sufficient abundance to provide the DM seen in the Universe today.  The masses of PBHs that explain the entirety of the DM range from    ($10^{17}-10^{26}$) g; heavier PBHs up to $1 \\, M_\\odot$ may provide still an interesting fraction of the DM.  We discuss the implications that PBH DM would have on the physics of the first stars, the so called Population III stars. These stars could range from $\\sim$ (1 - few 100) $M_\\odot$. First, we compare various possible heat sources due to PBHs with the ordinary heat from stellar fusion of the stars.  For the properties of the Pop III stars, we use results computed by Heger \\& Woosley.  Specifically, for a 100~$M_\\odot$  (10~$M_\\odot$) Pop. III star, we take the central temperature to be $1.2 \\times 10^8$~K ($9.6 \\times 10^7$~K), the central density  31~g/cm$^3$ ($226$~g/cm$^3$), the radius 7~$R_\\odot$ (1.2~$R_\\odot$), and the stellar fusion luminosity to be \\be \\label{eq:lstar} L_* &=& 6.5 \\times 10^{39} \\, {\\rm erg/s} \\,\\,\\,\\, (100 \\, M_\\odot) \\, , \\\\ L_* &=& 4.2 \\times 10^{37} \\, {\\rm erg/s} \\,\\,\\,\\, (10 \\, M_\\odot)  \\, . \\label{eq:lstar10} \\ee We find that the ordinary stellar fusion luminosity dominates over the heat sources due to PBHs, which include accretion onto the BHs, Hawking radiation, and the Schwinger mechanism.  Instead, we find the interesting result that PBHs inside the first stars may sink to the center and form a single BH, which may accrete very rapidly and swallow the whole star.  The phenomenon is relevant for PBHs heavier than about $10^{22}$~g, because the corresponding timescale for dynamical friction turns out to be shorter than the typical stellar lifetime, while it is less interesting or completely negligible for lighter BHs. So, for $M_{PBH} \\gtrsim 10^{22}$~g, the lifetimes of Pop. III stars may be  shortened, with implications for reionization of the Universe as well as for the first supernovae. In addition, since the stars are inside much larger haloes, they can in principle accrete even more matter (depending on the accretion mechanism).  Thus,  the end--products of the scenario are BHs of masses $10-10^5 \\, M_\\odot$.   These may be the seeds which produced the super--massive BHs seen at high redshifts; the Intermediate Mass Black Holes; as well as the black holes at the center of every normal galaxy today and whose origin is as yet uncertain. Possible mechanisms of production of superheavy BHs are reviewed in {Dokuchaev} {et~al.} (2007). In addition, although the PBH swallowing the star shortens the star's lifetime and its contribution to reionization, the newly formed hole can become a new, alternative source of ionizing photons.   The rest of the paper is organized as follows. In Section II, we briefly review the physics of PBHs, that is, how they can be formed in the early Universe and what current constraints on their cosmological abundance are. In Section III, we discuss the behavior of individual PBHs: how many of them are expected inside a single star (via adiabatic contraction), what is the luminosity due to accretion onto the PBHs, and what is the timescale for their size to double.  We also investigate alternative mechanisms for generating luminosity by these small PBHs.  Then in Section IV and V we turn to the most important part of the paper. We study the dynamical friction that pulls all the BHs into a single larger BH at the center of the star, and then watch this single large BH accrete the entire star surrounding it on a fairly rapid timescale. We conclude with a discussion in Section VI. Throughout the paper, we use units with $c = 1$. ",
        "conclusions": "Primordial black holes in the mass range $M_{PBH} \\sim 10^{17} - 10^{26}$~g are viable dark matter candidates. They may be produced in the early Universe by many mechanisms and so far there are no constraints on their possible abundance. Assuming that they make part of the cosmological dark matter, we expect that due to dynamical friction primordial black holes will make up a small but significant mass fraction of the first stars.  Primordial black holes with masses smaller than about $10^{22}$~g do not have a significant effect on the evolution of primordial stars, because their timescales for Bondi accretion and for dynamical friction are larger than the lifetime of a Main Sequence star of $10 - 100 \\; M_\\odot$. On the contrary, primordial black holes heavier than $10^{22}$~g might sink quickly to the center of the star by dynamical friction and form a larger black hole, which could swallow the whole star in a short time. So, Pop. III stars would likely have lived for a short time, with implications for the reionization of the Universe after the cosmic dark ages and  the nature of the first supernovae; in fact they may preclude any supernovae from the first stars.  Although the BH swallowing the star shortens the star's lifetime and its contribution to reionization, the newly formed hole can become a new, alternative source of ionizing photons. The 10--100~$M_\\odot$ BHs that form by swallowing the Pop. III stars may grow even larger: they reside in 1000~$M_\\odot$ of gas that are in excess of the Jeans mass and may fall into the BH.  Black holes of mass 1-1000 $M_\\odot$ may result.  Depending on the accretion mechanism at this point, the black hole may accrete  more matter and grow larger. The $10^6 \\; M_\\odot$ minihaloes of dark matter contain $\\sim 10^5 \\; M_\\odot$ of baryonic matter. This accretion from the minihalo, as well as from other haloes merging with the one containing the black hole,  would be from low density gaseous material ($\\rho \\sim 10^{-24}$~g/cm$^3$), which is considerably different from the accretion we considered earlier from  within the star ($\\rho \\sim 1$~g/cm$^3$).  In the case of accretion from the low density gas outside the star, feedback may become important.  As we have shown, the timescale for the Pop. III stars to become black hole can be much shorter than the lifetime of the Pop. III stars (3~Myr for a 100~$M_\\odot$ star), so that the feedback due to stellar heating and ionization of the medium surrounding the black hole may be minimal. However, the accretion may well be in a disk, with the accompanying radiation pressure as well as radiative feedback due to the accretion ({Silk} \\& {Rees} (1998); {Springel} {et~al.} (2005); {Ciotto} \\& {Ostriker} (2001); {Li} {et~al.} (2007); {Pelupessy} {et~al.} (2007); {Alvarez} {et~al.} (2008)) limiting the accretion rate.    A recent study ({Alvarez} {et~al.} (2008)) of the radiative feedback from the black hole accretion has been done for the case of $\\eta = 0.1$ and a Pop. III star that has undergone its full lifespan, and finds reduced accretion onto the BH; the story may be different here.   We have not studied these later stages. Since the end--products are $10-10^5\\; M_\\odot$ black holes, these objects may serve as seeds for Intermediate Mass Black Holes; the super--massive black holes which have been seen already at high redshifts ({Haiman} \\& {Loeb} (2001); {Volonteri} \\& {Rees} (2006)) and may be the progenitors of the super--massive black holes which are  in the center of every normal galaxy today.   Even if the primordial black holes do not explain the entirety of the dark matter in the Universe, they may still play a role in the first stars. Heavier primordial black holes than the ones studied here, i.e., primordial black holes with $M_{BH} >  10^{26}$~g,  are observationally constrained to be only a fraction of the total dark matter in the Universe, and yet could be important in the first stars. It would only take one such black hole to be pulled into the star via dynamical friction (timescale $\\sim 10^7$~yr for a 1~$M_\\odot$ black hole to get from 1~pc out into the center of the star (see Eq.~(\\ref{tau-df})) and to quickly eat up the whole star. In fact, a single massive primordial black hole would already have a major effect during the protostellar phase  while the molecular cloud is collapsing down into a protostar: the molecular cloud would already collapse into a black hole. In this case the fusion phase of a Pop. III star would be completely avoided. Another possibility would be to have the dark matter consist primarily of Weakly Interacting Massive Particles (WIMPs) but with a small component of primordial black holes. In that case there would be dark stars powered by WIMP annihilation ({Spolyar} {et~al.} (2008)), which would become black holes once the primordial black holes sink to the center of the dark star.  In principle, if the effects described in this paper do not take place, one could place bounds on the black hole abundances of various masses. For example, if primordial black holes swallowed primordial stars too quickly, the cosmological metal enrichment would be problematic and in absence of viable alternatives, the current allowed mass range $M_{PBH} \\sim 10^{17} - 10^{26}$~g could be further reduced to $\\sim 10^{17} - 10^{22}$~g."
    },
    "0812/0812.4029.txt": {
        "abstract": "{We derive quantum kinetic equations for scalar fields undergoing coherent evolution either in time (coherent particle production) or in space (quantum reflection). Our central finding is that in systems with certain space-time symmetries, quantum coherence manifests itself in the form of new spectral solutions for the dynamical 2-point correlation function. This spectral structure leads to a consistent approximation for dynamical equations that describe coherent evolution in presence of decohering collisions. We illustrate the method by solving the bosonic Klein problem and the bound states for the nonrelativistic square well potential. We then compare our spectral phase space definition of particle number to other definitions in the nonequilibrium field theory. Finally we will explicitly compute the effects of interactions to coherent particle production in the case of an unstable field coupled to an oscillating background.} ",
        "introduction": "% Quantum transport effects are gaining more and more interest in many applications in modern particle physics and cosmology. This is true in particular for the case of the electroweak baryogenesis~\\cite{EWBGrew,CJK,Earlier,KPSW,PSW}, but also for the leptogenesis~\\cite{buchmuller} or particle creation in the early universe~\\cite{Prokopec_partnumber} and during phase transitions~\\cite{brandenbeger}. We have recently developed new quantum transport equations for fermionic systems including nonlocal coherence, either in space (quantum reflection) or in time (coherent pair production)~\\cite{HKR1,HKR2}, including the effects of decohering collisions~\\cite{HKR2}. Here we will introduce a similar formalism for the scalar fields. As in the fermionic case, the coherence information is found to be encoded in new spectral solutions in the phase space of the dynamical 2-point function. The physical information of particle numbers or fluxes and of coherence is carried by a set of scalar functions that parametrize the spectral shells in the full 2-point correlation function.  Our approach can be summarized as follows. We first formulate Schwinger-Dyson equations for the 2-point correlation functions using the Closed Time Path (CTP) method. We solve the resulting Kadanoff-Baym (KB) equations for the 2-point Wightman function $\\prop^<$ in the noninteracting mean field limit in the mixed representation. The most general solution for $\\prop^<$ in this limit is a sum of singular spectral distribution functions corresponding to the usual mass-shell solutions with a dispersion relation $\\omega^2 = \\vec k^2 + m^2$, and a new coherence solutions living at shell $k_z= 0$ in the static planar symmetric case or at shell $k_0=0$ in the case of a spatially homogeneous system. We then turn back to the full KB-equations including collision terms and integrate them with an appropriate set of moments. On the adiabatic boundaries the lowest moments of $\\prop^<$ can be related to the spectral on-shell functions. From practical point of view the most important aspect of our method is that the singular shell structure reduces integrated equations of motion, including the collision terms, to a closed set of equations for spectral on-shell functions, or equivalently for a finite number of lowest moments of $\\prop^<$ (three in case of a single real scalar field).  Since the singular shell representation of the coherence is so crucial for our formalism, we have given several examples which illustrate their physical role. Firstly, we will solve the bosonic Klein problem. We will show that in the absence of the coherence shell the quantum nature of reflection is completely lost, but that the correct reflection and tunneling factors are recovered when the coherence shell is included. We will also show that the spectral phase space definition of particle number in our formalism is consistent with other definitions for nonequilibrium systems in the literature~\\cite{Berges}. In particular, our particle number, when applied to Bunch-Davies vacuum in the inflation, corresponds to the adiabatic particle number that remains always zero in a conformally coupled scalar theory~\\cite{BirrelAndDavis}. The Klein problem example also allows us to demonstrate how the on-shell functions are related to moment functions that must be used to formulate a dynamical problem with only an incomplete information about the variables defining the system. We will also consider production of unstable scalar particles by a coherent time dependent background potential and decoherence and thermalization of an initially highly correlated state.  This paper is organized as follows. In section \\ref{sec:formalism} we briefly review the basic CTP-formulation for the calculation of the 2-point function and in section~\\ref{sec:shell} we derive the spectral shell structure of the Wightman function in the mean field limit.  In section \\ref{sec:physical} we use our formalism to solve the bosonic Klein problem. We also derive an expression for our on-shell particle number in terms of moment functions and compare it with other definitions in the literature and apply it to the particle production during inflation. We also compute other measurable quantities such as energy density and the pressure. In section \\ref{sec:non-relativistic} we solve the nonrelativistic problem with a Schr\\\"odinger equation and show that there one obtains similar coherence solutions for the description of reflection in the planar symmetric case. We complete this section by solving the bound states of the square well potential with our formalism. In section \\ref{sec:spectral} we show that, just as with fermions, the coherence solutions are excluded from the spectral function by the spectral sum rule. In section \\ref{sec:interactions} we derive the dynamical (moment) equations for a scalar field including collisions for a spatially homogeneous system. In section \\ref{sec:numexample} we consider coherent production of unstable particles. Finally section \\ref{sec:discussion} contains our discussion and outlook. % ",
        "conclusions": "\\label{sec:discussion} In this paper we have derived quantum transport equations including nonlocal coherence effects for relativistic and nonrelativistic scalar fields in systems with particular spacetime symmetries. The time-dependent but homogeneous systems include for example particle production during phase transitions or during the inflation in the early universe. The static, planar symmetric problems include for example reflection off a potential, such as Klein problem, or off a mass wall induced by a phase transition front in the early universe.  The key observation leading to a calculable approximation scheme was the observation that in both of these geometries the 2-point correlator, written in the mixed representation, has new spectral solutions living on shell $k_0 = 0$ in the homogeneous case and on shell $k_z = 0$ in the static planar symmetric case. These solutions were interpreted to carry information on the nonlocal coherence between particle and antiparticle excitations in the former and between left- and right moving states in the latter case. We demonstrated the physical nature of these new spectral solutions by applying the formalism to exactly solvable models, such as relativistic Klein problem and bound states in one dimensional nonrelativistic potential well.  The nontrivial singularity structures described above were found as an approximation to collisionless equations to the lowest order in gradients. The core of our calculational scheme was the argument that these structures should provide a reasonable {\\em ansatz} for the 2-point function when relating the moments to physical on-shell functions at the boundary of the system and when computing collision terms in the moment expansion of the Kadanoff-Baym equation.  When the ansatz is introduced into the collision terms they become completely parametrized by the on-shell functions in the ansatz. Because the on-shell functions can be uniquely related to the lowest moment functions, the resulting moment equations  can be solved. Despite the simplicity of the resulting equations, they contain the essential information about the evolution of nonlocal quantum coherence under decohering interactions. Based on the nature of the approximation we argue that our method could be useful even when the background field is not necessarily slowly varying. Method requires the existence of an adiabatic boundary, or boundaries where the on-shell functions can be related to the moment functions. It is thus best suited to problems with localized disturbances in background field configurations with asymptotic adiabatic boundaries.  Our method provides a natural definition for the adiabatic particle number $n_{\\vec k}$ related to the value of the phase-space functions multiplying the singular mass-shells. This definition was applied to the particle number evolution during inflation where it was shown to correspond to the adiabatic particle number defined \\eg in ref.~\\cite{BirrelAndDavis}. Moreover, our particle number coincided with the definition of ref.~\\cite{Berges} for slowly varying fields. We also computed the particle number evolution in the presence of a driving background interaction in the form of a time-dependent mass term. This situation could model for example the particle creation during phase transitions or at the end of the inflation.  We then included decoherence assuming that the scalar particles created by the background fields were unstable. We found that the effect of interactions on particle number divided to two parts: first the existing particle number was suppressed by decays as expected and secondly interactions provided a friction term on the growth of the coherence. However, as the growth of $n_{\\vec k}$ was found to be completely controlled by coherence, the friction term turned out to be most efficient in reducing the particle number created by unit time.  Many of the results presented here were qualitatively similar to those derived earlier by us~\\cite{HKR1,HKR2} for fermions. However, details of the derivation were substantially different so as to warrant a complete independent treatment in this paper. It would be interesting to apply our formalism to study for example the effect of collisions or decays on the production of scalar fields in a realistic model for a parametric resonance. It should also be straightforward to extend our nonrelativistic formulation for example for 3D-cubic lattice potentials and study atoms in such lattices under the influence of external thermal noise."
    },
    "0812/0812.1530_arXiv.txt": {
        "abstract": "SIM-Lite is an astrometric interferometer being designed for sub-microarcsecond astrometry, with a wide range of applications from searches for Earth-analogs to determining the distribution of dark matter. SIM-Lite measurements can be limited by random and systematic errors, as well as astrophysical noise. In this paper we focus on instrument systematic errors and report  results from SIM-Lite's interferometer testbed. We find that, for narrow-angle astrometry such as used for planet finding, the end-of-mission noise floor for SIM-Lite is below 0.035\\ \\uas. ",
        "introduction": "The newly redesigned light version of the Space Interferometry Mission, SIM-Lite, is a long-baseline astrometric interferometer designed to reach sub-microarcsecond precision over the course of a five year mission. It has two modes of operation, one for global astrometry with 4~\\uas\\ end-of-mission accuracy, and the other for narrow-angle astrometry, with a single-epoch (1100s visit) accuracy below 1~\\uas\\ \\citep{Unwin08}. The redesign lowers costs primarily by shortening the baseline from $9~{\\rm m}$ to $6~{\\rm m}$ and replacing a guide interferometer with a telescope star tracker.  Narrow-angle astrometry is used to measure the orbital motions of non-luminous objects, like neutron stars and planets, by observing the motion of the parent star.  SIM-Lite is designed to make narrow angle measurements by alternately switching between a target star and anywhere from 3 to 6 reference stars within a 1 degree radius. The astrometric signal is the motion of the target star relative to the reference stars.  To take a concrete example, consider a candidate star that is observed 250 times over the course of a 5-year mission.  A peak in the joint periodogram power distribution (similar to a power spectral density) of the target star position constitutes the detection of a planet candidate. Detection of a planet with a 1\\% false-alarm probability requires a signal-to-noise ratio (SNR) of approximately 5.8. At 10 pc, the astrometric signature for an Earth orbiting a one Solar mass star at 1 AU is 0.3~\\uas, so the SNR requirement calls for a final instrument error of less than 50~nano-arcseconds (nas). This corresponds to a single-epoch error of less than 0.82~\\uas. For the nearest $\\sim60$ stars, a noise floor below 35 nas is needed to detect Earths in the habitable zone.  The key questions, therefore, are whether sub-microarcsecond single-epoch accuracy is attainable, and, equally important, whether the instrument systematic errors do average down well below the single-epoch accuracy. This paper reports testbed results relating directly to both the single-epoch and end-of-mission narrow angle accuracy.   The basic elements of a Michelson stellar interferometer are shown in Figure~\\ref{fig1}. Starlight is collected by two spatially separated siderostats. Each siderostat has an embedded fiducial that marks one end of the baseline and also acts as a retroreflector for the internal and external metrology laser beams. Along each arm, the starlight beam has an annular footprint while internal metrology occupies the center.  The quantity of interest is the delay $x$, which is the optical path difference (OPD) between two starlight beams originating at a star and terminating at the two fiducials. The delay is given by: \\begin{equation} x=\\vec{b}\\cdot \\hat{s} + C + \\eta \\label{eq-basicAstrom} \\end{equation} where $\\vec{b}$ is the baseline vector, $\\hat{s}$ is the unit vector to the star, and $\\eta$ represents measurement noise. $C$ is sometimes called the interferometer {\\em constant term} and represents the internal OPD when the metrology reading is zero. The baseline length $b$ is defined as the distance between the two fiducials and is monitored by external metrology. Thus, three measurements, the white light fringe delay, internal metrology and external metrology, form the basic ingredients of the astrometric angle.  The SIM-Lite instrument is described in detail elsewhere \\citep{SIM08, SIMEng08}. The fundamental SIM-Lite instrument is its ``science\" interferometer, operating in the visible range with $50~{\\rm cm}$ collectors and a $6~{\\rm m}$ baseline. The baseline vector is not stable at the microarcsec level, so guide interferometers observing bright (typically 7mag) guide stars are used to measure the motion of the baseline over time. A laser optical truss ties everything together at the microarcsec level. A detailed mathematical analysis of the SIM-Lite astrometric approach appears in  \\citet*{Mil03}. ",
        "conclusions": "For planet finding, the dominant error experienced by an astrometric interferometer is the thermally induced drift of the metrology with respect to the starlight. The testbed results show that, even in the presence of a harsher thermal environment than SIM-Lite will experience in space, chopping between the target and reference stars renders this error nearly ``white'' with a noise floor below 24 nas after 42 hours of averaging. This is below the 35 nas required to detect an Earth in the habitable zone of the nearest $\\sim60$ stars. While investigation of the next-level noise sources is continuing, these results suggest that the technology for astrometric detection of nearby Earths is at hand."
    },
    "0812/0812.2019_arXiv.txt": {
        "abstract": "{We discuss the model of magnetic field reconnection in the presence of turbulence introduced by us ten years ago. The model does not require any plasma effects to be involved in order to make the reconnection fast. In fact, it shows that the degree of magnetic field stochasticity controls the reconnection. The turbulence in the model is assumed to be subAlfvenic, with the magnetic field only slightly perturbed. This ensures that the reconnection happens in generic astrophysical environments and the model does not appeal to any unphysical concepts, similar to the turbulent magnetic diffusivity concept, which is employed in the kinematic magnetic dynamo. The interest to that model has recently increased due to successful numerical testings of the model predictions. In view of this, we discuss implications of the model, including the first-order Fermi acceleration of cosmic rays, that the model naturally entails, bursts of reconnection, that can be associated with Solar flares, as well as, removal of magnetic flux during star-formation.}    \\addkeyword{galaxies: magnetic fields- MHD} \\addkeyword{Sun: energetic particles} \\addkeyword{Sun: flares} \\addkeyword{Sun: magnetic fields}   \\begin{document} ",
        "introduction": " ",
        "conclusions": "The advantage of the model in Lazarian \\& Vishaniac (1999) is that, first of all, the reconnection is robust and is fast in any type of fluid, provided that the fluid is turbulent enough. The latter requirement is natural for most of astrophysical fluids (see Armstrong et al. 1995). At the same time, the model predicts that if the fluids are not turbulent initially, they should be prone to bursts of reconnection, which may provide an appealing explanation of Solar Flares.  \\subsection{Self-consistency of the model}  The high speed of reconnection given by equation (\\ref{main}) naturally leads to a question of self-consistency.  Is it reasonable to take the turbulent cascade suggested in GS95 when field lines in adjacent eddies are capable of reconnecting? It turns out that in this context, our estimate for $V_{rec,global}$ is just fast enough to be interesting.  We note that when considering the intersection of nearly parallel field lines in adjacent eddies the acceleration of plasma from the reconnection layer due to the pressure gradient is not $k_{\\|}V_A^2$, but rather $(k_{\\|}^3/k_{\\perp}^2)V_A^2$, since only the energy of the component of the magnetic field which is not shared is available to drive the outflow.  On the other hand, the characteristic length contraction of a given field line due to reconnection between adjacent eddies is only $k_{\\|}/k_{\\perp}^2$.  This gives an effective ejection rate of $k_{\\|}V_A$.  Since the width of the diffusion layer over a length $k_{\\|}^{-1}$ is just $k_{\\perp}^{-1}$, we can replace equation (\\ref{main}) with $V_{rec,global}\\approx V_A {k_{\\|}\\over k_{\\perp}}$. The associated reconnection rate is just \\begin{equation} \\tau^{-1}_{reconnect}\\sim V_A k_{\\|}, \\end{equation} which in GS95 is just the nonlinear cascade rate on the scale $k_{\\|}^{-1}$.  However, this result is general and does not involve assuming that GS95 is correct. As we discuss below, most of the energy liberated in reconnection goes into motions on length scales comparable to the dimensions of the reconnecting eddies, so this energy release will not short circuit the energy cascade described in GS95.  On the other hand, we can invert this argument to see that reconnection can play an important role in preventing the buildup of unresolved knots in the magnetic field. Such structures could play a major role in inhibiting the cascade of energy to smaller scales, flattening the energy spectrum relative to the predictions of GS95.  Our conclusion is that such structures will disappear as fast as they appear, supporting the notion that they play a limited role in the dynamics of MHD turbulence.  Finally, we note that if the magnetic field structure is driven by turbulence in another location, as when the footpoints of magnetic arcades are stirred by turbulent motions, then we can evaluate its effects in terms of the amplitude of field stochasticity and the scaling of structure anisotropy with scale. The actual turbulence may be balanced or imbalanced, have or not have ``polarization intermittency'' (see Beresnyak \\& Lazarian 2006, 2008) the robust nature of the reconnection implies that the reconnection will be sensitive to the amplitude of the induced field stochasticity, but not the details of the turbulent mixing process.  \\subsection{Turbulent diffusivity and dynamo}  To enable sustainable dynamo action and, for example, generate a galactic magnetic field, it is necessary to reconnect and rearrange magnetic flux on a scale similar to a galactic disc thickness within roughly a galactic turnover time ($\\sim 10^8$~years). This implies that reconnection must occur at a substantial fraction of the Alfv\\'en velocity. The preceding arguments indicate that such reconnection velocities should be attainable if we allow for a realistic magnetic field structure, one that includes both random and regular fields. This does solve one part of the problem of dynamo. The other part is related to magnetic helicity conservation (see Vishniac et al. 2003).  Interestingly enough, high turbulent diffusivity is also required for advecting heat in astrophysical plasmas, e.g. in clusters of galaxies (see Lazarian 2006). Turbulent reconnection enables eddies to transfer heat in the way similar to the advection of heat by turbulence in unmagnetized fluid (see Cho et al. 2003).  \\subsection{Dissipation of energy}  The usual assumption for energy dissipation in reconnection is that some large fraction of the energy given up by the magnetic field, in this case $\\sim \\rho V_A^2 L_x^3$, goes into heating the electrons.  This is not the case here.  Only a fraction, $\\sim 1/(k_{\\|}L_x)$ of any flux element is annihilated by Ohmic heating within the reconnection zone.  Over the entire course of the reconnection event the efficiency for electron heating is no greater than \\begin{equation} \\epsilon_e\\lesssim {\\eta/\\Delta\\over V_{rec,global}}={V_{rec,local}\\over V_{rec,global}}= {1\\over k_{\\|}L_x}. \\end{equation} or, with GS95 scaling substituted (see LV99) \\begin{equation} \\epsilon_e\\lesssim \\left({V_Al\\over\\eta}\\right)^{-2/5} \\left({v_l\\over v_A}\\right)^{8/5}\\left({l\\over L_x}\\right)^{4/5}. \\end{equation} The electron heating within the current sheet will not be uniform, due to the presence of turbulence, the intermittent presence of reconnected flux, and any collective effects we have neglected here.  To the extent that these are important they will also lower the electron heating efficiency by broadening the reconnection layer.  \\subsection{Summary}  The results above can be summarized as follows   $\\bullet$ Recent successful numerical testing of the model in Lazarian \\& Vishniac (1999), presented in a companion paper by Kowal et al. (2009, this volume) increase the appeal of the model and stimulate studies of its consequencies.  $\\bullet$ The aforementioned model of stochastic field reconnection provides justification for many of astrophysical simulations. If, on the contrary, the only reconnection model that works is the collisionless reconnection, this means that most of the numerical simulations, for instance, of interstellar medium are in error. Indeed, in the collisional environments the reconnection speed would be negligible, which cannot be achieved with the numerical simulations.  $\\bullet$ The model of magnetic field reconnection in the presence of weak magnetic field stochasticity is a natural generalization of the Sweet-Parker model of reconnection. Its many consequencies include bursts of reconnection, efficient first order Fermi acceleration of energetic particles, efficient diffusion of magnetic field in turbulent plasmas etc.  {\\bf Acknowledgments}. Research by AL is supported by the NSF grant AST 0808118 and the NSF Center for Magnetic Self-Organization and by EV is supported by the National Science and Engineering Research Council of Canada."
    },
    "0812/0812.2828_arXiv.txt": {
        "abstract": "We present results from a deep (1$\\sigma$ = 5.7~mJy~beam$^{-1}$ per 20.8~km~s$^{-1}$ velocity channel) \\atom{C}{}{12}O(1-0) interferometric observation of the central $60\\arcsec$ region  of the nearby edge-on starburst galaxy NGC 2146 observed with the Nobeyama Millimeter Array (NMA). Two diffuse expanding molecular superbubbles and one molecular outflow are successfully detected. One molecular superbubble, with a size of $\\sim1$~kpc and an expansion velocity of $\\sim50$~km~s$^{-1}$, is located below the galactic disk; a second molecular superbubble, this time with a size of $\\sim700$~pc and an expansion velocity of $\\sim35$~km~s$^{-1}$, is also seen in the position-velocity diagram; the molecular outflow is located above the galactic disk with an extent $\\sim2$~kpc, expanding with a velocity of up to $\\sim200$~km~s$^{-1}$. The molecular outflow has an arc-like structure, and is located at the front edge of the soft X-ray outflow. In addition, the kinetic energy ($\\sim 3 \\times 10^{55}$ erg) and the pressure ($\\sim 1 \\times 10^{-12\\pm1}$~dyne~cm$^{-2}$) of the molecular outflow is comparable to or smaller than that of the hot thermal plasma, suggesting that the hot plasma pushes the molecular gas out from the galactic disk. Inside the $\\sim1$ kpc size molecular superbubble, diffuse soft X-ray emission seems to exist. But since the superbubble lies behind the inclined galactic disk, it is largely absorbed by the molecular gas. ",
        "introduction": "\\label{sect-intro}  A starburst galaxy is characterized by strong star forming activity. Starburst phenomena usually occur in the central regions of galaxies, and produce large numbers of massive stars on short time scales with star formation rates tens or hundreds of times higher than those in normal galaxies. Through stellar winds and energetic supernova explosions, these massive stars generate a large amount of mechanical energy that can sweep the surrounding interstellar medium (ISM) from starburst regions, create ISM bubbles, and generate high velocity galactic winds. This hypothetical starburst evolution is supported by several numerical models (e.g., \\cite{tom88,str00,mar05}), and is actually based on various observational results. Shell-like structures and outflows have been observed in many starburst galaxies in optical emission lines and soft X-ray emission (e.g., \\cite{leh96,mar98,dah98,str04}). Shell-like structures are also often observed in nearby galaxies in the atomic hydrogen (H\\emissiontype{I}) line (e.g., \\cite{ten88}).  Evidence for these shell-like or outflowing features from starburst galaxies are, however, rarely seen using molecular tracers. The famous cases are the molecular outflows in M82 \\citep{nak87} and NGC 891 \\citep{han92}, and molecular bubbles in NGC 3628 \\citep{irw96}. The reasons for the rarity of such detections are mostly their diffuse and extended nature and the lack of sensitive detections in this wavelength range. Recent improvements of various instruments at millimeter-wave have resulted in progress in this field: The interferometric \\atom{C}{}{12}O(1-0) observations toward M82 succeeded in imaging molecular superbubbles and outflows (e.g., \\cite{nei98,wei99,mat00,wei01}), and revealed a self-induced starburst mechanism \\citep{mat05}. Molecular superbubbles and/or outflows have also been detected in NGC 253 \\citep{sak06}, NGC 4631 \\citep{ran00}, and NGC 4666 \\citep{wal04}. Since stars form from molecular gas, observing these peculiar features in molecular gas can provide information about the regulation and feedback processes of the starburst phenomena. Here we report a new example of molecular superbubbles and outflows from the starburst galaxy NGC 2146.  NGC 2146 is a nearby (17.2 Mpc, 1$\\arcsec$ = 80 pc; \\cite{tul88}), IR luminous ($1.2 \\times 10^{11}$~L$_\\odot$; \\cite{san03}), edge-on ($i = 63\\arcdeg$; \\cite{del99}) starburst galaxy. Its optical appearance shows disturbed features with peculiar spiral arms, which cannot be explained by a simple model of spiral structure \\citep{ben75,gre06}. The neutral hydrogen images display 100~kpc scale tidal tails, although there is no evidence of a nearby companion galaxy \\citep{fis76,tar01}. A large amount of molecular gas is concentrated around the central region, which is enough to support the starburst activities \\citep{jac88,you88}. The central region hosts several compact radio continuum sources, identified as supernova remnants, radio supernovae, and ultra-compact and/or ultra-dense H\\emissiontype{II} regions \\citep{tar00}. The soft X-ray images show kpc-scale outflows from the starburst region around the galactic center \\citep{inu05,del99,arm95}. All these characteristics are reminiscent of those present in the nearby starburst galaxy M82, where kpc-scale outflows have also been detected, together with a molecular superbubble. Given the similarity between the two galaxies, it is natural to expect that similar features can be revealed also in NGC 2146. Therefore, we have performed deep observations in the \\atom{C}{}{12}O(1-0) line toward the central region of NGC 2146 as a case study for molecular superbubbles and outflows. ",
        "conclusions": "\\label{sect-dis}  \\subsection{Soft X-ray Emission from the Large-Scale Molecular Bubble and Outflow} \\label{sect-xray}  NGC 2146 has a kpc-scale outflow from the galactic center detected at soft X-ray wavelengths \\citep{inu05,arm95}. Here, we compare our \\atom{C}{}{12}O(1-0) map with the soft X-ray image obtained with the Chandra X-ray Observatory \\citep{inu05}. Figure~\\ref{xray} shows the \\atom{C}{}{12}O(1-0) integrated-intensity contour map overlaid on the soft X-ray image. The soft X-ray image shows strong extended emission toward the northeastern part of the galaxy, and weak extended emission in the southwest. Between these two diffuse emission regions, there is almost no emission except for a number of strong point sources. The latter location corresponds to the strong CO emitting region, namely the galactic disk region, and therefore the lack of diffuse soft X-ray emission can be explained by absorption from the molecular gas in the galactic disk. Because the stronger X-ray emission appears in the northeastern region of the galactic disk rather than in the southwestern region, the most probable picture is that the southwestern edge of the CO emission is at the near side of the galactic disk and the northeastern one is at the far side. This scenario is indeed consistent with previous studies \\citep{arm95,gre00}.  Since the soft X-ray image traces the location of hot thermal plasma in this galaxy \\citep{arm95,del99,inu05}, the comparison between our \\atom{C}{}{12}O(1-0) map and the soft X-ray map gives the schematic image of the relative location between the molecular gas and the hot plasma. In figure~\\ref{xray}, the northeastern X-ray emission shows a conical structure, emanating from the center of the galaxy and spreading outwards, perpendicular to the major axis of the galaxy (color-scale features enclosed by two red shorter dashed lines). The molecular outflow, {\\it OF1}, is located around the front edge of the X-ray outflow. This fact and the good spatial correspondence suggest that these two outflows are physically connected.  We then compared the molecular gas outflow energy with that of the soft X-ray outflow. As reported in section~\\ref{sect-of} (and in Table~\\ref{tabbubl}), the energy of {\\it OF1} is calculated as $E_{\\rm OF1} \\sim 3.0 \\times 10^{55}$~erg. On the other hand, the energy of the thermal plasma outflow is estimated as $E_{\\rm X} = 5.1 \\times 10^{56}~f^{\\frac{1}{2}}$~erg (with their value re-scaled because of the different distance for NGC 2146 adopted by us; \\cite{inu05}), where $f$ is the volume filling factor for the X-ray emitting gas. If we adopt the standard value of 0.01 -- 0.1 for $f$ \\citep{arm95}, the thermal plasma outflow energy can be calculated as $(0.5-1.6)\\times10^{56}$~erg. The energy of the thermal plasma outflow is therefore comparable or an order of magnitude larger than the molecular outflow energy, depending on the volume filling factor for the X-ray emitting gas. The thermal plasma would therefore have enough energy to push {\\it OF1} out from the galactic disk.  We also compared the pressure of the molecular gas in the outflow with that of the soft X-ray outflow. The molecular gas pressure can be measured by $P_{\\rm mol} \\sim n({\\rm H_{2}}) k_{\\rm B} T$, where $n({\\rm H_{2}})$ is the H$_{2}$ gas number density, $k_{\\rm B}$ is the Boltzmann constant, and $T$ is the molecular gas temperature. Since our observations do not provide any density or temperature information, here we assume standard values, namely a molecular gas density and temperature traced by \\atom{C}{}{12}O(1-0) of $10^{2-3}$~cm$^{-3}$ and 10 -- 100 K, respectively. Thus the molecular gas pressure is estimated as $\\sim1.4\\times10^{(-12\\pm1)}$~dyne~cm$^{-2}$. The thermal plasma outflow pressure is estimated as $P_{\\rm X} = (1.0 - 3.2)\\times10^{-11}$~dyne~cm$^{-2}$ with the volume filling factor of 0.1 -- 0.01 \\citep{inu05}, the pressure of the thermal plasma outflow is therefore comparable or two orders of magnitude larger than that of the molecular outflow, depending on the volume filling factor for the X-ray emitting gas and the density and/or temperature of the molecular gas. This result further supports the capability of the thermal plasma to push {\\it OF1} out from the galactic disk. Since the energy and the pressure of the X-ray emission is larger than that of the molecular outflow, it is natural to consider that the expanding force from X-ray emission is still pushing the surrounding molecular gas.  The distribution, kinematics, energies, and pressure of the molecular gas and the hot thermal plasma outflows suggest that the hot thermal plasma pushes the molecular gas outward, and produces the large-scale outflows. The hot thermal plasma is caused by supernova explosions in the central starburst region of NGC 2146 (e.g., \\cite{arm95}) and, very likely, the molecular gas outflow is also caused by the intense starbursts at the center of NGC 2146 \\citep{gre00}. Furthermore, the arc-like feature of {\\it OF1}, with the conical structure of the hot thermal plasma outflow just behind {\\it OF1}, suggests that the molecular outflow could have had a bubble-like structure at the beginning of the expansion, and then been blown out as an arc-like feature as it evolved. Indeed, the molecular superbubble with hot gas inside observed in the central region of M82 \\citep{mat05} could well represent an analogue of {\\it OF1} at an early stage.  Two-dimensional hydrodynamical simulations (e.g., \\cite{tom88}) support the evolution of a molecular superbubble to an outflow. One of their models simulated the evolution of a superbubble under the ISM density of 100~cm$^{-3}$, which corresponds to that of molecular gas seen in \\atom{C}{}{12}O(1-0). Supernova explosions at a constant rate for a long time produce a vast amount of hot gas, and the hot gas accelerates and expands the surrounding ISM gradually. This model shows that the size of the molecular superbubble can expand to $\\sim$2 kpc on a timescale of $\\sim$10$^7$ years. Since the supernova rate (SNR) of this model adopts that of M82 (0.1 yr$^{-1}$), and NGC 2146 has a similar SNR (0.15 yr$^{-1}$; \\cite{tar00}), this model can be considered valid for NGC 2146. Furthermore, the values of the size and the timescale from this model are similar to our observations, and therefore our results are consistent with the model.  At the center of the expanding molecular bubble {\\it SB1}, diffuse soft X-ray emission, indicative of hot thermal plasma, seems to be present. Figure~\\ref{xray} show two soft X-ray images with different spatial resolutions: Figure~\\ref{xray}a is the higher spatial resolution image, where point sources are clearly visible, while it is more difficult to see the diffuse components. Figure~\\ref{xray}b is the lower spatial resolution image, where diffuse components are more clearly seen. Comparison of these two images can separate point sources and diffuse sources, and indeed at the center of {\\it SB1}, diffuse soft X-ray emission is detected. This situation is again similar to that observed in M82 and expected from hydrodynamical simulations, suggesting that the hot gas is embedded  inside of the expanding molecular bubble {\\it SB1}, pushing the surrounding molecular gas outward. A precise estimate of the energy of the soft X-ray component is, however, difficult since, in this region, the soft X-ray emission is absorbed by the foreground molecular gas and, therefore, any value derived for the energy would be significantly underestimated.    \\subsection{Configurations of Outflows and Bubbles} \\label{sect-conf}  Here we discuss the possible configuration of the outflows and bubbles in the central $60\\arcsec$ of NGC 2146 based on our CO and X-ray data, and the previously published optical spectra. A schematic diagram of the possible configuration of the molecular gas and other components is shown in figure~\\ref{conf} and, in the following paragraphs, its details are discussed.  As shown in section~\\ref{sect-model}, the molecular gas disk is well modeled by the standard rigid- and flat-rotation curves with an inclination of $63\\arcdeg$. Based on this simple regular galactic kinematics, it is safe to assume that the galactic disk is flat and has no warp in the central region. Comparison between the CO in the galactic disk and the X-ray data indicates that the outflow is directed toward us in the northeastern side, and the CO in the galactic disk in ``blocking'' the soft X-ray emission in the southwestern side, making that the near side (Sect.~\\ref{sect-xray}). Indeed, the long-slit optical spectra along the minor axis show that the ionized gas is blueshifted to the northeast and redshifted to the southwest, and the spectra taken in the central $2\\arcsec$ show a blue asymmetry in the emission lines, suggesting that the ionized gas outflow is flowing toward us in the northeast, and the redshifted outflow is blocked by the galactic disk around the galactic center \\citep{arm95}, which is consistent with the configuration we propose.  The soft X-ray outflow shows a conical structure with a large opening angle (see figure~\\ref{xray}), and the outflowing direction is inclined toward us. From the blue asymmetry in the optical spectra near the galactic center, the line-of-sight ionized gas outflow velocity toward us is estimated as $\\sim300-400$~km~s$^{-1}$ \\citep{arm95}. The redshifted side, on the other hand, shows only $\\sim100$~km~s$^{-1}$ (estimated from figure~11 of \\cite{arm95}), suggesting that the far side of the conical structure of the outflow is nearly perpendicular to the line-of-sight. {\\it OF1} is dominated by the redshifted gas, suggesting that most of the molecular gas is located at the far side of the conical structure. The line width of ionized gas increases with distance to the galactic center \\citep{arm95}, suggesting an acceleration as the gas moves away, consistent with the kinematics of {\\it OF1}.  {\\it SB1} is located in the southwestern part of the galaxy, indicating that {\\it SB1} is at the farside of the galactic disk. The radius of the molecular gas disk is about 2~kpc (about $25\\arcsec$) with an inclination of $63\\arcdeg$, while the central velocity of {\\it SB1} is somewhat blueshifted and the radius of {\\it SB1} is about $\\sim1$~kpc. Therefore, in order to be able to partially see {\\it SB1}, this has to be closer to us than the galactic center.  Summing up all the aforementioned information, we made a schematic diagram of the possible configuration of the molecular gas and the hot gas outflows in figure~\\ref{conf}. In the figure, the molecular gas is shown in red, X-ray emission in yellow, and optical emission line in light blue. The emission obscured by the molecular gas (and possibly dust that is mixed with the molecular gas) in front is shaded. The overall configuration of this schematic diagram is globally consistent with that presented in \\citet{gre00} but, by using our new observations, we have been able to describe the different structures and components presented in the inner region of NGC 2146 in greater detail.  In sections~\\ref{sect-sb} and \\ref{sect-of}, we derive the properties of {\\it SB1} and {\\it OF1}. Since the timescale of {\\it SB1} corresponds roughly to that of {\\it OF1}, it is possible that {\\it SB1} and {\\it OF1} originated from one starburst event. The difference in the structure ({\\it SB1} is still a bubble while {\\it OF1} is already outflowing) can be explained by the possible location where the starburst may have started up; if a starburst occured a bit above the galactic disk plane, an expansion upward is less frustrated than downward due to a less dense ISM. Since a bubble expands faster in less dense conditions \\citep{tom88}, an expansion in a lower density ISM leads to faster bubble evolution, and the bubble becomes an outflow earlier than that in a dense region. However, the possible central locations for {\\it SB1} and {\\it OF1} are different (figure~\\ref{m0m1}; see also figure~\\ref{conf}), it is difficult to say that {\\it SB1} and {\\it OF1} have the same explosion center. Thus it is safer to conclude that {\\it SB1} and {\\it OF1} may have formed in the same explosion period, but at different locations."
    },
    "0812/0812.2475_arXiv.txt": {
        "abstract": "The star HD 69830 exhibits radial velocity variations attributed to three planets as well as infrared emission at $8-35\\mu$m attributed to a warm debris disk.  Previous studies have developed models for the planet migration and mass growth \\citep{Alibert_et_al_05,Alibert_et_al_06} and the replenishment of warm grains \\citep{Wyatt_et_al_07a}.  In this paper, we perform n-body integrations in order to explore the implications of these models for: 1) the excitation of planetary eccentricity, 2) the accretion and clearing of a putative planetesimal disk, 3) the distribution of planetesimal orbits following migration, and 4) the implications for the origin of the infrared emission from the HD 69830 system.  We find that: i)  It is not possible to explain the observed planetary eccentricities ($\\sim 0.1$) purely as the result of planetary perturbations during migration unless the planetary system is nearly face-on. However, the presence of gas damping in the system only serves to exacerbate the problem again. ii) The rate of accretion of planetesimals onto planets in our n-body simulations is significantly different to that assumed in the semi-analytic models, with our inner planet accreting at a rate an order of magnitude greater than the outer ones, suggesting that one cannot successfully treat planetesimal accretion in the simplified manner of \\citet{Alibert_et_al_06}. iii ) We find that the eccentricity damping of planetesimals does \\emph{not} act as an insurmountable obstacle to the existence of an excited eccentric disk: All simulations result in a significant fraction ($\\sim$15\\%) of the total planetesimal disk mass, corresponding to $\\sim 25 \\Mearth$, remaining bound in the region $\\sim$1-9 AU, even after all three planets have migrated through the region iv) This swarm of planetesimals has orbital distributions that are size-sorted by gas drag, with the largest planetesimals ($\\sim 1,000km$), which may contain a large proportion of the system mass, preferentially occupying the highest eccentricity (and thus longest-lived) orbits. Although such planetesimals would be expected to collide and produce a disk of warm dust, further work will be required to understand whether these eccentricity distributions are high enough to explain the level of dust emission observed despite mass loss via steady state collisional evolution. ",
        "introduction": "Radial velocity surveys have discovered over 300 extrasolar planets and roughly 30 multiple planet systems (Butler et al.\\ 2006; \\mbox{http://www.exoplanets.org}, \\mbox{http://www.exoplanet.eu}). The multiple planet systems are particularly valuable for gaining insights into the processes of planet formation (e.g., \\citet{Lee02,Ford05}).  Among the known multiple planet systems, the {\\em Spitzer} Legacy Program FEPS \\citep{Meyer06} has identified four systems that emit an excess of infrared radiation (relative to that expected from the stellar photosphere), most likely due to reradiation of starlight absorbed by a dust disk.  In these systems, the properties of the dust disk can provide additional constraints on the system's formation and dynamical evolution (e.g., \\citet{Moro-Martin07}).  One of these systems is HD 69830, a bright K0V star located 12.6 pc away \\citep{Perryman97}. In addition to possessing 3 Neptune-mass planets \\citep{Lovis_et_al_06}, this exceptional system was the \\emph{only} system out of 84 FKG stars studied by \\citet{Beichman_et_al_05} which was found to display an excess of emission at 24$\\mu$m, indicative of \\emph{warm} dust (other systems were observable only at 70$\\mu$m, indicative of much cooler ($T<100$K) dust).  All three known planets orbit at less than 1 AU, so the putative dust disk is exterior to the known planets, but still close enough that gravitational perturbations from the planets might be effective at stirring a disk of planetesimals, resulting in collisions that could generate dust that would give rise to the observed IR excess.  Thus the HD 69830 system is particularly interesting due to a unique combination of radial velocity and IR excess observations that have revealed multiple planets and a close-in warm dust disk.     \\subsection{Planets: Radial Velocity Observations and Theoretical Models}\\label{Planets}  Based on HARPS radial-velocity measurements of HD 69830, \\citet{Lovis_et_al_06} identified planets ($M\\sin{i}=10.2$, $11.8$ \\& $18.1M_{\\oplus}$ respectively) orbiting close-in to the central star (semi-major axes of 0.08, 0.19 \\& 0.63 AU).  The planets' eccentricities ($0.10\\pm 0.04$, $0.13\\pm 0.06$ \\& $0.07\\pm 0.07$ respectively, \\citet{Lovis_et_al_06}) are large by solar system standards, but modest when compared to those of extra-solar planets \\citep{2007astro.ph..3163F}.  The host star is somewhat cooler ($T_{\\rm eff}\\simeq$ 5385 K) and less massive than the Sun ($M_{\\star}\\simeq 0.86 M_\\odot$).  Otherwise, it is quite similar to the Sun, with a nearly solar metallicity ([Fe/H]= $-0.05\\pm 0.02$) and age (4-10 Gyr; \\citet{Lovis_et_al_06}).  Since the known planets produce modest radial velocity perturbations ($2.2 - 3.5\\textrm{m s}^{-1}$) just large enough for detection, similar planetary systems could have often eluded detection by broad radial velocity planet searches on account of an insufficient number of velocity observations and/or velocity precision due to photon noise and/or radial velocity ``jitter''.  Recent results from the HARPS search for southern extra-solar planets and its discovery of a planetary system with 3 Super-Earths \\citep{Mayor08} bolsters the view that similar planetary systems might be quite common.  These radial velocity observations alone reveal an interesting planetary system.  Most of the known planets in multiple planet systems have masses roughly comparable to Jupiter (most likely due to detection biases). The unusual intermediate planetary masses revealed in the HD 69830 system have already inspired several theoretical investigations.  As a result, theorists have developed a detailed model for the mass growth and orbital migration of the observed planets through a series of papers \\citep{Alibert_et_al_05,Alibert_et_al_06,Lovis_et_al_06}.  In particular, \\citet{Alibert_et_al_06} found that the observed characteristics could be reproduced by starting with a gas disk in which the surface density, $\\Sigma$, is related to the disk radius, $a$, by $\\Sigma \\propto a^{-3/2}$ and normalized to $800 \\textrm{g/cm}^2$ at 5 AU (This surface density is around 4 times greater than the minimum mass solar nebula, giving $M_{disk} = 0.07M_{\\odot}$ ($0.07\\textrm{AU} \\rightarrow 30 \\textrm{AU}$)). The disk has a dust-to-gas ratio of 1/70. They then inserted 3 planetary embryos of mass $0.6 M_{\\oplus}$ at initial semi-major axes of 3, 6.5 \\& 8 AU and used a semi-analytic model to migration and grwoth of the embryos due to their interaction with the disk over the course of the disk lifetime, $\\tau_{disk} = 2$Myr. After accretion and migration, the 3 planets in the model were found to have shifted inwards to the observed semi-major axes, have total masses consistent with the minimum observed masses and possess \\emph{core} masses of $\\sim 10$, $\\sim 7.5$ \\& $\\sim 10 \\Mearth$ respectively.  While the \\citet{Alibert_et_al_06} model does an impressive job of matching the observed planet masses and semi-major axes, it is a 1-D semi-analytic model, using numerous assumptions and approximations to model the interaction between the protoplanetary disk and the growing and migrating planetary cores. This raises questions regarding the effects of the planets on any planetesimal disk and the feedback effects of disk evolution and ``shepherding'' on planetesimal accretion rates. In addition, the model does not address the eccentricity evolution of the planets and planetesimals.  In this paper, we build on this work, as we continue the quest to understand the formation and orbital evolution of this fascinating system by applying n-body methods to follow planetesimal evolution within the context of the semi-analytical model of \\citet{Alibert_et_al_06}.      \\subsection{IR Observations}\\label{IRobs}  Spitzer observations of the HD 69830 system reveal a strong infrared excess (relative to the stellar photosphere) between 8 \\& $35\\mu m$, but \\emph{no} significant excess at $70\\mu m$ \\citep{Beichman_et_al_05}.  This combination suggests that the system contains a disk of warm ($\\sim 400K$), small ($<1\\mu m$) crystalline silicate grains orbiting close to the star ($\\sim 1 \\textrm{AU}$). For these parameters, the collisional timescale (for $\\mu$m grains at $\\sim1$AU from HD 69830) is $\\sim400$ yr, while the Poynting-Robertson drag timescale is $\\sim700$ yr \\citep{Beichman_et_al_05}.  Thus, collisions in such a disk would grind down the grains until they became small enough to be removed via radiation pressure. Unless we happen to be observing the system at a very special time, this implies that the grains are replenished on a timescale $<1000$ yrs so as to sustain the observed IR excess.  Several models \\citep{Beichman_et_al_05,Lisse_et_al_07,Wyatt_et_al_07a} have been proposed to explain the origin of the dust responsible for the IR excess, including: \\begin{enumerate} \\item A massive cometary population, \\item The capture of a super-comet onto a circular orbit at $\\sim 1$ AU, \\item The steady-state evolution of a planetesimal belt at $\\sim 1$ AU, \\item A recent, large collision in a planetesimal belt at $\\sim 1$ AU, and \\item Recent dynamical instability which results in planetesimals from an outer belt being thrown inwards. \\end{enumerate}  In the first scenario, the observed dust would be released by pristine comets entering the inner region of the planetary system. \\citet{Beichman_et_al_05} noted the observed spectral signature was similar to that of comet Hale-Bopp, but that reproducing the emission using cometary ejecta similar to that of Hale-Bopp would require $\\sim10^6$ such comets per year to be delivered to the inner regions of the system.  If the putative comets were of the same mass as Hale-Bopp and the dust were to be sustained for $\\sim10^7$ yr, then this would imply $\\sim900M_{\\oplus}$ of comets entering the inner regions of the planet system, unfeasibly large for a residual Kuiper-Belt.  In the second scenario, \\citet{Beichman_et_al_05} suggested that a single object of the size of a large Kuiper Belt Object (e.g., Sedna) composed of ice and rock may have been scattered inwards to $\\sim 0.5$AU.  They estimate that such an object would have an evaporation timescale of $\\sim2$ Myr, allowing a reasonable chance of observation.  Given the masses and orbits of the three known planets, a Sedna-like object at $\\sim0.5$AU could not be dynamically stable, unless it were on a low eccentricity orbit. In order to have a reasonable chance of observing the IR excess from such a body, the dynamics of the system must be such that a Sedna-like object could have recently been perturbed into the inner system (presumably on a highly eccentric orbit) and then had its orbit circularized.  For the third scenario, \\citet{Wyatt_et_al_07a} considered the possibility that the dust originates in a primordial planetesimal belt that is coincident with the dust and which has evolved in a quasi-steady state. They found that collisional processing would have removed most of the belt's mass over the $>2$ Gyr age of the system and that the observed levels of dust are incompatible with this interpretation for a similar reason. One possible resolution to this is that planetesimals (and dust) may orbit with significant eccentricities, thus prolonging their survival, meaning that the emission at 1 AU could come from the inner edge of an extended disk when the planetesimals and dust are at pericentre (Wyatt et al, in preparation).  Similarly, for the fourth scenario \\citet{Wyatt_et_al_07a} showed that an {\\em in situ} planetesimal belt would be extremely unlikely to have undergone the single recent collision that would give rise to the observed IR excess. The results of \\citet{Lohne08} may slightly increase the probability of a collision having occurred, as their model predicts higher remnant masses at late times, but a significant increase in probability would seem to require unfeasibly high initial disk masses.  Finally, in the fifth scenario, \\citet{Wyatt_et_al_07a} have proposed a model in which the current IR excess is due to a delayed epoch of orbital instability among planetesimals in an outer disk, perhaps similar to that of the Late Heavy Bombardment (LHB) in our own solar system \\cite{2005Natur.435..466G}. Clearly this model then raises further issues such as how such a delayed instability might have been triggered and why there has been no detection of the source population of planetesimals.  All of these scenarios have unresolved problems, but their unifying feature is that they all depend to some extent on the existence of an extended distribution of planetesimals beyond the observed planets in the system. Fortunately the architecture of the planetary system sets constraints on how it formed, which in turn has implications for the planetesimal population remaining following planet formation. By combining planetary formation models (described in \\S \\ref{Planets}), with the N-body techniques described in \\S \\ref{Method}, we aim to develop an understanding of the remnant disk structure likely to populate the systems at the end of the planet formation process.      \\subsection{Outline} We perform dynamical simulations to investigate the formation and current state of the HD 69830 planetary system.  Using numerical methods described in \\S\\ref{Method}, we look in \\S\\ref{PlanetExcite} at the implications of the \\citet{Alibert_et_al_05} model for the planetary eccentricities and accretion rates.  In \\S\\ref{PlanetesimalExcite}, we perform additional simulations to model the interaction between the planets and a putative protoplanetary distribution of planetesimals in order to model the predicted planetesimal distribution after the migration and growth of planets has been completed. In \\S\\ref{Problems} we consider possible variations in the formation and orbital evolution of the HD 69830 planetary system.  In \\S \\ref{Discussion} we consider the planetesimal distributions resulting from our simulations and discuss their implications for explaining the observed dust emission. Finally, we summarize our conclusions and future prospects in \\S\\ref{Summary}. ",
        "conclusions": "\\label{Discussion} \\subsection{Planet Formation Models}\\label{PFM} Numerous studies \\citep{Kokubo06, 2006ApJ...644.1223R} show that the growth of embryos to isolation mass in a $\\Sigma \\propto a^{-3/2}$ disk will occur inside-out, that is, the timescale for growth to isolation in such a disk will be $\\tau_{iso} \\propto a^{2}$ and hence isolation will occur quickest in the inner disk. Therefore a migrating embryo which initially formed at 3 AU will be migrating through a region of space which is \\emph{not} occupied by a smooth solid disk, but rather by a series of massive embryos which have already grown to isolation. The core growth of the embryo must then take place via giant impacts, which will clearly have a significant impact on both the planetary semi-major axes and eccentricities which would be predicted at the end of the growth and migration regime.  In addition, we have seen in our work above that significant ``shepherding'' of material can occur in front of migrating protoplanets. This phenomena combined with the stirring up of the planetesimal disk by the passage of the first embryo means that the core growth rate for the first embryo is (at least) an order of magnitude greater than for the subsequent cores, a result that is significantly at odds with the assumptions of semi-analytic models.  We also find that embryos which have grown more massive \\emph{prior} to passage through MMR are prefered in the model, as this allows an easier, more natural explanation of the observed planetary eccentricities, although the subsequent gas damping of these eccentricities still presents problems.  These results serve to emphasise that a consistent model for planet formation and evolution is not possible using the current generation of semi-analytic models. In order to make progress, the semi-analytic growth models need to be combined with n-body models to ensure that both the mutual planet-planet interactions as well as the planet-planetesimal interactions are self-consistently modelled and able to be compared with observations.    \\subsection{Eccentric Planetesimal Disk}\\label{EPD} We find in \\S \\ref{PlanetesimalExcite} that the end product of our n-body simulations is the prediction that the excitation and retention of an eccentric planetesimal disk exterior to the three planets in HD 69830 is almost certain to take place (within the model parameters investigated). Significantly, we found that the mass of planetesimals remaining in this disk was large: after the planets had migrated and the gas disk dissipated, there would be $\\sim 25 - 45 \\Mearth$ of material present in the scattered disk (using the initial disk profile of \\citep{Alibert_et_al_06}).  The presence of this amount of material in the young system (few Myr) does not necessarily guarantee that significant amounts of mass will remain in the system at late times (few Gyr), as the steady-state collisional processing of the disk will remove a large amount of material. The results of \\citet{Wyatt_et_al_07a} indicate that low eccentricity planetesimal belts evolve to a mass that is \\emph{independent} of initial mass, and in the case of HD69830, if this material were placed in a ring at 1 AU, then after $2Gyr$ the collisional processing would result in just $10^{-5} \\Mearth$ of planetesimals remaining, which would be $1,000$ times less than required to explain the observed emission. This conclusion may be mitigated to some extent by the results of \\citet{Lohne08} which show a small dependence of dust mass at late times on initial planetesimal mass, but this would probably not be sufficient as there results imply that the initial planetesimal mass would have to be unrealistically large to explain the observed emission in HD 69830.  To work out how our planetesimals evolve we need to consider the collisional evolution of a population with a range of eccentricities and semimajor axes: This has yet to be achieved, however, work in preparation (Wyatt et al) already shows that increasing the eccentricity of a planetesimal belt from nearly circular to 0.9 while keeping the pericentre constant increases the amount of material remaining at late times by 1 to 2 orders of magnitude. Although certainly giving no guarantee, this increases confidence that the distributions we derived might be able to explain the observed emission.  However, our simulations also show that the size distribution of the planetesimals is extremely important, since it is only $> 100km$ planetesimals that remain on highly eccentric orbits following scattering by planet migration and subsequent gas damping. Thus, it is important to consider the fraction of the remaining $25-45\\Mearth$ of planetesimals that have radii $> 100km$. The planetesimals would be expected to have a bimodal population of the type seen in the simulations of \\citet{Morishima08}, who found that at $10^6$yrs, $\\sim 70\\%$ of the solid mass would be in the form of a few large embryos, with the rest residing in a population of planetesimals with a number density distribution given by $n(m)dm = m^{-2}dm$. Such a distribution has mass evenly spread at all sizes (on a log-scale), and thus we would anticipate that approximately $10\\%$ ($\\sim 2.5-4.5\\Mearth$) of the surviving planetesmal material would reside in bodies of size 100 - 1000 km, with a further $70\\%$ ($\\sim 18-32\\Mearth$) locked up in the very large embryos. In the \\citet{Alibert_et_al_06} model, from the initial $0.5 M_J = 160\\Mearth$ of solid material in the disk, only $\\sim 27\\Mearth$ ($15\\%$) is accreted onto the migrating cores, so the contribution of the large embryos to the scattered distribution could be significant. Thus there is likely to be a significant mass fraction in large scattered bodies, which, along with the size sorting of the planetesimals via gas drag, works to ensure that a large proportion of the remnant planetesimal mass would be locked up in large bodies which would be on the most eccentric orbits."
    },
    "0812/0812.2196_arXiv.txt": {
        "abstract": "We propose a simple model of supersymmetric dark matter that can explain recent results from PAMELA and ATIC experiments. It is based on a $U(1)_{B-L}$ extension of the minimal supersymmetric standard model. The dark matter particle is a linear combination of the $U(1)_{B-L}$ gaugino and Higgsino partners of Higgs fields that break the $B-L$ around one TeV. The dominant mode of dark matter annihilation is to the lightest of the new Higgs fields, which has a mass in the GeV range, and its subsequent decay mainly produces taus or muons by the virtue of $B-L$ charges. This light Higgs also results in Sommerfeld enhancement of the dark matter annihilation cross section, which can be $\\gs 10^3$. For a dark matter mass in the $1-2$~TeV range, the model provides a good fit to the PAMELA data and a reasonable fit to the ATIC data. We also briefly discuss the prospects of this model for direct detection experiments and the LHC. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.0522_arXiv.txt": {
        "abstract": "The excesses of the cosmic positron fraction recently measured by PAMELA and the electron spectra by ATIC, PPB-BETS, Fermi and H.E.S.S. indicate the existence of primary electron and positron sources. The possible explanations include dark matter annihilation, decay, and astrophysical origin, like pulsars. In this work we show that these three scenarios can all explain the experimental results of the cosmic $e^\\pm$ excess. However, it may be difficult to discriminate these different scenarios by the local measurements of electrons and positrons. We propose possible discriminations among these scenarios through the synchrotron and inverse Compton radiation of the primary electrons/positrons from the region close to the Galactic center. Taking typical configurations, we find the three scenarios predict quite different spectra and skymaps of the synchrotron and inverse Compton radiation, though there are relatively large uncertainties. The most prominent differences come from the energy band $10^4\\sim 10^9$ MHz for synchrotron emission and $\\gtrsim 10$ GeV for inverse Compton emission. It might be able to discriminate at least the annihilating dark matter scenario from the other two given the high precision synchrotron and diffuse $\\gamma$-ray skymaps in the future. ",
        "introduction": "The anti-matter particles such as positrons and antiprotons in cosmic rays (CRs) are very important in understanding the origin and propagation of CRs. These particles are usually produced by the CR nuclei that interact with the interstellar medium (ISM) when propagating in the interstellar space. Precise measurements of these particles will provide us valuable information about the primary CR sources and the interaction with matter. In addition, these secondary anti-matter particles have relative lower fluxes and characteristic spectra. Therefore they make themselves good objects to study the exotic origin of CRs, such as from dark matter (DM) annihilation or decay.  Recently the PAMELA collaboration released their first CR measurements on the positron fraction \\cite{Adriani:2008zr} and $\\bar{p}/p$ ratio \\cite{Adriani:2008zq}. The positron fraction of PAMELA data shows an uprise above $\\sim 10$ GeV till to $\\sim 100$ GeV, which exceeds the background estimation of the conventional CR propagation model. This result is consistent with previous measurements by, e.g., HEAT \\cite{Barwick:1997ig} and AMS \\cite{Aguilar:2007yf}. On the other hand, the $\\bar{p}/p$ ratio is compatible with the background. The results of PAMELA strongly indicate the existence of such sources that mainly generate leptonic particles at this energy range. The electron spectrum up to several TeV reported by ATIC collaboration also shows an obvious excess with an interesting bump around $300\\sim800$ GeV \\cite{Chang:2008zz}. In addition, the measurements of the electron spectrum by PPB-BETS \\cite{Torii:2008xu}, H.E.S.S. \\cite{Collaboration:2008aa,Aharonian:2009ah}, and most recently by Fermi \\cite{Collaboration:2009zk} all show the excesses of electron spectra, although they are not fully consistent between each other. The studies show that the PAMELA result on the positron fraction is consistent with the electron spectrum measured by ATIC or Fermi assuming the equal amount of the production of positrons and electrons (e.g., \\cite{Cholis:2008wq,Bergstrom:2009fa}).  One possible primary electron/positron source in the Galaxy is the pulsar. Pulsars and their nebulae are well known cosmic particle accelerators. Although the quantitative details of the acceleration processes are still open for study, early radio observations have established them as important high energy cosmic electron and positron sources (e.g., \\cite{Frail1997}). X-ray and $\\gamma$-ray observations show that some of the accelerated particles can reach an energy of a few tens of TeV, and there are indications that the particle distribution cuts off in the TeV energy range (e.g., \\cite{Pavlov2001,Lu:2002kx, Aharonian:2006xx}). These particles produce emission over a broadband frequency range in the source region, which has been observed and studied extensively (e.g., \\cite{Li:2007dg}). Some of these particles will escape into the ISM becoming high energy cosmic electrons and positrons. It is shown that one or several nearby pulsars will be able to contribute enough positrons to reproduce the PAMELA data \\cite{Hooper:2008kg,Yuksel:2008rf} as well as ATIC data \\cite{Profumo:2008ms}.  Another potential candidate, which has been widely discussed to solve the positron/electron excess problem, is the DM annihilation (e.g., \\cite{annidm}) or decay (e.g.,\\cite{decaydm,Yin:2008bs,Nardi:2008ix}).  It has been shown in many studies that the local measurements of the electrons/positrons can not provide enough power to discriminate the DM origin and the pulsar origin of the electrons/positrons. On one hand, the ATIC data suggest a sharp cutoff near $\\sim 600$ GeV. Such a feature can naturally appear for a DM origin when closing to the mass of DM particle, if the DM annihilates or decays directly to e$^+$e$^-$ pair, as shown in Ref. \\cite{Chang:2008zz}. However, the same thing is also expected if the observed excess is dominated by a single pulsar and/or its nebula and the accelerated high energy electrons and positrons have experienced significant energy loss as they propagate in the source and in the ISM \\cite{Profumo:2008ms}. The narrow peak resulting from local pulsars is even indistinguishable from that of DM origin \\cite{Pohl:2008gm}. On the other hand, if the electron spectrum is as smooth as that from Fermi, there are also degeneracies between the pulsar and DM interpretations. A continuously distributed pulsar population in the Galaxy might be able to generate smoother electron spectrum due to different energy losses of various pulsars. While for the DM scenario, if the electrons and positrons come from DM annihilation or decay, the spectrum can also be smooth, and is consistent with the Fermi measurements \\cite{Bergstrom:2009fa}.  In this work we study the photon emission associated with the different scenarios to see how the degeneracies between these models can be broken. Our point is that although the different scenarios can give almost the same local electrons/positrons, the spatial extrapolations of them might be significantly different, especially in the Galactic center (GC) region. The pulsars mainly concentrate in the Galactic plane, while the DM is spherically distributed in an extended halo and highly concentrated in the center of the halo. In addition, DM annihilation is proportional to $\\rho^2(r)$, different from that of DM decay ($\\propto \\rho(r)$). Therefore we anticipate the existence of remarkable differences in the skymaps of the synchrotron and inverse Compton (IC) emission from these extra primary electrons/positrons for the three scenarios.  We compare the differences of synchrotron as well as IC radiation among the pulsars, annihilating DM and decaying DM models, on the premise that they all reproduce the positron fraction and total electron + positron flux data. The spectral and directional features of the radiation are discussed in detail. We compare the results with the WMAP haze data \\cite{Finkbeiner:2003im,Dobler:2007wv} and EGRET diffuse $\\gamma$-ray data \\cite{Hunter:1997we}, and give the observable signals by future experiments such as Fermi/GLAST.  This paper is organized as follows. In Sec. II we will briefly present typical model configurations of the three scenarios which can fit the $e^\\pm$ data simultaneously. The discussions on the synchrotron and IC emission are given in Sec. III and IV respectively. In Sec. V we discuss the uncertainties of our predictions. Finally, we give a summary and draw the conclusion in Sec. VI. ",
        "conclusions": "To explain the recent measurements of the positron fraction by PAMELA and electron spectrum by ATIC, Fermi and HESS, several scenarios including the DM and pulsars are proposed. In this paper we studied the perspective to discriminate these models using the synchrotron and IC radiation generated by these electrons/positrons. The point is that although various kinds of models degenerate in the local environment\\footnote{High energy electrons lose energy fast and the detected electrons should indeed be ``local'', e.g., within $\\sim 1$ kpc \\cite{Maurin:2002uc}.} and give similar electrons/positrons spectra, their different spatial extension might be % revealed by the accompanied photon emission.  By properly adjusting the model parameters we first build three benchmark models which can reproduce the measured electron/positron data. We then studied the synchrotron and IC radiation from these primary electrons and positrons. We find that around the GC region, the synchrotron spectra for frequencies from $10^4 \\sim 10^9$ MHz and the IC $\\gamma$-ray spectra from several GeV to several hundred GeV are significantly different among the three scenarios, especially between annihilating DM scenario and others. For the directional profiles, the DM models show larger gradient near the GC, especially in the longitude profiles, than the pulsar model and the background. The photon emissions for pulsar model are very small and similar with the background, and will be the most difficult one to be detected. While for DM models, especially the annihilating DM model, there are strong observable signals. Finally we discuss the possible uncertainties of these conclusions. The major uncertainty comes from the DM inner profile. For the most conservative case with cored isothermal DM profile, the signals of DM scenarios become weaker and might be hard to be discriminated. However, the isothermal profile is generally disfavored by numerical simulations. Other uncertainties, such as the DM annihilation or decay channels, the propagation parameters and other astrophysical inputs, seems not to change the relative differences among the signals of the three scenarios significantly.  It should be pointed out that in this work the boost factor for the annihilating DM scenario is assumed to be universal, such as the Sommerfeld effect. It is also possible that the boost effect is spatially dependent, e.g., from the DM subhalos \\cite{Yuan:2006ju,Bi:2007ab,Bi:2006vk}. If so, the enhancement near the GC region will not be as important as at large radius in the halo. Actually, the calculation based on the N-body simulation indicates that the boost factor from DM substructure for anti-matter particles from DM annihilation in the Galaxy is generally negligible \\cite{Lavalle:1900wn,Lavalle:2008zb}. A recent calculation on the DM clumpiness boost factor taking into account the Sommerfeld effect still finds no remarkable enhancement effect, except for fine tuning to the strongly resonant case \\cite{Yuan:2009bb}. Therefore the conclusions in this work are generally held."
    },
    "0812/0812.0839_arXiv.txt": {
        "abstract": "We discuss a consolidation of determinations of the density of neutral interstellar H at the nose of the termination shock carried out with the use of various data sets, techniques, and modeling approaches. In particular, we focus on the determination of this density based on observations of H pickup ions on Ulysses during its aphelion passage through the ecliptic plane. We discuss in greater detail a novel method of determination of the density from these measurements and review the results from its application to actual data. The H density at TS derived from this analysis is equal to $0.087 \\pm 0.022$~cm$^{-3}$, and when all relevant determinations are taken into account, the consolidated density is obtained at $0.09 \\pm 0.022$~cm$^{-3}$. The density of H in CHISM based on literature values of filtration factor is then calculated at $0.16 \\pm 0.04$~cm$^{-3}$. ",
        "introduction": "\\label{sec:intro} With the plethora of detailed measurements available from heliospheric missions such as Voyager, Ulysses, SOHO, Cassini, ACE, Wind, and many others and in anticipation of the first mission dedicated to remote-sensing studies of the heliospheric interface IBEX, a consolidation of heliospheric parameters is an important task. Key parameters in this quest are the density of neutral interstellar hydrogen and helium at the nose of the termination shock. While a consensus was reached on the parameters of interstellar helium already a few years ago \\citep{mobius_etal:04a, mobius_etal:05a, witte:04, gloeckler_etal:04a, vallerga_etal:04a, lallement_etal:04a}, the hydrogen density remained more uncertain until very recently due to more complex effects on this species within the heliosphere. A dedicated team supported by the International Space Science Institute from Bern, Switzerland, has recently reached a consolidation of the density of neutral interstellar hydrogen at the nose of the termination shock based on different data sets and modeling. We discuss this consolidation, focusing mostly on the determination of the density at the termination shock based on hydrogen pickup ion (H PUI) observations by Ulysses during its passage through the ecliptic plane at aphelion of its orbit, performed using a novel analysis method. ",
        "conclusions": "\\label{sec:conclu} The density of neutral interstellar gas at the termination shock seems to be well constrained by observations performed using different techniques and appropriate modeling. The datasets used include Ulysses observations of H pickup ions, Voyager and Ulysses observations of solar wind slowdown, Cassini and Voyager observations of modulation in the heliospheric Lyman-alpha glow, and multi-year all-sky photometric observations of the glow by SWAN/SOHO. The densities obtained from analysis of these observations group in the range of 0.08 -- 0.09~cm$^{-3}$ and the uncertainty of this determination is about 20\\%. Extrapolation of this value to the CHISM based on analysis of heliospheric filtration factors from the literature brings the density of neutral interstellar hydrogen equal to $0.16 \\pm 0.04$~cm$^{-3}$."
    },
    "0812/0812.4592_arXiv.txt": {
        "abstract": "To foresee a solar flare neutrino signal we infer its upper and lower bound. The upper bound was derived since a few years by general energy equipartition arguments on observed solar particle flare. The lower bound, the most compelling one for any guarantee neutrino signal, is derived by most recent records of hard Gamma bump due to solar flare on January $2005$ (by neutral pion decay). Because neutral and charged  pions (made by hadron scattering in the flare) are born on the same foot, their link is compelling: the observed gamma flux \\cite{ref2} reflects into a corresponding one for the neutrinos, almost one to one. Moreover while gamma photons might be absorbed (in deep corona) or at least reduced inside the  flaring plasma, the secondaries neutrino are not. So pion neutrinos should be even more abundant than gamma ones. Tens-hundred MeV neutrinos may cross undisturbed the whole Sun, doubling at least their rate respect a unique solar-side for gamma flare. Therefore we obtain minimal bounds opening a windows for neutrino astronomy, already at the edge  of present but \\emph{quite within near future  Megaton neutrino detectors}. Such detectors are considered mostly to reveal cosmic supernova background or rare Local Group  (few Mpc) Supernovas events \\cite{ref3}. However  rarest (once a decade), brief (a few minutes) powerful solar neutrino \"flare\"  may shine and they may overcome by two to three order of magnitude the corresponding steady atmospheric neutrino noise on the Earth, leading in largest Neutrino detector at least to one or to  meaning-full few events clustered signals. The voice of such a solar anti-neutrino flare  component at a few tens MeVs may induce an inverse beta decay over a $vanishing$  anti-neutrino solar background. Megaton or even   inner ten Megaton Ice Cube detector at ten GeV threshold may also reveal  traces in hardest energy of solar flares. Icecube , marginally, too. Solar neutrino flavors may shine light on neutrino mixing angles. Not only on orbit satellites but even human  astronauts in Space   may exploit underground neutrino detectors for the prompt alert on (otherwise) fast and maybe lethal solar explosions. \\vspace{1pc} ",
        "introduction": "During largest solar flare, of a few minutes duration,  the particle flux escaping the corona eruption and hitting later on the Earth, is 3-4 order of magnitude above the common atmospheric CR background . If the flare particle interactions on the Sun corona is taking place as efficiently as  in terrestrial atmosphere, than their secondaries by charged pions and muons decays, are leading to a neutrinos fluency on Earth comparable to one day terrestrial atmospheric neutrino activity (upper Bound). One therefore may expect a prompt increase of neutrino signals of the order of one day integral events made by atmospheric neutrinos \\cite{ref0}. In present neutrino detectors the signal is just on the edge, but as long as the authors know, it has been never revealed \\cite{ref1}.  Sun density at the flare corona might be diluted and pion production maybe consequently suppressed (by a factor (0.1-0.05)). This may be the reason for the SK null detection. Indeed the low Gamma signals recently reported \\cite{ref2} confirm this suppressed signal, but just at the detection edge (See Fig.\\ref{01},\\ref{02}). Unfortunately the neutrino signal at  hundred MeV energies is rare while the one at ten MeV or below is polluted by Solar Hep neutrinos.  The expected signal is dominated by $10-30$ MeV neutrinos, that  might be greatly improved  by   anti-neutrino component via Gadolinium presence in next SK detectors \\cite{ref1}.  Our earliest (2006) upper limit estimate for  October - November $2003$ solar flares \\cite{ref0} and  the recent January $20th$ 2005 \\cite{ref1} exceptional flare were  leading to signal near unity for  Super-Kamiokande and to a few events above   unity for Megaton detectors (See Fig.\\ref{01},\\ref{02}). Our present estimate based on a lower (gamma) bound of neutrino signal, while below the upper ones, still confirms the order of magnitude and the near edge discovery (See Fig.\\ref{01},\\ref{02}). The recent peculiar solar flares as the October-November $2003$ and January $2005$ \\cite{ref2} were source of high energetic charged particles: large fraction of these {\\itshape{primary}} particles, became a source of both neutrons \\cite{ref1} and {\\itshape{secondary}} kaons, $K^{\\pm}$, pions, $\\pi^{\\pm}$ by their particle-particle spallation on the Sun surface \\cite{ref0}. Consequently, $\\mu^{\\pm}$, final secondaries muonic and electronic neutrinos and anti-neutrinos, ${\\nu}_{\\mu}$, $\\bar{\\nu}_{\\mu}$, ${\\nu}_{e}$, $\\bar{\\nu}_{e}$, $\\gamma$ rays, are released by the chain reactions $\\pi^{\\pm} \\rightarrow \\mu^{\\pm}+\\nu_{\\mu}(\\bar{\\nu}_{\\mu})$, $\\pi^{0} \\rightarrow 2\\gamma$, $\\mu^{\\pm} \\rightarrow e^{\\pm}+\\nu_{e}(\\bar{\\nu}_{e})+ \\nu_{\\mu}(\\bar{\\nu}_{\\mu})$ occurring on the sun atmosphere. There are two different sites for these decays (see \\cite{ref0}): A brief and sharp solar flare, originated within the $solar$ corona itself  and a diluted and delayed $ terrestrial$ neutrino flux, produced by late flare particles hitting the  Earth's  atmosphere. This latter delayed signal is of poor physical interest, like an inverse missing signal during the E. Forbush  phase. The main and first {\\itshape{solar}} flare neutrinos reach the Earth with a well defined directionality and within a narrow time range. The corresponding average energies $<E_{{\\nu}_{e}}>$, $<E_{{\\nu}_{\\mu}}>$  ( since  low solar corona densities) suffer negligible energy loss: $<E_{{\\nu}_{e}}> $ $\\simeq$ $50 MeV$, $<E_{{\\nu}_{\\mu}}> \\simeq$ 100 $\\div$ 200 MeV. The opposite occur to downward flare. In the simplest approach, the main source of pion production is $p+p\\rightarrow {{\\Delta}^{++}}n\\rightarrow p{{\\pi}^{+}}n$; $p+p\\rightarrow{{{\\Delta}^{+}}p}^{\\nearrow^{p+p+{{\\pi}^{0}}}}_{\\searrow_{p+n+{\\pi}^{+}}}$ at the center of mass of the resonance ${\\Delta}$ (whose mass value is ${m}_{\\Delta}=1232$ MeV). As  a first approximation and as a useful simplification after the needed boost of the secondaries energies  one may assume that the total pion $\\pi^+ $ energy is equally distributed, in average, in all its final remnants: ($\\bar{\\nu}_{\\mu}$, ${e}^{+}$, ${\\nu}_{e}$, ${\\nu}_{\\mu}$):${E}_{{\\nu}_{\\mu}} \\geq {E}_{{\\bar{\\nu}_{\\mu}}} \\simeq {E}_{{\\nu}_{e}} \\simeq \\frac{1}{4}{E}_{{\\pi}^{+}}$. Similar nuclear reactions (at lower probability) may also occur by proton-alfa scattering leading to: $p+n\\rightarrow {{\\Delta}^{+}}n\\rightarrow n{{\\pi}^{+}}n$; $p+n\\rightarrow {{{\\Delta}^{o}}p}^{\\nearrow^{p+p+{{\\pi}^{-}}}}_{\\searrow_{p+n+{\\pi}^{o}}}$. Here we neglect the ${\\pi}^{-}$ additional role due to the flavor mixing and the dominance of previous reactions ${\\pi}^{+}$. To a first approximation the flavor oscillation will lead to a decrease in the muon component and it will make the electron neutrino component a bit harder. Indeed the oscillation length (at the energy considered) is small with respect to the Earth-Sun distance: $ L_{\\nu_{\\mu}-\\nu_{\\tau}}=2.48 \\cdot10^{9} \\,cm \\left( \\frac{E_{\\nu}}{10^{9}\\,eV} \\right) \\left( \\frac{\\Delta m_{ij}^2 }{(10^{-2} \\,eV)^2} \\right)^{-1} \\ll D_{\\oplus\\odot}=1.5\\cdot 10^{13}cm$.  While at the birth place the neutrino fluxes by positive charged pions $\\pi^+$ are $\\Phi_{\\nu_e}$:$\\Phi_{\\nu_{\\mu}}$:$\\Phi_{\\nu_{\\tau}}$ $= 1:1:0$, after the mixing assuming a democratic number redistribution we expect $\\Phi_{\\nu_e}$:$\\Phi_{\\nu_{\\mu}}$:$\\Phi_{\\nu_{\\tau}}$ $= (\\frac{2}{3}):(\\frac{2}{3}):(\\frac{2}{3})$. Naturally in a more detailed balance the role of the most subtle and hidden parameter (the neutrino mixing  $\\Theta_{13}$)  may be deforming the present averaged  flavor balance. On the other side for the anti-neutrino fluxes we expect at the birth place: $\\Phi_{\\bar\\nu_e}$:$\\Phi_{\\bar\\nu_{\\mu}}$:$\\Phi_{\\bar\\nu_{\\tau}}$ $= 0:1:0$ while at their arrival (within a similar democratic redistribution): $\\Phi_{\\bar\\nu_e}$:$\\Phi_{\\bar\\nu_{\\mu}}$:$\\Phi_{\\bar\\nu_{\\tau}}$ $= (\\frac{1}{3}):(\\frac{1}{3}):(\\frac{1}{3})$. This neutrino  flux, derived by gamma one, hold $100$ s  duration and it is larger by two order of magnitude over the atmospheric one. ",
        "conclusions": ""
    },
    "0812/0812.4405.txt": {
        "abstract": "We use extensive new observations of the very rich $z \\sim 0.4$ cluster of galaxies A851 to examine the nature and origin of starburst galaxies in intermediate-redshift clusters. New {\\em HST} observations, \\tfm {\\em Spitzer} photometry and ground-based spectroscopy cover most of a region of the cluster about 10\\arcmin\\  across, corresponding to a clustercentric radial distance of about 1.6 Mpc. This spatial coverage allows us to confirm the existence of a morphology-density relation within this cluster, and to identify several large, presumably infalling, subsystems. We confirm our previous conclusion that a very large fraction of the starforming galaxies in A851 have recently undergone starbursts. We argue that starbursts are mostly confined to two kinds of sites: infalling groups and the cluster center. At the cluster center it appears that infalling galaxies are undergoing major mergers, resulting in starbursts whose optical emission lines are completely buried beneath dust. The aftermath of this process appears to be proto-S0 galaxies devoid of star formation. In contrast, major mergers do not appear to be the cause of most of the starbursts in infalling groups, and fewer of these events result in the transformation of the galaxy into an S0. Some recent theoretical work provides possible explanations for these two distinct processes, but it is not clear whether they can operate with the very high efficiency needed to account for the very large starburst rate observed. ",
        "introduction": "Thirty years of challenging observational work has confirmed the existence of rapid recent evolution of the star formation rate of galaxies in rich clusters (Butcher \\& Oemler 1978, BO), and has also provided support for BO's hypothesis that this evolution was caused by the transformation of spiral into S0 galaxies. However, despite this progress, no consensus has emerged concerning the nature of the process or processes driving the evolution of the spiral galaxy population. BO hypothesized that spirials faded after ram pressure stripping by the intracluster medium removed their gas supply. However, Dressler \\& Gunn (1983) soon showed that many of the blue cluster galaxies were undergoing starbursts, suggesting the working of processes more violent than the mere fading of star formation. Dressler \\& Gunn suggested that, rather than stripping the interstellar medium from a galaxy, the cluster ram pressure induced a starburst which consumed the galaxy's gas.  Since then, a number of other processes have been suggested, including starvation of a galaxy's star formation by the removal of the outer gas halos that are posited to surround spirals and replenish the disk gas by infall (Larson, Tinsley \\& Caldwell 1980), tidal shocks on a galaxy, either due to the cluster core (e.g. Byrd \\& Valtonen 1990, Henriksen \\& Byrd 1996), to unvirialized subclusters (Gnedin 2003), or to other galaxies (Richstone \\& Malmuth 1983, Icke 1985, Moore \\etal (1996), and galaxy-galaxy mergers (Dressler \\etal 1999, D99, van Dokkum \\etal 1999, Struck 2006).  These processes may be categorized in several ways, including method of gas removal, star formation history, and dependance on environment. In stripping and starvation, gas content and star formation rate go hand-in-hand; as gas is swept from a galaxy or the external supply is shut off, there is a monotonic decrease in star formation rate. On the other hand, mergers, ram pressure induced star formation, and tidal encounters (including harassment, the cumulative effect of many weak encounters [Moore \\etal 1996]) work at least in part by a temporary {\\em increase} in star formation rate, i.e. by a starburst which can consume much or most of the gas. The processes also differ in their dependance on environment. All rely on interactions of a galaxy with its surroundings, but harassment, ram pressure induced star formation and stripping work only in the hot, dense environment of a rich cluster core, while mergers are expected to be {\\em least} effective in such a hot environment.  One might suppose that distinguishing between these very different processes would be easy, but this has not been so.  A less-than-complete survey of 25 recent papers on the subject shows that 8 papers support starvation as the explanation of cluster-driven evolution, 4 papers support mergers, 3 support tidal encounters, 2 favor stripping, 2 favor ram pressure induced starbursts, 1 supports harassment, 3 favor a combination of mergers plus stripping, 1 favors stripping plus starvation and 1 favors stripping plus tidal encounters. Obviously, a consensus is not at hand.  In several previous papers (Poggianti \\etal 1999, P99, Dressler \\etal 2004, D04) we have presented evidence supporting the importance of starbursts in cluster galaxy evolution. (Definitions of ``starbursts'' vary. We shall mean by this term a significant temporary increase in the star formation rate of a galaxy above its long-term past average.) A companion paper to the present one (Dressler \\etal 2009, hereafter D09) uses $24 \\micron$ {\\em Spitzer} observations of A851 to further elucidate the connection between the optical spectral properties of cluster galaxies and the nature of the ongoing starbursts. D09 show that  galaxies with strong Balmer absorption lines and \\OII emission, which P99 and D04 have previously identified as starbursts, are indeed young starbursts buried under appreciable dust extinction, while galaxies with strong Balmer absorption lines but no \\OII emission, which P99 and D04 identify as post-starburst galaxies, divide into two very different classes. The majority are true post-starbursts, but a significant minority are declining starbursts in which optical emission from the remaining OB stars is completely hidden beneath dust. Together, these papers demonstrate that starbursts are a widespread phenomenon in galaxies at earlier epochs, both within and outside of clusters. They also show that the post-starburst history of many cluster galaxies is very different than that of field galaxies: most field galaxies apparently resume normal star formation after a burst but many cluster galaxies do not. These facts suggest that a combination of starbursts plus an additional, cluster-specific process form at least one of the mechanisms responsible for the evolution of the star formation in cluster galaxies. If correct, more information on the environmental dependance of the phenomenon would clearly be useful for pinpointing the cluster-specific post-starburst process, as well as for elucidating the cause of the starbursts.  In this paper we present new {\\em HST} imaging and ground-based spectroscopy of the cluster Abell 851, at a redshift $z = 0.41$. A851 is not a ``typical'' cluster; it is richer, has more substructure, and has a larger population of starbursting and disturbed galaxies than the average $z = 0.4$ cluster. However, there is every reason to think that the processes occurring at the time of observation in A851 typify those occurring--- at a somewhat lower rate--- in all clusters at that epoch. It therefore represents a particularly useful laboratory for studying the evolution of cluster galaxies.  The data presented here cover a much wider area than have most previous observations of intermediate redshift clusters, allowing us to probe further out into the cluster, beyond the cluster core and into the region where one might expect newly infalling galaxies and groups to predominate.  Thus they provide significant new information on the environmental dependance of the evolutionary processes, and the larger data set and the addition of {\\em Spitzer} IR photometry also allows us to strengthen our previous conclusions about the origin and role of starbursts. We will demonstrate that a majority of the galaxies in A851 with on-going or recent star formation have undergone significant starbursts, and that many are  associated with unmistakable mergers. The most recent starburst and merger events are concentrated in several kinematic and/or spatial subcomponents of the cluster, some of which we will suggest are currently infalling subclusters.  The paper is organized as follows.  In \\S2 we describe the data used in this paper, and in \\S3 we discuss the structure of the cluster derived from spatial and velocity distributions.  In \\S4 we revisit the morphology/density relation for A851.  In \\S5 we discuss the analysis of the spectra of A851 members, correlate their properties with other galaxy and cluster parameters, and compare them with those of field galaxies at the same epoch.  In \\S6 we examine the morphology and spatial distribution of the starbursting population, and in \\S7 try to deduce from these the origin of the starbursts. ",
        "conclusions": "We can summarize what we now know about A851 as follows:  \\begin{enumerate}  \\item A851 is a complex system of merging subclusters, many with their own hot X-Ray gas content.  \\item Despite this, A851 shows a well-established morphology-density relation, with early type galaxies concentrated in the densest lumps.  \\item In two infalling groups, an extremely high fraction of the starforming galaxies are undergoing starbursts. The starburst rate in the main cluster is significantly lower, except in the very center, but still much higher than in the field at the same epoch.  \\item The ratio of post-starburst objects to those currently undergoing starbursts is much higher in the main cluster than in either the infalling groups or the field.  \\item Most of the starbursts in the core of the main cluster are young, are buried behind very high dust extinction, and occur in galaxies which appear to be undergoing mergers . In contrast, these characteristics apply to a minority of the group and field starbursts.  \\item There is a strong radial gradient in the age of the starbursts in the main cluster. The youngest occupy the very center, with older starbursts and post-starbursts occurring at progressively greater cluster-centric distances.  \\end{enumerate}  These facts suggest that there are at least two causes of starbursts. In the cluster center, starbursts are the result of galaxy-galaxy mergers, or more precisely {\\em major} mergers in which the mass ratio between the two galaxies is less than a few, since only such events produce the morphological signatures seen in Fig.11. For reasons mentioned above, such major mergers are likely to be quite effective at removing a galaxy's gas, and therefore provide a simple explanation for the gas-free post-starburst objects found mostly in clusters. Furthermore, if Struck's (2005) arguments are to be believed, the cores of rich clusters provide an unexpectedly hospitable environment for such major mergers. Such a core-centered process explains the radial gradient in the ages of starbursts seen in A851.  There are two worries about this otherwise satisfying picture. Firstly, it is commonly believed that major mergers result in bulge-dominated galaxies, not disky objects like the post-starburst objects seen in Fig. 11. Secondly, Stuck's mechanism, at least, requires that the bursting galaxies be members of {\\em groups}. Given the small number of objects involved at any one time, one would expect them to be members of one or two groups. If the group(s) is infalling perpendicular to the line of sight one would expect a small velocity dispersion; if infalling along the line of sight one would expect a small velocity dispersion and a large velocity offset from the main cluster. Neither expectation is met by the central starbursts, whose velocity distribution is very similar to that of the entire cluster. These are serious problems. Nevertheless, the morphological evidence for major mergers is rather compelling, and the location of the bursts and the radial gradient in the ages of the bursting and post-bursting galaxies strongly points to a process sited in the cluster core.  The other primary site of starbursts appears to be infalling groups. The starbursting group members, while often peculiar, show much less compelling evidence for major mergers than do the core objects. Furthermore, group starbursts seems to produce gas-free post-starburst remnants much less frequently than do core starbursts. The absence of ram-pressure stripping in these low-density groups may be a partial explanation for this latter difference, but the available facts are consistent with a different mechanism driving starbursts in this environment. Either minor mergers, with smaller mass ratios, tidal encounters, or the cluster tides proposed by Bekki may provide a mechanism which is more consistent with the group starburst characteristics. Bekki's mechanism is particularly appealing because it can explain why the starburst rate in infalling groups is substantially higher than in the field- which is, after all, composed mostly of groups.  A major question is the extent to which the processes observed in A851 are typical of other intermediate-redshift clusters of galaxies. It cannot be ignored that a substantial fraction of the 25 recent papers mentioned in   \\S1 find evidence for different processes at work. For example, the very extensive studies of $z \\sim 0.4$ clusters by Moran \\etal (2005, 2007) provide compelling evidence for a gradual fading of star formation in some cluster galaxies. A851 is, admittedly, a rather unusual cluster, with more substructure and more disturbed galaxies than most. Nevertheless, astrophysical processes should be universal. It is reasonable to expect that what is observed occurring in the core of A851, and in its infalling groups, has also occurred during some phases of the building of all rich clusters. Whether the average rate at which these processes occur in most clusters is sufficient to entirely explain the transformation of spiral to S0 galaxies which characterizes the recent evolution of cluster galaxies, or whether it is only one of a number of processes working in concert is a question which only further observations of other clusters can answer."
    },
    "0812/0812.0558_arXiv.txt": {
        "abstract": "{Our understanding of the chemical evolution (CE) of the Galactic bulge requires the determination of  abundances in large samples of giant stars and planetary nebulae (PNe). Studies based on high resolution spectroscopy of giant stars in several fields of the Galactic bulge obtained with very large telescopes have allowed important progress.} {We discuss PNe abundances in the Galactic bulge and compare these results with those presented in the literature for giant stars.} {We present the largest, high-quality data-set available for PNe in the direction of the Galactic bulge (inner-disk/bulge). For comparison purposes, we also consider a sample of PNe in the Large Magellanic Cloud (LMC). We derive the element abundances in a consistent way for all the PNe studied. By comparing the abundances for the bulge, inner-disk, and LMC, we identify elements that have not been modified during the evolution of the PN progenitor and can be used to trace the bulge chemical enrichment history. We then compare the PN abundances with abundances of bulge field giant.} {At the metallicity of the bulge, we find that the abundances of O and Ne are close to the values for the interstellar medium at the time of the PN progenitor formation, and hence these elements can be used as tracers of the bulge CE, in the same way as S and Ar, which are not expected to be affected by nucleosynthetic processes during the evolution of the PN progenitors. The PN oxygen abundance distribution is shifted to lower values by 0.3~dex with respect to the distribution given by giants. A similar shift appears to occur for Ne and S. We discuss possible reasons for this PNe-giant discrepancy and conclude that this is probably due to systematic errors in the abundance derivations in either giants or PNe (or both). We issue an important warning concerning the use of \\emph{absolute} abundances in CE studies.} {} ",
        "introduction": "\\label{sec:intro}  The Galactic bulge is old, distant, and highly obscured and studies of its chemical composition include two types of stars that represent evolved stages of intermediate mass stars:  planetary nebulae (PNe) (e.g. G\\'orny et al. 2004), for which abundances are derived using the intensities of conspicuous emission lines, and giant stars (e.g. Lecureur et al. (2007) and Zoccali et al. (2008)), whose luminosities enable spectra of sufficient quality suitable for abundance determination to be obtained, using very large telescopes.  However, the use of evolved stars as test particles for probing the chemical evolution of the Galactic bulge requires some caution, since it is known that some elements have their abundances modified during stellar evolution. For example, intermediate-mass stars compete with massive stars in the enrichment of  the interstellar medium (ISM) in N and C (e.g. Chiappini et al. 2003, Henry 2004), and contribute to the enrichment of He. However, other elements (such as S and Ar) are unaffected by nucleosynthetic processes during the evolution of intermediate-mass stars and probe the chemical composition of the ISM at the time when the stars were born. The status of O and Ne is less clear, since their abundances can be affected by nucleosynthesis and mixing during stellar evolution, in amounts that  depend on metallicity and other properties (Charbonnel 2005,  Leisy \\& Dennefeld 2006, Pe\\~na et al. 2007).  In this paper, we use data sets both for PNe and giant stars to discuss the chemical evolution of the Galactic bulge. For PNe, we use a high quality sample obtained by merging the data from G\\'{o}rny et al. (2008) with those from G\\'{o}rny et al. (2004) and Wang \\& Liu (2007). All PN abundances were recomputed in a consistent way. A detailed comparison of our sample with other bulge PN samples in the literature was presented by G\\'orny et al. (2008). Our bulge PN sample constitutes the largest high-quality data-set of abundance measurements for bulge PNe available in the literature.  We complete a detailed analysis of the chemical abundances of these PNe with two main goals. First, we compare the properties of the PN population in an old, metal-rich\\footnote{The bulge appears to be metal-rich with respect to the halo and thick disk (the other two old components of the Milky Way). Within the central region, the bulge contains cold gas from which stars form until today.} environment (the Galactic bulge) with those in both metal-poor (LMC) and metal-rich (inner-disk) environments with ongoing star formation. We are then able to gain insights into processes occurring during the evolution of low- and intermediate-mass stars that can affect the abundances observed in the PN stage (such as hot-bottom burning and dredge-up), and infer their dependence on stellar mass and metallicity. Second, we attempt to improve our understanding of the formation and evolution of the Galactic bulge by studying elements that have not been modified during the evolution of the PN progenitor and, hence, can be compared with chemical evolution model predictions. After these elements have been safely identified, we compare the bulge PNe abundances with those of bulge giant stars quoted in the literature.  This paper is organized as follows. In Sects.~\\ref{sec:PN} and ~\\ref{sec:Stars}, we describe the PN and giant star samples on which we based our study, and discuss the main uncertainties in their abundance determinations. In Sect.~\\ref{sec:mixing}, we review current ideas about the main mixing processes occurring in low- and intermediate mass stars. In Sect.~\\ref{sec:PNeCE} we compare the properties of PNe in different environments. We present compelling evidence that both O and Ne are reliable chemical-evolution tracers for metal-rich, old populations, in contrast to their unreliability for metal-poor systems. In Sect.~\\ref{sec:StarsPNeCE}, bulge PN abundances are compared with field-giant bulge abundances. We start this section by recalling the recent results on the bulge formation obtained from the study of stars. We approach the question of whether PNe and giant stars provide similar answers concerning the bulge chemical evolution and discuss the biases involved. Section~\\ref{sec:conclusions} contains a summary of our main results. ",
        "conclusions": "\\label{sec:conclusions} Zoccali et al. (2008) compared the metallicity distribution of about 400 clump and giant stars in Baade's window, with that of about 200 giants at b$=- $6$^\\circ$ and b=$-$12$^{\\circ}$. The metallicity distributions exhibited a gradient in stellar populations on the metal-rich side, such that in Baade's window there is a metal-rich component at [Fe/H]$\\sim$+0.3, which becomes less evident at b$=-$6$^\\circ$. This result is based on iron abundances, and for oxygen the differences could be smaller (see below). In a preliminary kinematical study, Gom\\'ez et al. (in preparation) find a higher velocity dispersion for the metal-rich component by $\\sim$~20km/s, which could be interpreted as an indication of a different stellar population towards the inner regions.  Given that the Baade's window field seems to be more contaminated by the [Fe/H]$\\approx$+0.3 component mentioned above than the field at b=$-$6, and given that our bulge PN sample are projected across a wider area (see Fig.~\\ref{lb}), it is unclear in what proportion this new component is present in our PN sample. If we attempt to explain the 0.3~dex difference in oxygen (and probably in Ne and S as well) in terms of the properties of the stellar populations, we may conclude that our PN sample is essentially free of this metal-rich component. However, since [O/Fe] decreases with [Fe/H], the oxygen content in this metal-rich population will probably resemble that of the more metal-poor population.  To confirm that the metal-rich population in the Baade's window follows the same [O/Fe] relation as shown in Fig.~\\ref{figbulge}, it would be necessary to measure oxygen in the same stars. This would require very-high resolution (R=40000) data, such as that acquired for the UVES sample of Zoccali et al. (2006) for all of the metal-rich population found by Zoccali et al. (2008) in the Baade's window.  Checking for any difference in the V$_{lsr}$ velocity dispersion of our bulge PNe with b $<$ 4$^{\\circ}$ (88 objects) and b $>$ 4$^{\\circ}$ (78 objects), we measured 121~km/s and 101~km/s, respectively. Interestingly, the difference of 20~km/s is similar to that found for bulge giants in Baade's window with respect to other fields. However, for PNe this difference is not significant according to our statistical tests.  Finally, we note that significant samples of bulge objects are now available for which the metallicity distribution can be studied (RGB stars, red clump giants -- Zoccali et al. 2008 -- and in PNe -- this work). This opens the possibility of looking for additional constraints to stellar evolution models of low- and intermediate-mass stars, which would enhance our understanding of how metallicity and mass loss affect the metallicity distributions at different stellar evolution stages. It is also important to understand another discrepancy measured in the bulge: the metallicities of lensed turnoff stars (i.e. turnoff stars observed due to the amplification of their brightness by gravitational microlensing) appear to be systematically higher than giants, according to Cohen et al. (2008).\\\\  Our main conclusions are summarized below:\\\\  We have compared the properties of PNe in different systems (bulge, inner-disk, and LMC) by using the largest homogeneous sample of PN abundances presently available. We find that:  \\begin{itemize} \\item The Galactic bulge and inner-disk PN distributions of O, Ne, S, and Ar are shifted to higher values than those of the LMC, indicating that the bulge and inner-disk are more metal rich than the LMC (a result already known from other stellar studies). \\item Oxygen and neon in bulge PNe are close to their ISM values at the time of the PN progenitor formation, and hence can also be used as tracers of the bulge chemical evolution. \\item In LMC PNe, both oxygen and neon have been modified during the evolution of the PN progenitor. \\item Differences in metallicity appear to play a more dominant role in the mixing processes occurring in low- and intermediate-mass stars than the differences in the mass range of PN progenitors. \\end{itemize}  After identifying reliably the bulge PN abundances that can be used as tracers the bulge chemical evolution (O, Ne, Ar, S), we compared these abundances with those measured in giant stars. We found that:  \\begin{itemize} \\item The oxygen abundance distribution of bulge giant stars is shifted to higher values by 0.3~dex with respect to that of PNe. \\item A similar shift appears to exist for Ne and S (after converting the Mg and Si abundances of giant stars to those of Ne and S by adopting the solar values of Asplund et al. (2005), although this method is rather uncertain). \\item We discussed many reasons for the discrepancy between the abundances of PNe and giant stars, and concluded that the oxygen abundances in PNe (distributed over the entire bulge) and in giants (most of them in the Baade's window) do not convey the same evolutionary story. \\end{itemize}  After a thorough analysis, we conclude that the observed discrepancy between PN and giant star abundance distributions is probably due to systematic errors in the abundance derivations of either PNe or giant stars, or both. Our results constitute at least an important warning against a careless use of absolute abundances."
    },
    "0812/0812.0591_arXiv.txt": {
        "abstract": "Within the next decade, ground based gravitational wave detectors are in principle capable of determining the compact object merger rate per unit volume of the local universe to better  than $20\\%$ with more than $30$ detections. These measurements will constrain our models of stellar, binary, and star cluster evolution in the nearby  present-day and ancient universe. We argue that the stellar models are sensitive to heterogeneities (in age and metallicity at least) in such a way that the predicted merger rates are subject to an additional 30-50\\% systematic errors unless these heterogenities are taken into account. Without adding new electromagnetic constraints on massive binary evolution or relying on more information from each merger (e.g., binary masses and spins), as few as the $\\simeq 5$ merger detections could exhaust the information available in a naive comparison to merger rate predictions. As a concrete example immediately relevant to analysis of initial and enhanced LIGO results, we use a nearby-universe catalog to demonstrate that no one  tracer of stellar content can be consistently used to constrain merger rates without introducing a systematic error of order $O(30\\%)$ at 90\\% confidence (depending on the type of binary involved). For example, though binary black holes typically take many Gyr to merge, binary neutron stars often merge rapidly; different tracers of stellar content are required for these two types. \\optional{Specifically, we argue that though the cumulative blue light can be measured very accurately (at present, to within $30\\%$ in the local universe and to within \\editremark{20\\%} at large distances), systematic uncertainty associated with improperly incorporating emission from elliptical galaxies and other old populations could lead to a systematic error in any comparison between \\emph{merger rates alone} and stellar content of order  40\\%. } More generally, we argue that theoretical binary evolution can depend sufficiently sensitively on star-forming conditions -- even assuming no uncertainty in binary evolution model --  that the \\emph{distribution} of star forming conditions must be incorporated to reduce the systematic error in merger rate predictions below roughly $40\\%$. We emphasize that the degree of sensitivity to star-forming conditions depends on the binary evolution model and on the amount of relevant variation in star-forming conditions.  For example, if after further comparison with electromagnetic and gravitational wave observations future population synthesis models suggest all BH-BH binary mergers occur promptly and therefore are associated with well-studied present-day star formation, the associated composition-related systematic uncertainty could be lower than the pessimistic value quoted above. Further, as gravitational wave detectors will make available many properties of each merger  -- binary component masses, spins, and even short GRB associations and  host galaxies could be available -- many detections can still be exploited to create high-precision constraints on binary compact object formation models. ",
        "introduction": "Ground based gravitational wave detector networks (LIGO, described in \\citet{gw-detectors-LIGO-original}; and VIRGO, at the Virgo project website   \\nocite{gw-detectors-VIRGO-website} {\\tt www.virgo.infn.it}) are analyzing the results of a design-sensitivity search for the signals expected from the inspiral and merger of double compact binaries (here, NS-NS, BH-NS, BH-BH) \\citep[extending,for example, the search in][]{LIGO-Inspiral-s3s4}.  Sensitivity improvements in LIGO and other interferometers are expected over the next decade that will make multiple detections a near certainty.  For example, based on the short lifetime of the very massive black hole X-ray binary IC-10 X-1, \\cite{bhrates-Chris-IC10-2008} predict that even the current generation of interferometer has a good chance of detecting a (high-mass) merger.  With moderate improvements in detector sensitivity that will be in place by mid-2009 (``enhanced LIGO''), multiple detections are plausible.  More conservatively, theoretical calculations which explore a wide range of still-plausible assumptions [\\cite{PSgrbs-popsyn} (\\abbrvPSgrbs) and \\cite{PSellipticals} (\\abbrvPSellipticals)] predict that the advanced LIGO network is likely to detect several tens of mergers per year, allowing the merger rate per unit volume to be determined in principle to within $20\\%$. In fact, advanced LIGO can determine the merger rate per unit volume significantly more precisely \\citep[20\\%; see,e.g.][(\\abbrvCompanion)]{PSellipticals-Errors} than measurements have constrained the star formation history of the Milky Way and the distant universe \\citep[often at least $30\\%$; see,e.g.][and references therein]{2006ApJ...648..987P,sfr-HopkinsTrentham-MassLightTracers-Consistency-2008}, against which the predictions of \\cite{PSellipticals} and other theoretical models are normalized. As theoretical predictions can be no more precise than their input, even though a large number of merger detections are likely, advanced LIGO measurements cannot distinguish between different hypotheses about how merging binaries are produced if those merger rates differ by less than $O(30\\%)$, on the basis of the number of mergers alone.   With the ability to measure both the number and properties of merging compact binaries, LIGO has long been expected to provide invaluable assistance in better-constraining  hypotheses regarding compact binary formation \\cite[see,e.g.,][and references therein]{2008ApJ...676.1162S}. Given a systematic error of $\\varepsilon$ in any merger rate prediction, only approximately $1/\\varepsilon^2$ unique detections are needed to determine if reality is consistent with a model. Even optimistically assuming that three types of binaries can be distinguished (e.g., from their component masses) and their rates estimated at this level of accuracy, the fraction of \\emph{a priori} plausible models still consistent with the three merger rates provided by LIGO is comparatively large: $\\simeq [(\\log 1+\\epsilon)/2]^3$ (assuming two orders of magnitude uncertainty in the a priori plausible merger rate for each type of binary), at best comparable with the miniscule fraction of parameter space needed to constrain a high-dimensional model weakly.  For example, for the seven-dimensional binary evolution models compared with observations of double pulsars in \\cite{PSmoreconstraints}, each parameter would be constrained to just $[(\\log 1+\\epsilon)/2]^{3/7}\\simeq 0.3$ of its a priori range using merger rate estimates alone, comparable (albeit complementary) to the information provided by electromagnetic observations of double pulsars. On the contrary, had systematic errors been smaller, then detection of $n$ binaries of each of 3 types should imply an accuracy $(2\\ln 10)^{-3/7}n^{-3/14}\\simeq 0.2(n/100)^{-3/14}$ in each parameter. Furthermore, since this systematic uncertainty is introduced through our lack of knowledge about the nearby and ancient universe, even though third generation detectors such as the \\href{www.et-gw.eu}{Einstein Telescope} \\nocite{gw-detectors-ET-website} will harvest vastly more mergers, they will be similarly limited when comparing their observed merger rates with theoretical models that rely upon existing surveys of star formation. To take full advantage of the many mergers that in-construction and third-generation instruments will detect, compact-object theorists will need to compare the   distributions of binary parameters expected from theory (i.e., masses, spins) with observations.   In this paper we  estimate the limiting systematic error introduced into any theoretical prediction of binary compact object merger rates through the star formation history of the universe. We furthermore explain that the relevant uncertainty is not merely overall normalization of the nearby and even distant star formation history.  Instead, we argue that merger rates, particularly binary black hole merger rates, can  also be sensitive to the correlated distribution of age and metallicity of their progenitor star-forming regions. Though high-precision surveys and spectral energy distribution (SED) reconstructions of galaxies may precisely determine the mean star formation rate and metallicity by the epoch of advanced LIGO, \\citep[see,e.g.][and references therein]{araa-Renzini-2006,sfr-MeasurementsToSFR-Hopkins2003SDSSReview,sfr-HopkinsTrentham-MassLightTracers-Consistency-2008}, the more delicate analyses which estimate the \\emph{distribution} of star forming conditions, particularly those of low metallicity that are far more apt to produce massive black hole binaries, remain in their infancy \\citep[see,e.g.][]{sfr-ZEvolution-ByGalaxy-Panter2008}.  Given that advanced LIGO and future gravitational-wave detectors could observe mergers produced from binary stellar evolution out to as far as $z\\simeq 2$ (e.g., for an optimally oriented $30+30 M_\\odot$ binary black hole merger), an epoch of rapid star formation in massive galaxies, the relevant composition distribution needed to eliminate this systematic uncertainty is unlikely to be available in the near future. Equivalently, gravitational-wave interferometers will soon provide a uniquely accurate and potentially uniquely \\emph{biased} probe into the formation and evolution of high-mass stars in the early and low-metallicity universe.     \\subsection{Outline and relation to prior work} As discussed in more detail in \\S \\ref{sec:Catalog}, to account for local-universe inhomogeneities and to simplify the intrinsically mass-dependent results of gravitational wave searches, previous searches for gravitational wave inspiral and merger waveforms have ``normalized'' their result by the amount of blue light within the relevant time-averaged detection volume; see the discussion in \\cite{LIGO-Inspiral-s3s4} as well as the  considerably more detailed presentation in \\cite{mm-stats-LoudestEvents-Inspiral-FairhurstBrady2007a}.  By choosing to express results as a ``merger rate per unit blue light,''  however, the authors limit the accuracy of any attempt to compare merger rate predictions with their observations: as emphasized above, such a comparison is helpful only to the level of accuracy that ``mergers per unit blue light'' can be uniquely defined. The systematic  error so introduced is unlikely to seriously limit astrophysical comparisons once detections are available in the near future: initial and enhanced LIGO results, expected to have at best a handful of detections, will not reach this level of accuracy. But this composition-based systematic error is comparable to several formal uncertainties often quoted in relation to already-published upper limits \\optional{(\\editremark{pick some numbers}) } and is therefore already relevant to anyone attempting to constrain their models with existing observational upper limits.   In short, anyone planning on using or expressing results in this form should be aware of its limitations.  That being said, at present, gravitational-wave detectors survey only the nearby universe, where  uncertainties  in the distances to galaxies  dominate over photometric errors \\citep[74\\% vs 31\\%, respectively; see][]{LIGO-Inspiral-s3s4-Galaxies}.  At best,  the  cumulative asymptotic luminosity can be determined only to O(10\\%).  All of these uncertainties are comparable or greater than the uncertainty introduced into any astrophysical interpretation by assuming the number of mergers is proportional to blue light.     The uncertainties discussed in this paper therefore \\emph{bound below} the accuracy of any comparison between merger rate predictions and observations.          These limiting uncertainties arise because  most simple  prediction (or ``normalization'') methods build in an implicit assumption of \\emph{homogeneity} of star-forming conditions.  But because binary mergers are rare and exceptional events themselves and naturally arise more frequently from rare and exceptional conditions (e.g., old star formation or low metallicity), assuming homogeneity builds in systematic errors greater than the limiting uncertainty desired for advanced detectors. As outlined above and described in  \\S~\\ref{sec:Catalog},  we argue that dividing the rate of mergers by the amount of blue light in the detection volume oversimplifies the (implicit) inverse problem: predicting how many mergers should occur given an amount of blue light. To demonstrate that other bands give different yet potentially equally relevant normalizations, we introduce a multi-band galaxy catalog for the local universe.  In \\S \\ref{sec:Lags} we demonstrate that, after a starburst, different models of binary evolution and different types of binaries lead to different conclusions about the time-dependent ratio of mergers and light.  We use this tool to estimate the systematic error introduced by normalizing to blue light or, more generally, any single-band normalization. For advanced detectors, galaxy catalogs will not be available.  Nonetheless, as demonstrated in  \\S~\\ref{sec:Hetero} the gravitational wave detection rate should not be cavalierly normalized to the \\emph{mean} properties of the universe on large scales: exceptional circumstances (here, metallicity) can introduce systematic errors at least as large as the limiting uncertainty expected of advanced detectors. Though the exact magnitude of the effect cannot be determined without an equally exact theory of binary evolution, we estimate that even modestly reliable predictions could require fairly detailed input regarding the composition of the universe within reach. Finally, to clearly illustrate the effects summarized in this paper,  in \\S \\ref{sec:Multicomponent} we show that concrete, plausible predictions for a two-component universe cannot be well-modeled by a time-independent or homogeneous one. ",
        "conclusions": "Conclusions} In anticipation of an era of frequent binary coalescence detection and with the goal of divining the limiting astrophysical \\emph{measurement} uncertainties for future observations, in this paper we have examined the relevant  systematic errors intrinsic to proposed \\emph{absolute}  normalizations against which gravitational wave detections and upper limits can be compared.  In other words, we have examined the challenges associated with comparing just the \\emph{number} of binary merger detections with predictions.  We find that after a surprisingly small number of detections, either much more sophisticated models or richer data products (e.g., the observed mass distribution) will be needed to further constrain binary evolution.  For the nearby universe, relevant to initial and enhanced LIGO, we argue that the systematic error associated with using catalog-based normalizations has been  understated. Though the nominal accuracy of tracers of star formation inside a volume, such as blue light as adopted in \\abbrvGalaxy, can be comparable to the  systematic error target in LIGO ($15\\%$),  the relevant systematic error by adopting a normalization that does not trace old mass and remains the same for all binary types and evolution models -- the error introduced into any comparison between the number of detections and predictions  -- will be considerably larger ($\\simeq 40\\%$). This systematic error can be ameliorated but not eliminated by employing model-dependent normalizations. To provide a framework with which to calculate this two-band normalization, we introduced a multi-band galaxy catalog that extends the blue-light catalog presented in \\abbrvGalaxy; see Figures \\ref{fig:Cumulative} and \\ref{fig:MassFraction}. We recommend that this catalog and approach be applied to re-evaluate the astrophysical  systematic errors relevant to initial and enhanced LIGO upper limits.   Advanced detectors will probe the distant universe, for which a catalog is impractical.  Though merger rates can be compared against the \\emph{average} properties of the universe, we have demonstrated that treating the universe as homogeneous will introduce \\emph{at least} a 40\\% systematic error, because regions of different metallicity will have different relative probabilities of producing massive merging binaries.  We emphasize our estimate is conservative, assuming that the only variable in star formation is metallicity (e.g., no top-heavy IMFs or alternate modes of star formation) and that the universe was always homogeneous with a similar metallicity distribution to that observed at present. Because binary black hole detection rates in particular can be strongly influenced by metallicity variations (e.g., due to changes in the initial star-final black hole mass relation with metallicity) and because black holes are far more likely to be produced in the early universe in the epoch of peak star formation  in massive galaxies undergoing rapid metallicity evolution \\citep[binary merger delays for black holes are almost always long; based on results in ][the median merger delay for merging BH-BH binaries given \\emph{steady-state} star formation is $\\tau_{BBH}\\simeq 1-3\\unit{Gyr}$, depending on assumptions, while for NS-NS binaries it is almost always much smaller, $\\tau_{BNS}\\lesssim 0.3 \\unit{Gyr}$]{PSellipticals}, our estimate could significantly understate the relevant systematic uncertainty.     To summarize, we recommend the following:  (I) When interpreting LIGO data as constraints on merger rates,  unless composition distributions are explicitly incorporated into the predictive models, an additional systematic error of order $40\\%$ should be included to allow for fluctuations in composition and age between galaxies; for example, this revised uncertainty  will be used in \\abbrvCompanion{} to explore how advanced LIGO detections might constrain binary merger models.  (II) Future merger rate predictions should include metallicity evolution and distributions, to determine the most likely LIGO detection rates when low-metallicity environments are included.  For example, the forthcoming paper by Belczynski et al. \\nocite{popsyn-ChrisFutureMergerRateMetallicity2009} will  explore evolutionary scenarios over ensemble of metallicities in more detail. (III) To better assess all relevant systematic errors limiting comparisons between models and theory, more observational and theoretical work is needed to constrain the distribution of fluctuations, particularly IMF fluctuations early in the universe or in clustered star formation. (IV) Finally, to provide another handle with which to constrain binary evolution, future model constraint papers should describe how to compare the detected mass distribution with  highly model-dependent predictions. Given the immense computational requirements needed to both thoroughly and accurately explore the space of binary evolution models,  let alone globular clusters, and the relatively modest benefits that Moore's Law provides to a monte carlo simulation sampling a high dimensional space, a careful balance must be struck between accurately pinning down predictions for each model and thoroughly exploring the model space. The parameter-dependent detection efficiency $\\epsilon(D,m_1,m_2)$ and parameter-measurement-ambiguity functions provided in the gravitational-wave literature \\citep[see,e.g.][]{CutlerFlanagan:1994} will tell us how much we can learn about parameter distributions from LIGO and therefore determine where that balance will be struck.       When estimating systematic errors introduced by treating the universe as homogeneous, we have for simplicity assumed all mergers are produced only through binary evolution.   Interactions in globular clusters are expected to be an equally critical channel for forming merging double black hole binaries; see for example \\cite{2008ApJ...676.1162S} and references therein.   Though we have not performed a thorough exploration of parameter space, as we were able to do for binary evolution with \\texttt{StarTrack}, we expect this channel will be at least as sensitive to inhomogeneities as isolated binary evolution. More critically, this  competing channel may produce mergers that are indistinguishable from binary evolution.  The existence of such an unconstrained and often indistinguishable channel introduces yet another large systematic error into interpretation of binary compact object detection rates. Further study is critical, to determine not only the range of rates these models produce in a realistically heterogeneous universe but also methods with which to distinguish the randomly-oriented and equal-mass-biased mergers expected from this channel from the more aligned mergers expected from binary evolution."
    },
    "0812/0812.2846_arXiv.txt": {
        "abstract": "\\small{ For a large class of dark energy (DE) models, for which the effective gravitational constant is a constant and there is no direct exchange of energy between DE and dark matter (DM), knowledge of the expansion history suffices to reconstruct the growth factor of linearized density perturbations in the non-relativistic matter component on scales much smaller than the Hubble distance. In this paper we develop a non-parametric method for extracting information about the perturbative growth factor from data pertaining to the luminosity or angular size distances. A comparison of the reconstructed density contrast with observations of large scale structure and gravitational lensing can help distinguish DE models such as the cosmological constant and quintessence from models based on modified gravity theories as well as models in which DE and DM are either unified, or interact directly. We show that for current SNe data, the linear growth factor at $z=0.3$ can be constrained to $5 \\%$, and the linear growth rate to $6\\%$.  With future SNe data, such as expected from the JDEM mission, we may be able to constrain the growth factor to $2-3\\%$ and the growth rate to $3-4 \\%$ at $z=0.3$ with this unbiased, model-independent reconstruction method. For future BAO data which would deliver measurements of both the angular diameter distance and Hubble parameter, it should be possible to constrain the growth factor at $z=2.5$ to $9\\%$. These constraints grow tighter with the errors on the datasets.  With a large quantity of data expected in the next few years, this method can emerge as a competitive tool for distinguishing between different models of dark energy.  } ",
        "introduction": "Over the last decade, observations of Type Ia supernovae have shown that the expansion of the universe is currently accelerating \\citep{accl,accl1,accl2,accl3,accl4,accl5,accl6,accl7,union}. This remarkable discovery has led cosmologists to hypothesize the presence of dark energy (DE), a negative pressure energy component which dominates the energy content of the universe at present. Many theories have been propounded to explain this phenomenon, the simplest of which is the cosmological constant $\\ld$, with a constant energy density and the equation of state $w=-1$. Although $\\ld$ appears to explain all current observations satisfactorily, to do so its value must necessarily be very small $\\ld/8\\pi G \\simeq 10^{-47}$GeV$^4$. So, it represents a new small constant of nature in addition to those known from elementary particle physics, many of them being very small if expressed in the Planck units.  However, since it is not known at present how to derive $\\ld$ from these small constants and it is also unclear if DE is in fact time independent, other phenomenological explanations for cosmic acceleration have been suggested \\citep[see reviews][]{DE_review,rev2,rev3,rev4,rev5,rev6,ss06}. These are based either on the introduction of new physical fields (quintessence models, Chaplygin gas, etc.), or on modifying the laws of gravity and therefore the geometry of the universe (scalar-tensor gravity, $f(R)$ gravity, higher dimensional `Braneworld' models \\etc). The plethora of competing dark energy models has led to the development of parametric and non-parametric methods as a means of obtaining model independent information about the nature of dark energy directly from observations \\citep [see][and references therein]{star98,DE_recon,recon1,recon2,recon3,recon4,recon5,recon6,recon7,recon8,recon9,sss08,ss06}.  The next decade will see the emergence of many new cosmological probes. A large number of these are likely to make important contributions to the field of dark energy. The Sloan Digital Sky Survey began its stage III observations in 2008, and its Baryon Oscillation Spectroscopic Survey (BOSS) is expected to map the spatial distribution of luminous galaxies and quasars and detect the characteristic scale imprinted by baryon acoustic oscillations in the early universe \\citep{sdss}. The Joint Dark Energy Mission (JDEM) is expected to discover a large number of supernovae, and also provide important data on weak-lensing and baryon acoustic oscillations \\citep{jdem}. The Square Kilometer Array (SKA) will map out over a billion galaxies to redshift of about 1.5, and is expected to determine the power spectrum of dark matter fluctuations as well as its growth as a function of cosmic epoch \\citep{ska}. Important clues to the growth of structure will also come from current and future weak lensing surveys (CFHTLS, DES, JDEM, EUCLID, SKA, LSST), galaxy redshift-space distortions \\citep{guzo,percival, percival2} as well as galaxy cluster mass functions at different redshifts $z$ \\citep{vikh, rap}. With the wealth of data expected to arrive over the next several years, it is important to explore different methods of analyzing these datasets in order to extract the optimum amount of information from them. In this paper we explore the possibility of reconstructing the linearized growth rate of density perturbations in the non-relativistic matter component, $\\delta(z)$, taken at some fixed comoving scale much less than the Hubble distance, from datasets which have traditionally been used to explore only the smooth background universe, \\eg luminosity distance and angular diameter distance data.  In the case of {\\em physical} DE, the effective gravitational constant appearing in the equation for linear density perturbations in the matter component coincides with the Newton gravitational constant $G$ measured in the laboratory and using Solar system tests. If, additionally, there is no direct non-gravitational interaction between DE and DM in the physical reference frame, so that the DE energy-momentum tensor is covariantly conserved, the density contrast reconstructed in this manner should match that determined directly from observations of large scale structure. In this case the methods developed in this paper will provide an important consistency check on DE models such as $\\ld$ and Quintessence. On the other hand, {\\em geometrical} models of DE (Braneworlds, scalar-tensor gravity, etc.) usually predict a different growth rate for $\\delta(z)$ from that in general relativity (GR). Models where DE has a direct non-gravitational interaction with DM, or where DE and DM are unified, have a similar property even in the framework of GR. In this case, a reconstruction of the linearized density contrast from observations of standard candles/rulers will not match with $\\delta$ determined directly from large scale structure \\citep{f_dgp,bert,ishak,knox,chiba,MMG,haiman,hu,pol,koy}. Currently reconstructed values of the growth rate from galaxy redshift distortions \\citep{guzo,sdss2,sdss2-1,slaq} are not very constrictive, but future missions like Euclid \\citep{euclid} are expected to constrain the growth rate tightly. Therefore comparing the results from future supernova data, using the methods described in this paper, to those from future large scale structure data will help address important issues concerning the nature of gravity and dark energy.  This paper is organized as follows. In section~\\ref{meth}, we describe the reconstruction technique and the data used to test this method. Section~\\ref{res} shows the results and examines the dependence of the method on various factors such as the redshift distribution of the data and information on other cosmological parameters. The conclusions are presented in section~\\ref{concl}. ",
        "conclusions": "In this paper we have proposed a method for extracting growth parameters for dark energy models (within the spatially flat FRW universe) from observations that map the background universe, such as measures of the luminosity distance or the angular diameter distance. The method is model independent and unbiased. For current data, the growth factor $g(z)$ may be constrained to $\\sim 5 \\%$ at $z = 0.3$, while the growth rate $f(z)$ is constrained to $\\sim 6 \\%$. For future JDEM SNe data, we will be able to put constraints of the order of a few percent on the growth parameters, \\eg $2\\%$ on the growth factor and $3\\%$ on the growth rate at a redshift of $0.3$ , and $4\\%$ on the growth factor and $8\\%$ on the growth rate at a redshift of unity. In conjunction with the likelihood parameter estimation method, this method acts as an important consistency check on the accuracy of the priors on $\\om$ for SNe. With future probes like JDEM and BOSS taken in conjunction, it will lead to an unbiased estimation of the growth parameters upto a redshift of $z=2.5$.  It is well known that, in GR and for most DE models, the expansion history completely determines the linearized growth rate of density perturbations \\citep{star98,ss06} (the exact conditions for this are formulated at the beginning of Sec. II). Consequently, a comparison of the density contrast reconstructed from the expansion history would provide one more important consistency check for a large variety of DE models including the cosmological constant and quintessence. On the other hand, as explained in more detail at the beginning of Sec. II, any departure of the observed density contrast from that reconstructed using standard candles and rulers would almost certainly indicate that either there is an exchange of energy between DE and DM (so that the effective energy-momentum tensor of DE is not on its own covariantly conserved), or that cosmic acceleration is a consequence of modified, non-Einsteinian gravity. In modified gravity theories, such as Braneworld models, scalar-tensor and $f(R)$ gravity, etc., the linearized perturbation equation for $\\del$ does not follow the Newtonian form, Eq~(\\ref{eq:diff}) \\citep{f_dgp,brane_pert_DGP,brane_pert,fR,beps00,jain,koy,zhao,song1}. Hence the density contrast reconstructed using observations of standard candles/rulers via Eq~(\\ref{eq:int}) and the density contrast determined directly from observations of large scale structure, say, by weak lensing, galaxy redshift distortions or cluster abundances at different $z$ \\citep{guzo,percival,vikh,euclid}, are likely to differ.  In Sec.~\\ref{dgp} we show that one can obtain a strong signature of modified gravity by comparing results of this reconstruction method with future observations of galaxy redshift distortions using the DGP model as a toy example of modified gravity, where the growth factor $g(z)$ is scale-independent on small scales. However, as discussed in Sec. II, $g(z)$ often becomes scale-dependent both in modified gravity and in the case of direct DE--DM interaction (or their unification). Therefore, for further discrimination of DE models alternative to quintessence and the cosmological constant, measurement of $\\del$ at different comoving scales is required to determine if $g(z)$ is scale-dependent or not.  Future surveys such as JDEM are expected to deliver high quality data for both supernovae and weak lensing. Using such surveys it would then be possible to compare the reconstructed density contrast from standard candles (SNe) with the density contrast observed from gravitational clustering (lensing). Therefore, we hope that the techniques developed in this paper, combined with future observations, will help unravel the nature of that most enigmatic quantity -- dark energy."
    },
    "0812/0812.2594_arXiv.txt": {
        "abstract": "We discuss infrared {\\it Spitzer} observations of early type galaxies in the SAURON sample at 24, 60 and 170$\\mu$m. When compared with 2MASS $Ks$ band luminosities, lenticular (S0) galaxies exhibit a much wider range of mid to far-infrared luminosities then elliptical (E) galaxies. Mid and far-infrared emission from E galaxies is a combination of circumstellar or interstellar emission from local mass-losing red giant stars, dust buoyantly transported from the galactic cores into distant hot interstellar gas and dust accreted from the environment. The source of mid and far-IR emission in S0 galaxies is quite different and is consistent with low levels of star formation, 0.02 - 0.2 $M_{\\odot}$ yr$^{-1}$, in cold, dusty gaseous disks. The infrared 24$\\mu$m--70$\\mu$m color is systematically lower for (mostly S0) galaxies with known molecular disks. Our observations support the conjecture that cold dusty gas in some S0 galaxies is created by stellar mass loss at approximately the same rate that it is consumed by star formation, so the mass depletion of these disks by star formation will be slow. Unlike E galaxies, the infrared luminosities of S0 galaxies correlate with both the mass of molecular gas and the stellar H$\\beta$ spectral index, and all are related to the recent star formation rate. However, star formation rates estimated from the H$\\beta$ emission line luminosities $L_{H\\beta}$ in SAURON S0 galaxies are generally much smaller. Since $L_{H\\beta}$ does not correlate with 24$\\mu$m emission from dust heated by young stars, optical emission lines appear to be a poor indicator of star formation rates in SAURON S0 galaxies. The absence of H$\\beta$ emission may be due to a relative absence of OB stars in the initial mass function or to dust absorption of H$\\beta$ emission lines. ",
        "introduction": "In the discussion below we extend our earlier analysis of {\\it Spitzer} infrared observations of elliptical galaxies (E) (Temi, Brighenti \\* Mathews, 2007a) to include lenticular galaxies (S0). The data set we consider here is further restricted to 31 early type E (19) and S0 (12) galaxies that have been observed by both {\\it Spitzer} and SAURON. The SAURON sample is generally representative and has received much attention from observers at all wavelengths, allowing detailed comparisons of {\\it Spitzer} infrared observations with many different galactic attributes: stellar populations, nature and extent of cold gas, dust content, etc.  Our principal objective is to explore evidence of recent star formation in local S0 galaxies at mid and far infrared wavelengths and to contrast this with the apparent lack of measurable current star formation in the sample E galaxies. Many S0 galaxies in the {\\it Spitzer}-SAURON sample have attributes consistent with recent star formation: significantly larger mid and far infrared luminosities, large masses of molecular gas, and younger stellar ages as inferred from optical absorption line spectra. {\\it Spitzer}-SAURON S0 galaxies having the largest far-infrared luminosities are known to contain rotationally supported disks of molecular hydrogen with masses $\\gta 10^8$ $M_{\\odot}$. Using the 24$\\mu$ emission as a guide, stars are forming in many sample lenticular galaxies at very modest rates, 0.02 - 0.2 $M_{\\odot}$ yr$^{-1}$. For some of these S0 galaxies this level of star formation is comparable to the rate that gas is globally supplied from an old population of evolving stars, suggesting a quasi-steady state. In their optical spectra our sample S0 galaxiess exhibit Balmer absorption lines from young A-F stars, but the Balmer line emission from ionized gas associated with more massive O-B stars can be lower than expected. In addition, when star formation rates are estimated using standard empirical procedures, the rates determined from H$\\beta$ emission line luminosities in {\\it Spitzer}-SAURON S0 galaxies can be much less than those derived from 24$\\mu$m luminosity. The cold, dusty molecular disks in some S0 galaxies may originate from stellar mass loss, may be a vestige of a more substantial disk during a previous incarnation as a more normal spiral galaxy, or may have been acquired by a merger event. ",
        "conclusions": "Perhaps the most important conclusion from these {\\it Spitzer} observations of infrared emission from early type SAURON galaxies is that there are significant differences between S0 and E galaxies in the relative content of dust and cold gas and star formation rates. Only a small subset of S0 galaxies appear to share with E galaxies the same correlations between the 24$\\mu$m luminosity, the H$\\beta$ index, and molecular mass. Results from Sloan surveys of ``early type'' galaxies that combine S0 and E galaxies (e.g. Kaviraj et al. 2007) may therefore be misleading if applied only to E galaxies. It is likely that the source of cold gas in E and S0 galaxies is systematically quite different. Rotationally supported cold gas in S0 galaxies may be a relic of their previous incarnation as late type spirals.  The threshold for measurable star formation rates in the SAURON sample is about 0.02 $M_{\\odot}$ yr$^{-1}$. All S0 galaxies (except NGC 474, 1023, 7457 and 4382) exceed this value and all E galaxies (except NGC 4486 and possibly NGC 2768) have lower SFRs. We have shown that the 24$\\mu$m luminosity is a good measure of the SFR and that it correlates well with other related parameters such as the H$\\beta$ index, the mean stellar age determined from SAURON observations and the mass of molecular gas.  Positive residuals from the E galaxy correlation between $\\lambda L_{\\lambda}$ at 24$\\mu$m and $L_{Ks}$ represent the emission from interstellar dust alone and are a measure of the galactic star formation rate. It is remarkable that these residuals for S0 galaxies correlate so tightly with normalized far-IR luminosities $L_{60}/L_{Ks}$ and $L_{170}/L_{Ks}$. This suggests that the amount of dust, cold gas and old stars is rather finely tuned. The estimated mass loss rates from old stars in S0 galaxies are very nearly the same as the star formation rates derived from interstellar 24$\\mu$m emission. This supports the notion that the mass of cold, dusty gas in some SAURON S0 galaxies is in an approximate steady state -- gas is lost by star formation at nearly the same rate that new gas is ejected from the dominant old stellar population.  However, the H$\\beta$ emission line luminosity from S0 galaxies appears to be very low and unrelated to the star formation rate. (This is also true for the global $B - Ks$ colors of S0 galaxies in our sample.) Star formation rates determined from the mid-IR luminosity $(\\lambda L_{\\lambda})_{24}$ are often about an order of magnitude larger than those found from the H$\\beta$ emission line luminosity, using the standard SFR recipes. Perhaps star formation in cold gaseous disks in early type galaxies does not extend to massive OB stars because the disk column densities fall below the Krumholtz-McKee threshold of $\\sim1$ gm cm$^{-2}$. An alternative interpretation is that OB stars do indeed form but are deeply embedded in dust that absorbs H$\\beta$ line emission. If this is the explanation, radiation must still escape from slightly older A-F type stars since the star formation rate based on 24$\\mu$m emission correlates with the Lick H$\\beta$ index in the spectra of young $\\sim$A stars. These non-ionizing stars may be less concentrated toward the central plane of the dusty molecular disks, allowing their radiation to escape."
    },
    "0812/0812.1541.txt": {
        "abstract": "% Axisymmetric steady-state  weakly ionized Hall-MHD Keplerian thin disks are investigated by using asymptotic expansions in the small disk aspect ratio $\\epsilon$. The model incorporates the azimuthal and poloidal components of the magnetic fields in the leading order in $\\epsilon$. The disk structure is described by an appropriate Grad-Shafranov equation for the poloidal flux function $\\psi$ that involves two arbitrary functions of $\\psi$ for the toroidal and poloidal currents. The flux function is symmetric about the midplane and satisfies certain boundary conditions at the near-horizontal disk edges. The boundary conditions model the combined effect of the primordial as well as the  dipole-like magnetic fields. An analytical solution for the Hall equilibrium is achieved by further expanding the relevant equations in an additional small parameter $\\delta$ that is inversely proportional to the Hall parameter. It is thus found that the Hall equilibrium disks fall into two types: Keplerian disks with (i) small ($R_d\\sim\\delta^0$) and (ii) large (${R_d}\\gtrsim\\delta^{-k}$, $k>0$)  radius of the disk. The numerical examples that are presented demonstrate the richness and great variety of magnetic and density configurations that may be achieved under the Hall-MHD equilibrium. ",
        "introduction": "In the present study the Hall equilibrium of thin Keplerian disks embedded in 3D axially-symmetric magnetic fields is investigated. Protoplanetary disks have been the subject of numerous investigations in search of possible instabilities that may play a significant role in such yet not totally understood phenomena as the outwards transfer of angular momentum, and planet formation. However, in most astrophysics-related hydromagnetic stability studies which include accretion, jet, and wind effects the models for the underlying equilibrium states are based on the local approximation or a 'cylindrical disk' that is uniform in the axial direction.  A more realistic model of thin disks is adopted in \\cite{Regev 1983}, \\cite{Kluzniak and Kita 2000}, \\cite{Umurhan et al. 2006} for equilibrium nonmagnetized rotating disks (see also \\cite{Ogilvie 1997} for Keplerian disks of ideal MHD plasmas) and is based on an asymptotic approach in the small aspect ratio of the disk.  Traditionally, a big variety of geometrical configurations of axially symmetric steady-state equilibria are described through a few arbitrary functions of the poloidal magnetic flux. The latter satisfies the Grad-Shafranov (GS) type equations within both one-fluid and multifluid models of the plasmas (e.g. \\cite{McClements and Thyagaraja 2001, Thyagaraja and McClements 2006}, see also the earlier study by \\cite{Lovelace et al. 1986}, where in particular thin rotating disks have been discussed and references therein). Such solutions provide a base for their stability studies.  Although the magnetic field configuration in real disks is largely unknown and, as observations and numerical simulations indicate, toroidal magnetic component may be of the same order of magnitude, or some times dominate the poloidal field due to the differential rotation of the disk (see e.g. \\cite{Terquem and Papaloizou 1996, Papaloizou and Terquem 1997, Hawley and Krolik 2002, Proga 2003}). %In the initial stage of the disk rotation the poloidal field is commonly accepted as the dominant component. As shown in the present study there are three kinds of  possible  equilibrium in thin rotating disks which can not be reduced from one to the other: (i) pure poloidal, (ii) pure toroidal and (iii) mixed toroidal and poloidal magnetic field equilibria. The first case has been considered within the ideal MHD model by \\cite{Lovelace et al. 1986}, and \\cite{Ogilvie 1997}. The second case, i.e. the pure toroidal magnetic field has been investigated within the Hall MHD model by \\cite{Shtemler et al. 2007}. The present study is aimed to describe numerous axially-symmetric equilibria in 3D axially-symmetric magnetic fields that contain both toroidal and poloidal components, i.e. case (iii) within the Hall MHD model in thin disk approximation.  While most of recent research has focused on the magnetohydrodynamic (MHD) description of magneto-rotational instability (MRI), the importance of the Hall electric field to such astrophysical objects as protoplanetary disks has been pointed out by Wardle (1999). Since then, such works as \\cite{Balbus and Terquem 2001}, \\cite{Salmeron and Wardle 2003}, \\cite{Salmeron and Wardle 2005}, \\cite{Desch 2004}, \\cite{Urpin and Rudiger 2005}, \\cite{Rudiger and Kitchatinov 2005}, and \\cite{Pandey and Wardle 2008} have shed light on the role of the Hall effect in the modification of the MRIs. However, in addition to modifying the MRIs, the Hall electromotive force gives rise to a new family of instabilities in a density-stratified environment. That new family of instabilities is characterized by the Hall parameter  $\\pi_H$ of the order of unity ( $\\pi_H$  is the ratio of the Hall drift velocity  $V_{HD}$ to the characteristic velocity, where $V_{HD}=l_i^2\\Omega_i/H_*$,  $l_i$ and  $\\Omega_i$ are the inertial length and Larmor frequency of the ions, respectively, and  $H_*$  is the density inhomogeneity length (see \\cite{Huba 1991}). As shown by \\cite{Liverts and Mond 2004}, and \\cite{Kolberg et al. 2005}, the Hall electric field combined with radial density stratification gives rise to two modes of wave propagation. The first one is the stable fast magnetic penetration mode, while the second is a slow quasi-electrostatic mode that may become unstable for a short-enough density inhomogeneity length $H_*$. The latter is termed the Hall instability. Density stratifications also play a crucial role in the recently discussed starto-rotational instability. That instability is excited by the combined effect of vertical and radial density stratification.  The aim of the present work is, therefore to describe the various Hall equilibrium configurations of magnetic field and density in thin magnetized Keplerian disks as a first step towards a further study of their various density stratification-dependent instabilities. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The following simplifying assumptions are adopted throughout the present work: the plasma is treated in a 3D axially symmetric magnetic field within the Hall MHD model.  The problem inside the disk is treated within the thin-disk approximation, while the effects of the outer region is modeled by the boundary condition for the poloidal magnetic flux. Influence of the accretion, jet and wind effects are neglected in the thin-disk approximation (\\cite{Ogilvie 1997}). A near Keplerian region of the rotating disk is considered that starts at some radial distance away from a central body.  The influence of the latter on the Keplerian portion of the disk is modeled by a dipole-like contribution (singular at the disk axis) to the boundary condition for the poloidal magnetic flux at the near-horizontal disk edges. In addition, axial electrical currents that are localized in the central non-Keplerian regions of the disk are taken into account through the resulting toroidal magnetic field components within the rotating disk.  Influence of the interstellar (primordial) magnetic field is modeled by specifying the contribution of the latter on the near-horizontal disk edges.  The paper is organized as follows. The dimensional governing equations, the normalization procedure, and the resulting non-dimensional system are presented in the next Section. Section 3 contains an asymptotic analysis of the Hall-equilibrium in the small aspect ratio as well as the derivation the appropriate GS equation for the flux function and the corresponding boundary conditions at the disk edge. In section 4 an analytic model is developed for the Hall equilibrium with the aid of asymptotic expansions in a newly defined small parameter inversely proportional to the Hall parameter. Numerical examples for different regimes of the equilibrium are presented in section 5. Summary and discussion are in section 6. ",
        "conclusions": "3D axially-symmetric steady-state equilibria of weakly ionized polytropic Hall-MHD plasmas are investigated for the Keplerian portion of thin disks. Asymptotic expansions in small aspect ratio  $\\epsilon$ provide an efficient way to construct an equilibrium model for differentially rotating disks. There are three principle Hall equilibrium states which can not be reduced from one to the other, classified according to the following orientations of the magnetic field lines: (i) pure toroidal, (ii) pure poloidal, and (iii) mixed toroidal and poloidal equilibria. The three classes of equilibria differ from each other by the different ordering with respect to $\\epsilon$  of the various components of the magnetic field. In the present study the most interesting case is investigated that involves both toroidal and poloidal components of the magnetic field that are of the same zero order in $\\epsilon$. The disk structure is described by the GS equation for the poloidal flux   that involves two arbitrary functions  $I(\\psi)$ and  $\\phi(\\psi)$  for the toroidal electric current and electric potential, respectively. %magenta begin The arbitrariness of functions $I(\\psi)$ and $\\phi(\\psi)$  in steady-state problems reflects the uncertainty of the initial data in the general initial value problem. %%%magenta end The flux function   is symmetric about the midplane and satisfies certain boundary conditions at the horizontal edges of the disk (radial variations of the disk height are neglected). %magenta begin The boundary conditions for the magnetic flux $\\psi=\\Gamma(r)$ % at the horizontal disk edges express the joint effect of an primordial magnetic field, and a dipole-like magnetic field that reflects the influence of the central body on the Keplerian disk. % Solutions for different configurations of the magnetic field and density are obtained explicitly for trial linear approximations of the current  $I(\\psi)=I^{(0)}+ I^{(1)}\\psi$  flowing through a circular area in a plane   and of the electric potential   $\\phi(\\psi)=\\phi^{(1)}\\psi$   (where   $I^{(0)}$ is the total current concentrated along the disk axis), as well as several forms of the boundary magnetic flux $\\Gamma(r)$.  A particular noteworthy new feature of the present model is the finite radius of the rotating disks which is inherent for the Hall equilibrium. By using a small parameter  $\\delta$  proportional to the inverse Hall parameter with a small  coefficient of proportionality, it is established that the Hall equilibria disks fall into two types which are characterized by quite different orders in $\\delta$: Keplerian disks   with (i) small ($R_d\\sim\\delta^0$) and (ii) large (${R_d}\\gtrsim\\delta^{-k}$, $k>0$) radius of the disk. % For disks of the first family a finite radius of the disk  appears as a cut-off value at which the plasma density vanishes. % Disks of the second family  have  large  (finite or infinite) radii  due to non-uniformity of the asymptotic expansions in $\\delta$. %%magenta!!! %The value of $R_d$ for the first family is independent of the Hall parameter,  depends on the parameters of $I(\\psi)$,  %$J(\\psi)$ and $\\Gamma(r)$, and  in a wide range of their possible variation $R_d$ is additionally independent of the %plasma beta. The value of $R_d$ for the second family depends  on the Hall plasma $\\beta_H=\\beta\\pi_H$ additionally to %the parameters of $I(\\psi)$,  $J(\\psi)$ and $\\Gamma(r)$. %magenta end  % The method developed here allows to investigate analytically the equilibrium states with a large number of possible combinations of boundary conditions for the magnetic flux. Thus, it is demonstrated that all possible mutual orientations of the rotating axis, the primordial magnetic field, and the magnetic moment of the central body, as well as the relative strength of the two latter give rise to a great richness of possible topologies of the magnetic field lines. Such configurations includes geometries with O-points, with X-points, with material gaps and rings, with detached plasmoids, with magnetic islands, and with zero net induced magnetic flux. % Note that magnetic islands centered on O-points are typical for steady-state equilibria in Keplerian disk. For instance they have been simulated in thin disk approximation within MHD model for the pure poloidal magnetic field (\\cite{Lovelace et al. 1986}). On the other hand,  equilibria that contain X-points are more subtle and generally have a strong tendency to undergo a significant change in their topology due to possible reconnection processes (\\cite{Priest and Forbes 2000}). The latter however is commonly speculated to be responsible for widely observed astrophysical phenomena like jets and winds. % % %Particularly the X-point equilibria are susceptible to significant change in their topology due to possible reconnection processes (Priest and Forbes 2000) that may give rise to such phenomena as jets. Describing such processes requires an extensive stability analysis. Since the primordial poloidal flux on the disk is likely too small to be responsible for the observed total magnetic flux in disk systems (Colgate and Li 2001), the possible flux enhancing due to instability of axisymmetric Hall equilibria becomes an important issue. Additionally, it is natural then to incorporate multipoles of higher odd orders (Lepeltier and Aly 1999) instead of the primordial poloidal flux, in addition to the dipole adopted here, into to the boundary flux sources at the origin that describe the central object effect on the Keplerian disk. The stability study of the developed equilibrium disks may provide further selection of admissible magnetic configurations. For instance, development of X-point configurations may lead to an instability of such equilibria, since X-point equilibrium solutions are rather structurally unstable and a small perturbation of the equilibrium likely leads to an essential change in its form, such as magnetic lines reconnection and transition to the neighborhood stable (smooth) configuration. The present Hall equilibrium solution for thin Keplerian disks may be incorporated to the leading order in small aspect ratio into a more sophisticated problem accounting for accretion and jet inflow (similar to Ogilvie 1997)."
    },
    "0812/0812.0905_arXiv.txt": {
        "abstract": "{We present the first EVN parallax measurements of 6.7 GHz methanol masers in star forming regions of the Galaxy. The 6.7 GHz methanol maser transition is a very valuable astrometric tool, for its large stability and confined velocity spread, which makes it ideal to measure proper motions and parallaxes. Eight well-studied massive star forming regions have been observed during five EVN sessions of 24 hours duration each and we present here preliminary results for five of them. We achieve accuracies of up to 51 $\\mu$as, which still have the potential to be improved by more ideal observational circumstances.}   \\FullConference{The 9th European VLBI Network Symposium on The role of VLBI in the Golden Age for Radio Astronomy and EVN Users Meeting\\\\ September 23-26, 2008\\\\ Bologna, Italy}  \\begin{document} ",
        "introduction": "In astronomy, accurate distances are a very important quantity to measure. First, the determination of many physical properties of objects, such as the size-scale, the mass, the luminosity and the age depend on the distance. For example, the distance towards the Orion Nebula, determined by a trigonometric parallax of radio continuum stars \\cite{ment07}, turned out to be 10\\% smaller than previously assumed. The implications of this are a 10\\% drop in the masses, 20\\% fainter luminosities and 20-30\\% younger ages for the stars in the field. Moreover, one also needs distances on a large scale to study the structure of our Galaxy.  Unfortunately, accurate distance measurements are very difficult. The only method which allows an unbiased measure of distance is trigonometric parallax, which is simply a geometrical method, and therefore free of any assumptions. To achieve the stability needed for optical wavelengths, parallax measurements require space-borne observations. Historically, the first dedicated optical satellite for this purpose was Hipparcos \\cite{per95}, which had an accuracy of 0.8-2 mas, allowing it to measure distances to $\\sim$200 pc (with error bars of 20\\%). However, this successful experiment could reach only a tiny part of the Milky Way. A new optical astrometry satellite GAIA, to be launched in 2012, will be able to attain accuracies of 20 $\\mu$as, and therefore provide important results on Galactic-scale distances. Despite the high accuracy reachable in the optical, GAIA will suffer from dust extinction, especially towards the large and dense dust lanes in the Galactic plane. It is in these regions that radio astronomy can provide an important complement to GAIA; radio wavelengths do not suffer from dust extinction and VLBI phase referencing techniques can provide accuracies up to 10~$\\mu$as for obtaining accurate (with error bars of 10\\%) distances to 10~kpc (see for example \\cite{honma07},\\cite{xu06} and \\cite{reid08}). This is up to a 100-fold improvement on Hipparcos!  To measure the trigonometric parallax one requires very strong and compact sources. Masers are the ideal targets, found frequently in (dusty) star formation regions (SFRs) and AGB stars, objects whose distances one wants to know. Water masers at 22 GHz inhabit both these regions, while methanol masers (12 and 6.7 GHz) are exclusively associated with high mass SFRs. Water and 6.7 GHz methanol masers can be very strong, with peak fluxes of the order of a kilojansky, but because of their stability and confined velocity spread methanol masers are preferred for astrometry. Water masers are often associated with outflows in SFRs \\cite{hachi06}, which is manifested by a large velocity spread. Methanol masers are thought to be pumped by emission from warm dust, heated by a young stellar object and have velocity spreads of maximally $\\sim$20 kms$^{-1}$, but frequently much smaller \\cite{sobo07}. The frequency difference between water and methanol masers results in different sources of disturbance for each of the transitions. At higher frequencies the signal is more affected by tropospheric delay errors than at lower frequencies, where the ionosphere is typically more dominant. Also, the positional accuracy of the quasar one requires for phase referencing depends on frequency. At high frequencies one observes the more optically thick part of the quasar, while at lower frequencies one is dominated by the optically thin part and hence one can encounter more extended structure.   When the project began in 2006, the EVN was the only VLBI network that had receivers for observing the 6.7 GHz maser transition. It offered a unique opportunity to start a project measuring the parallaxes and proper motions towards eight well known SFRs with strong methanol masers. Accurate distances towards these targets will have an important impact on the study of massive star formation in these regions. ",
        "conclusions": "We have demonstrated that one can measure the trigonometric parallax of 6.7 GHz methanol masers using the EVN. Good results were obtained for sources of higher declinations ($\\delta$~>~25$^\\circ$). The accuracy reached in this experiment is $\\sim$50 -- 100~$\\mu$as, which is 10 -- 20 times better than Hipparcos. The accuracy of parallax measurements of the 6.7 GHz maser transition with the EVN can be improved by taking notice of the following. Our project was large; it consisted of eight masers, all at different R.A.s and declinations. In the five sessions, we did not manage to have the ideal parallax signature coverage for each maser (compare L1287 with ON1 in Figure~1), for which one prefers to sample the peaks. Often, as for L1287, some sessions were near the zero-signature of the parallax, decreasing the quality of the fit. Another problem encountered in the project was the large number (50\\%) of low-declination sources. We did not manage to observe them as long as their high declination counterparts. This, combined with their low elevations, resulted in a poor $u,v$ coverage.  As in previous parallax studies, the position uncertainties in declination are significantly worse than in Right Ascension (see for example L1287 in Figure 1). First, the errors introduced by the atmosphere are larger in declination. Second, the shape of the beam is more elongated in declination than in Right Ascension. For example, the source S269 has a beam of $7.7\\times3.7$~mas with a position angle of 55$^\\circ$.  Finally, we would like to touch upon the {\\sl geodetic-like} observations. We find an improvement in the calibration for most of the cases. However, in a few cases the corrections did not improve the calibration. We suspect that here the ionospheric delay plays a more important role, and actually we degrade the calibration by wrongly interpreting the determined delay as of tropospheric origin. In these cases, we did not apply the {\\sl geodetic-like} correction. A dual frequency observation mode, as employed in geodetic observations to calibrate the ionosphere, is not possible with most of the EVN antennas. Also switching receivers is impossible. Therefore, determining the S/N with or without a {\\sl geodetic-like} correction in each observation is the best option for now.  In conclusion, we have shown that 6.7 methanol masers, observed with the EVN, are a well suited tool for astrometry. The fact that the VERA array was recently equipped with 5 GHz receivers is a sign for a growing enthusiasm for 6.7 GHz methanol masers. The final astrometry of our sources will be used to reevaluate the SFRs' astrophysics in the light of the new distance determinations, and to study the 3D velocities of the SFRs relative to Galactic motion."
    },
    "0812/0812.1132_arXiv.txt": {
        "abstract": "{} {The main purpose of this paper is to study time delays between the light variations in different wavebands for a sample of quasars. Measuring a reliable time delay for a large number of quasars may help constraint the models of their central engines. The standard accretion disk irradiation model predicts a delay of the longer wavelengths behind the shorter ones, a delay that depends on the fundamental quasar parameters. Since the black hole masses and the accretion rates are approximately known for the sample we use, one can compare the observed time delays with the expected ones.} {We applied the interpolation cross-correlation function ($ICCF$) method to the Giveon et al. sample of 42 quasars, monitored in two ($B$ and $R$) colors, to find the time lags represented by the $ICCF$ peaks. Different tests were performed to assess the influence of photometric errors, sampling, etc., on the final result.} {We found that most of the objects show a delay in the red light curve behind the blue one (a positive lag), which on average for the sample is about +4 days (+3 for the median), although the scatter is significant. These results are broadly consistent with the reprocessing model, especially for the well-sampled objects. The normalized time-lag deviations do not seem to correlate significantly with other quasar properties, including optical, radio, or X-ray measurables. On the other hand, many objects show a clear negative lag, which, if real, may have important consequences for the variability models.} {} ",
        "introduction": "Although variability is an extensively studied and rather common feature of quasars, little is currently known about the exact nature of the processes driving the flux changes. Variations are observed in practically all energy bands, often with time delays between them. This fact suggests a common cause (process) connecting otherwise spatially separated regions (presumably parts of an accretion disk), where most of the corresponding wavelengths come from. Since the time lags are rather short for the expected distances, it is the speed of light that is perhaps the only way to casually connect these separate regions. Several competing ideas, often leading to different signs of the time lags between the bands have been discussed in the literature (Czerny et al. 2008):  \\begin{itemize} \\item {\\bf Reprocessing} of the central high-energy continuum into lower energy bands from the outer (and colder) regions of an accretion disk (Krolik et al. 1991). For this mechanism to work, the central X-ray emission must ``see'' the outer parts -- either because of elevation of the central source above the disk surface, of warping, or of flaring of the disk (or all three). Thus, a time lag between the variable X-rays and the lower wavelength (e.g. UV, optical, IR) will be observed with a relation between the lag and the wavelength, roughly scaled as $\\tau_{\\rm \\lambda} \\sim \\lambda^{4/3}$ (Sect. 4) for a standard accretion disk (Shakura \\& Syunyaev 1973). Therefore, if the reprocessing is responsible for the most of the optical variability, a delay between every two optical bands (say $B$ and $R$-bands) will be expected, with the shorter wavelengths leading (a positive lag). However, if the variable X-rays come not from the center, but from another location instead (e.g. from local flares above the disk surface), the sign of the lags can be either positive or negative, depending on the radial distance at which these flares predominantly occur. In this case, a connection between the signs of the lags and the fundamental quasar parameters could be searched, as the radial distance of the X-ray flares is probably governed by the structure of a corona above the disk, which in turn should depend on the quasar parameters.  \\item {\\bf A weak blazar component}, often assumed to exist even in radio-quiet sources (Czerny et al. 2008, and the references therein). Such a component can produce variable optical/IR synchrotron radiation, and respectively -- variable X-rays via the synchrotron self-Compton mechanism. These X-rays can, in turn, be reprocessed in a disk into optical/UV photons, leading to a complicated interband time-delay picture, in which the redder (synchrotron) variations will probably lead the bluer (reprocessed) ones.   \\item {\\bf Disk fluctuations}, drifting inward. This mechanism is sometimes invoked (Ar\\'evalo et al. 2008) to explain cases where the long-wavelength variation are leading. However, the viscous timescale at the optical band generation distance is much longer than the timescales of a few days, discussed in this paper.  \\end{itemize}  The situation can be very complicated, as several of the effects mentioned above can work in combination, or of course, an entirely different mechanism can be responsible for the interband relations. In either case, studying the lags between the continuum changes in different wavelengths for a large sample of quasars may help to clarify the situation, and therefore, to create a better picture of the quasar central engines.  Sergeev et al. (2005) studied the light curves of 14 nearby Seyfert galaxies, observed on $\\sim$60--150 epochs in 4 broad-band filters and found positive delays for most of the cases, varying between a fraction of a day and a few days. Their results are broadly consistent with the reprocessing model. Recently, similar results were reported by Liu et al. (2008) and Ar\\'evalo et al. (2008). Here we extend their works, applying a similar technique to a larger sample of quasars, optically monitored at the Wise observatory (Giveon et al. 1999) in two colors for several years. The sample is described in more detail in the next section. In Sect. 3 we describe the method we used to find the time delay between the bands, i.e. the interpolation cross-correlation function ($ICCF$) method. Next, the results for the time lags are presented and compared with the predictions of the reprocessing model. Finally, we discuss the reliability of the lag estimates and different implications for the central engine models of quasars, by comparing the deviations from the reprocessing model predictions with various quasar characteristics. ",
        "conclusions": "In this paper we analyze the light curves and study the time lags between the optical continuum bands for a large sample of quasars. In spite of the significant scatter, we show that the lags are broadly consistent with the reprocessing model, according to which the optical variations are largely due to the reprocessing of the central X-ray radiation in a surrounding thin accretion disk. There are also some indications that the central X-ray irradiating source may be located at some height (a few hundred Schwarzschild radii) above the disk plane, representing perhaps a (failed) jet base. The deviations of the observed lags from the expectations, assuming the reprocessing model, do not seem to correlate significantly with any other quasar properties and are probably due mostly to the scarce sampling. The paper also demonstrates that the broad-band optical monitoring of quasars could be a powerful tool to study the central engines, provided the light curve is well sampled. A small robotic telescope (e.g. 0.4--0.6m), dedicated to monitoring a large sample of brighter quasars once every 2--3 days, with an accuracy of 0.02 mag or better in 3 filters, might be able to clarify the role of reprocessing in quasar variability."
    },
    "0812/0812.4094_arXiv.txt": {
        "abstract": "\\noindent We report the detection of a candidate brown dwarf orbiting the metal-rich K dwarf HD~91669, based on radial-velocity data from the McDonald Observatory Planet Search.  HD~91669b is a substellar object in an eccentric orbit ($e$=0.45) at a separation of 1.2~AU.  The minimum mass of 30.6~\\Mjup\\ places this object firmly within the brown dwarf desert for inclinations $i\\gtsimeq$23\\degrees.  This is the second rare close-in brown dwarf candidate discovered by the McDonald planet search program. ",
        "introduction": "Brown dwarfs are commonly defined as substellar objects with masses between the deuterium-burning limit (about 13\\Mjup) and the hydrogen-burning limit (about 80\\Mjup).  Despite the large (several hundred \\ms) reflex radial-velocity signals produced by brown dwarfs in orbit around main-sequence stars, planet-search surveys monitoring thousands of stars have found remarkably few such objects in orbits with $a\\ltsimeq$5~AU.  By contrast, nearly 300 planets have been found by the radial-velocity method.  Early radial-velocity surveys by \\citet{campbell88} and \\citet{murdoch93} reported a distinct lack of candidate companions in the brown dwarf mass range, despite their relatively easy detectability.  \\citet{campbell88} noted that 7 of 15 targets showed velocity variability consistent with a distant companion, but the lack of astrometric variability led those authors to limit the companion mass range to $\\sim$1-9 \\Mjup.  In the nearly 20 years spanned by current radial-velocity planet searches, only a handful of brown dwarf candidates have been identified, e.g.~HD~114762b (11 \\Mjup; Latham et al.~1989, Cochran et al.~1991, Mazeh et al.~1996), HD~168443c (18 \\Mjup; Marcy et al.~2001), HD~202206b (17 \\Mjup; Udry et al.~2002), and HD~137510b (26 \\Mjup; Endl et al.~2004).  Recently, a transiting brown dwarf has been discovered by the \\textit{CoRoT} spacecraft.  This object, CoRoT-Exo-3b, has a mass of 21.7 \\Mjup\\ and a radius of 1.01 \\Rjup, in a remarkably close orbit at 0.05~AU \\citep{deleuil08}.  The steep rise in the planetary mass function toward lower masses \\citep{marcy05a} combined with the decline in lower-mass stellar companions \\citep{mazeh03} give evidence for the existence of a ``brown dwarf desert'': a paucity of brown dwarf companions orbiting within $\\sim$3-4~AU of solar-type stars \\citep{mb00}.  In contrast, there is no dearth of brown dwarfs orbiting at larger separations \\citep{gizis01}, or free-floating \\citep{reid99}.  \\citet{grether06} performed a detailed investigation of the planetary and stellar mass distributions in order to correct for biases and to assess the reality of the brown dwarf desert.  Defining close companions as those with orbital periods $P<$5~yr, \\citet{grether06} confirmed the deficit of close brown dwarf companions and estimated the driest part of the ``desert'' to lie at $M=31^{+25}_{-18}$ \\Mjup.  It is possible that selection biases curtail the discovery rate of brown dwarfs relative to planets, as measured by published results.  The push to detect lower-mass planets combined with the high value of telescope time lead planet search teams to discard targets which exhibit the large-amplitude radial-velocity variations induced by brown dwarf companions in close orbits.  However, a simple bias of this nature does not seem likely to explain the factor of $\\sim$100 deficit in the observed frequency of brown dwarf companions compared to planetary companions.  Since the radial-velocity method yields only a minimum mass for the companion, some number of brown dwarf candidates are likely to be low-mass stars orbiting at low inclinations.  Some authors have combined the radial-velocity orbits with astrometric data in order to determine the inclination and true mass of the companions.  In this way, the ``planets'' HD~38529c and HD~168443c were revealed to be brown dwarfs with masses of 37 \\Mjup\\ and 34~\\Mjup, respectively \\citep{reffert06}. \\citet{bean07} used \\hst astrometry to determine that the planet candidate HD~33636b, with a minimum mass of 9.3~\\Mjup, was in fact a star with a true mass of 142 \\Mjup\\ orbiting at a nearly pole-on inclination $i=4.1$\\degr.  These examples illustrate the importance of finding brown dwarf candidates with low values of m~sin~$i$ to increase the likelihood of their true mass remaining in the substellar mass range.  In this paper, we derive the parameters of the host star HD~91669 (\\S~3), and in \\S~4 we present evidence for its brown dwarf companion, and discuss possible constraints on its true mass. ",
        "conclusions": "We have presented radial-velocity observations indicating the presence of a rare brown dwarf candidate orbiting the metal-rich K dwarf HD~91669. HD~91669b is the second brown dwarf identified in the McDonald Observatory planet search program.  Like its predecessor HD~137510b \\citep{endl04}, this object has a relatively low mass and hence a high probability of having a true mass within the substellar regime.  We note that the rarity of brown dwarfs implies that the detection of two candidates in a sample of 250 stars is rather unlikely.  However, the McDonald Observatory program yields a detection rate of 0.8$\\pm$0.6\\%, compared to a rate of 0.7$\\pm$0.2\\% from the California \\& Carnegie Planet Search (7 candidates from $\\sim$1000 target stars: Vogt et al.~2002, Patel et al.~2007).  Within the limits of small-number stattics, the yields of brown dwarf candidates from these two programs are comparable."
    },
    "0812/0812.3950_arXiv.txt": {
        "abstract": "We study the prospects of finding the first quasars in the universe with ALMA and JWST. For this purpose, we derive a model for the high-redshift black hole population based on observed relations between the black hole mass and the host galaxy. We re-address previous constraints from the X-ray background with particular focus on black hole luminosities below the Eddington limit as observed in many local AGN. For such luminosities, up to $20\\%$ of high-redshift black holes can be active quasars. We then discuss the observables of high-redshift black holes for ALMA and JWST by adopting NGC~1068 as a reference system. We calculate the expected flux of different fine-structure lines for a similar system at higher redshift, and provide further predictions for high-$J$~\\COI lines. We discuss the expected fluxes from stellar light, the AGN continuum and the Lyman $\\alpha$ line for JWST. Line fluxes observed with ALMA can be used to derive detailed properties of high-redshift sources. We suggest two observational strategies to find potential AGN at high redshift and estimate the expected number of sources, which is between $1-10$ for ALMA with a field of view of $\\sim(1')^2$ searching for line emission and $100-1000$ for JWST with a field of view of $(2.16')^2$ searching for continuum radiation. We find that both telescopes can probe high-redshift quasars down to redshift $10$ and beyond, and therefore truely detect the first quasars in the universe. ",
        "introduction": "After the first detection of supermassive black holes at $z\\sim6$ \\citep{Fan01}, there has been a strong effort to find more sources at high redshift and to study their properties. Further high-$z$ quasi-stellar objects (QSOs) have been found in surveys like the Sloan Digital Sky Survey (SDSS)  \\citep{Fan01,Carballo06,Cool06,Fan06,Shen07,Inada08}, the Canada-France high-redshift quasar survey \\citep{Willott07}, the UKIDSS Large Area Survey \\citep{Venemans07} and the Palomar-QUEST survey \\citep{Bauer06}. They have been studied by follow-up observations with Chandra \\citep{Shemmer06}, the Subaru Telescope \\citep{Goto07}, with IR spectroscopy \\citep{Juarez07}, near-IR spectroscopy \\citep{Jurez08}, mid-IR observations \\citep{Stern07, Martinez08}, Swift observations \\citep{Sambruna07}, VLBA observations \\citep{Momjian07}, the Very Large Array as well as the Very Long Baseline Array \\citep{Momjian08}. Some of their host galaxies have also been detected in CO and submm emission \\citep{Omont96, Carilli02, Walter04, Weiss07, Riechers08b, Riechers08a}.  These observations furthermore inferred a number of relations between the mass of the central black hole and the properties of their host galaxy. The Magorrian relation connects the velocity dispersion and the mass of the galactic bulge with the black hole mass \\citep{Magorrian98, Ferrarese00, Gebhardt00,Graham01,Merritt01,Tremaine02,Haering04,Peterson08,Somerville08}. Quasars with black hole masses of the order $10^9\\,M_\\odot$ are observed with supersolar metallicities even at $z\\sim6$ \\citep{Pentericci02}, which is indicated in further studies as well \\citep{Freudling03,Maiolino03,Dietrich03a,Dietrich03b}. In the nearby universe, correlations have been found between black hole mass and metallicity \\citep{Warner03, Kisaka08}, with a mild slope of $0.38\\pm0.07$. This is plausible, as it is generally recognized that the overall mass of a galaxy is correlated with its metallicity \\citep{Faber73,Zaritsky94,Jablonka96,Trager00}. In this paper, we will in general assume that the black hole mass is a good indicator for the evolutionary stage of the quasar host galaxy, in particular with respect to the Magorrian relation and the expected metallicity.  In the local universe, active galactic nuclei (AGN) can be studied in much more detail and with higher resolution. An example that has been studied particularly well, in different wavelengths and with high resolution is the AGN NGC~1068. Its direct X-ray emission is shielded through absorption along the line of sight, but there is evidence for reflected X-ray photons that are scattered into the line of sight \\citep{Pounds06}. Its molecular disk has been studied through emission from different species, in particular molecular \\COI and fine-structure lines \\citep{Schinnerer00, Galliano03, Galliano05, Spinoglio05, Davies07, Davies07b, Poncelet07}. The continuum emission and the distribution of dust has been measured in further studies \\citep{Krips06, Mason06, Rhee06, Tomono06, Howell07, Hoenig08}. Such observations in particular allow to probe stellar and gaseous dynamics and the galactic gravitational potential \\citep{Gerssen06, Emsellem06, Das07}. Observations with the Green Bank Telescope (GBT) provide evidence for water maser activity \\citep{Greenhill95, Greenhill96, Kondratko06}. The inner radio jet can be detected with observations at $7$~mm wavelength \\citep{Cotton08}. Given this amount of detail, NGC~1068 thus provides an excellent test case for the inner structure of active galaxies.  The formation of quasars and their supermassive black holes is still one of the unresolved riddles of structure formation and cosmology. The simplest scenarios assume that they have grown from the remnants of the first stars, which are believed to be very massive \\citep{Abel02, Bromm04, Glover05}, and whose black hole remnants could grow further by accretion. Such a scenario however has problems, as the remnants of the first stars typically do not end up in the most massive quasars at redshift $z\\sim6$ \\citep{Trenti08}, and radiative feedback from the stellar progenitor can delay accretion as well \\citep{Johnson07, Alvarez08, Milosavljevic08}. In case of Eddington accretion, seed black holes of $\\sim10^5\\,M_\\odot$ are required in order to grow to the observed supermassive black holes at $z\\sim6$ \\citep{Shapiro05}.  Recently, it has also been discussed whether the first stars in the early universe were powered by dark matter annihilation rather than nuclear fusion \\citep{Spolyar08, Iocco08}. Such stars could reach masses of the order $1000\\ M_\\odot$ \\citep{FreeseBodenheimer08, IoccoBressan08} and were considered as possible progenitors for the first supermassive black holes. The evolution of such stars on the main sequence has been calculated by \\citet{Taoso08} and \\citet{Yoon08}. However, it was shown that such stellar models are highly constrained by the observed reionization optical depth \\citep{SchleicherBanerjee08a, SchleicherBanerjee08c}. Also, we note that such seeds would still require super-Eddington accretion to grow to the observed supermassive black holes at $z\\sim6$.  Therefore, alternative scenarios have been considered that lead to the formation of massive seed black holes by direct collapse \\citep{Eisenstein95, Koushiappas04, Begelman06, Spaans06, Dijkstra08}. However, for such scenarios, the situation is also not entirely clear. For instance, \\citet{Lodato06} argue that the gas in such halos should fragment if \\HzI cooling is efficient, and \\citet{Omukai08} and \\citet{Jappsen08} suggest that a non-zero metallicity will lead to fragmentation as well. Simulations by \\citet{Clark08} support these conjectures.  The black hole population at high redshift is further constrained by observations of the soft X-ray background. These constraints indicate that the population of high-redshift quasars was not sufficient to reionize the universe \\citep{Dijkstra04}, and suggest an upper limit to the black hole density of $\\sim4\\times10^4\\,M_\\odot\\,\\mathrm{Mpc}^{-3}$ \\citep{Salvaterra05}. The latter depends in particular on the adopted spectra and luminosity of the black hole, as well as their duty cycles, and it has been argued that larger black hole populations are conceivable as well \\citep{Zaroubi07}. It is therefore important to probe the number density of black holes and AGN activity as a function of redshift by direct observations.  In this work, we explore the possibility of searching for the first quasars in the universe with ALMA \\footnote{http://www.eso.org/sci/facilities/alma/index.html} and JWST \\footnote{http://www.stsci.edu/jwst/overview/}. ALMA is particularly relevant in light of recent predictions on the observability of the first miniquasars \\citep{Spaans08} and high-redshift Lyman $\\alpha$ galaxies \\citep{Finkelstein08}, while it is one of the main scientific goals of JWST to find the first galaxies at high redshift. In \\S \\ref{population}, we introduce a model for the black hole population which is based on observed properties and correlations between black holes and their host galaxy. In \\S \\ref{constraint}, we discuss how the number of active black holes is constrained by the unresolved X-ray background. In \\S \\ref{ngc1068}, we analyze the observed properties of NGC~1068 and derive further predictions for very high-$J$~\\COI lines. In \\S \\ref{observables}, we use these information to derive the relevant observables for ALMA and JWST and predict the number of detectable sources as a function of redshift. Further discussion and outlook is provided in \\S \\ref{discussion}. ",
        "conclusions": "\\label{discussion} In this paper, we have presented a model for the quasar population at high redshift, based on the known correlations between the black hole mass and its host galaxy. We have derived constraints on the fraction of active quasars from the SXRB and discussed the expected observables. We have based our discussion on using NGC~1068 as a template AGN, and calculated its observability at high redshift. Our focus lies in particular on high-$J$~\\COI lines, fine-structure lines of \\CII and \\OId, dust emission, stellar light, the AGN continuum and Lyman $\\alpha$ radiation. Above a metallicity of $10^{-2}$ solar, \\COId, \\CII and \\OI are the main coolants, and their detection can be used to derive an estimate for the heating rate and the impinging X-ray flux. A profile of the observed line width further allows to derive estimates for the black hole mass and probes the mass - luminosity relation at high redshift. We expect that galaxies containing a black hole in the mass range $10^6-10^7\\ M_\\odot$ can be observed at high redshift with ALMA and JWST. We predict $1-10$ high-redshift sources that can be found via molecular lines in an ALMA field of view of $(1')^2$, and $100-1000$ sources to be found by JWST using NIRCam and focusing on stellar emission.  One uncertainty in this discussion comes from the question whether the Maggorian relation is established already at high redshift. This seems fair in light of supersolar metallicities in high-redshift quasars. On the other hand, if our model overpredicts the number of high-redshift black holes, the X-ray constraints would be alleviated further, and a larger fraction of the black hole population could be active. In this sense, the number of active quasars might still be of the same order. Indeed, unless the population of black holes between redshifts $5$ and $6$ has accreted mass with high super-Eddington rates, there must be progenitors with masses of $\\sim10^6-10^7\\ M_\\odot$ at higher redshift. It is a secondary issue for our purposes whether the black hole would be directly visible in X-rays, or if the system is a hidden AGN and only visible in the excited high-$J$~\\COI lines.  In about three years, ALMA will become operational with initially $16$~antennas. This will provide sufficient sensitivity to detect high-$J$~\\COI lines, fine-structure lines and dust continuum emission from the brightest high-redshift quasars that are already known. As these systems are large, the high angular resolution of ALMA will provide detailed spatial information about the inner structure of such systems, and detailed diagnostics for the \\COI lines allow to infer the density and temperature distribution in the host galaxy \\citep{Meijerink07}.  Once ALMA has the full set of $50$~antennas, it reaches even higher sensitivity and it becomes possible to search for AGNs with less massive black holes. For such deep field observations, we recommend to focus initially on modest redshifts of $z\\sim5-6$, where we expect higher fluxes, more sources, and where one may be more confident that the Magorrian relation is well-established. This provides an important test to confirm the observability of high-redshift quasars, as well as important information about the progenitors of the most massive quasars in the universe. Such studies can then gradually be extended to higher redshifts and will provide a direct test for structure formation models in the high-redshift universe and of the Magorrian relation at different redshifts.  JWST can aid such surveys by providing additional sources that may be detected in stellar light, the AGN continuum or Lyman $\\alpha$ emission. In this case, ALMA does not need to glance into the dark sky, but can do pointed follow-ups of the known sources. This will speed up high-redshift surveys, and provide complementary information at different frequencies. It is particularly important at very high redshift beyond $z\\sim10$, where the probability of finding sources in the ALMA field of view decreases rapidly. The combination of these instruments can therefore truely detect the first quasars in the universe."
    },
    "0812/0812.1022_arXiv.txt": {
        "abstract": "We describe the creation of a set of artificially ``redshifted'' galaxies in the range $0.1<z<1.1$ using a set of $\\sim$100 SDSS low redshift ($v<7000$ km\\,s$^{-1}$) images as input. The intention is to generate a training set of realistic images of galaxies of diverse morphologies and a large range of redshifts for the GEMS and COSMOS galaxy evolution projects. This training set allows other studies to investigate and quantify the effects of cosmological redshift on the determination of galaxy morphologies, distortions and other galaxy properties that are potentially sensitive to resolution, surface brightness and bandpass issues. We use galaxy images from the SDSS in the $u$, $g$, $r$, $i$, $z$ filter bands as input, and computed new galaxy images from these data, resembling the same galaxies as located at redshifts $0.1<z<1.1$ and viewed with the Hubble Space Telescope Advanced Camera for Surveys (HST ACS). In this process we take into account angular size change, cosmological surface brightness dimming, and spectral change. The latter is achieved by interpolating a spectral energy distribution that is fit to the input images on a pixel-to-pixel basis. The output images are created for the specific HST ACS point spread function and the filters used for GEMS (F606W and F850LP) and COSMOS (F814W). All images are binned onto the desired pixel grids (0\\farcs03 for GEMS and 0\\farcs05 for COSMOS) and corrected to an appropriate point spread function. Noise is added corresponding to the data quality of the two projects and the images are added onto empty sky pieces of real data images. We make these datasets available from our website, as well as the code -- \\ferengi: ``Full and Efficient Redshifting of Ensembles of Nearby Galaxy Images'' -- to produce datasets for other redshifts and/or instruments. ",
        "introduction": "\\label{sec:intro} In the current era of observational astronomy the size of galaxy datasets means that number statistics starts to lose its spot as number one source of uncertainty in galaxy evolution studies. With the availability of large wide field galaxy surveys as the Sloan Digital Sky Survey \\citep[SDSS][]{york00} or 2dF Galaxy Redshift Survey, and the deep, space based high resolution projects like the Hubble Ultra Deep Field, GOODS, GEMS and STAGES projects and the Cosmic Evolution Survey (COSMOS) with their $10^4$ to $10^6$ galaxies, other error sources become vital to understand.  If we want to understand the buildup of galaxies with their intricately linked evolution in stellar and black hole mass, luminosities, colors, morphological types, and the alternations between interaction and relaxation, we have to understand what the tools we apply really measure. Any given galaxy will look different when viewed with different instruments or when located at different redshifts, due to cosmological dimming and changes in rest-frame bandpass.  This makes the qualitative or quantitative classification of e.g.\\ galaxy morphology non-trivial to compare for different redshifts or filters. As one example faint disks visible at one redshift will fade from the observer's view at larger distances, not only affecting eyeball-classifications of morphology but also automatic classifier software. Moreover, low surface brightness signs of interaction and structures, such as bars which are not prominent in the rest-frame ultraviolet, will not be visible at higher redshfits.  If the evolution of galaxies is to be studied via merger statistics, morphology-segregated evolution, star formation histories, and type-specific luminosity functions, all morphological classifiers have to be calibrated to deliver comparisons of similar quantities at all redshifts, taking into account bandpass-dependent properties and changes in signal-to-noise ratio.  From the analysis of HST survey data we know that the average surface brightness of the disc galaxy population fades with time \\citep[e.g.\\ ][]{lilly1998,barden2005}. The change is approximately 1-1.5mag depending on rest-frame bandpass over the redshift interval $0<z<1$. To some degree this surface brightness evolution counters the cosmological dimming and helps detecting low surface brightness features. Thus, if this effect is not taken into account, predictions about the recoverability of structural parameters or classifications are overly pessimistic.  For a few projects in the past, codes were created that would include some of these effects for specific applications or datasets \\citep[e.g.][]{abraham1996a,abraham1996b,giavalisco1996,bouwens1998, takamiya1999,burgarella2001,kuchinski2001,vdbergh2002,lisker2006}. However, to our knowledge no codes or datasets that include geometrical and cosmological bandpass shifting effects are currently publicly available. In the present article we describe the creation of artificial image data sets in the range $0.1<z<1.1$, computed from low-$z$ SDSS galaxies by artificial ``redshifting''. Since we refrain from adding any evolutionary models but purely apply cosmological changes in angular size, surface brightness and filter bandpass, this dataset shows exactly how such galaxies will appear when observed from cosmological distances, at which redshift certain features become undetectable, and any quantitative classifier procedure can be tested for dependency on redshift effects.  Our code \\ferengi, ``Full and Efficient Redshifting of Ensembles of Nearby Galaxy Images'', and the present datasets were originally created for the morphological classification of active and inactive galaxies in the GEMS \\citep{rix04}, STAGES \\citep{grayinprep}, and COSMOS \\citep{scov06} projects. The datasets for the HST ACS image characteristics of the GEMS and COSMOS project are freely available from our website\\footnote{Website for retrieval of simulated datasets and the code \\ferengi: \\url{http://www.mpia.de/FERENGI/}}. As many studies might benefit from such data we also make the \\ferengi\\ code available on the same webpage for others to use with different input samples, redshifts, and instrument characteristics.  In the following we describe the input data in Section~\\ref{sec:input}, the redshifting procedure including the basic cosmological formulae that enter, and the bandpass shift in Section~\\ref{sec:redshifting}, and the creation of realistic images from this information in Section~\\ref{sec:noise} including example results (Section~\\ref{sec:results}). Next, we present a series of tests characterising the robustness and accuracy of the code (Section~\\ref{sec:tests}). In Section~\\ref{sec:conclusions} we discuss limitations of this procedure.  All examples and cosmology-dependent numbers given here are computed assuming a flat universe with $\\Omega_\\Lambda=0.7$ and $h=H_0/(100\\,\\mathrm{km\\,s^{-1}\\,Mpc^{-1}})=0.7$. ",
        "conclusions": "\\label{sec:conclusions} We take from Figures~\\ref{fig:mean01} and \\ref{fig:mean02} that \\ferengi\\ has no unwanted systematic effects on measured morphological and photometrical parameters of simulated galaxies. The deviations from theoretical values are well within the uncertainties expected from, in this case, \\galfit, also in absolute terms. Both linear scales as well as magnitudes are well reproduced. This includes the bandpass shifts induced by redshifting and different filters, as is demonstrated in the right panels in Figures~\\ref{fig:mean02}: The \\kcorrect\\ module correctly interpolates between input filter images. We cannot exclude issues when {\\em extrapolating} \\kcorrect\\ templates, but we explicitly restrict the output range of redshifts to only use template {\\em interpolation}, between bands.  We make \\ferengi\\ publically available from our webpage, as are redshifted sets of images for redshifts $0.1\\le z\\le1.1$ and different HST ACS filters. Primarily these are created for the GEMS (F606W and F850LP at 0\\farcs03/pixel) and COSMOS (F814W at 0\\farcs05/pixel) projects, but the code can be modified and used for other filters. The user can extend the input sample of images (also provided on the webpage), to other objects, filter curves, and output redshifts, and in principle also to other data sources. We leave it up to the user to update \\ferengi\\ with respect to the $K$-correction module when including other input filter bands.  As stated, templates are used for interpolation between input filters. They can theoretically also be extrapolated beyond the input wavelength interval, if one has faith in the templates. However, we strongly advise against this for a different reason: Due to the limited S/N available for the pixel-by-pixel $K$-correction that is computed, a mild extrapolation might be acceptable, but the noise of the output pixels and systematic deviations will obviously increase with distance to the last supported wavelength. As various templates might be degenerate when interpolating, there might be striking differences on extrapolation, resulting in large errors in the output flux.  As a second caveat we note that stars in the input images are not treated separately in \\ferengi\\ or in the provided datasets. {\\it Any} flux in the input images, be it stars or back-/foreground galaxies, will be treated as if it were at the distance of the target and thus be redshifted. While distant background galaxies usually disappear, stars will be represented with the output PSF, resembling a faint dense star field (the density of stars is enhanced quadratically with decreasing linear scales as the simulation redshift increases). Of even greater concern are foreground (or closer background) galaxies. In the output, their physical properties are calculated falsely as their assumed distance is not correct. Whether the resulting images will in the end be reliable representations of reality, strongly depends on the quality of input images, the type of galaxies, and their correspondence with the templates used."
    },
    "0812/0812.2021_arXiv.txt": {
        "abstract": "We present the first results from the largest spectroscopic survey to date of an intermediate redshift galaxy cluster, the $z=0.834$ cluster \\rxjfull.  We use the colors of galaxies, assembled from a $D \\sim 12$~Mpc region centered on the cluster, to investigate the properties of the red-sequence as a function of density and clustercentric radius.  Our wide-field multi-slit survey with a low-dispersion prism (LDP) in the IMACS spectrograph at the 6.5~m Baade telescope allowed us to identify \\nldponly\\ new members of the cluster and its surrounding large-scale structure with a redshift accuracy of $\\sigma_z/(1+z) \\approx 1\\%$ and a contamination rate of $\\sim 2\\%$ for galaxies with \\ifil~$<23.75$~mag.  We combine these new members with the \\nspec\\ previously known spectroscopic members to give a total of \\ntot\\ galaxies from which we obtain a mass-limited sample of \\nmcut\\ galaxies with stellar masses \\mcutrange.  We find that the red galaxy fraction is $93 \\pm 3\\%$ in the two merging cores of the cluster and declines to a level of $64 \\pm 3\\%$ at projected clustercentric radii $R \\ga 3$~Mpc.  At these large projected distances, the correlation between clustercentric radius and local density is nonexistent.  This allows an assessment of the influence of the local environment on galaxy evolution, as opposed to mechanisms that operate on cluster scales (e.g. harassment, ram-pressure stripping).  Even beyond $R>3$~Mpc we find an increasing fraction of red galaxies with increasing local density.  The red galaxy fraction at the highest local densities in two large groups at $R>3$~Mpc matches the red galaxy fraction found in the two cores.  Strikingly, galaxies at intermediate densities at $R>3$~Mpc, that are not obvious members of groups, also show signs of an enhanced red galaxy fraction.  Our results point to such intermediate density regions and the groups in the outskirts of the cluster, as sites where the local environment influences the transition of galaxies onto the red-sequence. ",
        "introduction": "The environment of a galaxy has long been known to influence its properties, such as morphology, color, and star formation.  Numerous investigations showed that galaxy populations in rich cluster environments were dominated by red elliptical and S0 galaxies \\citep{wolf1901, smith1935, zwicky1942, oemler1974}.  \\citet{oemler1974}, in particular, postulated that more evolved clusters hosted a larger proportion of early-type galaxies, suggesting a link between galaxy properties and the buildup of clusters.  \\citet{dressler1980} subsequently studied a diverse set of clusters and made the key discovery that at a fixed local density, relaxed and un-relaxed clusters host a similar fraction of early-type galaxies and that this fraction increased at higher local densities.  Dressler's morphology-density relation (MDR) challenged the long-standing role of ram-pressure stripping \\citep{gunn1972} in transforming field spirals into cluster early-types galaxies, and instead pointed to the importance of local processes.  In recent years, studies have probed into the lower density environments that are characteristic of groups \\citep{postman1984,zabludoff1998} and the field galaxy population  \\citep{kauffmann2004, baldry2006}, and found the local environment to be an important factor there as well.  In general, these studies found the {\\em lowest} density environments to be comprised of lower mass, blue, late-type galaxies, a stark contrast to the types of galaxies found at higher densities.  Studies of galaxy properties and their dependencies on local environment are key to understanding the mechanisms that are responsible for galaxy evolution.  For example, in the high density cluster environment processes like ram pressure stripping \\citep{gunn1972} and harassment \\citep{moore1996} may suppress star formation and consequently evolve galaxies towards the red-sequence.  At intermediate densities that are characteristic of the group environment, galaxy-galaxy interactions and mergers are more probable, given the favorable relative velocities of galaxies \\citep{cavaliere1992}.   Early HST observations of distant clusters revealed tidal interactions of blue disk galaxies, implying a strong role for galaxy-galaxy interactions in forming present day cluster ellipticals \\citep{dressler1994,dressler1994b}.  An analysis of galaxy properties in different environments can therefore shed light on the processes responsible for the observed buildup of the red-sequence between $z \\sim 1$ and today \\citep{bell2004b, faber2007}.  For example, given that cluster galaxies comprise a small fraction of their $z \\sim 1$ sample, \\citet{cooper2006,cooper2007} argue against processes such as ram pressure stripping and harassment as the sole origin of the observed color-density trend found at low redshift.  Instead, they suggest that processes that form the high early-type fraction in $z \\sim 0$ groups of galaxies dominate red-sequence formation at higher redshifts.   The outskirts of {\\em massive} clusters present an ideal location to study the impact of environment on galaxy properties over a broad range of local densities: from the high density cluster core to the intermediate density groups in the outskirts and the low density surrounding field.  Furthermore, the findings of \\citet{holden2007} and \\citet{vanderwel2007b} suggest that the morphological makeup of galaxies changes dramatically at the intermediate densities found in the outskirts of clusters.  These mass-limited studies found that while in the cores of clusters the fraction of early-type galaxies is high and does not evolve significantly between $z \\sim 1$ and $z=0$, the early-type fraction (ETF) for galaxies in the low density field is lower and also does not evolve.  Because clusters grow by accreting galaxies from the field, and massive clusters are expected to grow by a factor of $\\sim 2-3$ in mass between $z \\sim 1$ and $z=0$ \\citep{wechsler2002}, the large reservoir of infalling galaxies from the field must transform in the outskirts of clusters in order to preserve the high ETF found in the high density core.  Thus, the outskirts of massive clusters provide an opportunity to assemble a large sample of galaxies in environments where we expect transformations to be occurring.  Studying the properties of galaxies in the environments of clusters at intermediate redshifts has several advantages.  At such redshifts, around $z \\sim 0.8$, at a lookback time of $\\sim 7$~Gyr, a number of clusters exist where the formation of a massive cluster out of the large scale structure can be witnessed.  This is an epoch of rapid buildup in the cluster population.  Observations of the evolution of the cluster mass function indicate that the {\\em number} of massive clusters increases by a factor of a few since such redshifts \\citep{vikhlinin2008}.  In addition, as noted above, simulations suggest that individual clusters grow in {\\em mass} by a factor of $\\sim 2-3$ over the same time period \\citep{wechsler2002}.  This changing landscape therefore allows one to study the properties of galaxy populations that will transform, evolve, and assemble into the present day cluster population.  On much smaller physical scales, the star forming activity in galaxies was also different at these intermediate redshifts.  In general, galaxies formed stars at a higher rate in the past \\citep{lilly1996, madau1998}, particularly higher mass galaxies that will have ceased star formation by the present epoch \\citep{cowie1996, noeske2007b}.  {\\em At a redshift of $z \\sim 0.8$, the contrast between star-forming and non-star-forming galaxies is accentuated, allowing us to better highlight those environments that influence star formation.}  We focus in this paper on a mass-limited sample for this important epoch of cluster mass growth.  Luminosity-limited samples selected in the rest-frame UV include many bright, low mass, star forming galaxies as a result of the large scatter in their mass-to-light ratios at a fixed galaxy stellar mass.  Selecting by stellar mass thus provides a more well-defined and homogeneous sample.  In this paper, we study the colors of galaxies in the $z=0.834$ cluster \\rxjfull\\ (hereafter \\rxj) and its outskirts.  This is a very rich cluster with extensive imaging and spectroscopic datasets that now cover a very wide area (out to $\\sim 6$~Mpc in radius).  As such it is an ideal target for a study of galaxy evolution and transformation over a wide range of densities, from the core to the field.  Our initial spectroscopic sample of members is currently $\\sim 750$ objects.  These galaxies are the focus of our first paper on this cluster.  Such samples are becoming large enough to provide robust constraints on cluster galaxy properties (e.g. mass functions, star formation rates, morphologies, environmental trends, etc).  The use of spectroscopic redshifts, as opposed to photometric ones, ensures that our membership list, and therefore results, do not suffer from the systematics that trouble photometric surveys.  Ultimately we expect to reach 2000-3000 spectroscopic members between two clusters at $z \\sim 0.8$.  The remarkable feat employs the new low-dispersion prism (LDP) in the Inamori Magellan Areal Camera and Spectrograph (IMACS) on the Baade 6.5~m telescope at Magellan (see below for more details), further enhancing our insight into galaxy evolution in a variety of environments at $z \\sim 0.8$.  In \\S~\\ref{observations} we describe the data employed in this wide-field observing campaign.  In \\S~\\ref{analysis} the derived properties of galaxies in our sample are discussed.  Galaxy colors and their trends with clustercentric radius and local density are presented in \\S~\\ref{results}.  The implications of these results are discussed in \\S~\\ref{discussion}.  We summarize our findings in \\S~\\ref{summary}.  Throughout this paper, magnitudes from Sloan filters will be on the AB system, and Johnson filters on the Vega system.  The use of both systems is a consequence of adhering to convention and making comparisons to previous work more manageable.  We assume a $\\Lambda$CDM cosmology with $H_0=70$~\\kmps\\ \\pmpc, $\\Omega_{M}=0.30$, and $\\Omega_{\\Lambda}=0.70$, resulting in a look-back time of 7.0~Gyr and an angular scale of 0.46~Mpc~arcmin$^{-1}$ at the cluster redshift, $z=0.834$. ",
        "conclusions": "\\label{summary} We studied the environmental dependence of galaxy colors in the $z=0.834$ galaxy cluster \\rxj\\ in order to investigate where galaxies transition from blue, late-type systems into red, early-type ones that dominate the central regions.  The cluster is a massive X-ray luminous system that is comprised of two cores.  The northern and southern cores have velocity dispersions $\\sigma=~$768~\\kmps\\ and $\\sigma=~$408~\\kmps, respectively, and overlapping projected virial radii of \\rtwoN~=~\\rtwoNval\\ and \\rtwoS~=~\\rtwoSval\\ and are separated in velocity along the line of sight by $\\Delta v \\sim 1600$~\\kmps.  With a low-dispersion prism (LDP) on IMACS, the multi-object spectrograph on the 6.5~m Baade (Magellan I) telescope, we have built one of the largest samples of {\\em spectroscopically} confirmed members for a cluster at intermediate redshift.  This was accomplished by targeting galaxies with \\ifil~$<23.75$~mag within a $D=27$\\arcmin\\ ($D \\approx 12$~Mpc at $z=0.834$) FOV encompassing the two cores.  Prior to this work, $\\sim$250 galaxies were known to be members of \\rxj\\ and its outskirts with redshifts in the range $0.80<z<0.87$.  We have added \\nldponly\\ new members.  The redshifts derived from the LDP show a scatter of $\\sigma = 0.018$, or 1\\% in $(1+z)$, about those derived from higher resolution spectra (\\S~\\ref{redshifts}), allowing us to localize galaxies within the superstructure of the cluster.  Stellar masses for galaxies were computed using the \\citet{bell2003} relation between rest-frame $B-V$ color and $M/L_{B}$ (\\S~\\ref{masses}).  We used a mass-limited sample of \\nmcut\\ galaxies with mass \\mcutrange\\ to conduct our study.  This mass threshold corresponds to the \\ifil-band magnitude along the red-sequence at which the spectroscopic completeness is $\\sim 50\\%$ (\\ifil~$\\sim 23$).  Because the spectroscopic completeness varies for galaxies of different mass, we applied completeness corrections when necessary (\\S~\\ref{completeness}).  We assigned galaxies to be blue or red based on their deviation from the red-sequence, with blue galaxies having \\ugz\\ colors $\\sim 0.12$~mag or more below the red-sequence.  This initial sample provides an opportunity to analyze cluster galaxy properties at intermediate redshift with substantial statistical significance, and is only a glimpse of what will be possible when our survey is complete and we are likely to have over $\\sim 2000$~members.  Our conclusions are the following: \\begin{enumerate} \\item The most massive galaxies (\\hirange) lie on the red-sequence at all clustercentric radii and all local densities, while less massive galaxies have a broad range in colors farther from the two cluster cores and in lower density environments.  In this massive cluster at $z \\sim 0.8$ most of the evolution due to environmental effects is occurring in the lower mass galaxies. \\item The fraction of galaxies that are on the red-sequence is at a maximum in the two cores of the cluster at $93 \\pm 3\\%$.  This red galaxy fraction declines at larger clustercentric radii until $R \\sim 3$~Mpc, beyond which it is roughly constant at $64 \\pm 3\\%$.  Our lowest density sample at $R > 3$~Mpc ($2~{\\rm Mpc^{-2}} < \\Sigma < 8~{\\rm Mpc^{-2}}$) has a red fraction of $45 \\pm 10\\%$.  These values are broadly consistent with the red galaxy fraction for the field at $z \\sim 0.8$. \\item At $R>3$~Mpc, where the correlation between clustercentric radius and density is nonexistent, we find that the red galaxy fraction increases with local density.  This implies that the local environment is a factor in galaxy evolution in the outskirts of the cluster. \\item The high density environments ($32~{\\rm Mpc^{-2}} < \\Sigma < 126~{\\rm Mpc^{-2}}$) in the outskirts ($R>3$~Mpc) have a red galaxy fraction that is similar to what is found in the two cores.  These galaxies are primarily located in two previously identified groups. \\item A large number of galaxies at $R>3$~Mpc are found in our intermediate density bin ($8~{\\rm Mpc^{-2}} < \\Sigma < 32~{\\rm Mpc^{-2}}$), where the red galaxy fraction is elevated relative to the value found for the lowest density bin ($2~{\\rm Mpc^{-2}} < \\Sigma < 8~{\\rm Mpc^{-2}}$).  These intermediate densities therefore represent a regime where galaxies above our mass-limit first begin to undergo transformations from star forming to quiescent as a result of the local environment.  Most of these galaxies are not associated with either of the large groups in the outskirts, but instead trace the surrounding filamentary structure. \\end{enumerate}  Exploring the outer parts of clusters provides opportunities to identify the regions where galaxies of different masses are growing in mass and/or transforming from blue, star-forming, late-type galaxies to red, quiescent, early-type galaxies.  The uniformity of the colors of the highest mass galaxies at $z \\sim 0.8$ suggests that the vast majority of their star formation has taken place already, even at this epoch, while the diverse colors of the somewhat lower mass sample indicate that the transformation process is still underway in the intermediate density regions in the outer parts of \\rxj. Studies of galaxies in such regions of transformation promise to provide great insight into the physical processes that drive galaxy evolution and build up the red-sequence population.  \\subsection{Future Work} We are in the process of expanding the present sample of galaxies in \\rxj\\ with the LDP and we are also targeting another cluster at $z \\sim 0.8$, \\msfull.  As we move forward, the LDP reductions and SED fitting are being refined to provide better constraints on redshifts and other galaxy properties.  We are also in the process of obtaining $K_s$-band imaging, which will allow us to improve our stellar mass estimates.  When completed, this dataset will yield a wealth of information concerning galaxies in overdense environments in and around clusters at intermediate redshift.  Between the two targeted fields, we expect to find 2000-3000 galaxies in the clusters and their outskirts, allowing us to study the environmental dependence of mass functions with $M^{\\star}$ and $\\alpha$ to significantly higher precision than at present, possibly as high as $\\pm 0.1$~dex and $\\pm 0.15$ respectively.  We will also explore stellar population parameters from SED fitting, an analysis of blue galaxies down to low masses, and many other exciting aspects of cluster galaxies at an epoch some 7-8~Gyr ago when clusters were rapidly building up."
    },
    "0812/0812.2398_arXiv.txt": {
        "abstract": "{} {Rotation periods of cometary nuclei are scarce, though important when studying the nature and origin of these objects. Our aim is to derive a rotation period for the nucleus of comet C/2004 Q2 (Machholz).} {C/2004 Q2 (Machholz) was monitored using the Merope CCD camera on the Mercator telescope at La Palma, Spain, in January 2005, during its closest approach to Earth, implying a high spatial resolution (50\\,km per pixel). One hundred seventy images were recorded in three different photometric broadband filters, two blue ones (Geneva U and B) and one red (Cousins I). Magnitudes for the comet's optocentre were derived with very small apertures to isolate the contribution of the nucleus to the bright coma, including correction for the seeing. Our CCD photometry also permitted us to study the coma profile of the inner coma in the different bands.} {A rotation period for the nucleus of P\\,=\\,9.1$\\pm$0.2\\,h was derived. The period is on the short side compared to published periods of other comets, but still shorter periods are known. Nevertheless, comparing our results with images obtained in the narrowband CN filter, the possibility that our method sampled P/2 instead of P cannot be excluded. Coma profiles are also presented, and a terminal ejection velocity of the grains v$_{\\rm gr}$\\,=\\,1609\\,$\\pm$\\,48\\,m\\,s$^{-1}$ is found from the continuum profile in the I band.} {} ",
        "introduction": "Comets are at the same time both the largest and the smallest celestial bodies in our solar system. While the nucleus of a comet is typically 1 to 15\\,km, the ion tail can easily extend over more than 1\\,AU \\citep[e.g.][]{jones00}. Cometary nuclei are chemically and morphologically complex bodies, consisting of non-volatile ices and frozen grains. When these nuclei approach the sun, gases start to sublimate from them, forming the cloud of gas and dust known as the coma. For comets in the inner solar system, this coma is much brighter than the nucleus, making the nucleus hidden from our view. Yet, information on the dynamics of this nucleus is important to study not only the nature and origin of these objects, but also more generally the origins of our solar system itself. In the first place, it is of a statistical importance to determine cometary periods, since a representative sample of periods could give important clues on possible statistical dependencies with other properties, like orbital period or mass, and with similar objects like asteroids. Moreover, information on the rotational state of a comet gives important clues on the recent history of the comet; fast rotation for example can point to strong recent activity in which the comet changed its own rotation state, or alternatively (but less likely) to a recent interaction with other solar system objects. Finally, rotational periods are also an important parameter in the development of sublimation models for nuclei \\citep{schleicher98}.  \\begin{table} \\caption{Basic properties of comet C/2004 Q2 (Machholz).}\\label{tab:basicprprts} \\begin{center} \\begin{tabular}{ll} \\hline\\hline \\multicolumn{2}{c}{\\bf Comet C/2004 Q2 (Machholz)}\\\\ \\hline Discovery            & August 27, 2004\\\\ Earth passage    & $\\Delta$\\,=\\,0.35\\,AU on January 5, 2005\\\\ Perihelion & $r_h$\\,=\\,1.21\\,AU on January 25, 2005\\\\ Orbital period               & $\\sim$117000 year \\\\ Inclination & 38.6$^{\\circ}$ with ecliptic\\\\ H$_2$O prod. rate & 2.6 10$^{29}$\\,molec.\\,s$^{-1}$ at $\\Delta$\\,=\\,0.394\\,AU ($\\dag$)\\\\ \\hline \\multicolumn{2}{l}{Refs.: Horizons On-Line Ephemeris System of JPL}\\\\ \\multicolumn{2}{l}{(NASA) except $\\dag$: \\citealt{bonev06}}\\\\ \\end{tabular} \\end{center} \\end{table}  In this paper, we present a study of the Oort Cloud comet C/2004 Q2 (Machholz). Some basic properties of this comet are summarised in Table~\\ref{tab:basicprprts}. It was discovered on August 27, 2004 by Don E. Machholz \\citep{machholz04}, and  reached a visual brightness of $\\sim$3.5 mag. at its closest approach to Earth, making it one of the most spectacular astronomical events of that year. The observations presented in this paper are taken around this close encounter, providing a high spatial resolution. In this paper, we concentrate both on the nucleus, and on the inner (10$^5$\\,km) coma. Our observational data consist of CCD images in three photometric bands. Our first objective is to derive a rotation period of the nucleus from our intensive photometric monitoring, while our second goal is to study the profiles of the inner coma. These profiles are relatively simple to determine, and directly lead to some interesting properties, like the terminal ejection velocity of the dust grains.  Published rotation periods of cometary nuclei are scarce \\citep[see][for a summary]{samarasinha04}, since the nucleus of a comet situated in the inner solar system is always hidden in an opaque and overwhelming coma. Several techniques have been developed, each with their own drawbacks. One possible technique is to search for structures in the inner coma that emerge on the nucleus, like fans or jets, and search for periodicities in their morphology. This technique was employed to search for the rotation state of C/2004 Q2 (Machholz) by \\citet{farnham07} and will be discussed later in this paper (Sect.~\\ref{subsect:rotltrtr}). The technique used in this paper is aperture photometry using extremely small aperture radii. The idea is that in the vicinity of the optocentre of the coma, a small but observable contribution of the nucleus to the total brightness is expected. In the formalism of \\citet{licandro00}, if the brightness measured in an aperture $\\rho$ around the optocentre is given by \\begin{displaymath} B(\\rho) = B_N + B_C(\\rho) \\end{displaymath} where $B_C(\\rho)$ is the brightness of the coma captured in a radius $\\rho$ and $B_N$ the nucleus brightness. The amplitude of the light curve is then \\begin{equation} \\Delta m (\\rho) = -2.5 \\log [{\\rm min}(B_{N}) + B_C(\\rho)] / [{\\rm max}(B_{N}) + B_C(\\rho)]. \\label{eq:licandroqtn} \\end{equation} Note that in this equation, we implicitly assume that $B_C(\\rho)$ is constant during the time span of the observations. In practice, changes in the activity can imply large coma variabilities, especially near perihelion. Nevertheless, from Eq. (\\ref{eq:licandroqtn}), it is easily seen that when taking large apertures, the contribution of the coma $B_C(\\rho)$ cancels out the amplitude of the light curve induced by the variability of the nucleus. Therefore, very small apertures around the comet's optocentre are usually employed to detect variability induced by the nucleus \\citep{jewitt85, jewitt90, meech93, meech97}. However, a problem arises when taking apertures smaller than the seeing disc of the observation. This problem was studied in detail by \\citet{licandro00}, who dubbed this problem the {\\em seeing effect}. They found that several published periods were in fact contaminated, or even attributable, to this seeing effect. In our study, we will take this effect into account.  \\begin{figure*} \\begin{center} \\resizebox{16 cm}{!}{\\rotatebox{-90}{\\includegraphics{9225fig1.ps}}} \\end{center} \\vskip -.8 cm \\caption{The low-resolution spectrum of comet C/2004 Q2 (Machholz) taken with the EFOSC2 instrument taken with the ESO 3.6\\,m at La Silla on 9 december 2004. See the text for the line identification. The passbands of the three photometric filters of the Merope instrument used in this paper are shown in red. It is obvious that these broad filters cover a mix of continuum and emission features, with the I filter having the most emission-free passband. The two spectra are ESO archive data, and are taken under programme 274.C-5020 (P.I. Weiler).}\\label{fig:lowrsltn} \\end{figure*}  The reduction of the data, as well as its correction for the seeing effect, are an important part of our work, since it will ultimately determine the quality of our analysis. Therefore, the technical part of our analysis will be a considerable fraction of this paper: The observations are described in Sect.~\\ref{sect:observtns}, while the reduction is discussed in Sect.~\\ref{sect:reductn}; Sect.~\\ref{sect:extrctn} is devoted to the extraction of the cometary magnitudes, including the correction for the seeing effect; Next, the variability analysis and its results are presented in Sect.~\\ref{sect:variabilitynlss}; In Sect.~\\ref{sect:comaprfls} we construct the coma profiles in the three filters that we used for this programme. The general discussion of the obtained results is found in Sect.~\\ref{sect:generaldscssn}; We end with the conclusions (Sect.~\\ref{sect:finalcnclsns}).  \\begin{table} \\caption{Observational log. For each night, the number of observations is given for each filter U, B and I. Also the heliocentric ($r_h$) and geocentric ($\\Delta$) distances are given (in AU, at 0:00\\,UT), together with the correction factor $f_a$ (see Eqn. \\ref{eqn:fa} in Sect.~\\ref{sect:distancecrrctn}).}\\label{tab:observationallg} \\begin{tabular}{ccccccc} \\hline\\hline \\multicolumn{7}{c}{Observational log}\\\\ \\hline & n(U) & n(B) & n(I) & $r_h$ & $\\Delta$ & $f_a$\\\\ \\hline 2005-Jan-08 & 5 & 5 & 5 &1.2337 & 0.3485 & 1.00\\\\ 2005-Jan-09 & 9 & 9 & 9 &1.2305 & 0.3500 & 1.00\\\\ 2005-Jan-10 &10& 9 & 8 &1.2274 & 0.3520 & 1.01\\\\ 2005-Jan-11 & 8 & 9 & 8 &1.2245 & 0.3546 & 1.02\\\\ 2005-Jan-12 & 8 & 8 & 8 &1.2219 & 0.3577 & 1.03\\\\ 2005-Jan-13 & 8 & 8 & 8 &1.2194 & 0.3613 & 1.05\\\\ 2005-Jan-14 & 6 & 6 & 6 &1.2171 & 0.3654 & 1.07\\\\ 2005-Jan-15 & 5 & 5 & 5 &1.2150 & 0.3699 & 1.09\\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "\\label{sect:finalcnclsns} Comet C/2004 Q2 (Machholz) was monitored using the Merope CCD camera on the Mercator telescope at La Palma, Spain, in January 2005, during its closest approach to Earth. 170 Images were recorded in three different bands. Our main goal was to find a rotation period of the nucleus through aperture photometry using very small apertures. This technique was already succesfully applied to other bright comets, but it was recently criticised by \\citet{licandro00}. These authors studied the effect of aperture radii smaller than the seeing disc, and concluded that many periods published in literature that make use of this technique are contaminated by this {\\em seeing effect}. In our analysis of comet C/2004 Q2 (Machholz), we took this effect into account by degrading our images to the least favourable seeing. Luckily, since the ambient conditions during our observational run were very good, only the images with a superb seeing had to be degraded slightly.  We searched in each filter and three aperture radii (1, 3 and 6 pixels) for frequencies between 0 and 20\\,c\\,d$^{-1}$ (i.e. periods down to 1.2\\,h), and this with different algorithms. After taking into account that a long term trend is in our data \\citep[which is possibly linked to a solar wind event, see][]{degroote08}, a period near 9.1$\\pm 0.2$\\,h was found. While this rotation period is in perfect agreement with the one of \\citet{sastri05}, it clearly deviates from the 17.6\\,h found by \\citet{farnham07}. However, the methodology of both methods could explain the possibility that our method in fact finds P/2 and not P.  Also coma profiles for the inner coma ($\\log \\rho$\\,$<$\\,4.6, $\\rho$ in km) were constructed for each photometric band. In the I-band, a clear $m$\\,=\\,$-$1 for $\\rho$\\,$<$\\,10$^{4}$\\,km and $m$\\,=\\,$-$1.5 for $\\rho$\\,$>$\\,10$^{4}$\\,km was found, in agreement with a {\\em steady state} coma without creation or destruction of particles. From this regime change, a terminal ejection velocity of the grains v$_{\\rm gr}$\\,=\\,1609\\,m\\,s$^{-1}$ was derived. In the other filters, smaller $m$-values were found, due to the contamination of strong gas features in those bands, while the I filter is relatively free from those emissions."
    },
    "0812/0812.3577_arXiv.txt": {
        "abstract": "\\noindent The general solution of the stationary Schr\\\"odinger equation for the associated Lam\\'e potentials with an arbitrary real energy is found. The supersymmetric partners are generated by employing seeds solutions for factorization energies inside the gaps. ",
        "introduction": "In quantum mechanics the exactly solvable models (ESM) are essential: since the complete physical information is encoded in few analytic expressions, they are ideal to test the convergence of numerical methods. Moreover, the ESM constitute the starting point for applying perturbation techniques. There are some other models, intermediate between the exactly solvable ones and those which can be solved just numerically, which are known nowadays as quasi-exactly solvable (QES). For them there exist analytical expressions encoding just partial physical information, e.g., for only a part of the Hamiltonian spectrum \\cite{us94}. Along the years several systems have been identified as belonging to the QES class, in particular, there was a dominant conviction that the associated Lam\\'e potentials were QES \\cite{rkp05}. However, very recently it was shown that the associated Lam\\'e potentials for $(m,\\ell)$ integers are exactly solvable, in the sense that the stationary Schr\\\"odinger equation admits analytic solutions for any value of the energy parameter \\cite{fg05,fg07}. The initial motivation to look for this result was the need to implement supersymmetric quantum mechanics for generating new exactly solvable models (periodic and asymptotically periodic) \\cite{df98,ks99,fnn00,fmrs02a,fmrs02b,sgnn03}. In order to implement non-singular transformations of general type, the explicit expressions for the Schr\\\"odinger solutions in the gaps as well as the band edge eigenfunctions were required \\cite{fmrs02a,fmrs02b}.  In the next section we will discuss two techniques to find the general solution for the associated Lam\\'e equation for an arbitrary value of the energy parameter: first an ansatz based procedure and then a systematic technique related to the well known Frobenius method. In section 3 we will apply the supersymmetry transformations for generating new periodic and asymptotically periodic potentials which are almost isospectral to the initial associated Lam\\'e potential. In section 4 our general results will be illustrated through the particular case characterized by $(m,\\ell) = (3,2)$. Our conclusions will be finally given at section 5. ",
        "conclusions": "We have shown that the associated Lam\\'e potentials for integers values of the parameter pair $(m,\\ell)$ are exactly solvable. When using as seeds Bloch-type solutions inside the gaps new exactly solvable periodic potentials are generated. On the other hand, for seeds chosen as general linear combinations of Bloch-type solutions, the SUSY partner potentials become asymptotically periodic, the corresponding spectra having bound states embedded into the gaps. These potentials are interesting since the new levels could work as intermediate transition energies for the electrons to jump between the energy bands.  \\begin{figure}[ht] \\epsfig{file=fig2a.eps, width=7cm} \\hskip1cm \\epsfig{file=fig2b.eps, width=7cm} \\caption{\\small Second-order SUSY partners (black curves) for the associated Lam\\'e potential (gray curves) with $m=3, \\ \\ell=2$, $k^2=0.9$. (a) Periodic case generated through two Bloch solutions $u_1(x) = \\psi_1^+(x)$, $u_2(x) = \\psi_2^+(x)$ for $\\epsilon_1 = 10$, $\\epsilon_2 = 10.1$. (b) Asymptotically periodic case in which the linear combinations $u_1(x) = \\psi_1^+(x) + \\psi_1^-(x)$, $u_2(x) = \\psi_2^+(x)-1.5 \\ \\psi_2^-(x)$ for the same $\\epsilon_{1,2}$ are used.} \\end{figure}  We would like to end up this paper by making a historical precision concerning SUSY techniques applied to periodic potentials (see also the discussion in \\cite{cjp08}). It is a fact that the recent interest on the subject \\cite{df98,ks99,fnn00,fmrs02a,fmrs02b,sgnn03,fg05,fg07,gin06,imv08,cjnp08} was catalyzed by Dunne and Feinberg discovery of the self-isospectrality for the Lam\\'e potentials with $m=1$, induced by the first-order SUSY transformation which employs as seed the ground-state eigenfunction. Self-isospectrality means, in particular, that the SUSY partner potential becomes just a displaced version of the initial one, by half the period in \\cite{df98}. Soon it was realized that the self-isospectrality in which the new potential becomes the initial one displaced by any real number arises as well for the Lam\\'e potential with $m=1$, the seeds employed being Bloch-type solutions associated to factorization energies in the gaps \\cite{fmrs02a,fmrs02b}. More general SUSY transformations, either of higher order or involving general solutions of the Schr\\\"odinger equation for a given factorization energy, have been also introduced \\cite{fmrs02a,fmrs02b,sgnn03}. However, it is worth to note that there are several interesting works, previous to Dunne and Feinberg paper, in which the SUSY techniques were applied to periodic potentials (see e.g. \\cite{tr89,bm85}). In particular, it is remarkable the work of Braden and Macfarlane in which the self-isospectrality for the Lam\\'e potential with $m=1$ was discovered for the first time \\cite{bm85}. This is a typical story of a discovery followed by a later rediscovery, which often arises in science. Our opinion is that both works are valuable, complementary to each other, and hence they are worth to be studied in detail."
    },
    "0812/0812.1966.txt": {
        "abstract": "In the absence of any compelling physical model, cosmological systematics are often misrepresented as statistical effects and the approach of marginalising over extra nuisance systematic parameters is used to gauge the effect of the systematic. In this article we argue that such an approach is risky at best since the key choice of function can have a large effect on the resultant cosmological errors.  As an alternative we present a functional form filling technique in which an unknown, residual, systematic is treated as such. Since the underlying function is unknown we evaluate the effect of every functional form allowed by the information available (either a hard boundary or some data). Using a simple toy model we introduce the formalism of functional form filling. We show that parameter errors can be dramatically affected by the choice of function in the case of marginalising over a systematic, but that in contrast the functional form filling approach is independent of the choice of basis set.  We then apply the technique to cosmic shear shape measurement systematics and show that a shear calibration bias of  $|m(z)|\\ls 10^{-3}(1+z)^{0.7}$ is required for a future all-sky photometric survey to yield unbiased cosmological parameter constraints to percent accuracy.  A module associated with the work in this paper is available through the open source {\\tt iCosmo} code available at {\\tt http://www.icosmo.org}. ",
        "introduction": "\\label{Introduction} %cosmology statistics Cosmology is entering a formative and crucial stage, from a mode in which data sets have been relatively small and in which the statistical accuracy required on parameters was relatively low, into a regime in which the data sets will be orders of magnitude larger and the statistical errors required to reveal new physics (for example modified gravity -- Heavens et al, 2007, Kunz \\& Sapone, 2007; massive neutrinos -- Kitching et al., 2008c, Hannestad \\& Wong, 2007, Cooray, 1999, Abazajian \\& Dodelson, 2003; dark energy -- Albrecht et al., 2006, Peacock et al., 2006) are smaller than any demanded thus far. The ability of future experiments to constrain cosmological parameters will not be limited by the statistical power of the probes used but, most likely, by systematic effects that will be present in the data, and that are inherent to the methods themselves.  The problem that we will address is how the final level of systematics, at the cosmological parameter estimation stage, should be treated. This problem is of relevance to all cosmological probes, some examples include weak lensing and intrinsic alignments (Heavens, Refregier \\& Heymans, 2000; Crittenden et al., 2000; Brown et al., 2002; Catelan, et al., 2001; Heymans \\& Heavens, 2003; King \\& Schneider, 2003; Hirata \\& Seljak, 2004, Bridle \\& King, 2007; Bridle \\& Abdalla, 2007); baryon oscillations and bias (e.g. Seo \\& Eisenstein, 2003), X-ray cluster masses and the mass-temperature relation (e.g. Pedersen \\& Dahle, 2007) to name a few.  The general thesis we advocate in this article is that the standard approach to systematics, that of assuming some parameterisation and fitting the extra parameters simultaneously to cosmological parameters (e.g. Kitching et al., 2008a; Bridle \\& King, 2007; Huterer et al., 2006; for weak lensing analyses), both misrepresents a systematic effect as a statistical signal and more importantly is not robust to the choice of parameterisation.  As an alternative we will present a method, `form filling functions', in which a systematic is treated as such: an unknown function which is present in the data. By exhaustively exploring the space of functions allowed by either data, simulations or theory the effect of a systematic on cosmological parameter estimation -- a bias in the maximum likelihood -- can be fully characterised. This is a natural extension of the work presented in Amara \\& Refregier (2007b). We will present this using a simple toy model to explain and demonstrate the essential aspects of the formalism. We will then apply this technique to the problem of shape measurement systematics in weak lensing.  We begin in Section \\ref{Alternative Approaches to Systematic Effects} by categorising the different approaches to systematics that can be taken, Section \\ref{The General Formalism} introduces the parameter estimation formalism and how systematics can be included in a number of alternative ways. We will then introduce a toy model that will then be used to introduce `form filling functions' in Section \\ref{Functional Form Filling} where we will also compare the standard approach to systematics to the one taken here. The application to weak lensing shape measurement systematics is presented in Section \\ref{An Application to Cosmic Shear Systematics} and conclusions will be discussed in Section \\ref{Conclusion}.  %---------------------------------------------- ",
        "conclusions": "\\label{Conclusion}  In this article we have presented a method that allows one to move beyond the tendency to treat a systematic effect as a statistical one. When a systematic is treated as a statistical signal extra parameters are introduced to describe the effect, and then these extra nuisance parameters are marginalised over jointly with cosmological parameters. We have shown that even in a very simple toy model such an approach is risky at best. The results being highly dependent on the choice of parameterisation, and in the limit of a large number of parameters any statistical signal on cosmological parameters can be completely diluted. One could use marginalisation if there exists a compelling physical theory for the systematic, however if the functional form is merely phenomenological or if it contains a truncated expansion then one should use caution. Even in the case that confidence is high with respect to the assumed functional form any residual must be investigated correctly.  As an alternative we advocate treating a systematic signal as such, an unknown contaminant in the data. We address the situation where we have at least some extra information on the systematic, either from some theory or simulation that may provide a hard boundary within which a systematic must lie, or from some external data set. We leave the case of `entangled systematics' (where there is no extra information and where the systematic depends on the cosmological parameters) to be investigated in Amara et al. (in prep). Throughout we have introduced each concept using a simple toy model.  Since the systematic is treated as a genuine unknown we must address every possible functional form that is allowed, by either the hard boundary or the extra systematic data. To do this we use complete basis sets and randomly sample from the space of coefficients until all functional behaviour has been experienced at every point with the hard boundary, or within a few $\\sigma$ of the data. We have shown that such ``functional form filling'' can be achieved by this technique using either Chebyshev, Fourier or tophat basis sets. By treating each function as a possible systematic, a bias in the maximum likelihood value is introduced whilst the marginal error stays the same. Throughout we have shown that as long as functional filling is achieved all conclusions on the magnitude of a systematic effect are independent of the choice of complete basis set -- though binning is highly inefficient in terms of computational time, and could be labelled as an unphysical basis set.  For the case of a hard boundary every function is given an equal probability and as such a maximum bias exists. For the case of extra systematic data we show that a probability can be assigned to each function, and bias, allowing for the production of robust systematic probability contours.  %lensing application We have made a first application of hard boundary functional form filling by addressing multiplicative systematics in cosmic shear tomography; we have left an application of the data bound for future work. We address a lensing systematic, that can result from shape measurement or PSF reconstruction inaccuracies, that has the same overall shape as the lensing power spectrum but is multiplied by some extra function. This is commonly represented using a multiplicative function $m(z)$ (Heymans et al, 2006; Massey et al., 2007) where $\\gamma^{\\rm sys}=m(z)\\gamma^{\\rm true}$. For a DUNE/Euclid type survey, we find that in order for the systematic on $w_0$ to yield a bias smaller than the marginal error the overall magnitude of $m(z)$ should be $m_0\\leq 8\\times 10^{-4}$ for a linear scaling in $(1+z)$, but as the magnitude of the redshift scaling is relaxed then the requirement on $m_0$ increases to $m_0\\leq 2\\times 10^{-3}$. The most recent shape measurement methods have $m\\sim 2\\times 10^{-3}$ (e.g. using the {\\tt lensfit} method Miller et al., 2007; Kitching et al., 2008b) and have a small scaling as a function of magnitude/size. The results shown here then, coupled with some currently available shape measurement techniques imply that constraints on cosmological parameters from tomographic weak lensing surveys should be robust to shape measurement systematics (with the caveat that PSF calibration and galaxy size dependence of the bias have been neglected here).  The techniques introduced here should have a wide application in cosmological parameter estimation: in any place in which there is some signal with some extra information on a systematic. For example baryon acoustic oscillations and galaxy bias, CMB and foregrounds, galaxy clusters and mass selection, supernovae and light rise times, weak lensing and photometric redshift uncertainties. A sophistication of these techniques could assign different weights/prior probabilities to particular functional forms, that are known \\emph{a priori} to have a large/small effect on cosmological parameter determination, such weights could come from either theoretical or instrumental constraints.  Cosmology is entering into a phase in which the statistical accuracy on parameters will be orders of magnitude smaller than anything achieved thus far. But we must take care of systematics in a rigorous way to be sure that our statistical constraints are valid. There exists the pervading worry that the functional forms used to parameterise systematics are not representative of the true underlying nature of the systematic and that something may have been missed. In such a scenario one should always add the warning that ``cosmological constraints are subject to the assumption of the systematic form''.  Here we have presented a way to address systematics in a way that requires no external assumptions, allowing for robust and rigorous statements on systematics to be made.  %----------------------------------------------"
    },
    "0812/0812.1434_arXiv.txt": {
        "abstract": "The methods and techniques for the slitless spectroscopy software aXe, which was designed to reduce data from the various slitless spectroscopy modes of Hubble Space Telescope instruments, are described. aXe can treat slitless spectra from different instruments such as ACS, NICMOS and WFC3 through the use of a configuration file which contains all the instrument dependent parameters. The basis of the spectral extraction within aXe are the position, morphology and photometry of the objects on a companion direct image. Several aspects of slitless spectroscopy, such as the overlap of spectra, an extraction dependent on object shape and the provision of flat-field cubes, motivate a dedicated software package, and the solutions offered within aXe are discussed in detail. The effect of the mutual contamination of spectra can be quantitatively assessed in aXe, using spectral and morphological information from the companion direct image(s). A new method named 'aXedrizzle' for 2D rebinning and co-adding spectral data, taken  with small shifts or dithers, is described. The extraction of slitless spectra with optimal weighting is outlined and the correction of spectra for detector fringing for the ACS CCD's is presented. Auxiliary software for simulating slitless data and for visualizing the results of an aXe extraction is outlined. ",
        "introduction": "A spectrum produced by a dispersing element, usually a prism or a grism, without a slit or slitlets has a number of characteristics that require specialized data reduction tools. Slitless spectroscopy is primarily a survey tool, since the spectra of all objects within a given area defined by the telescope-instrument-detector combination are recorded. While this survey aspect brings advantages, it also has some drawbacks. There can be no prior selection of target objects, and contamination, which is an overlapping of spectra, occurs frequently. The lack of defining slits implies that for extended objects the spectral resolution is controlled by the object size, more specifically the component of the object size in the dispersion direction. Moreover the shape of the objects, which in general are not aligned with the dispersion direction, must be taken into account during the spectral extraction. Since in slitless spectroscopy each pixel can receive radiation of any wavelength (within the sensitive range of the instrument), special techniques for flat-fielding and fringing are required. \\begin{deluxetable}{llccc} \\tablecaption{The HST slitless spectroscopic modes to which aXe has been applied.\\label{spec_modes}} \\tablehead{\\colhead{Instrument/Camera}&\\colhead{Disperser} & \\colhead{Wav. Range [\\AA]} & \\colhead{Dispersion [\\AA/pixel]} & \\colhead{FOV\\tablenotemark{1} [arcsecond]}} \\startdata ACS/WFC     & G800L     & $5500-10500$     &     $38.5$   & $202\\times 202$\\\\ ACS/HRC     & G800L     & $5500-10500$     &     $23.5$   & $29\\times 26$\\\\ ACS/HRC     & PR200L     & $1700-3900$     & 20[@2500\\AA]\\tablenotemark{2} & $29\\times 26$\\\\ ACS/SBC     & PR130L     & $1250-1800$     & 7[@1500\\AA]\\tablenotemark{2} & $35\\times 31$\\\\ ACS/HRC     & PR110L     & $1150-1800$     & 10[@1500\\AA]\\tablenotemark{2} & $35\\times 31$\\\\ NICMOS/NIC3 & G141      & $11000-19000$    & 80.0         & $51\\times 51$          \\\\ WFC3/UVIS\\tablenotemark{3} & G280      & $2000-4000$      & 13.0         & $160\\times160$\\\\ WFC3/IR\\tablenotemark{3} & G102      & $7800-10700$     & 25.0         & $123\\times137$\\\\ WFC3/IR\\tablenotemark{3}& G141      & $10500-17000$    & 47.0         & $123\\times137$\\\\ \\enddata \\tablenotetext{1}{Field-of-View} \\tablenotetext{2}{On account of the dispersion change with wavelength for prism spectra, the resolution is accompanied by the wavelength it refers to (see also Eqn.\\ \\ref{eq_prdisp})} \\tablenotetext{3}{Prior to the WFC3 installation in Servicing Mission 4, aXe was used in the reduction of ground based calibration data.} \\end{deluxetable}  Slitless spectroscopy has been employed for surveys from the ground such as for the Hamburg/ESO objective-prism survey survey \\citep{wi96}, the UK Schmidt Telescope (Clowes et al.\\ 1980; Hazard et al.\\ 1986) and the Quasars near Quasars survey \\citep{wo08}. Early efforts to use this technique for the determination of stellar radial velocities are described in \\cite{fe69}. However, the strong disadvantage of slitless spectroscopy from the ground is the high background, since the lack of a slit implies that each pixel receives the full transmission of the dispersing element. The background can be reduced by choosing a high resolving power disperser or adding a narrow band filter to the light path, such as for emission line studies on known lines (e.g. Salzer et al.\\ 2000; Jangren et al.\\ 2005).  Deep slitless surveys from the ground, such as searches for high redshift galaxies at $3 < z < 7$, are still valuable (see Kurk et al.\\ 2004), however space based observations are particularly effective, since the background is many orders of magnitude lower than from  earth and there are no strong and variable atmospheric absorption and emission components. For the Hubble Space Telescope (HST), the high spatial resolution (through lack of atmospheric seeing) brings a further gain, since the background contribution to the extracted spectrum is reduced. Although the background in slitless mode is larger by the ratio of the spectrum length to the slit width with respect to slit spectroscopy, HST slitless spectroscopy can challenge the faint limits of ground-based telescopes \\citep[e.g. ][]{ma05}.  Surveys with HST in slitless mode have included the NICMOS/Hubble Space Telescope Grism Parallel Survey of H$\\alpha$ emission line galaxies in the redshift range 0.7 to 1.9 with the G141 grism \\citep{mc99}; the Space Telescope Imaging Spectrograph Parallel Survey (Gardner et al.\\ 1998; Teplitz et al.\\ 2003) of [O~II] and [O~III] emission line galaxies $0.43<z<1.7$; and the Advanced Camera for Surveys GRAPES study of the Ultra-Deep Field \\citep{pir04} with, for example, detection of Ly-$\\alpha$ and Lyman-break galaxies to $z\\sim6$ \\citep{ma05}; the PEARS survey in the GOODS North and GOODS South fields \\citep{ma08} and searches for high-z supernovae to $z\\sim1.6$ (Riess et al.\\ 2004; Strolger et al.\\ 2004).  After exploratory attempts to use standard spectral extraction tasks in public data reduction packages (e.g., IRAF, IDL and MIDAS) for slitless spectroscopic data, it was decided to develop a dedicated extraction package.These standard packages encapsulate an enormous amount of expertise for slit spectroscopy. However the lack of solutions for specific problems in slitless spectroscopy together with the demand for the processing of hundreds to thousands of spectra per image mandated a dedicated software package.  This software package, called aXe, was originally developed for Advanced Camera for Surveys (ACS) grism and prism spectroscopy (see Table \\ref{spec_modes} and Pasquali et al.\\ 2006). However by design aXe was intended to be instrument-independent, and this aim was achieved by placing all the instrument specific parameters into a single data file called a {\\it configuration file}. In addition to ACS, aXe has subsequently been successfully applied to NICMOS slitless data \\citep{fr08} and data obtained during the ground calibration campaigns of Wide Field Camera 3 (WFC3, Bond et al.\\ 2007, Kuntschner et al. 2008b; see also Table \\ref{spec_modes}).  This paper is devoted to a description of the aXe spectral extraction package. aXe does {\\it not}\\/ work like a data reduction pipeline, which has a static set of products as output. Depending on the data and the scientific questions to be answered, the user can reduce the data following different recipes and methods, and the subsequent Sections illustrate all the methods and techniques implemented in the current version aXe-1.6. Section \\ref{sect:axe_strat} is devoted to the extraction of spectra from an individual slitless image. Section \\ref{sect:contamination} explains the various techniques to treat the contamination problem, followed by the introduction of a 2D-resampling technique which aims at co-adding individual spectra of the same object in Section \\ref{sect:axedrizzle}. The aXe implementation of optimal weighting is illustrated in Section \\ref{optimal_weighting}, followed by a discussion of the effects of detector fringing (Section \\ref{sect:fringing}). The structure of aXe and related software packages for simulation and visualization are presented in Section \\ref{sect_axe_software_package}. The paper closes with a summary and options for further refinement of the package. ",
        "conclusions": "aXe is not the only software for extracting HST slitless spectra. For ACS/WFC, an independent method has been described in \\cite{dr05}. {\\tt MULTISPEC} \\citep{app05} and the extraction method described in \\cite{ch99} concentrate on STIS slitless spectroscopy. Quickly after the first release in 2002, aXe developed into the standard reduction package for all ACS slitless modes. The methods implemented in aXe specifically for slitless spectroscopy (e.g. flat-fielding, 2D-resampling with aXedrizzle) have proven to produce science-ready and well calibrated data for a number of large programmes such as GRAPES \\citep{pir04}, PEARS \\citep{ma08} and searches for high-z supernovae (Riess et al. 2004; Strolger et al. 2004).  By generating instrument-specific configuration and calibration files, aXe can also be applied to slitless spectroscopic data from other HST instruments (e.g. NICMOS) and even from ground based instruments such as the Wide Field Imager (WFI) mounted on the ESO 2.2-m-telescope. With small modifications, aXe has even been used to successfully reduce data taken with FORS2 MXU masks (see Kuntschner et al.\\ 2005 and K\\\"ummel et al.\\ 2006). aXeSIM is being used to simulate data for future space missions such as SPACE (Cimatti et al.\\ 2008; now renamed to EUCLID) and SNAP (\\anchor{http://snap.lbl.gov/}{http://snap.lbl.gov/}) that plan to use slitless spectroscopy.  The ability within aXe to automatically extract spectra from slitless data was important in the NICMOS Hubble Legacy Archive (HLA) project \\citep{fr08} and aXe served as the central component of the completely automatic Pipeline for Hubble Legacy Archive Grism data (PHLAG) (K\\\"ummel et al.\\ 2007; K\\\"ummel et al.\\ 2008). Due to the modular setup of aXe, additional tasks which handle some features specific to NICMOS could comfortably be added in order to achieve optimal results leading to an archive of several thousand near-infrared spectra (\\anchor{http://www.stecf.org/archive/hla/}{http://www.stecf.org/archive/hla/}).  Besides instrument-specific adjustments also the core of the aXe extraction software is continually being improved and further developed, with releases of new versions at an approximately annual rate. New features and developments that are currently being implemented for aXe (versions 1.7 and 1.8) and aXeSIM (version 1.2) include spectral extraction based on a slit geometry optimized for slitless spectroscopy and flux conversion which compensates for the size of the direct object (developed in the NICMOS HLA project, see Freudling et al.\\ 2008); the identification and exclusion of cosmics and hot pixels within aXedrizzle, similar to MultiDrizzle \\citep{koe06} in direct imaging; the integration of aXeSIM into STSDAS.  The application of aXe to HST data is going to continue with the installation of the Wide Field Camera 3 (WFC3, Bond et al.\\ 2007) with its three slitless spectroscopic modes (see Tab.\\ \\ref{spec_modes}) during Servicing Mission 4. As for ACS, aXe is the recommended tool for reducing WFC3 slitless spectroscopic images. From the ground calibration data collected during the Thermal Vacuum testing, initial aXe configuration and calibration files are already available (Kuntschner et al.\\ 2008b and Kuntschner et al.\\ 2008c) that were, for example, used in generating the aXeSIM simulations shown in Fig.\\ \\ref{fig1:axesim}. Improved calibration will result from on-orbit data, enabling aXe to extract spectra from all slitless modes of WFC3."
    },
    "0812/0812.4211_arXiv.txt": {
        "abstract": "{We discuss the present performance and the future perspectives of VLBI in the 3\\,mm to 0.85\\,mm observing bands (so called mm-VLBI). The availability of new telescopes for mm-VLBI and the recent technical development towards larger observing bandwidth and higher data-rates now allow to image with 3mm-VLBI hundreds of sources with high dynamic range. At 1.3\\,mm, pilot VLBI studies have proven detectability of the brightest AGN, and the existence of ultra-compact regions therein. In the next few years global VLBI imaging will be established also at 1.3\\,mm wavelength. With an angular resolution in the $10-20$\\,$\\mu$ mas range, future 1.3\\,mm- and 0.8\\,mm VLBI will be an extraordinarily powerful astronomical observing method, allowing to image the enigmatic `central engines' and the foot-points of AGN-jets in greater detail than ever possible before. A sufficiently large number of telescopes is a prerequisite for global aperture synthesis imaging. Therefore a strong effort is needed to make more telescopes available for VLBI at short millimeter and sub-millimeter wavelengths. In this context, the further VLBI upgrade of both IRAM telescopes and the outfit of the APEX telescope in Chile in preparation for later mm-/sub-mm VLBI with ALMA is of high scientific importance. With a sufficiently large mm-VLBI network, the micro-arcsecond scale imaging of the post-Newtonian emission zone around the event horizon/ergosphere of nearby super-massive Black Holes (such as e.g. Sgr\\,A*, M\\,87, ...) should become possible within the next few years. }  \\FullConference{The 9th European VLBI Network Symposium on The role of VLBI in the Golden Age for Radio Astronomy and EVN Users Meeting\\\\ September 23-26, 2008\\\\ Bologna, Italy}  \\begin{document} ",
        "introduction": "The physical processes of energy release and the details of the jet launching, particle and jet acceleration in the nuclei of active galaxies (AGN) is still enigmatic and leads to many questions not yet answered. One way to address this problem, is the direct imaging of the central regions of these objects, with the highest possible angular and spatial resolution, and at wavelengths short enough to look through existing source intrinsic opacity barriers. The angular resolution of an interferometer is proportional to the observing wavelength ($\\lambda$) and inversely proportional to the separation ($D$) between its outer elements. This leads to millimeter VLBI (mm-VLBI), which observes at the shortest possible wavelength using telescopes on ground ($D <$ Earth diameter), and to space-VLBI utilizing baselines larger than the Earth's diameter through an orbiting radio antenna in space (VSOP2, RADIOASTRON). The planned space-VLBI missions will give an angular resolution of up to $\\sim 40$\\,$\\mu$as at 43\\,GHz for VSOP2 and up to $\\sim 10$\\,$\\mu$as at 22\\,GHz for RADIOASTRON. These numbers compare well with the angular resolution of ground based mm-VLBI: $\\sim 45$\\,$\\mu$as for present day global VLBI at 86\\,GHz (GMVA), and  $\\sim 17$\\,$\\mu$as, respectively $\\sim 11$\\,$\\mu$as for future VLBI at 230\\,GHz and 345\\,GHz. Owing to its larger collecting area and observing bandwidth, it is likely that future mm-/sub-mm VLBI will be more sensitive than the two planned space-VLBI missions (limitations of space-antenna size and down-link data rates), allowing to study more objects. Because of their different frequency coverage and  sensitivity space-VLBI and mm-VLBI are very complementary. Both can reach the smallest spatial scales, however at very different frequencies and opacities, facilitating the detailed study of the spectral characteristics of the corresponding emission regions. In the more distant future, these regions may be also investigated with near-infrared interferometry at the VLTI (GRAVITY, 2012: resolution $\\sim 1$\\,mas, astrometric accuracy $\\sim 10$\\,$\\mu$as, \\cite{Eisenhauer08}) and with NASA's X-ray interferometry mission (MAXIM, 2020: resolution $\\sim 100$\\,$\\mu$as, \\cite{Cash05}). The common driver of all these projects is the desire to see and directly image the central Black Hole and the post-Newtonian region around it. ",
        "conclusions": ""
    },
    "0812/0812.1328_arXiv.txt": {
        "abstract": "We consider the possibility that astrophysical Black Holes (BHs) can violate the Kerr bound; i.e., they can have angular momentum greater than BH mass, $J > M$. We discuss implications on the BH apparent shape. Even if the bound is violated by a small amount, the shadow cast by the BH changes significantly (it is $\\sim$ an order of magnitude smaller) from the case with $J \\le M$ and can be used as a clear observational signature in the search for super--spinning BHs. We discuss briefly recent observations in the mm range of the super--massive BH at the center of the Galaxy, speculating on the possibility that it might violate the Kerr bound. ",
        "introduction": "Black Holes (BHs) are quite strange objects, which are devoid of a true internal structure and are completely defined by a few parameters~\\cite{c-book, fn-book}. In the case of BHs which we may possibly find in our Universe, the number of these parameters reduces to three: the mass $M$, the charge $Q$, and the spin $J$. In this paper we pay close attention to the possible spin of the BH while we set $Q=0$. At present good BH candidates~\\cite{narayan} include super--massive objects ($10^5 - 10^9$~$M_\\odot$) at the center of galaxies and stellar mass objects ($5 - 20$~$M_\\odot$) in X--ray binary systems. In both cases, we can infer their mass from dynamical arguments, studying the Newtonian orbital motion of stars or gas around them~\\footnote{It is also possible that a third category of BHs exists, intermediate mass BHs ($10^3 - 10^4$~$M_\\odot$), but so far we do not have much information about them and, in particular, there are no clear measurements of their mass~\\cite{cmiller}.}. A challenge today is to measure the spin of these objects as may be experimentally feasible in the near future. Here the difficulty is that we need to probe the spacetime close to the horizon, because spin effects are absent in Newtonian gravity and are suppressed at small velocities and large distances.   An important feature is that BHs are expected to respect the Kerr bound $J \\le M$. This is just the condition to have a horizon. If the Kerr bound were violated, instead of a rotating BH we would have a naked singularity. To see this we can examine the 4 dimensional Kerr--Neumann or Reissner--N{\\\"o}rdstrom solutions. The position of the horizon is given by the expression~\\cite{mtw,lppt} \\be R_{H} = M + \\sqrt{M^2 - Q^2 - J^2} \\, , \\label{r-hor} \\ee where $Q$ and $J$ are, respectively, electric charge and angular momentum of BH. It is clear that in 4D spacetime the horizon cannot be formed if \\be M < \\sqrt{Q^2 + J^2} \\, . \\label{mass-limit} \\ee In the absence of a horizon, there would be naked singularities which are not allowed. Indeed if condition~(\\ref{mass-limit}) is fulfilled, the Kerr--Newmann metric makes it possible to reach the physical singularity at $r=0$ from some large $r$ in finite time without crossing any horizon. One could thus consider closed time--like curves and violate causality (see e.g. section~66c of~\\cite{c-book} or ref.~\\cite{carter}). For this reason, usually some kind of cosmic censorship is assumed and naked singularities are forbidden~\\cite{penrose}.   However, it seems reasonable that the singularity at the center of BHs has actually no physical meaning and  it is just the symptom of the breakdown of classical General Relativity. First, it is difficult to believe that all the matter can collapse into an infinitesimal volume. Second, this is usually the kind of pathology which is expected to be solved at the quantum level. On these general grounds, one is tempted to argue that actually there is no singularity at the center and that the Kerr bound may be violated. In particular~\\cite{horava} discusses possible origins of the breach of the Kerr bound in string theory.   One more comment is in order here. Since for $J > M$ there is no horizon, Robinson's theorem does not hold~\\cite{robinson} and thus (at least in classical General Relativity) the super--rotating object might not be described by the Kerr metric. In quantum gravity we simply do not know what happens.   The most promising approach to measure the spin of BHs is often believed to be the study of emission lines (notably the fluorescent iron K$\\alpha$ at 6.4~keV), where $J$ may be deduced by fitting the shape of the line~\\cite{reynolds}. The method has some weak points. In particular, one has to assume some emissivity function (usually modeled as a power law in the radius) and that there is no emission inside the Innermost Stable Circular Orbit (ISCO). Relaxing these assumptions, one can find quite different results~\\cite{reynolds2}. Another common approach is the X--ray continuum fitting method, which is also based on the fact that the inner edge of the accretion disk is presumably the ISCO~\\cite{shafee1, shafee2, liu}. Here we need to know the distance of the BH candidate, its mass, and the disk inclination angle, but in a few cases there are sufficiently reliable estimates of these quantities. For our purposes, if we want to test the Kerr bound, both approaches do not look suitable, because there is no clear difference between a Kerr BH near extremality and one which violates the Kerr bound by a small amount: the radius of the ISCO is a continuous function of $J$, while we would like to observe some physical quantity which is discontinuous at $J = M$.   In this work we study the apparent shape of a super--spinning BH and we claim that the observation of its shadow could be used to test the Kerr bound. The shadow of the BH is the non--illuminated area seen by an observer if the BH is in front of a planar light source. In realistic situations, it is usually unlikely to have a bright source of this kind. Nevertheless, something very similar to the shadow can be observed if the BH is surrounded by an emitting medium (typically the accreting gas) which is optically thin (and this is always possible for enough high frequencies). Here an arbitrarily small violation of the Kerr bound makes the horizon disappear and changes significantly the apparent shape of the BH: now only the photons reaching the center of the BH are lost, while all the others, with turning points at finite distances from the center (or at distances larger than some scale coming from new physics), are not captured and can therefore come back to infinity. ",
        "conclusions": "In classical General Relativity, BH spinning more rapidly than the Kerr bound, i.e., spinning with $J > M$, would imply the existence of a naked singularity and the violation of causality. However, if quantum gravity effects can resolve the singularity, causality can be restored and we do not need the Kerr bound. Then super--spinning black holes may exist.   In this paper we have discussed how we may observationally identify a black hole which violates the Kerr bound. The key point is the absence of the horizon, which leads to a very different apparent shape for the black hole. By shape we mean essentially the cross section for capturing photons. If the black hole is surrounded by accreting gas which is optically thin, as we believe is the case for the black hole at the center of the Galaxy for sub--millimeter wavelengths, we can presumably see something similar to the black hole shadow. For the standard $J \\le M$, the precise measurement of the black hole spin is difficult because the image size is always about 10~$M$. On the other hand, the observational difference between BH with $J<M$ and with $J>M$ should be quite dramatic. The test of the Kerr bound can instead be relatively easy, because we have just to be able to distinguish an image of apparent size about 10~$M$ (for the case where the Kerr bound is satisfied) from one of about 2~$M$(where the Kerr bound is violated). A more detailed study of the picture is  necessary; in particular, the black hole image has to be evaluated in a more realistic framework, assuming some astrophysical model for the emitting region surrounding the black hole. This will be the subject of another work.   The possible violation of the Kerr bound does not strictly imply the breakdown of classical General Relativity, since the theory does not require $J \\le M$, but is surely something which would be unexpected in the standard framework and which demands new physics."
    },
    "0812/0812.2601_arXiv.txt": {
        "abstract": "Recently, deeper understanding of QCD emerges from the study of the AdS/CFT correspondence. New results include the properties of quark-gluon plasma and the confinement/deconfinement phase transition, which are both very important for the scenario of the QCD phase transition in the early universe. In this paper, we study some aspects of how the new results may affect the old calculations of the cosmological QCD phase transition, which are mainly based on the studies of perturbative QCD, lattice QCD and the MIT bag model.  \\vspace{0.5cm} \\noindent \\emph{Key words}: Confinement/deconfinement phase transition; Cosmological phase transition; Finite temperature QCD; Gauge/string duality\\\\ \\emph{PACS}: 98.80.Cq, 11.25.Tq, 25.75.Nq ",
        "introduction": "}  Phase transitions can produce relics, affect the anisotropies of the universe, or have other observable consequences; hence, it is very important in astrophysics. A particularly important phase transition is the QCD confinement/deconfinement phase transition, in which the deconfined quark-gluon plasma (QGP) phase transits to the confined hadronic phase. By assuming this phase transition is first order and that it has nonzero surface tension, it suffers chronologically the processes of supercooling, reheating, bubble nucleation, and may produce relics such as quark nuggets. For the up to date reviews of the cosmological QCD phase transition, see~\\cite{Schwarz:2003du,Boyanovsky:2006bf}.  Recently, deeper understanding of QCD emerges from the study of the anti-de Sitter/conformal field theory (AdS/CFT) correspondence~\\cite{Aharony:1999ti}. In its prototype version, type IIB superstring theory on $AdS_5 \\times S^5$ is dual to $\\mathcal{N} = 4$ $U(N_\\mathrm{c})$ super Yang-Mills (SYM) theory in $(3+1)$-dimensional spacetime~\\cite{Maldacena:1997re}. Generally speaking, conventional quantum field theories make sense only in the perturbative regions, where the 't Hooft coupling $\\lambda_4 = g_\\mathrm{YM}^2 N_\\mathrm{c}$ is small; however, the dual gravity theory is easy to handle when the supergravity (SUGRa) description becomes reliable, that is, in the strong coupling region. Hence, we can use the AdS/CFT correspondence to study field theory in the region where perturbative approaches are not applicable. It is also believed that the generalization of the AdS/CFT correspondence can realize some more realistic systems, such as QCD-like theories or QCD itself, which have running coupling constants (hence are not conformal), fundamental matter, and reproduce some phase transitions. Recent reviews on the connection between string theory and QCD can be found in~\\cite{Peeters:2007ab,Mateos:2007ay}.  New observational results from the relativistic heavy ion collider (RHIC)~\\cite{Gyulassy:2004zy} data tell us that the shear viscosity of the hot plasma is very small~\\cite{Teaney:2003kp}; thus the QGP at temperature $T \\gtrsim T_\\mathrm{dec}$ should in fact be strongly coupled~\\cite{Shuryak:2003xe,Shuryak:2004cy}, rather than asymptotic free as we used to think about it, where $T_\\mathrm{dec}$ is the critical temperature of the confinement/deconfinement phase transition. Hence, all phenomenological applications of QGP, which are based on perturbative QCD or the MIT bag model~\\cite{Chodos:1974je,DeGrand:1975cf}, should be reconsidered. These applications include the neutron stars/quark stars and the cosmological QCD phase transition. AdS/CFT provides an excellent tool to study them. Because of its strong interactive nature, it can explore the properties of QGP and the confinement/deconfinement phase transition, in both the high temperature and high baryon number density regions. However, in this paper, we will limit our focus on the property of high temperature region, which is important for the cosmological QCD phase transition.  The organization of this paper is as follows. In Sec.~\\ref{sec:QGP}, we present the results of QGP and the confinement/deconfinement phase transition from AdS/CFT. In Sec.~\\ref{sec:cosmology}, we study how these new results affect the conventional scenarios of the cosmological QCD phase transition, including the nucleation rate, the supercooling scale and the mean nucleation distance. We summarize our results in Sec.~\\ref{sec:conclusion}. We will set $\\hbar = c = k = 1$ throughout this paper. ",
        "conclusions": "}  In this paper, we discussed some implication of the new AdS/CFT results to the cosmological QCD confinement/deconfinement phase transition. We limit our discussion to the homogeneous nucleation case. The values of the hydrodynamical quantities, like the shear viscosity $\\eta$ or the bulk viscosity $\\zeta$, can significantly lower the prefactor $(\\kappa/2\\pi) \\Omega_0$ of the nucleation rate formula compared to the old estimations; however, they can hardly affect other characteristic parameters of this process, such as the supercooling scale $\\Delta = 1 - T_\\mathrm{f}/T_\\mathrm{dec}$ or the main nucleation distance $d_\\mathrm{nuc}$. The new EoS's, which differ from the MIT bag model, can affect the phase transition scenario mainly in two ways. (i) As most of these EoS's are comparatively more weakly first order than the bag model, it is not adequate to treat their latent heat $L_\\mathrm{h}$ as a constant. For some not-very-small supercooling, the effective latent heat is always much smaller. Hence, $d_\\mathrm{nuc}$ enhances comparing to the old estimation $d_\\mathrm{nuc} \\propto \\sigma^{3/2}/L_\\mathrm{h}$~\\cite{Fuller:1987ue} when $\\sigma$ becomes larger or $L_\\mathrm{h}$ becomes smaller. In a large acceptable parameter space of $\\sigma$ and $L_\\mathrm{h}$, $d_\\mathrm{nuc}$ is not as small as people used to think as about $\\simeq 2~\\mathrm{cm}$~\\cite{Christiansen:1995ic} for the homogeneous nucleation case. (ii) The high temperature phase should have an intrinsic maximum supercooling scale $\\Delta_<$ based on a Hawking-Page type phase transition. This is in contrast with the old belief that the range of supercooling is caused by impurities or perturbations. We discussed the possibility that this maximum supercooling scale is saturated in the cosmological QCD phase transition, which may happen when this phase transition is extremely weakly first order. If it happens, the nucleation distance $d_\\mathrm{nuc}$ can be increased tremendously. However, it is unlikely to be so; because to get an appropriate $d_\\mathrm{nuc}$ for our universe (that is, to help understand the surviving of the quark nuggets, or to get the appropriate initial conditions of the BBN), $L_\\mathrm{h}$ needs to be fine tuned.  Some related works are listed as below for comparison reasons. The nucleation rate and also some of its cosmological applications, base on the holographic RS I model, are discussed in~\\cite{Creminelli:2001th,Kaplan:2006yi}. In this model, a ``Planck brane'' and a ``TeV brane'' are added to the $AdS_5 \\times S^5$ spacetime with a dual CFT. The ``Planck brane'' makes a UV cutoff hence adds a $(3+1)$-dimensional gravity; the ``TeV brane'' makes an IR cutoff, and the standard model fields in it are understood as bound states out of the strong interacting CFT~\\cite{ArkaniHamed:2000ds}. When at finite temperature, to make a lower free energy, the low temperature phase is as in the RS I model, but the high temperature phase favors an AdS-Schwarzschild solution (duals to the free CFT gas); hence, our universe should suffer a phase transition at some $T_\\mathrm{c}$ lower than the Fermi scale. To ensure that the phase transition is completed thus for avoiding an empty universe, we need a strong upper bound for $N_\\mathrm{c}$ of the dual CFT field. This model has already been discussed in Sec.~\\ref{subsubsec:large-N_constraint}, where we pointed out that an upper limit of $N_\\mathrm{c}$ may be universal for some large-$N_\\mathrm{c}$ theories which suffer some phase transitions.  The phase transition of an AdS/CFT model, in which a $(2+1)$-dimensional field theory is dual to some (confined) AdS soliton or some (deconfined) black 3-brane metric compactified in a brane dimension, is discussed in~\\cite{Horowitz:2007fe}. The supercooling and the rapid reheating (hadronization) after it, are considered. Notwithstanding, in the large-$N_\\mathrm{c}$ limit, the slowly hadronized phase at the temperature $T_\\mathrm{dec}$ do not happen in their model. To begin at some supercooling temperature $T_0 > 0$, the residual deconfined regions after the phase transition still hold the energy portion larger than $1/4$. In that model, the supercooling scale is given by hand, and a lower limit $T_0 = 0$ ($\\Delta = 1$) is considered. Comparing to that work, what we do in this paper is calculating $\\Delta$ explicitly within some physical environments (what we use is the cosmological QCD phase transition). We use some AdS/CFT models more pertinent to the $(3+1)$-dimensional QCD than theirs.  In addition, an interesting relation between the KSS bound and strange quark stars, is shown in~\\cite{Bagchi:2007ir}. The authors argued that, the surface of quark stars at the temperature $T \\sim 80~\\mathrm{MeV}$, has already saturated the KSS bound.  The question which parallels to the topic we discussed in this paper, is how the RHIC results of strong interacting QGP and the AdS/CFT correspondence can affect the research of neutron stars and quark stars. The difference is that the deconfined QGP in quark stars is mainly caused by its high chemical potential, rather than caused by their high temperature in RHIC or the early universe. A lot of AdS/CFT models for finite chemical potential have already been constructed; although just as in the finite temperature case, they are mainly studied in the large-$N_\\mathrm{c}$ limit. We will leave these issues to the follow-up studies."
    },
    "0812/0812.5021_arXiv.txt": {
        "abstract": "{We discuss bulk viscosity due to non-leptonic processes involving hyperons and Bose-Einstein condensate of negatively charged kaons in neutron stars. It is noted that the hyperon bulk viscosity coefficient is a few order of magnitude larger than that of the case with the condensate. Further it is found that the hyperon bulk viscosity is suppressed in a superconducting phase. The hyperon bulk viscosity efficiently damps the r-mode instability in neutron stars irrespective of whether a superconducting phase is present or not in neutron star interior.}  \\FullConference{10th Symposium on Nuclei in the Cosmos\\\\ July 27 - August 1, 2008\\\\ Mackinac Island, Michigan, USA}  \\begin{document} ",
        "introduction": "Neutron stars pulsate in various modes due to its fluid perturbation. Those oscillatory modes are classified according to different restoring forces. The study of Coriolis restored r-mode in neutron stars has generated tremendous interest in recent times \\cite{Kok}. The mode becomes unstable due to gravitational radiation. R-modes of rotating neutron stars are important sources of detectable gravitational waves. The $l=m=2$ r-mode frequency ($\\omega_r$) is related to the angular frequency of the compact star as $\\omega_r = \\frac{2m}{l(l+1)} \\Omega$ \\cite{Kok}. The calculation of $\\Omega$ depends on the equation of state (EoS) as it will be shown in Figure 3. Therefore, r-modes, if detected, would shed light on the EoS and composition of matter in neutron star interior.  Depending on the nature of strong interaction, different novel phases with large strangeness fraction such as, hyperon matter, Bose-Einstein condensate of negatively charged kaons and quark phase may appear in neutron star interior. The r-mode instability could be effectively suppressed by bulk viscosity due to non-leptonic processes including hyperons \\cite{Jon1,Jon2} in neutron star interior. Here we discuss bulk viscosity coefficients and the corresponding damping time scales due to non-leptonic weak processes involving hyperons and negatively charged kaons. Next we demonstrate the effect of bulk viscosity on the damping of the gravitational radiation driven r-mode instability in neutron stars. Further we discuss the influence of antikaon condensate on the hyperon bulk viscosity. ",
        "conclusions": "We have investigated bulk viscosity due to non-leptonic processes involving $\\Lambda$ hyperons and $K^-$ condensate in neutron stars. It is noted that the hyperon bulk viscosity coefficient is a few orders of magnitude higher than the antikaon bulk viscosity coefficient. Hyperon bulk viscosity efficiently damps the r-mode instability.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=6cm,height=6cm]{Fig4.eps} \\caption{Hyperon bulk viscosity coefficient in $K^-$ condensed matter is shown as a function of normalised baryon density.} \\end{center} \\end{figure}"
    },
    "0812/0812.1755.txt": {
        "abstract": "In this work, considering the impact of a supernova remnant (SNR) with a neutral magnetized cloud we derived analytically a set of conditions that are favorable for driving gravitational instability in the cloud and thus star formation. Using these conditions, we have built diagrams of the SNR radius, $R_{SNR}$, versus the initial cloud density, $n_c$, that constrain a domain in the parameter space where star formation is allowed.  This work is an extension to previous study performed without considering magnetic fields (Melioli et al. 2006). The diagrams are also tested with fully 3-D MHD radiative cooling simulations involving a SNR and a self-gravitating cloud and we find that the numerical analysis is consistent with the results predicted by the  diagrams.  While the inclusion of a homogeneous magnetic field approximately perpendicular to the impact velocity of the SNR with an intensity $\\sim 1 \\; \\mu$G within the cloud results only a small shrinking of the star formation zone in the diagram relative to that without magnetic field, a larger magnetic field ($\\sim 10\\;\\mu$G) causes a significant shrinking, as expected. Though derived from simple analytical considerations these diagrams provide a useful tool for identifying sites where star formation could be triggered by the impact of a SN blast wave. Applications of them to a few regions of our own galaxy (e.g., the large CO shell in the direction of Cassiopeia, and the Edge Cloud 2 in the direction of the Scorpious constellation) have revealed that  star formation in those sites could have been triggered by shock waves from  SNRs  for specific values of the initial neutral cloud density and the SNR radius. Finally, we have evaluated the effective star formation efficiency for this sort of interaction and found that it is generally smaller than the observed values in our own Galaxy (sfe $\\sim$ 0.01$-$0.3). This result is consistent with previous  work in the literature and also suggests that the mechanism presently investigated, though very powerful to drive structure formation, supersonic turbulence and eventually, local star formation, does not seem to be sufficient to drive $global$ star formation in normal star forming galaxies, not even when the magnetic field in the neutral clouds is neglected. ",
        "introduction": "Essentially all present day star formation takes place in molecular clouds (MCs; e.g. Blitz 1993; Williams, Blitz \\& McKee 2000). It is likely that the MCs are relatively transient, dynamically evolving structures produced by compressive motions in the diffuse HI medium of either gravitational or turbulent origin, or some combination of both (e.g., Hartmann et al. 2001; Ballestero-Paredes et al. 2007). In fact, observations of supersonic line-widths in the MCs support the presence of supersonic turbulence in these clouds with a wealth of structures on all length-scales (Larson 1981; Blitz \\& Williams 1999; Elmegreen \\& Scalo 2004; Lazarian \\& Esquivel 2003). Recent numerical simulations in periodic boxes have shed some light on the role of the turbulence in the evolution of the MCs and the star formation within them. They suggest that the $continuous$ injection of supersonic motions, maintained by internal or external driving mechanisms (see below), can support a cloud  $globally$ against gravitational collapse so that the net effect of turbulence seems to inhibit collapse and this would explain the observed low overall star formation efficiencies in the Galaxy (Klessen et al. 2000; Mac Low \\& Klessen 2004; Vazquez-Semandeni et al. 2005). On the other hand, the supersonic turbulence is also able to produce density enhancements in the gas  that may allow $local$ collapse into stars in both nonmagnetized (Klessen et al. 2000;  Elmegreen \\& Scalo 2004) and magnetized media (Heitsch et al. 2001; Nakamura \\& Li 2005).  What are the possible sources of this turbulence in the MCs? Suggested candidates for an internal driving mechanism of turbulence include feedback from both low-mass and massive stars. These later, in particular, are major structuring agents in the ISM  in general (McCray \\& Snow 1979),  initially through  the production of powerful winds and  intense ionizing radiation and, at the end of their lives through the  explosions as supernovae (SNe). It is worth noting however, that  MCs with and without star formation have similar kinematic properties (Williams, Blitz \\& McKee 2000). External candidates include galactic spiral shocks (Roberts 1969; Bonnell et al. 2006) and again SNe shocks (Wada \\& Norman 2001; Elmegreen \\& Scalo 2004). These processes seem to have sufficient energy to explain the kinematics of the ISM and can generate the observed velocity dispersion-sizescale relation (Kornreich \\& Scalo 2000). Other mechanisms, such as magnetorotational instabilities, and even the expansion of HII regions and fluctuations in the ultraviolet (UV) field apparently inject energy into the ambient medium at a rate which is about an order of magnitude lower than the energy that is required to explain the random motions of the ISM at several scales, nonetheless the relative importance of all these injection mechanisms upon star formation is still a matter of debate (see, e.g., Joung \\& Mac Low 2006; Ballesteros-Paredes et al. 2007; Mac Low 2008 for reviews).  In this work, we focus on one of these driving mechanisms $-$ supernova explosions produce large blast waves that strike the interstellar clouds, compressing and sometimes destroying them, but the compression by the shock may also trigger local star formation (Elmegreen \\& Lada 1977; see also Nakamura \\& Li 2005 and references therein).   There has been extensive analytical and numerical  work exploring the role  of the SNe  in generating supersonic turbulence and  multi-phase structures in the ISM (e.g.,  Cox \\& Smith 1974; McKee \\& Cowie 1975; McKee \\& Ostriker 1977; Kornreich \\& Scalo 2000; Mac Low \\& Klessen 2004; Vazquez-Semadeni et al. 2000; de Avillez \\& Breitschwerdt 2005) and even at the galactic scale,  in the generation of large-scale outflows like galactic fountains and winds (e.g., de Gouveia Dal Pino \\& Medina-Tanco 1999; de Avillez 2000; Heckman et al. 2001; de Avillez \\& Berry 2001; Melioli et al. 2008a, 2008b). The collective effect of the SNe is likely to be the dominant contributor to the observed supersonic turbulence in the Galaxy (Norman \\& Ferrara 1996; Mac Low \\& Klessen 2004), however, it is possible that it inhibits $global$ star formation rather than triggering it (Joung \\& Mac Low 2006). Here, instead of examining the global effects of multiple SNe explosions upon the evolution of the ISM and the MCs, we will explore  the $local$ effects of these interactions. To this aim, we will consider a supernova remnant (SNR) either in its adiabatic or in its radiative phase impacting with an initially homogeneous diffuse neutral cloud and show that it is possible to derive analytically a set of conditions that can constrain a domain in the relevant parameter space where these interactions may lead to the formation of gravitationally unstable, collapsing structures. In Melioli et al. (2006; hereafter Paper I), we gave a first step into this direction by examining interactions involving a SNR and a non-magnetized cloud and then applying the results to a SF region in the local ISM where  the young stellar association of $\\beta$ Pictoris was born. Presently, besides including the effects of the magnetic fields, we will extend this analysis to few other SF regions with some indication of recent-past interactions with SN shock fronts (e.g., the Edge Cloud 2,  Yasui et al. 2006, Ruffle et al. 2007;  and the \"Great CO Shell\", Reynoso \\& Mangum 2001). Finally, we will also test our analytically derived $SF$ $domain$ with 3D MHD simulations of SNR interactions with self-gravitating neutral clouds.  In Section \\ref{SNR-Cloud interactions}, we consider the equations that describe the interaction between an expanding SNR and a cloud. In Section \\ref{Conditions to collapse}, we obtain analytically the set of constraints from these interactions with both non-magnetized and magnetized clouds that may lead to star formation (SF) and then build diagrams where these constraints delineate a domain in the parameter space which is appropriate for SF. In Section \\ref{MHD simulations}, we describe radiative cooling 3D MHD simulations of these interactions which  confirm the results obtained from the analytical study, and \\ref{Applications} we apply these analytical results to few examples of  regions of the ISM that present some indication of SF due to a recent SNR-cloud interaction.  In Section \\ref{SFE}, we briefly discuss the role  of these interactions to the global efficiency of star formation in the Galaxy and in Section \\ref{Conclusions}, we draw our conclusions. ",
        "conclusions": "\\label{Conclusions}  We have presented here a study of isolated interactions between SNRs and diffuse neutral clouds focusing on the determination of the conditions that these interactions must satisfy in order to lead to gravitational collapse of the shocked cloud material and to star formation, rather than to cloud destruction. A preliminary study of these interactions neglecting the effects of the magnetic field in the cloud had been  performed previously by Melioli et al. (2006). We have presently incorporated these effects and derived a new set of conditions for the interactions. A first condition determines the Jeans mass limit for the compressed cloud material in the impact. A second one establishes the penetration extent of the SNR shock front inside the cloud before being stalled due to radiative losses. It must have energy enough  to compress as much cloud material as possible before fainting.  A third constraint establishes the condition upon the same shock front under which it will $not$ become too strong to destroy the cloud completely. We have then built diagrams of the radius of the SNR as a function of the cloud density where this set of constraints delineate a domain within which star formation may result from these SNR-cloud interactions (Section 3). As expected, we find that an embedded magnetic field in the cloud normal to the shock front with an intensity of 1 $\\mu$G inhibits slightly the domain of SF in the diagram when compared to the non-magnetized case. The magnetic field plays a dominant role over the Jeans constraint causing a drift of the allowed  SF zone to higher cloud densities in the diagram. When larger intensities of magnetic fields are considered (5-10 $\\mu$G), the shrinking of the allowed SF zone in the diagrams is much more significant. We must emphasize however that, though observations indicate typical values of $B_c \\simeq 5-10 \\mu$G for these neutral clouds, the fact that we have assumed uniform, normal fields in the interactions have maximized their effects against gravitational collapse. We should thus consider as  more realistic the result obtained when an effective $B_c \\simeq 1 \\mu$G was employed. These diagrams derived from simple analytical considerations provide a useful tool for identifying sites where star formation could be triggered by the impact of a SN blast wave.   {\\bf We have also performed fully 3D MHD numerical simulations of the impact between a SNR and a self-gravitating cloud for different initial conditions (Section 4) tracking the evolution of these interactions and identifying the conditions that have led either to cloud collapse and star formation or to complete cloud destruction and mixing with the ambient medium. We have found the numerical results to be consistent with those established by the SNR-cloud density diagrams, in spite of the fact that the later have been derived from simplified analytic theory. We remark that the radiative cooling in the MHD numerical simulations has been considered through the adoption of an approximate polytropic pressure equation with $\\gamma_{eff} \\sim 1.2$ (see Spaans \\& Silk 2000). We are presently implementing a more realistic cooling function in our  Godunov-MHD code in order to test, e.g.,  the re-expanding cloud case and its validity under more realistic situations.}   We have applied the results above to a few examples of regions in the ISM with some evidence of  interactions of the sort examined in this work. In paper I, the application of the results of the $SF$ diagram for a non-magnetized cloud (as in Figure \\ref{fig:diag1muG}) to the conditions around the young stellar association of $\\beta-$Pictoris in our local ISM had led us to conclude that this stellar association could have originated from past cloud-SNR interaction only under very restrict conditions, i.e., with  a cloud with radius $\\sim$ 10 pc and density $\\sim$ 20 cm$^{-3}$ and a SNR with a radius $\\sim$ 42 pc. However, in the present work we find that  with the inclusion of an effective magnetic field in the cloud with an intensity of only 1 $\\mu$G this interaction is unlikely to produce that stellar association (Figure 2, cross in the third panel from top to bottom), at least not for the set of initial conditions proposed in the literature for that system (see also Melioli et al. 2006). In the case of the expanding Great CO Shell$-$O9.5 star system, we find that local star formation could have been induced in this region if, at the time of the interaction, the SNR that probably originated this expanding shell was still in the adiabatic phase and a radius between $\\sim$ 8 pc $-$ 29  pc impinged a magnetized cloud with  density around 30 $ \\rm cm^{-3}$ (Figure \\ref{fig:compdiag50pcpts}). Another example is the SF region near the Edge Cloud 2. This is one of the most distant cloud complexes from the galactic center where external perturbations should thus be rare. But the recent detection of two young associations of T-Tauri stars in this region could have been formed from the interaction of a SNR in the radiative phase with a cloud,  if the interaction started $\\lesssim 10^6$ yr,  the SNR had a radius $R_{SNR} \\simeq$ $46$ pc $-$ $84$ pc and the magnetized cloud a density around $n_c\\sim14\\; \\rm cm^{-3}$ (Figure \\ref{fig:compdiag20pcRadN2c}).  Finally, though in this study we have focused on isolated interactions involving SNRs and clouds, we  used the results of the diagrams to  estimate the contribution of these interactions to global star formation. Our evaluated effective star formation efficiency for this sort of interaction is generally smaller than the observed values in our own Galaxy (sfe $\\sim$ 0.01$-$0.3) (Figures \\ref{fig:sfe10pc} and \\ref{fig:sfe10pcrad}). This result seems to be consistent  with previous analysis (e.g., Joung \\& Mac Low 2006) and suggests that these interactions are powerful enough to drive structure formation, supersonic turbulence (see, e.g., simulation of Figure \\ref{fig:dmag10}) and eventually \"local\" star formation, but they do not seem to be sufficient to drive $global$ star formation in our galaxy or in other normal star forming galaxies, not even when the magnetic field in the neutral cloud is neglected. In conclusion, the small size of the allowed SF domain in the diagrams and the results for the estimated sfe indicate that these interactions must lead more frequently to the destruction of the clouds, rather than to their gravitational collapse."
    },
    "0812/0812.2938_arXiv.txt": {
        "abstract": " ",
        "introduction": "}  The study of almost all astronomical objects relies on an understanding of their hydrodynamics.  Indeed, for many such objects the involved hydrodynamics are complex enough to require numerical modeling.  The process of numerical modeling usually proceeds by writing partial differential equations describing the behavior of a continuous medium as an equivalent set of algebraic equations for a finite set of discretized elements.  This discretization can generally be performed in two different ways.  In an Eulerian approach, one discretizes the spatial domain into volumes termed grid cells. The fluid is considered to move through this fixed background grid.  By contrast, in a Lagrangian approach the fluid is discretized into fluid elements (or `particles') which can then move freely according to their initial velocities, and only their interactions need to be modeled.  Lagrangian methods work well in situations with large background flows where Eulerian methods would spend the bulk of their time advecting the (uninteresting) balanced flow, accumulating numerical errors with the numerous iterations required. Lagrangian methods have a large dynamic range in length but not in mass, achieving good spatial resolution in high-density regions but performing poorly in low-density regions.  In addition, the usual implementations of Lagrangian methods, based on Smoothed Particle Hydrodynamics (SPH), do not easily allow the higher spatial accuracy that grid methods can employ nor do they capture shocks as accurately as grid methods. By contrast, Eulerian methods provide a large dynamic range in mass but not in length. In general they are also computationally faster by several orders of magnitude, easier to implement, and easier to parallelize.  The RAPID code (Rapid Algorithm for Planets In Disks), which we present here, uses an Eulerian approach, adapted for a cylindrical grid.  While we focus on planet-disk interactions in this paper, the code is intended for the general study of accretion disks containing a dominant central mass.  In such systems the gas disk surrounding the central object is of a small enough mass that its self-gravity may be ignored.  Such disks will have a roughly Keplerian velocity profile resulting from the mass of the central object.  In order to obtain higher algorithm efficiency in the presence of this Keplerian flow, we make use of a FARGO-type algorithm \\citep{masset00}.  The algorithm's underlying strategy is to subtract off the bulk flow, which can be considered simply a translation of grid quantities, leaving the dynamically important residual velocity. RAPID is second-order accurate in space and time. Advection is accomplished through a nonlinear Total Variation Diminishing (TVD) scheme, which helps to control spurious oscillations.  Time-stepping is accomplished through a standard Runge-Kutta scheme.  Operator splitting is used to account for multiple dimensions and source terms such as those due to gravitational potentials and viscosity.  In Section \\ref{sec:eulerian} we outline the fluid equations to be solved and discuss considerations of angular momentum important for accuracy on cylindrical grids.  In Section \\ref{sec:numericalmethod} we discuss the details of the RAPID algorithm.  We provide results of basic hydrodynamical tests in Section \\ref{sec:basictests} and demonstrate the code's performance on typical planetary disk setups in Section \\ref{sec:planettests}.  Conclusions are presented in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have developed a new, efficient, parallelized hydrodynamic code for studying accretion disk processes. The current incarnation is optimized to study planetary disk-planet interactions.  A FARGO-type algorithm is implemented to help alleviate CFL time step restrictions imposed by the rapidly rotating inner disk region.  Open-MP directives are also implemented to obtain faster computations on shared memory machines. Parallelization on distributed memory machines (such as with Message Passing Interface (MPI) protocols) requires further development.  We have shown that the RAPID code performs comparably to the well-established piece-wise para\\-bolic method (PPM) on standard hydrodynamic tests.  The largest difference observed between the two algorithms is the level of pre- and postshock oscillations that are allowed.  PPM codes flatten such oscillations quite stringently.  We note that this procedure is not necessarily physically motivated and may lead to a decrease in the amount of structure present in some simulations.  We have also compared how results from our code differ depending on the choice of flux limiter. The amounts of diffusion and dispersion vary quite substantially, but the relative amounts amongst limiters are  observed to be qualitatively consistent on all tests.  In addition, we presented a large series of comparisons on the standard protoplanet problem, showing results that are consistent with ensemble results from a wide variety of other codes documented in the protoplanet comparison project \\citep{valborro06}.  In particular we confirmed the existence of libration islands at the $L_4$ and $L_5$ points, an asymmetry in the density of those islands, the sign and magnitude of the torque exerted on the planet by the disk and the growth and merging of vortices outside the planet's orbit at low viscosity.  We found that the large vortex which forms outside the plan\\-et's orbit causes substantial torque oscillations on the planet.  These oscillations correspond to repeated passes of the vortex by the planet.  Increasing the viscosity beyond $\\nu \\ge 10^{-5}$ damps the formation of the vortex thereby removing the oscillatory signature from the torque.  The existence of additional \\textit{vortex lines} to the inside of the planet's orbit are demonstrated.  We have shown that the \\textit{FARGO}-like fast-advection algorithm reduces the required simulation time by a factor of $6.5$ in standard planet-disk setups.  We also illustrated that using the inertial angular momentum rather than angular velocity as a solution variable decreases the numerical viscosity present in the simulations by an order or magnitude or more.  This finding supplements previous work by \\citet{kley98}.  In addition, the choice of inertial angular momentum as a solution variable conserves the total angular momentum on the grid to higher precision.  We determined the level of numerical viscosity present with\\-in the code to be $\\nu < 10^{-6}$ ($\\alpha < 10^{-3.5}$), enabling the simulation of scenarios with Reynolds numbers on the order of $Re = UL \\times 10^6 $."
    },
    "0812/0812.3339_arXiv.txt": {
        "abstract": "We present a set of numerical simulations addressing the effects of magnetic field strength and orientation on the flow-driven formation of molecular clouds. Fields perpendicular to the flows sweeping up the cloud can efficiently prevent the formation of massive clouds but permit the build-up of cold, diffuse filaments. Fields aligned with the flows lead to substantial clouds, whose degree of fragmentation and turbulence strongly depends on the background field strength. Adding a random field component leads to a ``selection effect'' for molecular cloud formation: high column densities are only reached at locations where the field component perpendicular to the flows is vanishing. Searching for signatures of colliding flows should focus on the diffuse, warm gas, since the cold gas phase making up the cloud will have lost the information about the original flow direction because the magnetic fields redistribute the kinetic energy of the inflows. ",
        "introduction": "\\label{s:motivation}  Evidence is accumulating that star formation follows rapidly upon molecular cloud formation (e.g. \\citealp{2001ApJ...562..852H} and \\citealp{2007RMxAA..43..123B} for the solar neighborhood; \\citealp{2003ApJS..149..343E} for M33; \\citealp{2007ApJ...668.1064E} in the context of M51). This rapid onset suggests that the clouds need to acquire high, non-linear density enhancements during their formation, since massive, finite clouds are highly susceptible to global gravitational collapse which could overwhelm small-scale fragmentation necessary for (local) star formation \\citep{2004ApJ...616..288B}. Thus to understand the initial conditions for star formation, we need to understand the formation of the parental clouds.  \\citet{1999ApJ...527..285B} and \\citet{2001ApJ...562..852H} proposed that the build-up of clouds in large-scale, converging flows of diffuse atomic gas could explain the crossing time problem, i.e. the observation that the stellar age spreads in a large number of local star forming regions are substantially smaller than the lateral crossing time \\citep{2001ApJ...562..852H,2007RMxAA..43..123B}. In this picture, there need not be a causal connection between star formation events in the plane perpendicular to the large-scale flows (see also \\citealp{2000ApJ...530..277E}). Rapid star formation is a necessary requirement for this scenario to work.  Numerical models of flow-driven cloud formation (we give only an early and the most recent numerical work of each group, namely \\citealp{1999A&A...351..309H} and \\citealp{2008A&A...486L..43H}; \\citealp{2000ApJ...532..980K} and \\citealp{2008ApJ...687..303I}; \\citealp{2005ApJ...633L.113H} and \\citealp{2008ApJ...674..316H}; \\citealp{2006ApJ...643..245V} and \\citealp{2007ApJ...657..870V}) have identified the thermal and dynamical instabilities that are responsible for the rapid fragmentation of the  nascent cloud (see \\citealp{2008ApJ...683..786H} for an assessment of the roles of the physical processes). Despite these promising successes, many questions about the physics at play remain unanswered, among one of the most pressing is the role of magnetic fields during the cloud formation process.  The role of magnetic fields in the flow-driven cloud formation scenario has been largely envisaged as one of ``guiding the flows'' to assemble the clouds, whether in form of the Parker instability along galactic spiral arms \\citep{1966ApJ...145..811P,1967ApJ...149..517P,1974A&A....33...73M}, in a generally turbulent interstellar medium (ISM) \\citep{1995ApJ...455..536P,2001ApJ...562..852H}, or during the sweep-up of gas in spiral shocks \\citep{2006ApJ...646..213K,2008MNRAS.383..497D}. Based on the models of Passot et al., Hartmann et al. suggested that the field orientation with respect to the flows selects the locations of cloud formation, namely that clouds will only form if the fields are aligned with the flows. A perpendicular field will reduce the compression of the post-shock gas and thus will limit the strong cooling and the thermal instability (TI, \\citealp{1965ApJ...142..531F}) necessary for the rapid flow fragmentation and the build-up of high-density contrasts \\citep{2004ApJ...616..288B,2008ApJ...674..316H,2008ApJ...683..786H}.  Given sufficiently high strengths, fields aligned with the inflows can suppress the dynamical instabilities responsible for the generation of turbulence, namely the non-linear thin shell instability (NTSI, \\citealp{1994ApJ...428..186V}; for a magnetic version see \\citealp{2007ApJ...665..445H}) and the Kelvin-Helmholtz instability (KHI, e.g. \\citealp{1961hhs..book.....C}, more recently \\citealp{2008MNRAS.385.1494K}, and for numerical studies \\citealp{2008ApJ...678..234P}). Yet magnetic fields are intrinsically three-dimensional, and already two-dimensional models by \\citet{2008ApJ...687..303I} show that even for fields perpendicular to the inflow, cold (albeit diffuse) clouds can form.  Thus, three-dimensional models of flow-driven cloud formation including magnetic fields are needed. \\citet{2008A&A...486L..43H} present a first approach to the problem, modeling the formation of a cloud in converging, perturbed flows, including fields and self-gravity. Here, we focus on the effects of magnetic field strength and orientation on the early stages of flow-driven cloud formation. We work in the ideal MHD limit (i.e. we do not explicitly consider ambipolar drift or resistivity), and we do not include gravity in the models.  All our models start out with field strengths below equipartition with the kinetic energy of the inflows. At a factor of $4.3$ below equipartition -- corresponding to an absolute field strength of $5\\mu$G at flow densities and velocities of $1$~cm$^{-3}$ and $16$~km~s$^{-1}$ -- , the fields already suppress the dynamical instabilities (and thus the generation of turbulence) leading to slab-like molecular clouds, while weaker fields -- at $2.5\\mu$G, corresponding to a factor of $17$ below equipartition -- lead to clouds closely resembling the hydrodynamical case, albeit with more coherent filaments. Fields at $0.5\\mu$G perpendicular to the inflows suppress the build-up of massive clouds in the collision plane, while they lead to the formation of diffuse, cold filaments perpendicular to the mean field, reminiscent of the cold HI clouds discussed by \\citet{2003ApJ...586.1067H}. A tangled field allows the assembly of substantial column densities in regions where the lateral field component is small or vanishing. Our results are consistent with the notion that magnetic fields select the location of cloud formation \\citep{2001ApJ...562..852H}. ",
        "conclusions": "\\subsection{The Role of  Magnetic Fields for Cloud Formation}  The field strength, and the orientation of the mean magnetic field with respect to the flows sweeping up the gas play a crucial role for the flow-driven formation of molecular clouds (Fig.~\\ref{f:polmap}). If the fields are dynamically dominant, the only chance to build up substantial clouds is by channeling the flows along the fields. This is the situation shown in model X50, and it is also borne out by simulations of molecular cloud formation in galactic spiral arms \\citep{2006ApJ...646..213K}, where the clouds tend to be oriented perpendicularly to the large-scale field (also possibly visible in the models by \\citealp{2008MNRAS.383..497D}), until sufficient material has been accumulated that they decouple dynamically from the large-scale field. Similarly, for the sweep-up of material by e.g. H{\\small{II}}-regions or supernova shells, one would expect the densest clouds to appear at the locations where the field is perpendicular to the shell (see Fig.~2 of \\citealp{2006ApJ...641..905H}, although the effect might be less clear in a highly turbulent environment, see Fig.~10 of \\citealp{2004ApJ...617..339B}).  \\citet{2001ApJ...562..852H} point out -- based on simulations by \\citet{1995ApJ...455..536P} -- that dynamically weak (but not necessarily ordered) fields would lead to a general selection effect for the formation of molecular clouds. Since $\\beta_{ram}>1$, the flows stretch out the fieldlines, leading to a natural alignment. In this picture, clouds form in the bends of large-scale fields (see Figs. 4 \\& 5 of \\citealp{2001ApJ...562..852H}). Such a scenario is to some extent addressed by model XR25, where  varying field orientations entail a local selection effect, picking out the formation sites of molecular clouds over e.g. a broad shock front. Note that while the field is dynamically weak ($\\beta_{ram}=8.5$, $\\beta_{th}=4.3$) in the initial conditions (and in the inflows) of model XR25, $\\beta_{th} < 1$ within the cloud (Fig.~\\ref{f:prssprof}).  This selection effect comes about because already a small oblique component can be amplified sufficiently to withstand the compression, preventing the high densities needed for cloud formation \\citep{2001ApJ...562..852H}. Such a situation is addressed in the extreme by model Y05. A field perpendicular to the sweeping-up flow can suppress the formation of massive clouds, although the three-dimensional situation is much less clear-cut than its one-dimensional counterpart (see e.g. \\citealp{1980ARA&A..18..219M}; \\citealp{2004ApJ...612..921B}). In one dimension, a density increase by a factor of $100$ from e.g. $n=1$~cm$^{-3}$ to $100$~cm$^{-3}$ would entail the same factor for the magnetic field strength since $B\\propto n$. Figure~\\ref{f:babsdens} shows this is not the case in three dimensions. The weak perpendicular field (model Y05) has been amplified by a peak factor of $\\approx 30$, while the density has increased by up to a factor of $300$. Generally, our models show a weak correlation of field strength with density over the whole thermal range, from the WNM to the CNM, consistent with observations of the field-density relation in the WNM and CNM (\\citealp{1986ApJ...301..339T}; HT05), and with numerical results (e.g. \\citealp{2005A&A...436..585D,2008A&A...486L..43H}). For models X50 and X25 a weak correlation between field and density is not overly surprising. For model Y05, the decorrelation\\footnote{The seemingly strongly correlated B(n) for $n<1$~cm$^{-3}$ in model Y05 does not affect the argument. These are a few regions (low mass fraction) at the edges of the expanding slab, subjected to numerical reconnection.} is a consequence of the fact that material is still free to move along the field lines perpendicular to the original inflow \\citep{2007ApJ...665..445H}, thus leading to the build-up of higher-density filaments perpendicular to the field (but aligned with the inflow), see Figure~\\ref{f:polmap}. Also, other effects, such as the acceleration of magnetic field transport by turbulence (\\citealp{2002ApJ...567..962Z,2002ApJ...570..210F,2004ApJ...603..165H} for ion-neutral drift, \\citealp{1999ApJ...517..700L} for reconnection), or a decorrelation due to MHD waves \\citep{2003A&A...398..845P} could explain the observed weak correlation.  Magnetic fields will rarely be completely uniform. Model XR25 tests the more general case of a uniform field at $2.5\\mu$G and a random component of equal size, consistent with (although slightly lower than) observational estimates for magnetic field strengths in the diffuse gas \\citep{1996ASPC...97..457H,2004Ap&SS.289..293B,2006ChJAS...6b.211H}. \\citet{2007ApJ...663L..41G} showed in a two-dimensional numerical experiment that pre-existing  perturbations in the inflows can lead to substantial magnetic field amplification due to a rippling of the shockfront and subsequent fieldline stretching. We observe a similar effect in model XR25, although our Mach numbers are substantially lower (their study addressed the propagation of a supernova shock front). \\citet{2008A&A...486L..43H} perturb the velocities of the inflows and find only a modest increase of the field strength. Clearly, the initially tangled field leads to rather different dynamics in the forming cloud (Figs~\\ref{f:polmap}, \\ref{f:prsstime}).  Models X25 and X50 demonstrate that not only the field orientation will play a role during cloud formation (see model Y05), but also the field strength, since all the instabilities involved have threshold limits for the field strength -- at least in two dimensions. It might well be that the stronger field in model X50 suppressing the formation of filaments could be offset by higher inflow speeds or substructures in the flows. This remains to be studied.  \\begin{figure*} \\includegraphics[width=\\textwidth]{f9.eps} \\caption{\\label{f:babsdens}Magnetic field strength against volume density for models as indicated in the plots. The colors denote temperatures, and the intensity the mass fraction. Generally, there is no clear correlation between field strength and density. The steeper of the two dashed lines denotes $B\\propto n$, the flatter one $B\\propto n^{1/2}$. Dotted lines denote the initial conditions.} \\end{figure*}  \\subsection{Turbulence and Thermal States}\\label{ss:turbtherm} Fields aligned with the inflows tend to suppress the NTSI, and thus lead to an approximate equipartition between the spatial components of the kinetic energy in the cold and in the thermally unstable gas (Fig.~\\ref{f:prssprof}, bottom couple of rows). For comparison, the hydrodynamical model H has the bulk of the kinetic energy in the (flow-aligned) $x$-component. Magnetic fields may play an important role to isotropize highly directional flows. Thus, searching for observational signatures of flow-driven cloud formation should focus on the warm, diffuse gas phase, since the inflow signature will be erased in the cold dense gas.  Fractions of thermally unstable gas (Fig.~\\ref{f:vrmsunm}) for flow-aligned fields (models X50, X25) are lower than values observed for diffuse CNM clouds \\citep{2003ApJ...586.1067H}. A lateral field component results in a thermally unstable gas fraction of $\\approx 50$\\%, consistent with observations. Based on these findings, one could feel tempted to extend the above argument about the selection effect introduced by magnetic fields: not only could magnetic fields control the locations of molecular cloud formation, but they also could lead to ``failed'' molecular clouds, i.e. diffuse atomic hydrogen clouds, if there is a non-negligible field component perpendicular to the sweeping-up flow (see also \\citet{2008ApJ...687..303I} for a similar argument based on two-dimensional simulations).  \\subsection{Ordered vs. Random Component}\\label{ss:components}  Another observational constraint is given by the ratio of the ordered over the unordered (or turbulent) field component. The observational evidence points to the components being of similar magnitude (e.g. \\citealp{1996ASPC...97..457H}; \\citealp{2004Ap&SS.289..293B}; \\citealp{2004ASSL..315..277B}; \\citealp{2006ChJAS...6b.211H}; see also discussion in HT05). A direct comparison to our models is hampered by the fact that in order to see the varying component, sufficiently large scales need to be addressed, which is why HT05 argue that their observed median field strength of $6\\mu$G should be identified with the {\\em total} magnetic field strength. Likewise, it is not obvious that the components should be of equal magnitude locally everywhere. Bearing this limitation in mind, it is clear from Figure~\\ref{f:fieldtime} that only for models X25 and Y05 the components are comparable.  \\subsection{Gravity}  We deliberately left out self-gravity in our simulations, in order to get a clearer view of the role of magnetic fields during the early cloud formation phase. Thus, our clouds are only confined by the ram pressure of the inflows, and at later stages, the dynamics of the clouds are probably underestimated since gravity as a source of turbulence is missing (e.g. \\citealp{2008MNRAS.385..181F}). As a result of the restricted physics, a comparison of our models with observations is only meaningful for models where gravity is not expected to play a role, i.e. for model Y05 addressing the formation of diffuse HI clouds. For all other models, we expect gravity to be relevant during the cloud formation process \\citep{2008ApJ...689..290H}."
    },
    "0812/0812.4519_arXiv.txt": {
        "abstract": "Our work focuses on a comprehensive orbital phase dependent spectroscopy of the four High Mass X-ray Binary Pulsars (HMXBPs) 4U~1538-52, GX~301-2, OAO~1657-415 \\& Vela~X-1. We hereby report the measurements of the variation of the absorption column density and iron-line flux along with other spectral parameters over the binary orbit for the above-mentioned HMXBPs in elliptical orbits, as observed with the Rossi X-ray Timing Explorer (RXTE) and the BeppoSAX satellites. A spherically symmetric wind profile was used as a model to compare the observed column density variations. Out of the four pulsars, only in 4U 1538-52, we find the model having a reasonable corroboration with the observations, whereas in the remaining three the stellar wind seems to be clumpy and a smooth symmetric stellar wind model appears to be quite inadequate in explaining the data. Moreover, in GX 301-2, neither the presence of a disk nor a gas stream from the companion was validated. Furthermore, the spectral results obtained in the case of OAO~1657-415 \\& Vela~X-1 were more or less similar to that of GX~301-2. ",
        "introduction": "{\\bf 4U~1538--52} was first detected with the UHURU satellite (Giacconi et al. 1974). The spin period and orbital period of the pulsar were first estimated to be 529 s and 3.73 days, respectively, from the observations with Ariel 5 and OSO-8 (Davision et al. 1977). Using data from the RXTE, the eccentricity of the binary orbit was calculated to be $\\sim$ 0.18 (Mukherjee et al. 2006). The X-ray spectrum has a prominent iron K-line (Makishima et al. 1987) and a pulse phase-dependent cyclotron resonance feature at 20 keV (Clark et al. 1990). It has a moderate X-ray luminosity of $\\sim$ 10$^{36}$ erg s$^{-1}$. The optical counterpart was found to be an early B type supergiant star (QV~Nor) with H$\\alpha$ emission lines (Parkes et al. 1978). The mass-loss rate and the terminal wind velocity were estimated as $\\sim$10$^{-6}$ M$_{\\odot}$ yr$^{-1}$ and $\\sim$1000 km s$^{-1}$ respectively.  {\\bf GX 301--2 (4U 1223-62)} was discovered by White et al. (1976). Using data from BATSE observations, the orbital period and eccentricity of the binary system were determined as $\\sim$41.5 days and $\\sim$0.46, respectively (Koh et al. 1997). The companion star Wray~977 has a B1 Ia+ spectral classification with a mass of 39 M$_{\\odot}$ (Kaper et al. 2006). The mass-loss rate and terminal velocity of the stellar wind are 10$^{-5}$ M$_{\\odot}$ yr$^{-1}$ and 305 km s$^{-1}$, respectively. GX~301-2 shows a variable X-ray luminosity in the range (2-400) $\\times$ $10^{35}$ erg s$^{-1}$, depending on the amount of the stellar wind captured, which in turn depends on the density and velocity of the wind.  {\\bf OAO~1657--415} was discovered by the Copernicus satellite (Polidan et al. 1978). White \\& Pravdo (1979) detected a $\\sim$38~s pulsation period for the neutron star. A steady spin-up time scape of 125 yr with short-term fluctuations are observed from the period measurements from RXTE and BATSE (Baykal 2000). BATSE observations also aided the discovery of X-ray eclipses by the stellar companion and the determination of a $\\sim$ 10.44 day orbital period (Chakrabarty et al. 1993). They in turn used the orbital parameters to infer that the companion is a supergiant of spectral class B0--B6. OAO~1657-415 is unique among the known HMXBs in that it appears to occupy a transition region between mass transfer via a stellar wind and Roche lobe overflow (Chakrabarty et al. 1993).  {\\bf Vela X-1} is also an eclipsing HMXBP with an orbital period of $\\sim$8.96 days (Barziv et al. 2001). It exhibits X-ray pulsations with a pulse period of 283 s (McClintock et al. 1976). The companion is an early-type primary star HD 77581, which is a massive B0.5Ib-\u2013type supergiant (Brucato \\& Kristian 1972). The inferred mass-loss rate of the stellar wind of the companion is of the order of 10$^{-7}$ M$_\\odot$ yr$^{-1}$ (Sako et al. 1999). Moreover, the terminal wind velocity has been determined to be 1700 km s$^{-1}$ (Dupree et al. 1980). The pulsar orbits about the center of mass of the system at a distance of only about 0.6 stellar radii from the surface of the supergiant, which has a radius of 53.4 R$_\\odot$. This implies that the neutron star is deeply embedded within the influence of the stellar wind. The typical X-ray luminosity of Vela X-1 is 4 $\\times$ 10$^{36}$ erg s$^{-1}$, but large flux variations have been observed (Kreykenbohm et al. 1999).  We observed 4U~1538-52 with RXTE from 2003-07-31 to 2003-08-07 covering out of eclipse phases for two binary orbits. We also used the archival data from BeppoSAX obtained between 1998-07-29 to 1998-08-01, covering one binary orbit. GX~301-2 was observed by RXTE first from 1996-05-10 to 1996-06-15 and secondly from 2000-10-12 to 2000-11-19. For OAO~1657-415, we analyzed the archival data as observed with RXTE from 1997-10-31 to 1997-11-11. In addition to these, we also analyzed an archival BeppoSAX Medium Energy Concentrator Spectrometer (MECS) observation taken on 2001-08-14 for $\\sim$104 ks exposure. Vela~X-1 was observed with the Proportional Counter Array (PCA) of RXTE from 2005-01-01 to 2005-01-09. For the RXTE data, we took Standard-2 data products of the PCA and extracted the source and background spectra using the tool saextrct v 4.2d. The BeppoSAX data products were extracted from the MECS and Low Energy Concentrator Spectrometers (LECS) using circular regions of radius 4$'$ and 8$'$ respectively. The background subtracted source spectra were analyzed with the spectral analysis package XSPEC v 11.2.0. Since PCU0 had lost the propane layer, we did not use the data from it in any further analysis for Vela~X-1. This was not a major issue for such a highly luminous pulsar since it did not affect the statistics.  For 4U~1538-52, the spectral model used was a simple power law along with a line-of-sight absorption, a high energy cut off and a Gaussian line with center energy $\\sim$6.4 keV. However in case of GX~301-2, the  Partial Covering Absorber Model (PCAM, Endo et al. 2002) with a high energy exponential cutoff and two Gaussian functions for iron K$_\\alpha$ and K$_\\beta$ lines was found to fit the RXTE data well compared to other models. The PCAM is described as two different power law components with the same photon index but different normalizations, being absorbed by different column densities (N$_{H1}$ \\& N$_{H2}$) respectively. The analytical form of the PCAM that we have used for spectral fitting is : $N(E) = {e^{-\\sigma(E)N_{\\mathrm H1}}}(S_{1}+S_{2}e^{-\\sigma(E)N_{\\mathrm H2}}) {E^{-\\Gamma}}$, where $N(E)$ is the intensity, $\\Gamma$ is the photon index, $N_{\\mathrm H1}$ and $N_{\\mathrm H2}$ are the two equivalent hydrogen column densities, $\\sigma$ is the photo-electric cross-section, $S_{1}$ and $S_{2}$ are the respective normalizations of the power law. For OAO~1657-415 \\& Vela~X-1 too, we observed that the PCAM was providing the best fit for both MECS and the RXTE spectra. Moreover, since RXTE is a non--imaging instrument and OAO~1657-415 lies under the influence of the Galactic ridge emission, we had to explicitly incorporate the ridge emission as a separate background spectral component. We put that in XSPEC as a sum of Raymond-Smith plasma and a power-law with appropriate normalizations (Valinia \\& Marshall 1998).  \\begin{figure} \\vskip 9.5 cm \\special{psfile=fig1.ps vscale=54 hscale=55 hoffset=-10 voffset=320 angle = -90.0} \\caption{Variation of Column Density versus Orbital Phase. The dashed line represents the model for inclination angle of 65$^\\circ$, the dashed-dotted line for 75$^\\circ$ and the dotted line for 85$^\\circ$. Diamonds, filled squares, and circles denote measurements from the observations with RXTE in 1997, BeppoSAX in 1998, and RXTE in 2003 respectively.} \\end{figure} ",
        "conclusions": "{\\bf 4U~1538-52} : The photon index, iron-line flux, cut-off energy and the e-folding energy measured with the pulse average spectrum taken over 2-3 ks did not show any substantial variation along the orbit which suggests that the continuum X-ray spectrum of the pulsar is hardly affected during its revolution. An important detection was a notably variation in N$_{H}$. It shows a smooth variation over orbital phase, increasing gradually by an order of magnitude as the pulsar approaches eclipse (Fig 1 : mid-eclipse is defined by phase zero). At orbital phases far from the eclipse, the column density has a value of $\\sim$ 1.5 $\\times$ 10$^{22}$ atoms cm$^{-2}$. Moreover, we compare the observed column density profile with a model estimated by assuming a spherically symmetric Castor, Abbott \\& Klein (CAK 1975) wind from the companion star. The velocity profile of the line-driven wind is: $v_{\\mathrm wind} =  {v_{\\infty}}\\sqrt{1-\\frac {{R_\\mathrm c}}{r}}$, where $v_{\\infty}$ is the terminal velocity for the stellar wind, R$_{\\mathrm c}$ is the radius of the companion and $r$ is the radial distance from center of the companion star. The column density profiles were derived using a numerical integration along the line of sight from the pulsar to the observer. With a mass-loss rate of $\\sim$10$^{-6}$ M$_\\odot$ yr$^{-1}$ and $v_\\infty$ $\\sim$1000 km s$^{-1}$, the model calculations of N$_H$ for different inclination angles when superposed on the observed values, it shows fairly reasonable agreement for three different inclination angles 65$^\\circ$, 75$^\\circ$ and 85$^\\circ$ respectively (Fig. 1), indicating that a CAK wind from the companion star may produce the observed orbital dependence of the column density for a certain range of the orbital inclination ($>$65$^\\circ$). We note here that, in Fig. 1, we have not done any fitting of the measured column densities and the model calculated column densities at different orbital phases are rather plotted together. The phase resolved column density values for the wind density model used here depend on the rate of mass loss, the terminal velocity and the inclination angle of the binary orbit. The data presented here, however, is not suitable for such a detailed analysis but only shows the consistency of the model with the observations.  \\begin{figure} \\vskip 10.9 cm \\special{psfile=fig2.ps vscale=63 hscale=99 hoffset=-50 voffset=-60 angle = 0.0} \\caption{The figure shows the N$_H$ variation for GX~301-2 with the uppermost panel depicts the model variation, the middle panel showing the observed variation of N$_{\\mathrm H1}$ and the lower panel shows the observed variation of N$_{\\mathrm H2}$. The error-bars shown correspond to 90$\\%$ confidence interval.} \\end{figure}  {\\bf GX~301-2} : As in 4U~1538-52, the continuum parameters do not show any orbital modulation in GX~301-2. In this case, the variation of N$_{H1}$ \\& N$_{H2}$ with orbital phase was not smooth. The values were very high with a large variation throughout the binary orbit (from 10$^{22}$ to 10$^{24}$ atoms cm$^{-2}$), indicating a clumpy nature of the stellar wind (Fig. 2). It is also seen that the covering fraction (defined as  Norm2/[Norm1+Norm2]) remains substantially high almost throughout the orbit which means that there is dense and clumpy material present throughout. This is also supported by the detection of a strong Compton recoil component detected with the Chandra grating spectrum and its successful reproduction by Monte Carlo simulations (Watanabe et al. 2003). Thus it is clear that the observed variation in column density cannot be explained by a spherically symmetric CAK wind only, indicating the presence of strong inhomogeneities in the wind. There are two proposed models in this regard : a gas-stream model by Leahy (1991) and the equatorial disk model by Pravdo \\& Ghosh (2001). However, the orbital dependence of the absorption column densities measured by us is very different from their predictions. Now, for the the iron K-line fluxes, we obtained peaks near periastron (Fig. 3). The line equivalent width had a correlation with the column density (N$_{H2}$), suggesting that most of the iron line is produced by the local clumpy matter surrounding the neutron star.  \\begin{figure} \\begin{center} \\includegraphics[angle=-90,width=12 cm]{fig3.ps} \\caption{The figure shows the variation of the iron-line flux with the orbital phase in GX~301-2. The upper and lower panels depict the flux corresponding to the 6.4 ${\\mathrm keV}$ and 7.0 ${\\mathrm keV}$ iron-lines respectively.} \\end{center} \\end{figure}  {\\bf OAO~1657-415 \\& Vela~X-1} : The average values of the free spectral parameters measured here over the out-of-eclipse phases of the binary orbit; viz. photon index, e-folding energy and cut-off energy, do not exhibit any modulation due to the revolution of both the pulsars. These results seem to be in conformity with that obtained for GX~301-2 \\& 4U~1538-52. The column density (N$_{H2}$) of the material that absorbs the primary X-ray emission is found to be high for both the pulsars (Fig 4). A large variation of the column densities throughout the out-of-eclipse phases (from 10$^{22}$ to 10$^{24}$ atoms cm$^{-2}$) characterizes their X-ray spectra. There were instances of moderate to high values of covering fraction. Here, it should be mentioned that the number of out-of-eclipse measurements for OAO~1657-415 are probably not sufficient to arrive at a strong conclusion.  We do not attempt to compare the variations of the observed column densities with the CAK model variation as we had done in case of 4U~1538-52 and GX~301-2 since it is clear from the measurements that such a comparison would definitely rule out the said model. In this context, we may point out that the spectral measurements for these two pulsars are similar in case of GX~301-2 (Mukherjee \\& Paul 2004). For GX~301-2 though, formation of clumped blobs of matter could be accounted for the unusually low wind velocity and high mass loss rate, but for the present case, such an inference is probably not possible. Moreover, unlike GX~301-2, we do not observe any systematic orbital modulation for the Iron-line fluxes. Furthermore, the iron-line equivalent width also does not show any definite correlation with N$_{H2}$.  \\begin{figure} \\centering \\includegraphics[width=3in,height=4in,clip,angle=-90.0]{fig4a.ps} \\includegraphics[width=3in,height=4in,clip,angle=-90.0]{fig4b.ps} \\caption{The figure shows the variation of N$_{H2}$ with orbital orbital phase for OAO~1657-415 (top) and Vela~X-1 (bottom). It may be observed that sufficient number of observations are lacking for OAO~1657-415. The error-bars shown correspond to 90$\\%$ confidence interval.} \\end{figure}   In view of the above exposition, we may conclude as follows :  \\begin{enumerate}  \\item The PCAM appears to be somewhat generic for pulsars which have variable column density, at least when observed with RXTE or BeppoSAX.  \\item For highly luminous supergiant binaries, the CAK Model of stellar wind does not suffice to describe the column density variations. This is corroborated by the results of GX~301-2, OAO~1657-415 \\& Vela~X-1 in this work and of 4U~1700--37 (Haberl et al. 1989) and 4U~1907+09 (Roberts et al. 2001).  \\item Moderately luminous pulsars probably validate the spherical wind model, as in the case of 4U~1538-52 (this work) and X~1908+075 (Levine et al. 2004). \\end{enumerate}"
    },
    "0812/0812.1750_arXiv.txt": {
        "abstract": "The diffuse plasma that fills galaxy groups and clusters (the intracluster medium, hereafter ICM) is a by-product of galaxy formation. The present thermal state of this gas results from a competition between gas cooling and heating. The heating comes from two distinct sources: gravitational heating associated with the collapse of the dark matter halo and additional thermal input from the formation of galaxies and their black holes.  A long term goal of this research is to decode the observed temperature, density and entropy profiles of clusters and to understand the relative roles of these processes. However, a long standing problem has been that cosmological simulations based on smoothed particle hydrodynamics (SPH) and Eulerian mesh-based codes predict different results even when cooling and galaxy/black hole heating are switched off.  Clusters formed in SPH simulations show near powerlaw entropy profiles, while those formed in Eulerian simulations develop a core and do not allow gas to reach such low entropies. Since the cooling rate is closely connected to the minimum entropy of the gas distribution, the differences are of potentially key importance.  In this paper, we investigate the origin of this discrepancy. By comparing simulations run using the \\gadget SPH code and the \\flash adaptive Eulerian mesh code, we show that the discrepancy arises during the idealised merger of two clusters, and that the differences are not the result of the lower effective resolution of Eulerian cosmological simulations.  The difference is not sensitive to the minimum mesh size (in Eulerian codes) or the number of particles used (in SPH codes).  We investigate whether the difference is the result of the different gravity solvers, the Galilean non-invariance of the mesh code or an effect of unsuitable artificial viscosity in the SPH code.  Instead, we find that the difference is inherent to the treatment of vortices in the two codes.  Particles in the SPH simulations retain a close connection to their initial entropy, while this connection is much weaker in the mesh simulations.  The origin of this difference lies in the treatment of eddies and fluid instabilities. These are suppressed in the SPH simulations, while the cluster mergers generate strong vortices in the Eulerian simulations that very efficiently mix the fluid and erase the low entropy gas. We discuss the potentially profound implications of these results. ",
        "introduction": "There has been a great deal of attention devoted in recent years to the properties of the hot X-ray-emitting plasma (the intracluster medium, hereafter ICM) in the central regions of massive galaxy groups and clusters.  To a large extent, the focus has been on the competition between radiative cooling losses and various mechanisms that could be heating the gas and therefore (at least partially) offsetting the effects of cooling.  Prior to the launch of the {\\it Chandra} and {\\it XMM-Newton} X-ray observatories, it was generally accepted that large quantities of the ICM should be cooling down to very low energies, where it would cease to emit X-rays and eventually condense out into molecular clouds and form stars (Fabian 1994). However, the amount of cold gas and star formation actually observed in most systems is well below what was expected based on analyses of {\\it ROSAT}, {\\it ASCA}, and {\\it Einstein} X-ray data (e.g., Voit \\& Donahue 1995; Edge 2001). New high spatial and spectral resolution data from {\\it Chandra} and {\\it XMM-Newton} has shown that, while cooling is clearly important in many groups and clusters (so-called `cool core' systems), relatively little gas is actually cooling out of the X-ray band to temperatures below $\\sim10^7$ K (Peterson et al.\\ 2003).  It is now widely believed that some energetic form of non-gravitational heating is compensating for the losses due to cooling. Indeed, it seems likely that such heating goes beyond a simple compensation for the radiated energy. As well as having a profound effect on the properties of the ICM, it seems that the heat input also has important consequences for the formation and evolution of the central brightest cluster galaxy (BCG) and therefore for the bright end of the galaxy luminosity function (e.g., Benson et al.\\ 2003; Bower et al.\\ 2006, 2008).  The thermal state of the ICM therefore provides an important probe of these processes and a concerted theoretical effort has been undertaken to explore the effects of various heating sources (e.g., supermassive black holes, supernovae, thermal conduction, dynamical friction heating of orbiting satellites) using analytic and semi-analytic models, in addition to idealised and full cosmological hydrodynamic simulations (e.g., Binney \\& Tabor 1995; Narayan \\& Medvedev 2001; Ruszkowski et al.\\ 2004; McCarthy et al.\\ 2004; Kim et al.\\ 2005; Voit \\& Donahue 2005; Sijacki et al.\\ 2008).  In most of these approaches, it is implicitly assumed that {\\it gravitationally-induced} heating (e.g., via hydrodynamic shocks, turbulent mixing, energy-exchange between the gas and dark matter) that occurs during mergers/accretion is understood and treated with a sufficient level of accuracy.  For example, most current analytic and semi-analytic models of groups and clusters attempt to take the effects of gravitational heating into account by assuming distributions for the gas that are taken from (or inspired by) non-radiative cosmological hydrodynamic simulations --- the assumption being that these simulations accurately and self-consistently track gravitational processes.  However, it has been known for some time that, even in the case of identical initial setups, the results of non-radiative cosmological simulations depend on the numerical scheme adopted for tracking the gas hydrodynamics.  In particular, mesh-based Eulerian (such as adaptive mesh refinement, hereafter AMR) codes appear to systematically produce higher entropy (lower density) gas cores within groups and clusters than do particle-based Lagrangian (such as smoothed particle hydrodynamics, hereafter SPH) codes (see, e.g., Frenk et al.\\ 1999; Ascasibar et al\\. 2003; Voit, Kay, \\& Bryan 2005; Dolag et al.\\ 2005; O'Shea et al.\\ 2005).  This may be regarded as somewhat surprising, given the ability of these codes to accurately reproduce a variety of different analytically solvable test problems (e.g., Sod shocks, Sedov blasts, the gravitational collapse of a uniform sphere; see, e.g., Tasker et al.\\ 2008), although clearly hierarchical structure formation is a more complex and challenging test of the codes.  At present, the origin of the cores and the discrepancy in their amplitudes between Eulerian and Lagrangian codes in cosmologically-simulated groups and clusters is unclear. There have been suggestions that it could be the result of insufficient resolution in the mesh simulations, artificial entropy generation in the SPH simulations, Galilean non-invariance of the mesh simulations, and differences in the amount of mixing in the SPH and mesh simulations (e.g., Dolag et al.\\ 2005; Wadsley et al.\\ 2008). We explore all of these potential causes in \\S 4.  Clearly, though, this matter is worth further investigation, as it potentially has important implications for the competition between heating and cooling in groups and clusters (the cooling time of the ICM has a steep dependence on its core entropy) and the bright end of the galaxy luminosity function.  And it is important to consider that the total heating is not merely the sum of the gravitational and non-gravitational heating terms.  The Rankine-Hugoniot jump conditions tell us that the efficiency of shock heating (either gravitational or non-gravitational in origin) {\\it depends on the density of the gas at the time of heating} (see, e.g., the discussion in McCarthy et al.\\ 2008a).  This implies that if gas has been heated before being shocked, the entropy generated in the shock can actually be amplified (Voit et al.\\ 2003; Voit \\& Ponman 2003; Borgani et al.\\ 2005; Younger \\& Bryan 2007).  The point is that gravitational and non-gravitational heating will couple together in complex ways, so it is important that we are confident that gravitational heating is being handled with sufficient accuracy in the simulations.  A major difficulty in studying the origin of the cores in cosmological simulations is that the group and cluster environment can be extremely complex, with many hundreds of substructures (depending on the numerical resolution) orbiting about at any given time .  Furthermore, such simulations can be quite computationally-expensive if one wishes to resolve in detail the innermost regions of groups and clusters (note that, typically, the simulated cores have sizes $\\la 0.1 r_{200}$).  An alternative approach, which we adopt in the present study, is to use idealised simulations of binary mergers to study the relevant gravitational heating processes.  The advantages of such an approach are obviously that the environment is much cleaner, therefore offering a better chance of isolating the key processes at play, and that the systems are fully resolved from the onset of the simulation.  The relatively modest computational expense of such idealised simulations also puts us in a position to be able to vary the relevant physical and numerical parameters in a systematic way and to study their effects.  Idealised merger simulations have been used extensively to study a variety of phenomena, such as the disruption of cooling flows (G{\\'o}mez et al.\\ 2002; Ritchie \\& Thomas 2002; Poole et al.\\ 2008), the intrinsic scatter in cluster X-ray and Sunyaev-Zel'dovich effect scaling relations (Ricker \\& Sarazin 2001; Poole et al.\\ 2007), the generation of cold fronts and related phenomena (e.g., Ascasibar \\& Markevitch 2006), and the ram pressure stripping of orbiting galaxies (e.g., Mori \\& Burkert 2000; McCarthy et al.\\ 2008b).  However, to our knowledge, idealised merger simulations have not been used to elucidate the important issue raised above, nor have they even been used to demonstrate whether or not this issue even exists in non-cosmological simulations.  In the present study, we perform a detailed comparison of idealised binary mergers run with the widely-used public simulations codes \\flash (an AMR code) and \\gadget (a SPH code).  The paper is organised as follows.  In \\S 2 we give a brief description of the simulation codes and the relevant adopted numerical parameters.  In addition, we describe the initial conditions (e.g., structure of the merging systems, mass ratio, orbit) of our idealised simulations.  In \\S 3, we present a detailed comparison of results from the \\flash and \\gadget runs and confirm that there is a significant difference in the amount of central entropy generated with the two codes.  In \\S 4 we explore several possible causes for the differences we see. Finally, in \\S 5, we summarise and discuss our findings. ",
        "conclusions": "In this paper, we set out to investigate the origin of the discrepancy in the entropy structure of clusters formed in Eulerian mesh-based simulations compared to those formed in Lagrangian SPH simulations.  While SPH simulations form clusters with almost powerlaw entropy distributions down to small radii, Eulerian simulations form much larger cores with the entropy distribution being truncated at significantly higher values. Previously it has been suspected that this discrepancy arose from the limited resolution of the mesh based methods, making it impossible for such codes to accurately trace the formation of dense gas structures at high redshift.  By running simulations of the merging of idealised clusters, we have shown that this is not the origin of the discrepancy. We used the \\gadget code (Springel 2005b) to compute the SPH solution and the \\flash code (Fryxell et al.\\ 2000) to compute the Eulerian mesh solution. In these idealised simulations, the initial gas density structure is resolved from the onset of the simulations, yet the final entropy distributions are significantly different. The magnitude of the difference generated in idealised mergers is comparable to that seen in the final clusters formed in full cosmological simulations.  A resolution study shows that the discrepancy in the idealised simulations cannot be attributed to a difference in the effective resolutions of the simulations.  Thus, the origin of the discrepancy must lie in the code's different treatments of gravity and/or hydrodynamics.  We considered various causes in some detail. We found that the difference was {\\em not} due to: \\begin{itemize} \\item{The use of different gravity solvers. The two codes differ in that \\gadget uses a TreePM method to determine forces, while \\flash uses the PM method alone. The different force resolutions of the codes could plausibly lead to differences in the energy transfer between gas and dark matter. Yet we find that the dark matter distributions produced by the two codes are almost identical when the mesh code is run at comparable resolution to the SPH code.} \\item{Galilean non-invariance of mesh codes. We investigate whether the results are changed if we change the rest-frame defined by the hydrodynamic mesh. Although we find that an artificial core can be generated in this way in the mesh code, its size is much smaller than the core formed once the clusters collide, and is not enough to explain the difference between \\flash and \\gadget. We show that most of this entropy difference is generated in the space of $\\sim 1$ Gyr when the cluster cores first collide.} \\item{Pre-shocking in SPH. We consider the possibility that the artificial viscosity of the SPH method might generate entropy in the flow prior to the core collision, thus reducing the efficiency with which entropy is generated later. By greatly reducing the artificial viscosity ahead of the core collision, we show that this effect is negligible.} \\end{itemize}  Having shown that none of these numerical issues can explain the difference of the final entropy distributions, we investigated the role of fluid mixing in the two codes. Several recent studies (e.g., Dolag et al.\\ 2005; Wadsley et al.\\ 2008) have argued that if one increases the amount of mixing in SPH simulations the result is larger cluster entropy cores that resemble the AMR results.  While this is certainly suggestive, it does not clearly demonstrate that it is the enhanced mixing in mesh simulations that is indeed the main driver of the difference (a larger entropy core in the mesh simulations need not necessarily have been established by mixing).  By injecting tracer particles into our \\flash simulations, we have been able to make an explicit comparison of the amount of mixing in the SPH and mesh simulations of clusters.  We find very substantial differences. In the SPH computation, there is a very close relation between the initial entropy of a particle and its final entropy. In contrast, tracer particles in the \\flash simulation only show a close connection for high initial entropies. The lowest $\\sim5$\\% of gas (by initial entropy) is completely mixed in the \\flash simulation. We conclude that mixing and fluid instabilities are the cause of the discrepancy between the simulation methods.  The origin of this mixing is closely connected to the suppression of turbulence in SPH codes compared to the Eulerian methods. This can easily be seen by comparing the flow structure when the clusters collide: while the \\flash image is dominated by large scale eddies, these are absent from the SPH realisation (see Figure~\\ref{instabilities}). It is now established that SPH codes tend to suppress the growth of Kelvin-Helmholtz instabilities in shear flows, and this seems to be the origin of the differences in our simulation results (e.g., Agertz et  al.\\ 2007).  These structures result in entropy generation through mixing, an irreversible process whose role is underestimated by the SPH method.  Of course, it is not clear that the turbulent structures are correctly captured in the mesh simulations (Iapichino \\& Niemeyer 2008; Wadsley et al.\\ 2008). The mesh forces fluids to be mixed on the scale of individual cells. In nature, this is achieved through turbulent cascades that mix material on progressively smaller and smaller scales: the mesh code may well overestimate the speed and effectiveness of this process.  Ultimately, deep X-ray observations may be able to tell us whether the mixing that occurs in the mesh simulations is too efficient.  An attempt at studying large-scale turbulence in clusters was made recently by Schuecker et al.\\ (2004).  Their analysis of {\\it XMM-Newton} observations of the Coma cluster indicated the presence of a scale-invariant pressure fluctuation spectrum on scales of 40-90 kpc and found that it could be well described by a projected Kolmogorov/Oboukhov-type turbulence spectrum.  If the observed pressure fluctuations are indeed driven by scale-invariant turbulence, this would suggest that current mesh simulations have the resolution required to accurately treat the turbulent mixing process. Alternatively, several authors have suggested that ICM may be highly viscous (eg., Fabian et al.\\ 2003) with the result that fluid instabilities will be strongly suppressed by physical processes. This might favour the use of SPH methods which include a physical viscosity (Sijacki \\& Springel 2006).  It is a significant advance that we now understand the origin of this long standing discrepancy.  Our work also has several important implications.  Firstly, as outlined in \\S 1, there has been much discussion in the recent literature on the competition between heating and cooling in galaxy groups and clusters.  The current consensus is that heating from AGN is approximately sufficient to offset cooling losses in observed cool core clusters (e.g., McNamara \\& Nulsen 2007).  However, observed present-day AGN power output seems energetically incapable of explaining the large number of systems that do not possess cool cores\\footnote{Recent estimates suggest that $\\sim50$\\% of all massive X-ray clusters in flux-limited samples do not have cool cores (e.g., Chen et al.\\ 2007).  Since at fixed mass cool core clusters tend to be more luminous than non-cool core clusters, the fraction of non-cool core cluster in flux-limited samples may actually be an underestimate of the true fraction.} (McCarthy et al.\\ 2008). Recent high resolution X-ray observations demonstrate that these systems have higher central entropies than typical cool core clusters (e.g., Dunn \\& Fabian 2008).  One way of getting around the energetics issue is to invoke an early episode of preheating (e.g., Kaiser 1991; Evrard \\& Henry 1991). Energetically, it is more efficient to raise the entropy of the (proto-)ICM prior to it having fallen into the cluster potential well, as its density would have been much lower than it is today (McCarthy et al.\\ 2008).  Preheating remains an attractive explanation for these systems.  However, as we have seen from our idealised merger simulations, the amount of central entropy generated in our mesh simulations is significant and is even comparable to the levels observed in the central regions of non-cool core clusters.  It is therefore tempting to invoke mergers and the mixing they induce as an explanation for these systems.  However, before a definitive statement to this effect can be made, much larger regions of parameter space should be explored. In particular, a much larger range of impact parameters and mass ratios is required, in addition to switching on the effects of radiative cooling (which we have neglected in the present study).  This would be the mesh code analog of the SPH study carried out by Poole et al. (2006; see also Poole et al.\\ 2008).  We leave this for future work. Alternatively, large cosmological mesh simulations, which self-consistently track the hierarchical growth of clusters, would be useful for testing the merger hypothesis.  Indeed, Burns et al.\\ (2008) have recently carried out a large mesh cosmological simulation (with the \\enzo code) and argue that mergers at high redshift play an important role in the establishment of present-day entropy cores.  However, these results appear to be at odds with the cosmological mesh simulations (run with the \\art code) of Nagai et al.\\ (2007) (see also Kravtsov et al.\\ 2005).  These authors find that most of their clusters have large cooling flows at the present-day, similar to what is seen in some SPH cosmological simulations (e.g., Kay et al.\\ 2004; Borgani et al.\\ 2006). On the other hand, the SPH simulations of Keres et al.\\ (2008) appear to yield clusters with large entropy cores.  This may be ascribed to the lack of effective feedback in their simulations, as radiative cooling selectively removes the lowest entropy gas (see, e.g., Bryan 2000; Voit et al.\\ 2002), leaving only high entropy (long cooling time) gas remaining in the simulated clusters.  However, all the simulations just mentioned suffer from the overcooling problem (Balogh et al.\\ 2001), so it is not clear to what extent the large entropy cores in clusters {\\it in either mesh or SPH} simulations are produced by shock heating, overcooling, or both.  All of these simulations implement different prescriptions for radiative cooling (e.g., metal-dependent or not), star formation, and feedback, and this may lie at the heart of the different findings.  A new generation of cosmological code comparisons will be essential in sorting out these apparently discrepant findings.  The focus should not only be on understanding the differences in the hot gas properties, but also on the distribution and amount of stellar matter, as the evolution of the cold and hot baryons are obviously intimately linked.  Reasonably tight limits on the amount of baryonic mass in the form of stars now exists (see, e.g., Balogh et al.z 2008) and provides a useful target for the next generation of simulations.  At present, merger-induced mixing as an explanation for intrinsic scatter in the hot gas properties of groups and clusters remains an open question.  Secondly, we have learnt a great deal about the nature of gas accretion and the development of hot gas haloes from SPH simulations of the universe.  Since we now see that these simulations may underestimate the degree of mixing that occurs, which of these results are robust, which need revision?  For example, Keres et al.\\ (2005) (among others) have argued that cold accretion by galaxies plays a dominant role in fuelling the star formation in galaxies. Is it plausible that turbulent eddies could disrupt and mix such cold streams as they try to penetrate through the hot halo? We can estimate the significance of the effect by comparing the Kelvin-Helmholtz timescale, $\\tau_{\\rm KH}$, with the free-fall time, $\\tau_{\\rm FF}$.  The KH timescale is given by (see, e.g., Nulsen 1982; Price 2008)  \\begin{equation} \\tau_{\\rm KH} \\equiv \\frac{2 \\pi}{\\omega} \\end{equation}  \\noindent where  \\begin{equation} \\omega = \\frac{2 \\pi}{k} \\frac{(\\rho \\rho')^{1/2} v_{\\rm rel}}{(\\rho+\\rho')} \\end{equation}  \\noindent and $\\rho$ is the density of the hot halo, $\\rho'$ is the density of the cold stream, $k$ is the wave number of the instability, and $v_{\\rm rel}$ the velocity of the stream relative to the hot halo.  If the stream and hot halo are in approximate pressure equilibrium, this implies a large density contrast (e.g., a $10^4$ K stream falling into a $10^6$ K hot halo of a Milky Way-type system would imply a density contrast of $100$).  In the limit of $\\rho' \\gg \\rho$ and recognising that the mode responsible for the destruction of the stream is comparable to the size of the stream (i.e., $k \\sim 2 \\pi / r'$), eqns.\\ (7) and (8) reduce to:  \\begin{equation} \\tau_{\\rm KH} \\approx \\frac{r'}{v_{\\rm rel}}\\biggl(\\frac{\\rho'}{\\rho}\\biggr)^{1/2} \\end{equation}  Adopting $\\rho'/\\rho = 100$, $r' = 100$ kpc, and $v_{\\rm rel} = 200$ km/s (perhaps typical numbers for a cold stream falling into a Milky Way-type system), we find $\\tau_{\\rm KH} \\sim 5$ Gyr.  The free-fall time, $\\tau_{\\rm FF} = R_{\\rm vir}/V_{\\rm circ}(R_{\\rm vir})$ [where $R_{\\rm vir}$ is the virial radius of main system and $V_{\\rm circ}(R_{\\rm vir})$ is the circular velocity of the main system at its virial radius], is $\\sim 1$ Gyr for a Milky Way-type system with mass $M_{\\rm vir} \\sim 10^{12} M_\\odot$.  On this basis, it seems that the stream would be stable because of the large density contrast in the flows.  It is clear, however, that the universality of these effects need to be treated with caution, as the free-fall and KH timescales are not vastly discrepant.  High resolution mesh simulations (cosmological or idealised) of Milky Way-like systems would provide a valuable check of the SPH results.  Finally, the SPH method has great advantages in terms of computational speed, effective resolution and Galilean invariance. Is it therefore possible to keep these advantages and add additional small scale transport processes to the code in order to offset the suppression of mixing?  Wadsley et al.\\ (2008) and Price (2008) have presented possible approaches based on including a thermal diffusion term in the SPH equations. Although the approaches differ in their mathematical details, the overall effect is the same. However, it is not yet clear how well this approach will work in cosmological simulations that include cooling (and feedback), since the thermal diffusion must be carefully controlled to avoid unphysical suppression of cooling in hydrostatic regions (e.g., Dolag et al.\\ 2005). One possibility might be to incorporate such terms as a negative surface tension in regions of large entropy contrast (Hu \\& Adams 2006).  An alternative approach is to combine the best features of the SPH method, such as the way that it continuously adapts to the local gas density and its flow, with the advantage of a Riemann based method of solving the fluid dynamic equations (e.g., Inutsuka 2002).  Clearly, there is a great need to find simple problems in which to test these codes: simple shock tube experiments are not sufficient because they do not include the disordered fluid motions that are responsible for generating the entropy core.  Idealised mergers represent a step forward, but the problem is still not sufficiently simple that it is possible to use self-similar scaling techniques (e.g., Bertschinger 1985, 1989) to establish the correct solution . One  possibility is to consider the generation of turbulent eddies in a fluctuating gravitational potential. We have begun such experiments, but (although the fluid flow patterns are clearly different) simply passing a gravitational potential through a uniform plasma at constant velocity does not expose the differences between SPH and Eulerian mesh based methods that we see in the idealised merger case. We will tackle the minimum complexity that is needed to generate these differences in a future paper."
    },
    "0812/0812.1799_arXiv.txt": {
        "abstract": "{} {We investigate the origin of a flux increase found during a transit of TrES-1, observed with the HST. This feature in the HST light curve cannot be attributed to noise and is supposedly a dark area on the stellar surface of the host star eclipsed by TrES-1 during its transit. We investigate the likeliness of two possible hypothesis for its origin: A starspot or a second transiting planet.} {We made use of several transit observations of TrES-1 from space with the HST and from ground with the IAC-80 telescope. On the basis of these observations we did a statistical study of flux variations in each of the observed events, to investigate if similar flux increases are present in other parts of the data set.} {The HST observation presents a single clear flux rise during a transit  whereas the ground observations led to the detection of two such events but with low significance.  In the case of having observed a starspot in the HST data, assuming a central impact between the spot and TrES-1, we would obtain a lower limit for the spot radius of 42000 km. For this radius the spot temperature would be 4690 K, 560 K lower then the stellar surface of 5250 K. For a putative second transiting planet we can set a lower limit for its radius at 0.37 R$_J$ and for periods of less than 10.5 days, we can set an upper limit at 0.72 R$_J$.} {Assuming a conventional interpretation, then this HST observation constitutes the detection of a starspot. Alternatively, this flux rise might also be caused by an additional transiting planet. The true nature of the origin can be revealed if a wavelength dependency of the flux rise can be shown or discarded with a higher certainty. Additionally, the presence of a second planet can also be detected by radial velocity measurements. } ",
        "introduction": "TrES-1 was the first planet transiting a bright star that was discovered by a wide-field survey \\citep{2004ApJ...613L.153A}, and it has become a frequently observed and well characterized exoplanet. It orbits with a period of 3.03 days a bright K0V star with an apparent magnitude of V=11.79, approximately 157 pc away from us. A detailed light curve analysis done by \\citet{2007ApJ...657.1098W} allowed to lower the errors of the star/planet parameters. \\citet{2005ApJ...626..523C} obtained the first direct detection of light emitted from an exoplanet using the Spitzer Space Telescope by an observation of the TrES-1 secondary eclipse. \\citet{2004ApJ...616L.167S,2006AJ....131.2274S} analyzed the chemical composition of the host star and compared it with other known planet-hosting stars. They found no chemical peculiarity compared to the other stars. \\\\  With the aim to acquire a very precise transit light curve of TrES-1, our group observed in the year 2005 several transit events with the Hubble Space Telescope (GO-10441). The most relevant finding has been a flux increase during one transit, which is clearly a feature coming from the TrES-1 system itself, for a preliminary description see \\citet{2007prpl.conf..701C}. Coverage of numerous further transits events was obtained between 2004 and 2007 with the IAC 80-cm telescope (IAC-80). The motivation for this coverage was for one a long-term study of transit-timing variations and later, a follow-up of the HST observations. The analysis of the flux-rises detected in these observations constitutes the main part of this paper, whereas the results from the timing study will be reported elsewhere. \\\\  Both on TrES-1 \\citep{2007prpl.conf..701C} and on other systems (HD209458, by \\citeauthor{2003ApJ...585L.147S} 2003 based on data from \\citeauthor{2001NewA....6...51D} 2001; HD189733 by \\citeauthor{2007A&A...476.1347P} 2007), there have been reports of potential stellar spots detected from flux-rises during planetary transits. In starspots, intense magnetic fields can suppress the convection of heat to the surface, and hence, that area will radiate less light, causing a dark area on the stellar surface. Starspots could cover up to 55 \\% of the stellar surface \\citep{1998ApJ...501L..73O} and are generally 500 to 2000 Kelvin cooler then the surrounding area. The possibility to detect these during transits was first discussed by  \\citet{2003ApJ...585L.147S}. Their simulations showed that a starspot fully or partly occulted by the transiting planet leaves a short brightness increase in the light curve of a transit. The flux variation feature from a starspot depends on two parameters: Its duration depends on the transit path-length across the spot and across the stellar surface, on the rotational velocity of the parent star, and on the projected orbital speed of the transiting planet. The ratio of transit path-lengths may be interpreted in first order, assuming a central impact,  as an area ratio between spot and the transiting planet \\citep{2003ApJ...585L.147S}. The feature's amplitude for another, depends on the temperature difference between the normal stellar surface and the spot as well as on the fraction of the planet's shadow that grazes over the spot, unless the projected planet comes to lie fully within the starspot, in which case only the temperatures are relevant.\\\\  For example, \\citet{2007A&A...476.1347P} found clear evidence of starspots on the star HD189733 by observing a transit of its planet with the HST at a high signal-to-noise ratio. They found an increase in flux of 0.001 mags and attributed this to the planet occulting a starspot. On TrES-1, there has been a previous ground-based search for spots during transits by \\citet{2007ApJ...657.1098W} with negative results. They gave an upper 2 $\\sigma$ limit for spots of $f < 4 x 10^{-4}$, where $f$ is the product of area ratio and the intensity contrast.\\\\  In this paper, we present a second hypothesis for the origin of the short flux-rise observed with HST, occurring if two planets transit simultaneously, and one of the planets occults the other: In a system of two transiting planets with different orbits, the known one has a short period (planet A) and the hypothetical other one has a much longer period (planet B). Assuming that the transit of the short-period planet, i. e. planet A, is being observed and that the centers of both transits coincide more or less in time, then planet B can be considered as moving slowly across the stellar disc. When planet A is eclipsing planet B; it is causing in this moment a flux increase in the combined transit light curve. Due to the short observation span, the ingress and egress of the much longer transit events of planet B might neither have been observed nor detected, see Figure \\ref{model_TrES1_2_full} for an arbitrary example light curve. For a light curve of two transiting planets, the flux variation depends on the area ratio and on the impact parameter between both planets, whereas the duration depends on the area ratio and on the respective orbital velocities of both planets. \\\\  In Section 2 we describe the HST and IAC-80 observations and data reductions. We identified an increase in flux in one HST transit observation. Diagnostics to differentiate between the two origins of the flux rise are presented in Section 3, where we evaluate its color dependency from red to blue and then further investigate the flux variability of the other IAC-80 light curves. A discussion of the results will take place in Section 4.\\\\ ",
        "conclusions": "\\label{DiscCon}  We observed transits of TrES-1 in 2004 and 2005 with the HST and identified an intriguing flux rise during one transit. We further followed-up transits of TrES-1 during the years 2004 to 2007 from the ground, finding two further possible flux rises, and used the color information in the HST data to assess the origin of this feature. We can identify two possible causes, one is the possibility of a circular starspot occulted by TrES-1 during transit and the other is the detection of a second smaller transiting planet in an orbit larger than TrES-1. \\\\  A good fit for the feature's shape can be obtained by using a circular dark patch on the stellar surface and neglecting deformation of a circular disk due to projection on a sphere. Assuming further that the dark patch and TrES-1 have a central impact on each other and using the orbital velocity for TrES-1 of $v=140$ km/s, then, in a time interval of 10 min., TrES-1 crosses 84000 km, which gives us a lower limit for the diameter of the occulted patch. The real diameter depends on the stellar rotation (for the starspot hypothesis) or on the orbital velocities of the transiting planets (for the planet-planet hypothesis). \\\\  The host star of TrES-1 is an active star \\citep{2004ApJ...616L.167S} which is favouring the spot scenario. In that case, we calculated first the spot radius. Since the parent star of TrES-1 is slowly rotating with a rotational velocity of $I_S=1.08$ km/s \\citep{2005ApJ...621.1072L}, corresponding to a rotation period of 38.4 days, the stellar rotation can be neglected and under the assumption of a central impact between TrES-1 and spot, the spot radius is simply half the diameter as obtained previously, i. e.  42000 km. An effective temperature of the stellar surface of 5250 K was obtained by \\citet{2004ApJ...616L.167S} and combined with the equation 1 given by \\citet{2003ApJ...585L.147S}, we estimated the spot temperature to be approximately 4690 K for a starspot with the radius of 42000 km. These values are perfectly reasonable for a large sunspot.\\\\  To generate a realistic model that shows that the observed HST transit light curve might as well arise from two occulting planets, we used ``Universal Transit Modeller'' (UTM, \\citeauthor{2008arXiv0807.3915D} 2008). UTM is an IDL-routine permitting simulations of transit light curves of multiple bodies, including rings and moons, and the fitting of observed light curves. We fitted first the summed HST light curve without considering the flux rise part. As fit parameters we used the radii of the host star and of the planet TrES-1, respectively; the limb darkening coefficient, assuming a quadratic limb darkening law, the inclination and the distance. We then used the best fit parameters and introduced a second transiting planet, but in an expanded orbit. However, we lack the determination of an unique period for the putative second transiting planet due to the singularity of the observed event and we further have no significant evidence in the IAC-80 data to derive a period. \\\\  Hence, we can only give as an example a model of two transiting planets, see Figures \\ref{HST_visits} and \\ref{model_TrES1_2_full}. Figure \\ref{model_TrES1_2_full} shows a complete example model, including ingress and egress of an arbitrary putative second planet. The introduced second planet has an arbitrary period of 300 days and radius of 4.25 R$_{\\oplus}$. Its transit duration is 10 hours and the depth is 2.7 mmags. It is very unlikely that a complete transit of this second planet would have been observed, because the transit duration is too long for nightly observations and it is even more difficult to keep a sufficient high photometric precision from ground over such a long time span due to changing atmosphere, extinction and airmass. All IAC-80 observations were done in time spans shorter than 10 hours and even under ideal conditions, i. e. at low airmass, the differential photometry applied to our transit observations did not correct completely for extinction. Therefore, as noted in Section \\ref{reduction}, we corrected the extinction by subtracting a second order polynomial fit from the observed light curves which removed all possible hints to a second planet's transit from the light curve's off-transit part.\\\\  \\begin{figure} \\centering \\includegraphics[angle=90,scale=0.34]{model_TrES1_2_full.ps} \\caption{Example of the same planet-planet transit model as shown in Figure \\ref{HST_visits}, for TrES-1 and an arbitrary additional transiting planet, but showing the model over longer time span, so that the ingress and egress of a possible second planet are visible. The vertical solid lines indicate the usual observation span of 6 hours centered at the mid-transit of TrES-1. We note that during this 6 hour time span the apparent transit depth of TrES-1 is identical with or without a second planet.  \\label{model_TrES1_2_full}} \\end{figure}  Whereas a mutual occultation would be an extremely rare event, limits for the presence of transits from a second planet can be derived from size-estimates of this planet. Its minimum size can be obtained from the amplitude of the mutual transit event and assuming that TrES-1 and the other planet have a central impact on each other. In this case, the flux increase is directly related to the area of overlap of the projection of the two planets into the star, and is independent of the period. For TrES-1 we have a transit depth of 23 mmags and the flux rise height is 2.7 mmags, and hence we obtain an area of overlap of $1/8.5$ of the cross section of TrES-1. Knowing the radius of TrES-1, 1.081 R$_J$ \\citep{2007ApJ...657.1098W}, we obtained a minimum size of a second planet of 0.37 R$_J$. However, the shape of the flux rise indicates a more grazing impact between TrES-1 and the hypothetical second planet, favouring a bigger radius.\\\\  The original STARE and PSST data of the field where TrES-1 has been discovered cover a time span of 96 days \\citep{2004ApJ...613L.153A}. This is a time span long enough to possibly detect a putative second planet on a longer orbit, but no transit has been detected in this data. Therefore, based on an assumption of a standard deviation of 10 mmags in the TrES-1 discovery data, we can set an upper limit for the size of a possible second planet. In the limiting case, a transit with a depth of 10 mmags would not have been detected. This 10 mmags corresponds to a transit depth of a planet with a radius of 0.72 R$_J$, which constitutes an upper detection limit for the radius of a possible second planet. \\\\ \\\\ For a putative long-period outer planet and due to the limited and gaped observations of the TrES-1 field the detection depends mostly on the probability of having observed a transit. Therefore, we further constrain the period in which this upper radius limit is valid. We applied a statistical analysis to a time vector which contains the time of observation. We established 2000 trial periods between a period of 4 days and 100 days on the STARE/PSST discovery data. Within each trial period we set randomly 300 epochs and tested how often we could recover a transit signal. We set a detection limit for a possible planet at periods where the recovery rate drops below 95 \\%. For the original STARE and PSST data this occurs at a period of 10.5 days. Hence, for periods of less than 10.5 days, we can set an upper limit for the radius of 0.72 R$_J$ for a possible second planet. However, larger planets in longer periods are possible.  \\\\ \\\\ Based on the conventional interpretation of flux-rise features during transits, we may conclude that a starspot has been detected in the HST data and possibly as well in the ground-based IAC-80 data. We also presented an alternative hypothesis of a mutual planet-planet occultation, which would be a  rare but not impossible event. The verification or exclusion of this hypothesis is difficult unless a wavelength-dependency of the flux-rise features can be demonstrated, or until the presence of the secondary planet can either be verified or excluded, be it from observations of its own transits or from radial velocity measurements."
    },
    "0812/0812.3080_arXiv.txt": {
        "abstract": "We present evidence that the star-forming region NGC~346/N66 in the Small Magellanic Cloud is the product of hierarchical star formation, probably from more than one star formation event. We investigate the spatial distribution and clustering behavior of the pre-main sequence (PMS) stellar population in the region, using data obtained with {\\sl Hubble Space Telescope}'s Advanced Camera for Surveys. By applying the nearest neighbor and minimum spanning tree methods on the rich sample of PMS stars previously discovered in the region we identify ten individual PMS clusters in the area and quantify their structures. The clusters show a wide range of morphologies from hierarchical multi-peak configurations to centrally condensed clusters. However, only about 40 per cent of the PMS stars belong to the identified clusters. The central association NGC~346 is identified as the largest stellar concentration, which cannot be resolved into subclusters. Several PMS clusters are aligned along filaments of higher stellar density pointing away from the central part of the region. The PMS density peaks in the association coincide with the peaks of [{\\ion{O}{3}}] and  8~$\\mu$m emission. While more massive stars seem to be concentrated in the central association when considering the entire area, we find no evidence for mass segregation within the system itself. ",
        "introduction": "Star formation is generally thought to be clumped into a hierarchy of structures, from small multiple systems to giant stellar complexes and beyond \\citep[e.g.][]{ef+el98}, and the interaction between gravity and supersonic turbulent motions has been proposed as a major player in the formation of hierarchical structures \\citep{maclow+klessen04,elm06,ballesteros07,mckee+ostriker07}. The interstellar medium has also been observed to be arranged in a hierarchical structure \\citep{scalo85,vazquez04}, sometimes described as fractal, from the scales of the largest giant molecular clouds (GMCs) down to individual cores and embedded clusters, which are sometimes hierarchical themselves.  The velocity structure in such turbulent regions is suggested to be scale-free, meaning that the same physical processes and the same velocity-separation relations occur over a wide range of absolute scales \\citep[e.g.][]{elm06}. According to this hypothesis, large clumps originate from the large velocities that compress the gas on large scales, and small clumps result from the compression of the gas characterized by small velocities on small scales. Considering that stars are formed by the collapse of such hierarchical structured clumps, they are themselves hierarchically structured. Indeed, stars are mostly formed in groups \\citep[e.g.][]{sta+pal05}, which can be seen as the bottom parts of this hierarchy \\citep[e.g.][]{elm00}.  Naturally, the structure of such systems, i.e.\\ the spatial distribution of their members, may hold important information on the mechanism and the initial conditions of their formation. However, a thorough investigation of the structural behavior of star formation in environments different than our Galaxy was not possible, until recently, due to observational limitations. In this paper we present such an investigation for the brightest star-forming region in the Small Magellanic Cloud (SMC).  \\begin{figure} \\epsscale{1.1} \\plotone{ngc346_fig1.eps} \\caption{Color-composite image of the three partially overlapping ACS/WFC observed fields of NGC 346, covering a total area of $5\\arcmin \\times 5\\arcmin$. North is up, and east is to the left. Image Credits: NASA, ESA and A. Nota (STScI/ESA).\\label{fig:ngc346-ima}} \\end{figure}  LHA 115-N66 or in short N66 \\citep{henize56} is a very active star-forming region in the SMC with a diameter of more than 9\\arcmin\\ at a distance of about 60\\,kpc \\citep[e.g.][]{laney+stobie94}. This \\ion{H}{2} region is also known as DEM~S~103 \\citep{davies76} or NGC~346 referring to the bright stellar association located in its center. Being the largest and most luminous \\ion{H}{2} region in the SMC, NGC~346/N66 has been mapped in several wavelengths. Observations of its area have been taken in the X-rays with {\\em XMM-Newton} and {\\em Chandra} \\citep{naze02, naze04}, in the UV with the {\\em International Ultraviolet Explorer} \\citep{deboer+savage80} and the {\\em Far-Ultraviolet Spectroscopic Explorer} \\citep{danforth03}, in the CO(2$-$1) line with the SEST telescope \\citep{contursi00}, and in the infrared with the {\\em Infrared Space Observatory} \\citep{contursi00} and more recently with the {\\em Spitzer Space Telescope} \\citep{bolatto07, simon07}. The region has also been covered by observations of the radio continuum \u00b8\\citep{ye91, reid06} and of the 21-cm atomic line \\citep{staveleysmith97}. Maps of the ionized gas have also been made in H$\\alpha$ \\citep[e.g.][]{kennicutt88, lecoarer93}, and [{\\ion{O}{3}] with NTT \\citep[e.g.][]{rubio00}.   N66 is characterized by weak molecular emission and is also deficient in \\ion{H}{1}. It hosts, though, a variety of discrete H$_2$ emission peaks, which coincide with the main features of the ionized gas, and with compact embedded young clusters, where candidate young stellar objects (YSOs) have been identified. The central part of the whole \\ion{H}{2} region is characterized by an oblique bright emission region, extending from southeast to northwest, known as the bar of N66 (Fig.~\\ref{fig:ngc346-ima}). This area is dominated by a strong UV radiation field related to the central association NGC~346. Indeed this system hosts the largest sample of spectroscopically confirmed OB stars in the whole SMC \\citep{massey89}, which have been the subject of several previous investigations \\citep[e.g.][]{niemela86,walborn00,evans06, hunter08}. NGC~346 ionizes the remaining molecular cloud to the southwest of the bar, triggering the development of a photo-dissociated region seen as a curved front eating into the cloud.  Recent observations with the {\\em Advanced Camera for Surveys} (ACS) onboard the {\\em Hubble Space Telescope} (HST) provided substantial information on the stellar content of this magnificent star-forming region by revealing a plethora of low-mass pre-main sequence (PMS) stars. This discovery clearly indicates that NGC~346/N66 also hosts low-mass star formation. Our photometric data set delivers more than 98\\,000 stars in total, with almost 8000 being low-mass PMS stars \\citep{gouliermis06}. The subsequent investigation of their spatial distribution showed that these stars are not homogeneously distributed across the region, being grouped in discrete concentrations, apart from the association NGC~346. Most of these ``subclusters'' are found to coincide with bright H$\\alpha$ emission, and the H$_2$ emission peaks \\citep{h08}. Low-mass stars with masses \\lsim~2-3~M$_\\odot$ remain in their PMS phase for $\\sim$~20-30~Myr \\citep[e.g.][]{sta+pal05,briceno07}. As a consequence the rich sample of low- and intermediate-mass PMS stars, found with ACS in NGC~346/N66, offers a unique opportunity to investigate clustered star formation within a time-scale of the order of $\\sim$~10~Myr and length-scales of 1 to 100~pc in a detail never achieved before for a star-forming region outside our Galaxy.   In this paper we analyze the spatial distribution of the young PMS population unveiled by HST/ACS in the bright SMC star-forming region NGC~346/N66 and compare the results to current theoretical scenarios of star formation. For this purpose, we make use of our deep ACS photometry of this region, and apply two complementary statistical methods to identify all the compact PMS clusters and analyze their structures. The data set used here is described in Section~\\ref{sec:data}, the statistical methods are explained in Section~\\ref{sec:methods} and the results of our structural analysis of the PMS stars is presented in Section~\\ref{sec:analysis}. A discussion on the star formation process revealed from our analysis is given in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Theoretical Concepts}  With the advancement of infrared detectors and the ability to do wide-area surveys, it has been recognized that a large fraction of all stars in the Milky Way form in clusters and aggregates of various size and mass scales and that isolated or widely distributed star formation is the exception rather than the rule \\citep{lada03}. It is therefore highly interesting to study the clustering behavior in neighboring galaxies as well. The complex hierarchical structure of molecular clouds provides a natural explanation for clustered star formation. Molecular clouds vary enormously in size and mass. In small, low-density, clouds stars form with low efficiency, more or less in isolation or scattered around in small groups of up to a few dozen members. Denser and more massive clouds may build up stars in associations and clusters of a few hundred members.  This appears to be the most common mode of star formation, at least in the solar neighborhood \\citep{adams01}. The formation of dense rich clusters with thousands of stars or more is rare.  The cluster NGC 346, therefore, is a very extreme case of clustered star formation.  Star formation is governed by the complex interplay between gravitational attraction in molecular clouds and any opposing processes such as supersonic turbulence, thermal pressure and magnetic fields.  There are two main competing models that describe the evolution of the cloud cores.  It was proposed in the 1980's that cores in low-mass star-forming regions evolve quasi-statically in magnetically subcritical clouds \\citep{shu87}. Gravitational contraction is mediated by ambipolar diffusion \\citep{Mou76, Mou91} causing a redistribution of magnetic flux until the inner regions of the core become supercritical and go into dynamical collapse.  Measurements of the Zeeman splitting of molecular lines in nearby cloud cores, however, indicate magnetic field strengths that fall below the critical value, in some cases only by a small margin, in many cases however, by factors of many \\citep{crutcher99, bourke01,crutcher08}. Magnetic fields therefore seem not to be the dominant agent regulating stellar birth.  This has lead to the suggestion that the supersonic turbulence ubiquitously observed in molecular clouds is the main physical process governing star formation \\citep{maclow+klessen04,ballesteros07,mckee+ostriker07}. Supersonic turbulence plays a dual role. On large scales it can support clouds against contraction, however, on small scales it can provoke localized collapse. Turbulence establishes a complex network of interacting shocks, where dense cores form at the stagnation points of convergent flows. The density can be large enough for gravitational collapse to set in. However, the fluctuations in turbulent velocity fields are highly transient.  The random flow that creates local density enhancements can disperse them again.  For local collapse to actually result in the formation of stars, high density fluctuations must collapse on time scales shorter than the typical time interval between two successive shock passages.  Only then are they able to `decouple' from the ambient flow and survive subsequent shock interactions.  The shorter the time between shock passages, the less likely these fluctuations are to survive. Hence, the timescale and efficiency of protostellar core formation depend strongly on the wavelength and strength of the driving source \\citep{klessen00b,Heitschetal2001,Vazquez03}, and accretion histories of individual protostars are strongly time varying \\citep{Klessen2001b,sk04}.  Interstellar turbulence is observed to be dominated by large-scale velocity modes \\citep{maclow00b,Ossenkopfetal2001,Ossenkopf02}. This implies it is very efficient  in sweeping up molecular cloud material, creating massive coherent structures which can go into large-scale gravitational collapse. In the extreme case, dense clusters with many thousand stars build up on one or two crossing times \\citep{Hartmann01,elmegreen00,elmegreen07}. Numerical simulations of gravoturbulent cloud fragmentation indicate that in such cases star clusters would build up in a hierarchical fashion with sub-clusters which evolve dynamically and merge together to build up the final, centrally-condensed cluster \\citep{klessen00,clarke00,bonnell03,clark05b,BonnellBate2006}, with stellar feedback strongly influencing this process \\citep[][]{krumholz07}. \\citet{sk06} investigated the temporal evolution of a large set of numerical simulations of gravoturbulent fragmentation and showed that the \\Q\\ parameter of all resulting clusters increases with time. This is confirmed by observations of embedded clusters in nearby molecular clouds, where the younger Class 0/1 sources show substantially lower \\Q\\ values than the more evolved Class 2/3 objects \\citep{skf08}.    \\subsection{The Structure of NGC 346/N66}  Within our investigation of the spatial distribution of PMS stars in the NGC~346/N66 region using the nearest neighbor and minimum spanning tree methods an increased NN density of PMS stars is found in the central association and in two filamentary structures, one extending to the north-east from the center of N66, and another at the southern edge of the observed area, extending from east to west. We identify ten individual clusters of PMS stars in the region, with the central association being the largest. While several PMS density peaks are found within the association NGC~346, the system could not be resolved into individual subclusters. Our methods show evidence of centrally concentrated UMS stars and PMS stars in a hierarchical distribution. The PMS density peaks of the association coincide with peaks of both [{\\ion{O}{3}] and 8~$\\mu$m emission, indicating that the system is in an early evolutionary stage.  The identified PMS clusters in the whole area of NGC~346/N66 show a wide range of morphologies from hierarchical multi-peak configurations to centrally condensed clusters. Only about 40\\% of the PMS stars in the investigated area belong to the ten identified clusters. The remaining 60\\% of PMS stars is found concentrated in small (statistically insignificant) groups and mostly dispersed over the entire area. This fraction of clustered PMS stars is comparable to that found in Galactic star-forming regions. In their SCUBA survey of the Perseus molecular cloud \\cite{hatchell05} found that 40~-~60\\% of the protostars and prestellar cores are located in small clusters ($<$~50~M$_\\odot$) and isolated objects. Investigations in other Galactic star-forming regions like Ophiuchus, Orion and Monoceros yield a percentage of distributed populations between 20 and 60\\% with a ``typical'' value of 25\\% \\citep{carpenter00, allen07}. However, there are uncertainties due to the incompleteness of the samples and the definition of the clusters. Moreover, the clustering properties themselves depend on the age of the population considered and on the masses of the stars. As a consequence, it is very difficult to compare and interpret these numbers.  Based on the analysis of \\Q\\ versus magnitude, both PMS and UMS stars show roughly the same degree of central concentration over the entire magnitude range. When we consider the entire area, \\Q\\ is found to decrease with magnitude, providing evidence that both PMS and UMS stars are centrally concentrated in respect to the whole observed region of NGC~346. This, especially for the bright stars, is expected considering that the central association contains the largest fraction of stars. However, when we consider only the central region of the area, \\Q\\ does not seem to depend strongly on magnitude, and therefore no evidence of mass segregation in the central association is found.  Whether the clusters in NGC~346 have formed coevally \\citep{sabbi07} or in a sequential process over several Myr \\citep{massey89, rubio00} is still under debate. As young clusters seem to evolve from a hierarchical configuration to a centrally concentrated one with time, the fact that the identified clusters in NGC~346 show a wide range in their structures with \\Q\\ values ranging from 0.65 to 0.93, may indeed indicate that they are at different evolutionary stages. However, the absolute \\Q\\ values are not correlated with the cluster age, but probably depend on a variety of additional factors like initial conditions or turbulent energy in the cloud \\citep{skf08}, and therefore this is no direct proof that the clusters have formed at different times. However, the rather high \\Q\\ values of the southern clusters (H, I, J) around or above 0.8 agree with other age indicators. While the peaks of the central concentration and the clusters to the north of it coincide with the peaks of dust emission (indicated by the  8~$\\mu$m observations) and gas (traced by [{\\ion{O}{3}}]), the southern clusters appear to be associated with less amounts of gas and dust, suggesting an older age of these clusters. \\citet{h08} find that the northern clusters comprise a younger PMS population (not older than 2.5\\,Myr) than the other clusters. Recently, \\citet{gouliermis08} presented evidence that these clusters are the product of triggered star formation. In combination, all these indications strongly suggest that the identified PMS clusters are not the product of a single star formation event.  In conclusion our investigation of the spatial distribution of PMS stars in the region of NGC~346/N66 using the nearest neighbor and minimum spanning tree methods showed that this region appears to be the product of hierarchical star formation, where the individual clusters show different structures according to their size, evolutionary stage, position in the cloud and possible additional internal or environmental factors. The overall structure and statistical properties of the region are consistent with numerical calculations of gravoturbulent cloud fragmentation \\citep[e.g.][]{klessen00b,clark05a,BonnellBate2006}. NGC~346/N66 is a unique case, where the complexity of star formation can be observed in its entire length-scale."
    },
    "0812/0812.1566_arXiv.txt": {
        "abstract": "We introduce a new statistical technique for extracting the inhomogeneous reionization signal from future high-sensitivity measurements of the cosmic microwave background temperature and polarization fields.  If reionization is inhomogeneous, then the optical depth to recombination will be a function $\\tau(\\n)$ of position on the sky.  Anisotropies in $\\tau(\\n)$ alter the statistics of the observed CMB via several physical mechanisms: screening of the surface of last scattering, generation of new polarization via Thomson scattering from reionization bubbles, and the kinetic Sunyaev-Zel'dovich effect. We construct a quadratic estimator  $\\htau_{\\ell m}$ for the modes of the $\\tau$-field. This estimator separates the patchy reionization signal from the CMB in the form of a noisy map, which can be cross-correlated with other probes of reionization or used as a standalone probe. A future satellite experiment with sufficient sensitivity and resolution to measure the lensed B-modes on most of the sky can constrain key parameters of patchy reionization, such as the duration of the patchy epoch or the mean bubble radius, at the $\\sim 10$\\% level. ",
        "introduction": "Upcoming generations of cosmic microwave background (CMB) experiments will make precise measurements of secondary anisotropies on small scales $(2000\\lesssim\\ell\\lesssim 10000)$ in temperature and E-mode polarization and will also have sufficient sensitivity to measure secondary B-modes, for example the ones arising from gravitational lensing of the primary E-mode.  Secondary anisotropies are caused by fluctuations in the distribution of matter after recombination and are generated by gravitational lensing, inverse Compton scattering of photons by the hot intracluster medium (the thermal Sunyaev-Zel'dovich effect), and the Doppler effect produced by Thomson scattering of photons from radially moving electrons (the kinetic Sunyaev-Zel'dovich effect (kSZ)).  Secondary anisotropies are also generated during the inhomogeneous, or patchy, phase of reionization. Analytic studies and numerical simulations suggest that the period of reionization was more complex than a sudden change in the ionization fraction \\cite{Zahn:2006sg,Furlanetto:2004nt,Barkana:2000fd}. Reionization is an inhomogeneous process and the inhomogeneities contribute  to the small scale CMB temperature and polarization anisotropies. In contrast to other secondaries, such as gravitational lensing \\cite{Hu:2001fa,Hu:2001kj,Zaldarriaga:1998ar,Okamoto:2003zw,Hirata:2002jy,Hirata:2003ka,Smith:2007rg,Hirata:2008cb} or the Sunyaev-Zel'dovich effect (SZ) \\cite{Sunyaev:1970er,Sunyaev:1980vz}, the potential science returns from patchy reionization have not been extensively studied in the context of the CMB.  There are currently almost no observational constraints on the evolution of the ionization fraction during the epoch of reionization. The Gunn-Peterson trough in high-redshift quasar spectra has been observed, showing that the transition from partial reionization to a fully ionized universe ($\\overline x_e\\approx 1$) took place at redshift $z\\sim 6$ \\cite{Becker:2001ee,White:2003sc,Fan:2005es}. Large-scale E-mode polarization measurements from five-year WMAP data show a total optical depth to recombination $\\tau=0.087\\pm 0.017$. If reionization is assumed instantaneous, this would imply a transition redshift $z_{\\rm rei}=11.0\\pm 1.4$ with $68 \\%$ confidence \\cite{Dunkley:2008ie}. Combining these two observations, we therefore have good indirect evidence for patchy reionization but no detailed information.  Predictions from simulations tell us that in order to allow the contributions of the first stars \\cite{Ricotti:2004xf}, the beginning of reionization could go until $z\\sim 30$. However, because of a lack of experimental data, this prediction is still uncertain. The large-scale E-mode is sensitive to reionization history, but only contains information about the (spatial) average ionization fraction $\\overline x_e(z)$ during reionization, not the size or morphology of the ionized regions. The WMAP data are not sensitive enough to place strong constraints on the reionization history beyond determination of the total optical depth $\\tau$, but future large-scale E-mode measurements from Planck or CMBpol can constrain up to $\\sim 5$ principal components or redshift bins in $\\overline x_e(z)$ \\cite{Mortonson:2007hq,Hu:2003gh}.  The patchy reionization signal is small and, therefore, will be difficult to separate from the other secondaries using the power spectrum alone. The patchy contribution to the temperature power spectrum is smaller than the sum of the contributions from lensing and low-redshift kSZ on all angular scales \\cite{Gnedin:2000gr,Zahn:2005fn,Hu:1999vq,Iliev:2006un,McQuinn:2005ce,Santos:2003jb}. The thermal SZ signal from galaxy clusters leads to a larger signal than the one caused by the kSZ effect, but it can be separated from the other secondaries due to its non-blackbody frequency dependence. In polarization, currently favored models produce B-mode polarization power spectra with an amplitude ($\\sim 0.01$$\\mu$K)  that is significantly lower than the B-modes coming from gravitational lensing ($\\sim 0.1$$\\mu$K) \\cite{Gruzinov:1998un,Liu:2001xe,Mortonson:2006re,Dore:2007bz}.  In this paper, we will propose an estimator which isolates the patchy reionization signal in the CMB. If reionization is patchy, then the optical depth to last scattering will be a 2D field $\\tau_{\\ell m}$ rather than a constant $\\tau$. Our estimator $\\htau_{\\ell m}$ will reconstruct each mode of this field, within statistical noise, using the small change in the CMB statistics which is induced by the mode. This construction was inspired by the well-known lens reconstruction estimator \\cite{Hu:2001fa,Hu:2001kj,Okamoto:2003zw,Hirata:2002jy,Hirata:2003ka}, which reconstructs each mode $\\phi_{\\ell m}$ of the CMB lens potential in an analogous way. We will show that our new estimator $\\htau_{\\ell m}$ isolates the patchy signal, in the sense that its expectation value is simply the underlying field $\\tau_{\\ell m}$, with no contribution from the Gaussian part of the CMB.  In \\S\\ref{sec:reionization_model} and ~\\S\\ref{sec:reionization_and_the_CMB}, we describe our semi-analytic modeling of reionization and its effect on the CMB. We will split the signal from patchy reionization into three effects: direction-dependent screening of the acoustic peaks, polarization generated by Thomson scattering from reionization bubbles, and temperature anisotropy generated by radial motion of the bubbles (i.e. the kSZ effect).  In \\S\\ref{sec:simple_estimator} and \\S\\ref{sec:quad_est_realistic} we will construct the estimator $\\htau_{\\ell m}$, by analogy with the quadratic estimator for lens reconstruction. We will show that the first two effects from \\S\\ref{sec:reionization_and_the_CMB} (screening and Thomson scattering) have an algebraic form where the quadratic estimator construction applies, but the third effect (kSZ) does not, because the kSZ anisotropy is proportional to the velocity field during reionization, and this field is not highly correlated with the primary CMB. Consequently, our estimator $\\htau_{\\ell m}$ is not very sensitive to the kSZ effect, but should be a near-optimal statistic for extracting the screening and Thomson signals. Although a complete treatment which formally extracts all the signal-to-noise is quite complex (\\S\\ref{ssec:pca}) we show that a simple estimator (\\S\\ref{ssec:simple_estimator}) contains all the information in practice.  In \\S\\ref{sec:simulation} we demonstrate our estimator with Monte Carlo simulations. We make simplifying assumptions, such as simulating the lensed component of the $B$-mode as if it were a Gaussian field. In \\S\\ref{sec:forecasts} we present our forecasts. A future all-sky $B$-mode experiment can constrain key reionization parameters, such as the mean radius of the ionized bubbles or the duration of reionization, at the $\\sim 10\\%$ level. It might be possible to improve this measurement making use of a delensing procedure or cross-correlating the $\\tau$-map with large-scale structure. We conclude in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In this paper we have introduced a new technique for extracting the patchy reionization signal from the cosmic microwave background. In a fixed realization of the $\\Delta\\tau$-field, the two-point function of the CMB acquires terms which are linear in $\\Delta\\tau$, and this allows us to construct a quadratic estimator $\\htau_{\\ell m}$ for the modes of the $\\Delta\\tau$-field. This construction is formally similar to the lens reconstruction estimator $\\hphi_{\\ell m}$ which estimates modes of the CMB lens potential $\\phi$ using the induced two-point function of the observed CMB.  In principle the general construction in this paper can use all combinations of the \\{T,E,B\\} fields, and can distinguish transverse modes of the $\\Delta\\tau$-field at different redshifts using the redshift dependence of the induced CMB two-point function. However, we find that in practice, the redshift dependence is so weak that only one 2D field $(\\Delta\\tau)_{\\ell m}$ can be reconstructed, and that nearly all the signal-to-noise comes from the $EB$ quadratic estimator (with contributions from Thomson and screening effects which are comparable in signal-to-noise). This analysis also implies that it is possible to write the estimator in a simple form (\\S\\ref{ssec:simple_estimator}) which ignores subtleties like the redshift dependence.  It is worth emphasizing that the quadratic estimator framework depends on having a large cross-correlation between the new CMB anisotropy generated by patchy reionization and the primary CMB. This is why our estimator is sensitive to the Thomson and screening effects during patchy reionization: in both of these cases, the new CMB anisotropy is highly correlated to the large-scale reionization E-mode, or to the acoustic E-mode, respectively. However, the quadratic estimator is only sensitive to the kSZ effect in temperature via its cross-correlation to the Doppler effect on large scales. Another way of saying this is that the quadratic estimator can only ``see'' the (Doppler, kSZ) cross-correlation but not the (kSZ, kSZ) power spectrum, since the two are $\\bigoh(\\Delta\\tau)$ and $\\bigoh(\\Delta\\tau^2)$ effects respectively. For this reason, we consider it a separate unsolved problem, not addressed in this paper, to construct a statistic which can separate the kSZ signal in temperature from other sources of small-scale power such as CMB lensing or the primary anisotropy.  In this paper we have done a complete analysis of the signal-to-noise that can be achieved using a quadratic estimator construction, concluding with a Fisher matrix forecast for a simple five-parameter reionization model (\\S\\ref{ssec:fisher}).  In this model, there are parameter degeneracies that leave only two well-constrained power spectrum observables: an ``amplitude'' observable $A$ which is mainly sensitive to the duration of patchy reionization, and a ``peak location'' observable $R_{\\rm eff}$ which is mainly sensitive to the size distribution of bubbles.  Let us conclude by commenting on some possible continuations of this work which are outside the scope of a first paper.  One such extension is to increase the realism of the reionization model, and to study prospects for constraining a more complex parameter set. It is unclear whether the ``two-observable'' picture that we obtained in our Fisher matrix analysis of the five-parameter model will be qualitatively different in a more complete model.  Another natural continuation is to study prospects for cross-correlation with other flavors of cosmological data which probe the epoch of inhomogeneous reionization. A nice feature of the quadratic estimator framework (as compared to other statistics such as the CMB power spectrum) is that the patchy reionization signal is extracted in the form of a map $\\htau_{\\ell m}$ which can be cross-correlated to other datasets. This cross-correlation technique is part of future work, and can be used either to boost signal-to-noise for a first detection (e.g. \\cite{Smith:2007rg,Hirata:2008cb} in the context of CMB lens reconstruction) or to make the measurement more robust by reducing the range of systematic effects which may contaminate the auto power spectrum of the reconstruction.  Finally, a continuation of this work which is currently in preparation is the implementation of the estimator in more realistic simulations which include the non-Gaussian statistics of the lensed B-mode. The simulations in this paper (\\S\\ref{sec:simulation}) show that our signal-to-noise forecasts can be achieved in toy Monte Carlo simulations in which the lensed B-mode has the correct power spectrum but the statistics are treated as Gaussian. Including non-Gaussianity is an important step: it can potentially bias the estimator $\\htau_{\\ell m}$ (which reconstructs the modes of the $\\Delta\\tau$-field by using the induced non-Gaussian signature) but also presents the opportunity for improving signal-to-noise by combining $\\tau$ reconstruction with delensing."
    },
    "0812/0812.2490_arXiv.txt": {
        "abstract": "We follow the bright, highly energetic afterglow of {\\it Swift}-discovered GRB\\,080721 at $z=2.591$ out to 36 days or 3$\\times$10$^{6}$~s since the trigger in the optical and X-ray bands. We do not detect a break in the late-time light curve inferring a limit on the opening angle of $\\theta_j \\ge 7.3^{\\circ}$ and setting tight constraints on the total energy budget of the burst of $E_{\\gamma} \\ge 9.9\\times 10 ^{51}$ erg within the fireball model. To obey the fireball model closure relations the GRB jet must be expanding into a homogeneous surrounding medium and likely lies behind a significant column of dust. The energy constraint we derive can be used as observational input for models of the progenitors of long gamma-ray bursts: we discuss how such high collimation-corrected energies could be accommodated with certain parameters of the standard massive star core-collapse models. We can, however, most probably rule out a magnetar progenitor for this GRB which would require 100\\% efficiency to reach the observed total energy. ",
        "introduction": "The discovery \\citep{VanParadijs,Costa} of afterglows to long-duration gamma-ray bursts (GRBs) showed them to occur in star-forming galaxies at high redshifts. Many aspects of the observed GRB behaviour could be explained reasonably well by the relativistic fireball models, in which the prompt emission is largely produced by shocks internal to the outflow, and the long-lived afterglow by the outflow impacting and shocking the external ambient medium \\citep{MeszarosRees,Sari}. However, it was early appreciated that if they emitted isotropically then the radiative energies implied for some bursts would be very large (in excess of a Solar rest mass in the case of GRB\\,990123, for example; \\citealt{Kulkarni}). Since no known mechanism can produce high-energy photons with efficiency approaching 100\\%, the total explosive energy required would be even greater, implausibly large for a stellar core-collapse powered event.  The energetics argument led to the expectation that the GRB outflow must be confined to a jet, reducing the overall energy requirements.  In that case a sufficiently massive core collapsing to a black-hole might produce an outflow which could pierce a hydrogen-stripped envelope to still produce a relativistic jet \\citep{MacFadyen}. This picture received support from the observation of concurrent supernova events at GRB sites for some (low-redshift) low-luminosity bursts \\citep{Galama,Hjorth,Stanek,Pian}. The class of GRB-supernovae could be explained by the explosion of a Wolf-Rayet star, and it has been widely assumed that this progenitor must account for most if not all long-GRBs.  A strong prediction of all collimated models is that the observed light curve should exhibit an achromatic break, visible in both optical and X-ray light curves, when the relativistic outflow slows to the point that the Doppler beaming angle becomes wider than the opening angle of the jet \\citep{Rhoads}: the later the break time, the wider the jet. In the pre-{\\it Swift} era a number of bursts with good afterglow observations showed breaks in the optical, at times ranging from a few hours to a few days. Furthermore, it was contended that the implied jet opening angles anti-correlated with the isotropic-equivalent $\\gamma$-ray energies $E_{\\rm iso}$ \\citep{Frail} -- in other words, corrected for collimation nearly all bursts seemed to be accessing a rather standard energy reservoir (to within about an order of magnitude), potentially usable as standard candles at high redshift. While the total rotational energies available from a collapsing core may be of order 10$^{54}$ erg, the maximal extraction of energy is unlikely to exceed 10$^{51}$--10$^{52}$ erg \\citep{Paczynski,Popham}. With this in mind, the standard energy reservoir with energies of $\\sim$10$^{51}$ erg appeared very attractive.  Unfortunately the X-ray data available pre-{\\it Swift} only provided definitive evidence of (single) breaks for a few bursts. One of the expectations of {\\it Swift} was that many more bursts would be seen with clear breaks in the X-ray coincident with optical breaks. After more than four years of operations this is not what has been found. In fact, the situation is much more confusing, with X-ray light curves exhibiting complex behaviour including flares and plateaux \\citep{Nousek,OBrien,Willingale}. When late breaks (beyond a few hours) are seen they rarely coincide in time or degree with optical breaks, and are hard to reconcile with simple fireball model predictions \\citep[e.g.][]{Panaitescu06}, though there are notable exceptions. It is possible that from X-ray data alone a number of jet breaks, particularly those at late times or those which turn over rather smoothly, may go undetected through misinterpretation of their light curves. These so-called `hidden' breaks may be revealed on comparison of X-ray light curves with optical light curves \\citep{Curran}, but the emphasis remains on X-ray light curves for break detection owing to the large numbers of well-sampled X-ray light curves now available.  Most crucially, in many GRBs no temporal break is seen at all to late times, implying much more energy than expected. The most extreme example to-date is that of GRB\\,060729, which continued a smooth X-ray power-law decline for 125 d, implying an opening angle greater than 28 degrees \\citep{Grupe}. This source continues to be monitored in the X-rays with {\\it Chandra} and was still detected one year after the GRB event (Grupe et al. 2009). The light curve now shows a possible break at $\\sim$290 d: the nature of this break is not yet known. While the modest luminosity and low redshift of GRB\\,060729 ($E_{\\rm iso}\\sim 7 \\times 10^{51}$ erg, \\citealt{Grupe}) make it unsuitable to test collapsar models to their limits, such observations provide a crucial guide to further theoretical developments. If such high collimation-corrected energies continue to be inferred, then the situation looks bleak for simple core-collapse models as the explanation for the high-luminosity long-GRBs. In fact, it would suggest a different progenitor is required for high- and low-luminosity long-bursts. On the other hand, if late breaks are seen, although the jet model is not proven this could provide evidence for an upper envelope to the energy tapped by GRBs which may still be consistent with that available in principle from core-collapse, but would require a new understanding of how such high efficiencies of radiative emission could be achieved.  In this paper we follow a high redshift GRB with a high isotropic energy out to late times to search for a jet break and thereby measure the total energy budget of the prompt $\\gamma$-ray emission. In Section 2 we describe the $\\gamma$-ray, X-ray and optical observations and analyses. We derive the GRB fluence in Section 3. In Section 4 we characterise the afterglow spectra and light curves and compare the results to the fireball model in Section 5 discussing the possibility of an early jet break. In Section 6 we calculate limits on the jet opening angle and collimation-corrected energy of this GRB. In Section 7 we note the caveats associated with our results, discuss the energetics of this GRB in relation to current popular progenitor models and compare this GRB to other bursts of interest. In Section 8 we conclude. ",
        "conclusions": "We have observed the bright, highly energetic afterglow of GRB\\,080721 out to 36 days since the trigger in the optical and X-ray bands. We conclude that no jet break is present in the late-time light curve and we rule out a jet-break origin for the early light curve break, inferring a limit on the opening angle of $\\theta_j \\ge 7.22^{\\circ}$ and setting constraints on the total energy budget of the burst of $E_{\\gamma} \\ge 9.88\\times 10 ^{51}$ erg within the fireball model. To obey the fireball model closure relations the GRB jet must be expanding into a homogeneous medium and be extincted in the optical bands by approximately 0.6 magnitudes. The energy constraint we derive can be used as observational input for models of the progenitors of long gamma-ray bursts. We can likely rule out a magnetar progenitor for this GRB as this would require close to 100\\% efficiency to reach the observed total energy. Such high collimation-corrected energies could be accommodated with certain parameters of the standard massive star core-collapse models. One of the key observational parameters in distinguishing between various core-collapse models is the metallicity, diagnostic only if it lies outside of the typical range of $0.01<Z<0.1$. The metallicity we measure for GRB\\,080721 unfortunately does not allow a distinction between models to be made.  The occurrence of such highly energetic or narrowly beamed GRBs is rather rare, and for those which have the required measurements to allow a study such as this one on GRB\\,080721 we estimate there may be a handful per year among the 100--130 well localised GRBs currently triggering operational satellites. Measurement of the high energy spectral peak, securing the redshift and monitoring the light curves for as long as possible (preferably in multiple bands) are all crucial, and these data have shown that a full understanding of the blastwave physics can be difficult to achieve even with a good quality, well-sampled, broad-band dataset. To determine whether these sources are true outliers from the GRB population, perhaps requiring different or more extreme progenitors, will require dedicated efforts to build up small samples to compare with each other and with the growing database of more typical GRBs."
    },
    "0812/0812.4339_arXiv.txt": {
        "abstract": "Formation of primordial black holes (PBHs) requires a large root-mean-square amplitude of density fluctuations, which generate second-order tensor perturbations that can be compared with observational constraints. We show that pulsar timing data essentially rules out PBHs with $10^{2-4}\\msolar$ which were previously considered as a candidate of intermediate-mass black hoes and that PBHs with mass range $10^{20-26}$ g may be probed by future space-based laser interferometers. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.1346_arXiv.txt": {
        "abstract": "{GX~339--4 is a well-known microquasar. In this contribution we show the obtained results with the INTEGRAL and XMM-Newton observatories of the outburst undertaken on 2007. The observations cover spectral evolution from the hard, soft intermediate states to the high/soft state. Spectral hardening correlated with the appearance of an skewed Fe line is detected during one of the observations during the soft intermediate state. In all spectral states joint XMM/EPIC-pn, JEM-X, ISGRI and SPI data were fit with the hybrid thermal/non-thermal Comptonization model (EQPAIR). With this model a non-thermal component seems to be required by the data in all the observations. Our results imply evolution in the coronal properties, the most important one being the transition from a compact corona in the first observation to the disappearance of coronal material in the second and re-appearance in the third. We discuss the results obtained in the context of possible physical scenarios for the origin and geometry of the corona and its relation to black hole states.}  \\FullConference{7th INTEGRAL Workshop\\\\ September 8-11 2008\\\\ Copenhagen, Denmark}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0812/0812.3669_arXiv.txt": {
        "abstract": "Proposed  \\HI  structure observatories for studying the epoch of reionization ($z \\approx$ 6--15)  and dark energy ($z \\approx$ 0--6) envision compact arrays with tens of thousands of antenna elements. Fully correlating this many elements is computationally expensive using XF or FX correlators, and has led some groups to reconsider direct imaging/FFT correlators. In this paper we develop a variation of the direct imaging correlator we call the MOFF correlator. The MOFF correlator shares the computational advantages of a direct imaging correlator, while avoiding a number of its shortcomings. In particular the MOFF correlator makes no constraints on the antenna arrangement or type, provides a fully calibrated output image including widefield polarimetry and non-coplanar baseline effects, and can be orders-of-magnitude more efficient than XF or FX correlators for compact radio cosmology arrays. ",
        "introduction": "A new generation of radio observatories are being developed to measure the history and evolution of our universe using deep radio surveys of large scale structure. The MWA, LOFAR, and PAPER are being built to observe the Epoch of Reionization (EoR) through large scale \\HI emission at redshifts of 6--11, and much larger second generation EoR observatories are being planned. Additionally, at redshifts of 0--6 \\HI structure experiments can trace large scale structure over much of the observable universe \\citep{Wyithe:2007p4,Mao:2008p196,Chang:2008p439,Loeb:2008p3284}, enabling a number of precision cosmography and dark energy measurements. In particular the dark energy baryon acoustic oscillation (BAO) signal has inspired the development of a number of new instrument concepts, including the FFT Telescope \\citep{Mao:2008p196}, the Canadian Hydrogen Intensity Mapping Experiment, the Cylinder Radio Telescope, the high frequency MWA5000 \\citep{Wyithe:2008p2214}, the Lunar Array for Radio Cosmology, and the Dark Ages Radio Interferometer.  The \\HI cosmology and second generation Epoch of Reionization instruments are similar in their instrumental characteristics, featuring compact arrangements of many thousands of antennas. The many short baselines and wide field of view maximizes the sensitivity of power spectrum measurements over all scales \\citep{Morales:2005p796,MoralesArecibo08}, however, it also puts enormous processing demands on the correlator system. A radio correlator measures the cross-power correlation between all antenna pairs in many narrow frequency channels, and for a modern FX correlator the computation scales as the square of the number of antennas. The correlator under construction for the MWA requires 15.5 trillion complex multiplies and adds per second (Tcmacs) for 512 antennas and 31 MHz of bandwidth. For many of the proposed radio cosmology instruments this quickly scales into the peta-flop regime, making the correlator a dominant cost for these arrays. This has driven some concepts such as the FFT telescope to use  direct imaging correlators, despite their significant shortcomings, and constrain the antenna layout to a regular array and filled aperture.  In this paper we introduce the MOFF (Modular Optimal Frequency Fourier) correlator concept. The MOFF correlator incorporates the gridding and calibration usually associated with post-correlation image processing into the correlation process. This work expands on direct imaging correlator concepts by \\citet{Daishido:2000p3323} and \\citet{RogersDirectImaging}, with particular emphases on data calibration and modular design.  For the compact antenna layouts of proposed radio cosmology telescopes, the MOFF correlator is very efficient, makes no constraints on the antenna placement, and produces a provably optimal data product. ",
        "conclusions": "\\label{MOFFreview}  The MOFF correlator concept has a number of advantages over traditional correlator designs for envisioned radio cosmology observatories. The advantages we have detailed include: \\begin{itemize} \\item Calibrated output image, equivalent in quality to FX visibilities (\\S\\ref{SHsec}). \\item More efficient than FX or XF correlators for compact arrays with very large numbers of antennas, sometimes by orders-of-magnitude (\\S\\ref{computation}). \\item No constraints on antenna placement or type (\\S\\ref{MOFF}). \\item Modular installation and reduced cabling costs (\\S\\ref{modular}). \\end{itemize} A couple of features we have not discussed but which follow directly from the MOFF design are: \\begin{itemize} \\item Electric field image can be tapped for pulsar studies and transient detection. Each pixel is effectively a calibrated tied-array beam. \\item Hybrid modes are possible, where a MOFF correlator is used for all of the antennas in a compact station and the electric field image from selected pixels is sent to an FX correlator for high angular resolution interferometry between the stations. Similar hybrid modes may be useful for phased-array feeds, where a MOFF correlator would turn the image back into calibrated spatial Fourier modes for widefield FX correlation [citation of eigenvalue correlation]. \\end{itemize}  As with any design there are some disadvantages as well. The primary disadvantage of the MOFF correlator is that the calibration must be applied as part of the correlation process---one cannot re-calibrate the data during post-processing. For very large cosmology arrays, the overwhelming number of antenna pairs may make saving raw visibility data impractical anyway. The MWA does not save the visibility information for this reason and instead performs realtime calibration and imaging, and future cosmology arrays will probably do the same. However, this may limit the usefulness of the MOFF concept in certain applications.  A second disadvantage is that the holographic antenna gain patterns needed for calibration (\\S\\ref{SHsec}) cannot be obtained directly from the image produced by the MOFF correlator.  Effectively the antenna beam patterns have been co-added during the correlation process (most easily seen in Equation \\ref{m1}), and one needs the individual visibilities to measure the holographic antenna gain response. As part of real world implementations of the  MOFF correlator, digitized signals from selected antennas will need to be pulled out of the MOFF correlator (Figure \\ref{moffFig} row b) so holographic gain patterns can be formed offline."
    },
    "0812/0812.2772_arXiv.txt": {
        "abstract": "{} {We consider the polarization arising from scattering in an envelope illuminated by a central anisotropic source.  This work extends the theory introduced in a previous paper (Al-Malki \\etal\\ 1999) in which scattering polarization from a spherically symmetric envelope illuminated by an anisotropic point source was considered. Here we generalize to account for the more realistic expectation of a non-spherical envelope shape.} {Spherical harmonics are used to describe both the light source anisotropy and the envelope density distribution functions of the scattering particles.  This framework demonstrates how the net resultant polarization arises from a superposition of three basic ``shape'' functions:  the distribution of source illumination, the distribution of envelope scatterers, and the phase function for dipole scattering.} {Specific expressions for the Stokes parameters and scattered flux are derived for the case of an ellipsoidal light source inside an ellipsoidal envelope, with principal axes that are generally not aligned. Two illustrative examples are considered:  (a) axisymmetric mass loss from a rapidly rotating star, such as may apply to some Luminous Blue Variables, and (b) a Roche-lobe filling star in a binary system with a circumstellar envelope.} {As a general conclusion, the combination of source anisotropy with distorted scattering envelopes leads to more complex polarimetric behavior such that the source characteristics should be carefully considered when interpreting polarimetric data. } ",
        "introduction": "For many stars linear polarization is produced mainly from scattering of starlight by circumstellar matter (Kruszewski \\etal\\ 1968; Serkowski 1970; Dyck \\etal\\ 1971; Shawl 1975).  This polarization can be used as a diagnostic of the geometry of the circumstellar envelope and of the light source(s) (e.g., Shakhovskoi 1965; Serkowski 1970; Brown \\& McLean 1977; Brown \\etal\\ 1978; Rudy \\& Kemp 1978; Simmons 1982, 1983; Friend \\& Cassinelli 1986; Clarke \\& McGale 1986, 1987).  In many models for circumstellar scattering polarization, the illuminating sources are treated as isotropic point sources (e.g., Brown \\& McLean 1977; Brown \\etal\\ 1978; Rudy \\& Kemp 1978; Shawl 1975; Simmons 1982). The effect of the finite size of the star as a light source has been studied (Cassinelli, Nordsieck, \\& Murison 1987), including the effects of limb darkening (Brown \\etal\\ 1989) and stellar occultation (Brown \\& Fox 1989; Fox \\& Brown 1991; Fox 1991).  Variable polarization is also revealing about a system, as evidenced in a recent study by Elias, Koch, \\& Pfeiffer (2008).  With recent emphasis on clumped wind flows of hot stars (e.g., Hamann, Feldmeier, \\& Oskinova 2008), models have been developed to interpret variable polarization from hot star winds (Richardson, Brown, \\& Simmons 1996; Brown, Ignace, \\& Cassinelli 2000; Li \\etal\\ 2000; Davies \\etal\\ 2007).  Relatively little has been done to explore the anisotropy of the source illumination and its consequences for interpreting polarimetric observations.  There are some exceptions, such as a study of gravity-darkening effects for the polarization from Be~stars by Bjorkman \\& Bjorkman (1994), the influence of star spots for polarizations from pre-main sequence stars by Vink \\etal\\ (2005), and more recently calculations of source anisotropy for interpreting polarimetric behavior observed in the post-red supergiant star HD~179821 by Patel \\etal\\ (2008).  In a previous paper (Al-Malki \\etal\\ 1999; hereafter Paper I), we modeled the polarization arising from Thomson and Rayleigh scattering explicitly for an arbitrary anisotropic (point) light source.  As a proof of concept, the circumstellar envelope was taken as spherically symmetric to explore the polarization signals that arise solely from the properties of the source.  An upper limit to the polarization of about 20\\% was derived in the limiting case of a disk-like star when viewed edge-on.  This paper generalizes the results of Paper~I by allowing both the point light source and the particle density distribution to be functions of arbitrary angular form.  As in Paper~I, we restrict ourselves to optically thin envelopes and single Thomson or Rayleigh scattering mechanisms. Although the effects of a finite-sized illuminator will be important for obtaining quantitatively a more accurate polarization amplitude (e.g., Ignace, Henson, \\& Carson 2008), the point source approximation is certainly adequate for exploring polarization trends of generalized source and envelope geometries.  In Sect.~2, we develop the theory required to evaluate the polarization, based on Paper I and the discussion of Simmons (1982). Applications to specific cases of astrophysical interest are explored in Sect.~3, with emphasis given to a treatment of the star and the circumstellar envelope as ellipsoidal. A discussion of these results and the need for ongoing modeling is given in Sect.~4.  \\begin{figure*}[t] \\centering{\\epsfig{file=fig1.eps,angle= -90,width=15cm}} \\caption{Illustration of the ellipsoidal envelope used in our models. The scattering envelope is shown in cross-section as shaded, with an equatorial width ranging from an inner radius of $R_2$ to an outer one of $R_1$.  Along the pole, the inner and outer boundaries are scaled by the factor $A_{\\rm r}$.  Although the theory allows for fully triaxial ellipsoids, our specific applications are for axisymmetric ellipsoids (ranging from oblate to prolate). \\label{fig1}} \\end{figure*} ",
        "conclusions": "In this work we derive the polarization arising from an anisotropic point light source within an arbitrary shaped envelope. We specifically considered Thomson and Rayleigh scattering mechanisms. The mathematical analysis made use of the properties of spherical harmonics, which can readily be generalized to more complicated cases than those discussed here (e.g., spots on a star). By using the first few spherical harmonics, an approximation was derived to represent distortions of stars and envelopes from spherical by varying amounts. The anisotropy properties of the star and envelope can be reduced to representations as products of multipole contributions.  We considered applications to rotating stars and binary systems.  For rotating stars we sought to explore the polarizations resulting from an approximate representation of enhanced bipolar wind flow illuminated by a rapidly rotating and gravity darkened star. This scenario has relevance to some LBV winds.  We also considered a highly distorted star such as can occur in a binary when one of the components is Roche lobe filling and dominates the luminous output of the system.  In this case the star's illuminating characteristics were modeled as a prolate ellipsoid rotating about an axis perpendicular to its symmetry axis. Scenarios involving an oblate circumstellar envelope that was aligned or inclined with respect to the star's rotation axis were considered.  The theory presented here is fairly general; however, there remain a number of improvements that should be pursued.  Foremost is inclusion of the star's finite size, because of both occultation and the finite depolarization effects.  An initial attempt at this has been made by Ignace \\etal\\ (2008), who consider scattering polarization from a structured giant star chromosphere. In addition, our theory assumes that the radial and angular descriptions of the envelope are separable, which may not always be the case.  Including these modifications is far from trivial but may be necessary in order to accurately represent complex systems.  Certainly, the growth in the size and capability of telescopes and modern instrumentation, along with the increased importance of spectropolarimetry in ascertaining source geometries (e.g., supernova ejecta -- Wang \\etal\\ 2001; Leonard \\etal\\ 2006) motivates an effort to generalize further our approach.  \\appendix"
    },
    "0812/0812.0631_arXiv.txt": {
        "abstract": "Recent observational studies of intermediate-age star clusters (SCs) in the Large Magellanic Cloud (LMC) have reported that a significant number of these objects show double main-sequence turn-offs (DMSTOs) in their color-magnitude diagrams (CMDs). One plausible explanation for the origin of these DMSTOs is that the SCs are composed of two different stellar populations with age differences of $\\sim 300$ Myr. Based on analytical methods and numerical simulations, we explore a new scenario in which SCs interact and merge with star-forming giant molecular clouds (GMCs) to form new composite SCs with two distinct component populations. In this new scenario, the possible age differences between the two different stellar populations responsible for the DMSTOs are due largely to secondary star formation within GMCs interacting and merging with already-existing SCs in the LMC disk. The total gas masses being converted into new stars (i.e., the second generation of stars) during GMC-SC interaction and merging can be comparable to or larger than the masses of the original SCs (i.e, the first generation of stars) in this scenario. Our simulations show that the spatial distributions of new stars in composite SCs formed from GMC-SC merging are more compact than those of stars initially in the SCs. We discuss both advantages and disadvantages of the new scenario in explaining fundamental properties of SCs with DMSTOs in the LMC and in the Small Magellanic Cloud (SMC). We also discuss the merits of various alternative scenarios for the origin of the DMSTOs. ",
        "introduction": "Recent photometric studies of intermediate-age SCs in the LMC have discovered objects exhibiting extremely unusual main-sequence turn-offs (MSTOs) in their CMDs. The first indications that such clusters might exist date back several years to observations obtained with terrestrial facilities -- Bertelli et al. (2003) demonstrated that the LMC cluster NGC 2173 apparently possesses an unusually large spread in colour about its MSTO, while Baume et al. (2007) obtained a similar result for NGC 2154.   More recent studies based on deep precision photometry from the {\\it Hubble Space Telescope} ({\\it HST}) Advanced Camera for Surveys (ACS) have revealed the truly peculiar nature of many intermediate-age LMC clusters. Most striking are the relatively massive clusters NGC 1846, 1806, and 1783 which possess double MSTOs (hereafter referred to as DMSTOs for convenience) on their CMDs (Mackey \\& Broby Nielsen 2007, MBN07; Mackey et al. 2008; Milone et al. 2008; Goudfrooij et al. 2008). Despite this, the remaining CMD sequences for these clusters (i.e., their red-giant branches, subgiant branches, main sequences, and red clumps) are all extremely narrow and well-defined. This suggests that none of the clusters possesses a significant line-of-sight depth or internal dispersion in $[$Fe$/$H$]$, or suffers from significant differential extinction. The simplest explanation is that the observed DMSTOs in each of these clusters represent two distinct stellar populations with differences in age of $\\sim 200 - 300$ Myr. It should be noted, however, that apart from the study of Mucciarelli et al. (2008) who found no significant star-to-star dispersion in $[\\alpha/$Fe$]$ for $6$ stars in NGC 1783, no constraints on the possibility of internal variations in chemical abundances other than $[$Fe$/$H$]$ presently exist for these SCs.   Milone et al. (2008) studied a large sample of $16$ intermediate-age LMC clusters, and found that $11$ of these possess CMDs exhibiting MSTOs which are not consistent with being simple, single stellar populations. Of these $11$, four clearly show DMSTOs (the three described above plus NGC 1751) while the remainder possess more sparsely populated CMDs that make the precise morphologies of their peculiar MSTOs difficult to ascertain. All are consistent with being DMSTOs, or alternatively with possibly being more smoothly distributed intrinsic broadenings. Nonetheless, all $11$ clusters again exhibit very narrow sequences across the remainder of their CMDs, suggesting that their peculiar MSTOs are due to internal age dispersions of $\\sim 200-300$ Myr. Milone et al. (2008) measured  the ratio between  the  upper and the  lower MSTO population in   NGC1846, NGC1806 and   in  the less-populated NGC1751 and found that the inferred young populations may comprise up to $\\sim 70$ per cent of the stars in the central regions of these SCs.   In addition to the above objects, several younger LMC clusters have been identified as possibly possessing irregular CMDs (e.g., Santiago et al. 2002, Gouliermis et al. 2006). It is further worth noting that one intermediate-age SMC cluster, NGC 419, also likely possesses a DMSTO (Glatt et al. 2008).   The results described above suggest that, contrary to expectation, DMSTOs may be a common feature among intermediate-age SCs of the Magellanic Clouds (MCs). If the origin of such DMSTOs is indeed closely associated with an age difference between a first and second generation of stars in these SCs, the above observations raise the following three key questions: (i) why the second generation of stars in a SC with a DMSTO can be formed about 300 Myr after the formation of the first generation of stars, (ii) why the total mass of the second generation can be comparable or even larger than that of the first, and (iii) how the two different populations can now be in the same SC. Although the origin of the observed multiple stellar populations in the Galactic globular clusters (GCs) has been observationally and theoretically discussed in terms of GC formation scenarios (e.g., Piotto 2008), the above fundamental questions so far have not been discussed by theoretical models of SC formation in the LMC.    \\begin{table*} \\centering \\begin{minipage}{175mm} \\caption{Model parameters for numerical simulations of GMC-SC mergers} \\begin{tabular}{ccccccc} {Model no \\footnote{For models M3 - M32, the ranges of model parameters investigated in the present study are shown in parentheses: the left and the right numbers are the minimum and maximum values.}} & {$M_{\\rm s}$ \\footnote{The total mass of stars in a SC  in units of $10^{4}$  ${\\rm M_{\\odot}}$. The total gas mass in a GMC is $5\\times 10^4 {\\rm M}_{\\odot}$ for all models. }} & {$R_{\\rm s}$ \\footnote{The size of a SC in units of pc. The size of a GMC is set to be 11.9 pc for all models in the present study.}} & {$y_{\\rm g}$ \\footnote{Initial $y$-position of a GMC with respect to the center of a SC. This corresponds to the impact parameter of the GMC-SC collision.}} & {$|u_{\\rm g}|$ \\footnote{The absolute magnitude of the initial velocity of a GMC along the $x$-axis in units of km s$^{-1}$, with respect to the center of a SC. For all models, $u_{\\rm g} \\le 0$.}} & {Star formation \\footnote{``YES'' (``NO'') means that the star formation model is (is not) included in the simulation. }} & Comments \\\\ M1 &  5.0  &  10  & 6.0  & 2.2 &  NO & the standard model \\\\ M2 &  5.0  &  10  & 6.0  & 2.2 &  YES &  the star formation model  \\\\ M3-32 &  (1.0, 20.0)  &  (5, 25)  & (0, 23.8)   & (0, 12.8)  & NO \\& YES  &  \\\\ \\end{tabular} \\end{minipage} \\end{table*}  The purpose of this paper is to explore a new scenario in terms of the above three fundamental questions on the origin of the observed DMSTOs. In the new scenario, an already formed SC (i.e., the first generation of stars) can merge or interact with a GMC so that a second generation of stars is formed in the GMC. The second generation can then merge and mix dynamically with the SC to form a new SC with two stellar populations of different ages. Thus, the origin of the DMSTOs in SCs of the LMC results from merging between SCs and GMCs in the new ``GMC-SC-merger''  scenario. The typical age difference ($t_{\\rm gap}$) of the two populations in SCs required for explaining the DMSTOs corresponds to the typical time scale of merging between SCs and GMCs of the LMC in the new scenario. We mainly discuss (i) whether star formation is likely to be enhanced during GMC-SC merging, (ii) the mass fraction of the second generation of stars relative to the first in a SC and its dependence on parameters of GMC-SC merging, and (iii) time scales of GMC-SC merging in the LMC.   The plan of the paper is as follows: In the next section, we describe our numerical models for GMC-SC merging in the LMC. In \\S 3, we present numerical results mainly concerning the final distributions of gas and new stars in the remnants of GMC-SC mergers. In \\S 4, we discuss the origin of the DMSTOs of SCs in a more general way. In this section, we also discuss the advantages and disadvantages of two alternative scenarios in explaining the physical properties of SCs with DMSTOs. We summarize our conclusions in \\S 5. ",
        "conclusions": "We have numerically investigated how GMCs evolve during GMC-SC merging in order to better understand the origin of the DMSTOs observed in intermediate-age SCs in the MCs. We have mainly investigated mass fractions of gas within 10 pc of the remnants of GMC-SC mergers for variously different model parameters. We summarise our principle results as follows:   (1) Gas initially in GMCs can be accumulated within SCs during GMC-SC merging so that high-density compact gaseous regions that are significantly flattened are formed in the central regions of the SCs. The mass fractions of gas ($f_{\\rm g}$) within the central 10 pc of the merger remnants (i.e., the newly formed SCs) can be as large as 0.5 for models with $s_{\\rm g} \\sim 1$ (i.e., major mergers).   (2) There is an optimum $M_{\\rm s}$ (or $s_{\\rm g}$) for a given set of model parameters, for which $f_{\\rm g}$ is a maximum. For example, the model with $M_{\\rm s}=M_{\\rm g}$ (i.e., $s_{\\rm g}=1$) shows the maximum $f_{\\rm g}$ in models with $M_{\\rm g}=5 \\times 10^4 {\\rm M}_{\\odot}$ and $y_{\\rm g}=0.5 R_{\\rm g}$ (=6 pc). The remnants of major GMC-SC mergers are likely to show larger $f_{\\rm g}$ in the present study.   (3) As long as GMC-SC merging occurs, $f_{\\rm g}$ does not depend strongly on $y_{\\rm g}$ and $|u_{\\rm g}|$. The models with larger $y_{\\rm g}$ and $|u_{\\rm g}|$, for which GMCs can interact with SCs without merging, show only very small $f_{\\rm g}$ (an order of $10^{-2}$). These results mean that merging is essential for the formation of new SCs with large $f_{\\rm g}$.   (4) If star formation is included in GMC-SC merging, compact, flattened star clusters composed of new stars can be formed in the central regions of the merger remnants. Since the simulated SCs-within-SCs appearances (i.e., doubly nested SCs) are not observed in the MCs, some later dynamical evolution processes need to transform the doubly nested SCs into normal-looking (well-mixed) objects with standard (King-type) density profiles.  (5) The time scale of GMC-SC merging ($t_{\\rm m}$) in the LMC can be similar to the typical age difference of $\\sim 3 \\times 10^8$ yr between the component stellar populations implied by observations of DMSTOs in SCs in the LMC. However $t_{\\rm m}$ depends strongly on the number of GMCs, the sizes of GMCs, and the velocity dispersions of GMCs and stars in the LMC a few Gyr ago, all of which remain observationally unclear.   Based on these results, we have pointed out that the observed possibly large fractions of the second generations of stars in SCs with DMSTOs in the LMC can be due to the past GMC-SC merging. We have also suggested that time lags between SC formation and the subsequent GMC-SC merging can be responsible for the possible age differences between the first and the second generations of stars in the SCs with DMSTOs in the LMC. The adopted numerical code does not allow us to investigate whether the simulated doubly-nested SCs can evolve into normal SCs with canonical radial density profiles due to internal dynamical processes. We plan to investigate this question in our future studies using the appropriate numerical codes (e.g., {\\tt NBODY4})."
    },
    "0812/0812.3542_arXiv.txt": {
        "abstract": "{There are only a few tens of open clusters for which ellipticities have been determined in the past.} {In this paper we derive observed and modelled shape parameters (apparent ellipticity and orientation of the ellipse) of 650 Galactic open clusters identified in the \\ascc catalogue. } {We provide the observed shape parameters of Galactic open clusters, computed with the help of a multi-component analysis. For the vast majority of clusters these parameters are determined for the first time. High resolution (``star by star'') N-body simulations are carried out with the specially developed $\\phi$GRAPE code providing models of clusters of different initial masses, Galactocentric distances and rotation velocities.} {The comparison of models and observations of about 150 clusters reveals ellipticities of observed clusters which are too low (0.2 vs. 0.3), and offers the basis to find the main reason for this discrepancy. The models predict that  after $\\approx 50$ Myr clusters reach an  oblate shape with an axes ratio of $1.65:1.35:1$, and with the major axis tilted by an angle of  $q_{XY} \\approx 30^\\circ$ with respect to the Galactocentric radius due to differential rotation of the Galaxy. } {Unbiased estimates of cluster shape parameters requires reliable membership determination in large cluster areas up to 2-3 tidal radii where the density of cluster stars is considerably lower than the background. Although dynamically bound stars outside the tidal radius contribute insignificantly to the cluster mass, their distribution is essential for a correct determination of cluster shape parameters. In contrast, a restricted mass range of cluster stars does not play such a dramatic role, though deep surveys allow to identify more cluster members and, therefore, to increase the accuracy of the observed shape parameters.} ",
        "introduction": "Our current project aims at studying the properties of the local population of Galactic open clusters. The sample contains 650 open clusters and cluster-like associations identified in the all-sky compiled catalogue of 2.5 million stars \\ascc \\citep{kha01}. For each cluster, a combined spatio-kinematic-photometric membership analysis was performed \\citep[]{starcat} and a homogeneous set of different cluster parameters was derived \\citep[][]{clucat,newclu}. In two recent papers \\citep[][]{clumart,clumart1} tidal radii and masses of open cluster were determined and their relation to cluster ellipticity was briefly discussed. In the present study we discuss in detail all the issues related to the shape of local clusters.  The shape of a star cluster, as well as its size and mass are the most important dynamical parameters. They are predesignated already at the stage of the cluster formation, when the cluster keeps the memory of the size and shape of the parent cloud. In addition to internal processes (self-gravitation and rotation) a considerable role in cluster shaping is played by external forces, e.g. by the Galactic tidal field, which is acting in two ways \\citep[cf.][]{wiel74,wiel85}: a) stretching the cluster into an ellipsoid directed towards the Galactic centre, and b) producing cluster tails outpouring from the ellipsoid endpoints (cluster Lagrangian points). The tails are composed of stars lost by the cluster, which lead and/or trail the cluster along its orbit due to differential rotation of the Galactic disk \\citep[see e.g.][]{chura06a}. The encounters of star clusters with giant molecular clouds randomize the regular effect of the Galactic field \\citep[][]{giea06}.  Additionally, molecular clouds produce a screening effect, when they  partly overlap the clusters.  Most of the above reasons lead to the violation of cluster's sphericity with different effects on the cluster core and corona. % The analysis of such violations and their comparison with predictions of N-body models can shed light on the dynamical history of open clusters. We assume in the following that open clusters can be represented by triaxial ellipsoids.  The violations of the sphericity of globular clusters are evident and their ellipticities were determined for the first time already at the beginning of the last century. For example, \\citet{sheso27} published the estimates of ellipticity of 75 globular clusters. Currently, shape parameters are determined for one hundred Galactic globular clusters \\citep[]{white87}, i.e. according to Harris' on-line catalogue\\footnote{ http://www.physics.mcmaster.ca/resources/globular.html} for about two thirds of the known objects in the Galaxy.  The current status of the shape measurements for open clusters is much poorer. Until now, indications of a flattening were obtained for a few tens of open clusters only. \\citet{ramer98a, ramer98b} and \\citet{adea01,adea02} found evidence of a flattening in the Pleiades and in Praesepe. Their results support the predictions of Wielen's model that explains the ellipsoidal form of clusters by tidal coupling with the Galaxy. In contrary, the flattening found in the Hyades \\citep{oo79} deviates from theoretical expectations. Recently, \\citet{chen04} published morphology parameters, including cluster ellipticities, for 31 Galactic open clusters, located preferentially in the anticentre direction of the Galaxy.  With this paper we fill this gap in the parameter list of open clusters and determine the shape parameters for all 650 open clusters in our sample \\citep{clucat,newclu}. For a comparison with these observed shape parameters we carried  out a set of high resolution (``star by star'') N-body dynamical calculations of cluster models with different initial angular momenta located at different Galactocentric distances in the Milky Way.  The paper has the following structure. In Sec.~\\ref{sec_bas-eq} we provide the basic equations used in the present work; in Sec.~\\ref{sec_num-mod} we describe the N-body models; in Sec.~\\ref{obs_mod-res} we compare the observed shape parameters with model calculations and with literature data; in Sec.~\\ref{sec_concl} we summarize the results. ",
        "conclusions": "\\label{sec_concl}  Based on a multicomponent analysis of coordinates of the most probable cluster members we have determined shape parameters of 650 Galactic open clusters. Ellipticities and orientation angles complete the list of morphological and dynamical parameters (core sizes and apparent cluster radii, indicators of mass segregation, tidal radii) which were determined and analysed in a series of previous papers \\citep[][]{clusim,clumart,clumart1}.  We have carried out high resolution N-body simulations with the specially developed $\\phi$GRAPE code on the parallel GRAPE systems developed at ARI Heidelberg and MAO Kiev. The set of 27 cluster models (3 initial masses $\\times$ 3 Galactocentric distances $\\times$ 3 rotation parameters) with a maximum number of particles of 40404 was evolved during 1 Gyr. The cluster particles are initially distributed according to a Salpeter IMF in the mass range $0.08...8.0~m_\\odot$.  For each particle, the decision on its cluster membership is made by comparing its potential and kinetic energies.  The calculations of the apparent shape parameters of the modelled and observed clusters were carried out with the same technique. The selection of suitable models was based on initial cluster masses that were estimated from the earlier studies of our cluster sample \\citep[][]{clumart,clumart1}. This enabled us to realize a self-consistent approach for comparison of model and observed shape parameters using the data on 152 of the most populated Galactic clusters from our sample.  N-body calculations show that all the models loose more than about 50\\% of their initial mass during the first Gyr of the evolution due to two body encounters only. A model cluster older than $\\approx 50$~Myr keeps an oblate shape, with the major axis $a$ tilted with respect to the Galactocentric radius at an angle of $\\approx$30-40 degrees. The model counterparts corresponding to 152 observed clusters have an ellipticity distribution with a peak at $e \\approx 0.3..0.4$, and, on average, the ellipticity becomes even larger for rotating model clusters.  However, the observed clusters show a significantly lower ellipticity, typically of about $e=0.2$. A comparison of observed cluster shapes with the detailed models allowed us to realize that lower ellipticities are only apparent. We explain this disagreement  by a bias in observations due to underestimated cluster sizes $r_2$ based on the \\ascc data and used in the determination of cluster ellipticities. According to the models, the ellipticity of the central region of a cluster is small, it increases if outer layers are taken into account, and approaches asymptotically the largest value at 2-3 tidal radii $r_t$. Particles outside the tidal radius contribute negligibly to the cluster mass, however, their spatial distribution is decisive for the determination of the cluster ellipticity. Due to a relatively low density of cluster members above the background  in outer cluster regions, real clusters are usually observable only within distances smaller than $r_t$ from the cluster centres.  Though the apparent cluster sizes $r_2$ based on the \\ascc are twice as large as  previously published in the literature, the ratio $r_2/r_t$ was found to be, on average, only 0.5 \\citep{clumart1}.  In contrast, the deepness of the survey does not play such a dramatic role for the ellipticity determination. Although  deeper surveys offer a possibility to identify  more numerous cluster members of lower masses and, consequently, to increase the random accuracy of the ellipticity determination, this alone is not sufficient to avoid  the ``restricted area bias''.  We adopt the assumption of \\citet{wiel74,wiel85} on the preferential shape and orientation of cluster ellipsoids in the Galaxy and derive simple formulae giving a relation between the apparent ellipticity and spatial location of clusters of our sample. Due to relatively high systematic and random errors in observed ellipticities, however, we could not use the present observations for estimating the parameters of the reference ellipsoid. According to N-body calculations, the cluster ellipsoid has an axes ratio of $a:b:c=1.65:1.35:1$, and it is tilted by an angle $q_{XY} \\approx 30^\\circ$ with respect to the Galactocentric radius.  In order to test the models with the real observations, the ellipticity determination for clusters must be based on sufficiently large areas (up to two to three tidal radii) around the clusters. From this point of view, equidensity methods are hardly useful to determine real cluster sizes, especially, due to a low density contrast of cluster members above the background at large distances from the cluster centres. The membership determination with kinematic and photometric criteria fits the problem better but requires surveys with higher accuracy of proper motions than it is in the \\ascc."
    },
    "0812/0812.2478_arXiv.txt": {
        "abstract": "We look for a non-Gaussian signal in the WMAP 5-year temperature anisotropy maps by performing a needlet-based data analysis. We use the foreground-reduced maps obtained by the WMAP team through the optimal combination of the W, V and Q channels, and perform realistic non-Gaussian simulations in order to constrain the non-linear coupling parameter $\\fnl$. We apply a third-order estimator of the needlet coefficients skewness and compute the $\\chi^2$ statistics of its distribution. We obtain $-80<\\fnl<120$ at 95\\% confidence level, which is consistent with a Gaussian distribution and comparable to previous constraints on the non-linear coupling. We then develop an estimator of $\\fnl$ based on the same simulations and we find consistent constraints on primordial non-Gaussianity. We finally compute the three point correlation function in needlet space: the constraints on $\\fnl$ improve to $-50<\\fnl<110$ at 95\\% confidence level. ",
        "introduction": "With the increasing amount of high-quality observations performed in the last decade (\\cite{Hinshaw2008WMAP5,Acbar2008,CBI2007,Quad2008Wu,Quad2008Pryke,Quad2008Hinderks,Masi2006,Maxipol2007Johnson}), statistical tests of the CMB temperature anisotropy pattern are getting more and more accurate.  This has made possible to test one of the basic tenets of the standard cosmological scenario, i.e.\\ that the primordial density perturbations follow a Gaussian distribution. This is a definite prediction of the simplest inflationary models \\citep{Guth1981,Sato1981,Linde1982,AlbrechtSteinhardt1982}: the detection of primordial deviations from Gaussianity would be a smoking gun for more complicated implementations of the inflationary mechanism, such as those of multi-fields \\citep{LythWands2002,LindeMukhanov2006,AlabidiLyth2006}, ekpyrotic \\citep{Mizuno2008,Khoury2002PhDT} or cyclic scenarios \\citep{SteinhardtTurok2002,LehnersSteinhardt2008}.  When dealing with the search for a non-Gaussian statistics in real data, two major issues have to be addressed. One has to do with the statistical tools used to analyse the data and detect deviations from Gaussianity: not only can these deviations be very subtle and elusive, but they could be generated by processes that are not directly related to the primordial density perturbations --- such as unremoved astrophysical foregrounds \\citep{Cooray2008WL,Serra2008} or instrumental systematics. The other issue is theoretical, and relates to the way the non-Gaussianity is parameterised: while there is only one way to realize a Gaussian distribution, non-Gaussian statistics can be produced in countless ways. One then has to assume a non-Gaussian parameterisation which relates in some sensible way to an underlying early universe scenario.  The latter issue is usually addressed by introducing a parameter $\\fnl$, which quantifies the amplitude of non-Gaussianity as a quadratic deviation with respect to the primordial Gaussian gravitational potential $\\Phi_{\\rm L}$, i.e.: \\begin{equation} \\Phi(x)=\\Phi_{\\rm L}(x)+\\fnl \\left[ \\Phi^2_{\\rm L}(x) -\\langle \\Phi_{\\rm L}^2(x)\\rangle \\right] \\end{equation} The major advantage of this parameterisation is that, regardless of the specific underlying early universe model, it can represent the second-order approximation of any non linear deviation from Gaussianity. For an excellent review on this topic see \\cite{Bartolo2004NGreview}.  From the point of view of data analysis, a number of techniques have been proposed in the past few years to quantify the level of deviation from Gaussian statistics in the data. The most used one in harmonic space is the bispectrum \\citep{Luo1994,Heavens1998,SpergelGoldberg1999,KomatsuSpergel2001,Cabella2006IntBis}. The bispectrum is defined as the three-point correlation function, and an estimate of $\\fnl$ through the bispectrum requires the sum over all the triangle configurations. Since this is extremely time-consuming, regardless whether the computation is performed in harmonic space or pixel space, \\cite{KomatsuSpergelWandelt2005} have proposed a fast cubic estimator based on the Wiener filter matching, which reduces considerably the computational challenge. This estimator has been further developed by \\cite{CreminelliEtAl2005} introducing a linear correction which allows the correct treatment of the anisotropic noise, and finally optimised \\citep{YadavKomatsuWandelt2007} and extended to polarisation measurements \\citep{YadavEtAl2007}. Recently, \\cite{YadavWandelt2007} applied the cubic estimator to the WMAP 3-year data, finding a detection of a primordial non-Gaussian signal at more than 2.5 sigma. An indication of a primordial non-Gaussian signal has been also found by the WMAP collaboration in the analysis of the 5-year dataset, although with a lower confidence level  \\citep{WMAP5Komatsu2008}. An interesting discussion on optimal and sub-optimal estimators can be found in \\cite{SmithZaldarriaga2006}. Se for further details \\cite{YuLu2008} and \\cite{Babich2005}.  Concerning the methods in pixel space, \\cite{DeTroia2007B2K,Acbar2008,Hikage2006MinkFunc,Curto2008Archeops,Curto2007Archeops} applied Minkowski functionals to several CMB datasets; \\cite{Cabella2005LocCurv} applied local curvature on WMAP 1-year data and \\cite{Monteserin2006} developed scalar statistics using the Laplacian as a tool to test Gaussianity. Alternative indicators based on skewness and kurtosis have been proposed by \\cite{Bernui2008}. Tests based on wavelets were applied to WMAP 1-year data \\citep{Vielva2004,Mukherjee2004}, and WMAP 5-year data by \\cite{Curto2008waveNG} and \\cite{McEwen2008}. Wavelets have been also used in CMB studies \\citep{Martinez-G2002} to identify anomalies in WMAP data \\citep{McEwen2008,Wiaux2008Morph,Cruz2008,Cruz2007ColdSpotW3,Vielva2007ColdSpot,Wiaux2006,Cruz2005ColdSpot,Vielva2004,Pietrobon2008AISO}, denoising \\citep{Sanz1999Denoising}, point sources extraction \\citep{Cayon2000PointS,Gonzalez-Nuevo2006PointS}. Very recently \\cite{Vielva2008} constrained primordial non-Gaussianity by means of N-point probability distribution functions.  In this work, we constrain the primordial non-Gaussianity by applying for the first time a novel rendition of spherical wavelets called {\\em needlets} to the WMAP 5-year CMB data. The needlet construction is discussed by \\cite{NarcowichPetrushevWard2006,Baldi2006,Baldi2007,Marinucci2008NEE,Geller2008Mat}. Needlets have a number of interesting properties which make them a promising tool for CMB data analysis: they live on the sphere, without relying on any tangent plane approximation, are quasi-exponentially localised in pixel space and can be easily constructed from a filter function with a finite support in harmonic space. Combined with the specific shape of this function, this guarantees a tiny level of correlation among needlets, which are only marginally affected by the presence of sky cuts. An exhaustive discussion can be found in \\cite{Pietrobon2006ISW} and \\cite{Marinucci2008NEE}. Recently, the needlet formalism has been extended to the the polarisation field, as discussed by \\cite{Geller2008SpinNee}. Here we try to test whether needlets can lead to interesting constraints on the non-Gaussian amplitude, when compared to previous estimates of $\\fnl$.  This paper is organised as follows: in Sec.~\\ref{sec:need_form} we review the needlets formalism; in Sec.~\\ref{sec:sims_res} we describe the statistical analysis of WMAP 5-year data and the constraints on $\\fnl$; we summarise our conclusions  in Sec.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  Primordial non-Gaussianity is becoming one of the keys to understand the physics of the early Universe. Several tests have been developed and applied to WMAP data to constrain the non-linear coupling parameter $\\fnl$. Recently, different methods \\citep{YadavWandelt2007,Curto2008waveNG} brought to different constraints on $\\fnl$ using similar datasets, WMAP 3-year and WMAP 5-year respectively. The two approaches have been shown to have the same power in constraining primordial non-Gaussianity, while they obtained different best fit values for $\\fnl$. Whereas this might be due to the masks applied to the datasets, it certainly underlines the complexity and difficulty of measuring $\\fnl$. The next generation of experiments will provide data with excellent angular resolution and signal-to-noise ratio which will be decisive to confirm or confute the measurements of  $\\fnl$ of the references above. In this respect, it will be even important to constrain $\\fnl$ with different methods in order to get a more robust detection or to spot spurious presences of non-Gaussian signal. Moreover, integrated estimators, not based on Wiener filters, are differently sensitive to the non-linear coupling and can be useful to address exotic non-Gaussian models which predict high values of $\\fnl$ and whose bispectrum evaluation, for instance, may require prohibitive computational time due to the convolution in the bispectrum formula.  In this work, we constrained the primordial non-Gaussianity parameter $\\fnl$ by developing the needlets formalism and applying it to the WMAP 5-year CMB data. We estimated $\\fnl$ to be $20$ with $-30<\\fnl<70$ and $-80<\\fnl<120$ at 1 and 2 sigma respectively, then consistent with the Gaussian hypothesis. We performed two different analyses, the $\\chi^2$ statistics and an estimator based on the skewness of the primordial non-Gaussian sky, finding an excellent agreement between the two results. Needlets have been proven to be a well understood tool for CMB data analysis, sensitive to the primordial non-Gaussianity. Since the skewness of the needlets coefficients is mainly sensitive to the equilateral triangle configurations, we improved our estimator computing the three point correlation function in needlet space which indeed recovers the signal due to squeezed triangle configurations. We obtain $\\fnl=30\\pm40$ at $68\\%$ confidence level, consistent with the previous analysis. Our constraints are slightly broader than those achieved by \\cite{Curto2008waveNG} and not in contrast with the values found by \\cite{YadavWandelt2007} since we were limited by a smaller range of multipoles, whereas the tighter constraints on $\\fnl$ crucially depend on the maximum multipole considered.  Our limits on $\\fnl$ are quite promising for future experiments like Planck, whose sensitivity and angular resolution will be enormously improved."
    },
    "0812/0812.0011_arXiv.txt": {
        "abstract": "A new era of directly imaged extrasolar planets has produced a three-planet system \\citep{2008M}, where the masses of the planets have been estimated by untested cooling models.  We point out that the nominal circular, face-on orbits of the planets lead to a dynamical instability in $\\sim$$10^5$~yr, a factor of at least $100$ shorter than the estimated age of the star.  Reduced planetary masses produce stability only for unreasonably small planets ($\\lesssim 2$~$M_{\\rm Jup}$).  Relaxing the face-on assumption, but still requiring circular orbits while fitting the observed positions, makes the instability time even shorter.  A promising solution is that the inner two planets have a 2:1 commensurability between their periods, and they avoid close encounters with each other through this resonance.  That the inner resonance has lasted until now, in spite of the perturbations of the outer planet, leads to a limit $\\lesssim 10 M_{\\rm Jup}$ on the masses unless the outer two planets are \\emph{also} engaged in a 2:1 mean-motion resonance.  In a double resonance, which is consistent with the current data, the system could survive until now even if the planets have masses of $\\sim20$~$M_{\\rm Jup}$.  Apsidal alignment can further enhance the stability of a mean-motion resonant system.  A completely different dynamical configuration, with large eccentricities and large mutual inclinations among the planets, is possible but finely tuned. ",
        "introduction": "\\label{sec:intro}  The method of direct imaging for the discovery of extrasolar planets has yielded spectacular first results over the last several years \\citep{2004C, 2008L, 2008M, 2008K, 2009Lagrange}.  Direct imaging is a method for discovering planets located far from their host stars, an as-yet unexplored region of parameter space, and it promises new opportunities to characterize the planets using their own radiation.  However, because the gravitational influence of directly-imaged planets is not measured and the astrometric orbital arcs obtained so far are short, determining the planetary masses and orbital architectures of these systems is challenging.  In the newly-discovered planetary system HR 8799 ($=$ HD 218396), three planets have been imaged at projected separations of 24, 38, and 68 AU from their host star \\citep{2008M}.  The best current estimate of their masses is derived from the planetary luminosities, measured in the infrared.  Because these planets are young and massive, they are still radiating prodigiously as they contract, cool, and become more gravitationally bound.  The masses are estimated using untested models of this contraction and cooling process.  One class of such models, the ``hot-start'' models, provides the largest luminosity possible at a certain mass and age, given assumptions about opacities in the planetary atmosphere.  Hot-start models have initially extended envelopes and a large entropy per baryon; even hotter models converge to a common track after a few Myr \\citep{2002B}.  Therefore, for a given age and luminosity, these models should provide a lower limit on the mass.  For HR 8799, the lower-limit masses are $5$-$11$, $7$-$13$, and $7$-$13$~$M_{\\rm Jup}$ for planets b, c, and d, respectively, based on a rather uncertain stellar age of $30$-$160$~Myr\\footnote{ This estimate is given by \\cite{2008M} based on four age indicators: Galactic space motion, main-sequence fitting, stellar pulsations, and the massive debris disk.  The first is circumstantial but consistent with the quoted ages; the others suggest an age $\\lesssim 100$~Myr.  That the star has reached the main sequence suggests that it is at least several $10$s of Myr old. }, which is presumably also roughly the age of the planets.  The following simple calculation illustrates why a lower mass limit can be inferred from a planet's contraction luminosity.  For HR 8799, the planetary luminosities have been measured to be $L \\simeq 10^{-5} L_\\odot$, and radii of $R \\simeq 1.2 R_{\\rm Jup}$ were derived from the objects' temperatures, measured by fitting photometry with a variety of synthetic spectral energy distributions \\citep{2008M}.  Because the objects are cooling, they were more luminous in the past, so they have radiated at least $L t_{\\rm age} \\gtrsim 4 \\times 10^{43}$~erg.  Their current binding energy, which supplied this luminosity, is $\\simeq G M^2 R^{-1} \\simeq 3 \\times 10^{43}(M/M_{\\rm Jup})^2$~erg, where the radius is roughly independent of the mass for Jupiter-mass objects.  Consequently, $M > 1 M_{\\rm Jup}$.  Cooling models also take into account that $L$ diminishes with time, and thus arrive at a considerably larger mass.  Whether this larger calculated mass is a robust lower limit depends on the accuracy of the model.  Recently, \\cite{2009D} measured the dynamical masses for a system of brown dwarfs (both of mass $\\approx 57 M_{\\rm Jup}$) and showed that cooling models overpredict the component masses by $\\sim$25\\%.  If energy is lost during the process of planet formation, then an even larger planet mass would be needed to generate the currently observed luminosity.  For example, in the planetary core-accretion models of \\cite{2007M}, considerable luminosity is radiated in the accretion stream and shock, and that energy is not internalized by the planet.  At the end of formation the planet has less gravitational potential energy to later supply its luminosity.  The integrated luminosity since formation would not account for the planet's current binding energy, so the mass needed to supply an observed luminosity at a given age may be much bigger.  The HR 8799 system has survived an order of magnitude longer than the primordial gas disk, which, if typical of disks of A stars, lasted $\\lesssim3$~Myr \\citep{1993H,2005H}.  The system has therefore had time to dynamically evolve in the absence of gas. Though the planets orbiting HR 8799 are separated by tens of AU, the inferred minimum masses of the planets were large enough that their mutual gravitational interactions are important.  For example, a planet with mass $M_p = 10$~$M_{\\rm Jup}$ orbiting a star of mass $M_* = 1.5$~$M_\\odot$ at semi-major axis $a = 40$~AU dominates gravitational dynamics within its Hill radius of size $R_H = a(M_p/3M_*)^{1/3} = 5$~AU.  Because $R_H$ is a large fraction of the planetary separation, gravitational interactions among the planets can substantially modify the dynamical evolution of the system.  In fact, the nominal orbits reported in the discovery paper \\citep{2008M} are unstable.  We integrated the Newtonian equations of motion of the proposed system using the Bulirsh-Stoer (BS) algorithm of the {\\slshape Mercury} \\citep{1999C} package (version 6.2), with an accuracy parameter of $10^{-12}$.  The planets are assigned circular, face-on orbits, and we used the nominal masses for all four bodies: $7$, $10$, $10$~$M_{\\rm Jup}$ for planets b, c, and d, respectively, and $1.5$~$M_\\odot$ for the star\\footnote{For the specific initial conditions of this and other integrations herein, see Tables~\\ref{tab:logI} and \\ref{tab:logII}.}.  Figure~\\ref{fig:ae} shows the results for the semi-major axis and maximum radial excursion of each planet as a function of time.  A close encounter between planets c and d at $0.298$~Myr (i.e., they enter within one Hill radius of one another) leads to a brief interval of strong scattering which ejects planet b at $0.316$~Myr (i.e., it reaches $>500$~AU with positive energy, and is removed from the simulation).  Planets c and d swap orbits and finish in a stable configuration, with no further semi-major axis evolution, but they exhibit a regular secular eccentricity cycle with a period of $1.5$~Myr.  This evolution is not unique in its details since the orbital evolution is chaotic.  However, qualitatively similar evolutions are common for simulated planetary systems constructed to match the discovery data: instability usually sets in well before the star's age of $\\gtrsim 30$~Myr.  \\begin{figure} \\epsscale{1.1} \\plotone{f1.eps} \\caption{ Semi-major axis, periapse, and apoapse for the three planets as a function of time in a numerical integration of the nominal model of the system, which has circular, face-on orbits, and planetary masses $M_b = 7 M_{\\rm Jup}$, $M_c = 10 M_{\\rm Jup}$, and $M_d = 10 M_{\\rm Jup}$.  An instability occurs only $3\\times10^5$~yrs into the integration, in which planets b and c suffer a close encounter, suggesting the planetary masses or orbits of the nominal model are in error. \\vspace{0.1 in}} \\label{fig:ae} \\end{figure}  The goal of this paper is to determine orbits that are consistent with the astrometric data, the inferred planetary masses, \\emph{and} with dynamical stability over the system's age.  Neglecting stability considerations, there is a large amount of freedom in fitting orbits to the discovery data, because (1) the measured astrometric arcs cover only $\\sim$2\\% of the middle orbit and $\\sim$1\\% of the outer orbit, (2) the velocity of the inner planet is almost entirely unconstrained, and (3) the line-of-sight positions and velocities of the planets relative to the star are unknown.  {\\it A priori}, two classes of orbital architectures are possible---those in which the planets occupy roughly coplanar orbits and those with large mutual inclinations.  Since planets form in disks, it is likely that they initially occupy nearly coplanar orbits, and systems that remain stable indefinitely are likely to stay roughly coplanar.  Alternatively, the system may not be indefinitely stable. While old compared to the lifetime of the protoplanetary disk, the current age of the planetary system is probably less than one tenth the main-sequence lifetime of the star ($\\sim$$1.5$~Gyr; \\citealt{1967I}).  Without further analysis, it is thus possible that the planets are in the process of scattering off of one another, currently have large eccentricities and mutual inclinations, and will not be stable over the lifetime of the star.  In fact, current models predict that planetary systems undergo periods of strong mutual excitation, perhaps generically leading to the ejection of planets (e.g., \\citealt{lld98}, \\citealt{gls04}, \\citealt{2009SM}, \\citealt{2009VCF}).  \\begin{deluxetable*}{lcccc} \\tablecaption{Astrometric constraints\\label{tab:nom}} \\tabletypesize{\\normalsize} \\tablewidth{0pt}  \\tablehead{ \\colhead{Planet} & \\colhead{[$x_E$, $x_N$]  (AU)} & \\colhead{$s$ (AU)} & \\colhead{[$v_E$, $v_N$]  (10$^{-3}$ AU day$^{-1}$)} & \\colhead{$v_p$  (10$^{-3}$ AU day$^{-1}$)} }  \\startdata b\t&[$60.16(6), 31.50(6)$] & $67.91(6)$&\t[$1.49(0.14),  -2.18(0.14)$] & $ 2.64(0.14)$\\\\ c\t&[$-25.90(6), 27.76(6)$] &$37.97(6)$& [$2.15(0.14), 2.47(0.14)$]& $ 3.27(0.14)$\\\\ d\t&[$-8.45(6), -22.93(6)$] &$ 24.44(6)$& [$-3.5(2.7),  0.0(2.7)$] & --- \\enddata \\vspace{0.1 in} \\tablecomments{   \\scriptsize{Sky-projected positions and velocities of each planet relative to the star, found by a rectilinear least-squares fit to the astrometry of Table 1 of \\cite{2008M}.  Positions are at the epoch of 2008 Aug. 12, velocities assume no detectable orbital acceleration.} \\vspace{0.2 in} }  \\end{deluxetable*}   To explore these possibilities systematically, we take the following approach. We start with restrictive assumptions about the orbital architecture of the system, and we then progressively relax those assumptions.  At each stage, we find parameters that maximize the stability, and we finally argue that a resonant configuration is most likely for the system to have survived to its current age.  In \\S\\S\\ref{sec:astrom}--\\ref{sec:planarei}, we assume that the orbits of the planets are close to coplanar.  In \\S\\ref{sec:astrom} we discuss astrometric constraints on the orbits and show that, although the data are consistent with circular and coplanar orbits, the system orientations that fit the data best do not generate stable orbits.  In \\S\\ref{sec:masses} we determine what planetary masses would be needed for circular, coplanar orbits to be stable and argue that they are too low given the observed luminosities.  Having thus ruled out circular, coplanar solutions, we next allow the planetary eccentricities to vary.  Since the inner planet's eccentricity is unconstrained by the data, we first scan over non-circular orbits for the inner planet while keeping the outer two planets on circular orbits (\\S\\ref{sec:stable}).  The suggestive results of this experiment led us to our preferred configuration for the planetary system: a mean motion resonance between the inner two planets (\\S\\ref{sec:mmr}).  An initial exploration of the parameter space of possible resonant orbits shows that if the outer two planets are also in a mean motion resonance, the system could be stable even if the companion masses are twice as large as the nominal masses.  In \\S\\ref{sec:montecarlo}, we allow all the orbital parameters to vary via a Monte-Carlo method. We confirm that mean-motion resonance is the most likely reason the planetary system has survived.  We also find a scattering-type configuration that is stable for $30$~Myr, but we argue that it is unlikely.  In \\S\\ref{sec:conclude} we discuss our conclusions. ",
        "conclusions": "\\label{sec:conclude}  We have investigated the orbital stability of the newly-imaged planetary system HR 8799.  The nominal orbital model and masses are not stable.  In fact, no model with circular, coplanar orbits that also fits the astrometry well is stable, regardless of the inclination and orientation of the system on the sky.  To overcome this problem by reducing the planetary masses, values $\\lesssim 2$~$M_{\\rm Jup}$ are required. This can happen if the cooling models under-predict the luminosity, though that is difficult to understand, as even hot-start models cannot produce the observed luminosity at such low masses (see \\S\\ref{sec:intro}).  Such masses would be plausible if the system is considerably younger than expected, yet the star has reached the main sequence \\citep{2008M}.  Our favored solution is that a 2:1 resonance between the inner two planets preserves stability.  Assuming that the inner pair of planets are in resonance, two qualitatively different configurations are possible: \\begin{itemize} \\item{The outer two planets are not in resonance. } This configuration remains stable in the perturbing presence of planet b only if the planetary masses are $\\lesssim 10$~$M_{\\rm Jup}$ (Fig.~\\ref{fig:massfactor}). (It is also possible, given a less likely system orientation, that only the outer resonance is active.  This configuration leads to a similar mass limit.) \\item{The outer pair of planets are also in 2:1 resonance.  } This solution fits all the current data for the system.  At the nominal masses, the system can easily survive for the age of the star.  In fact, the planetary masses could be up to $\\sim1.9$ times bigger than their nominal values without violating stability constraints (Fig.~\\ref{fig:massfactor}).  This value is similar to the maximum planetary masses that can stably exist at the fixed point of the 2:1 resonance \\citep[fig. 8]{2003B}.  It will be very interesting to find a test of this hypothesis.  The Laplace angle can also librate in this system, though this is not a requirement for stability. \\end{itemize} Strong apsidal alignment, while unnecessary for system stability, allows planetary survival in a weaker 2:1 mean-motion resonance.  One final type of system architecture is, in principle, consistent with stability of the HR 8799 system.  The planets may be hierarchically spaced, with large eccentricities and mutual inclinations, but inhabit phases of their orbits that look closely-packed at the moment.  Such systems fit the current data but are finely tuned both in their orbital parameters and in the time at which we are viewing the system.  It is possible that other stabilizing resonant configurations exist, but were missed because they occupy small regions of phase space.  The 2:1 mean-motion resonance dominates our randomly-generated stable systems and would naturally yield the observations without fine tuning, so we consider it to be the most likely stabilizing mechanism.  This study brings up several issues for future observations of HR 8799, and directly-imaged multiplanet systems in general, as follows.  First, it serves as the first test of hot-start cooling models for exoplanets.  They barely pass the test if only planets d and c are in resonance, and they comfortably pass the test if all three planets are in resonance.  We hope more detailed dynamical studies of this system will sharpen this test as more data are collected.  If the doubly-resonant configuration can somehow be verified, it would considerably weaken this test.  We have not yet directly used dust observations as a dynamical constraint.  The spectral energy distribution reveals a massive debris disk surrounding the planetary system, with an orbital radius of $\\gtrsim 66$~AU \\citep{2006WA}.  Given that planet b is observed at $a_b \\gtrsim 68$~AU, we expect that the inner edge of the debris disk must be $\\gtrsim 90$~AU, and that future measurements and modeling will find that orbital radius to be plausible and even preferred.  Such a model could in turn serve as a complementary test of planet b's mass, in analogy to the test of the mass of Fomalhaut b \\citep{2009C}.  Perhaps other directly-imaged systems will fortuitously arrange for complementary tests.  Second, we found evidence of a mean motion resonance at very large orbital separations, much farther than those found by the radial velocity technique, of which there are many (e.g., \\citealt{2001Marcy,2004Mayor}).  Sometimes resonant identification is based on stability arguments in those systems (e.g., \\citealt{2005C}), as it is here.  The most commonly invoked evolutionary mechanism for trapping planets into resonance with one another is convergent migration in the protoplanetary disk.  The properties of migration in a disk with multiple massive planets deserve further investigation to determine the conditions under which convergent migration and resonance capture are possible at the locations of the HR 8799 planets. We verified that if planets d and c were initially placed in circular orbits exterior to the 2:1 resonance, they are stable to collisions or ejections for $\\gtrsim 30$~Myr, so getting into the resonance without first becoming unstable is not a problem in this case.  One difficulty with differential migration is that any additional migration, after the resonance is reached, efficiently increases the eccentricities.  That this requires implausible fine-tuning in the absence of eccentricity damping by the gas disk has been discussed for the 2:1-resonant GJ876 system \\citep{2002LP}. In a trial integration, we introduced a force to simulate outward migration of the inner planet, following \\cite{2002LP}, with a timescale $a/(da/dt) = 10^7$~yr and no eccentricity damping, starting from double-2:1 resonance at the nominal masses (fig.~[\\ref{fig:boundres}], table~\\ref{tab:logI}).  The planetary eccentricities rapidly increased and the system began scattering after a semi-major axis change of $\\sim15$\\%, illustrating its fragility.  We expect calculations of migration into resonance will be very interesting for this system.  We also showed how perturbations by a third planet tend to disrupt a mean-motion resonance (when only the inner sub-system is in resonance).  This mechanism adds to a growing list of ways to disrupt resonances among planets, including turbulent fluctuations in a protoplanetary disk \\citep{2008A}, tidal dissipation \\citep{2007T}, and scattering of planetesimals \\citep{2006M,2007Morby}.  Third, it may seem surprising that dynamical stability arguments are needed to correctly solve the orbits of the first directly-imaged multiplanet system.  However, this situation is also common for multiplanet systems discovered by radial velocity.  For instance, \\cite{2005V} and \\cite{2006Lee} have found that dynamical stability can constrain orbital parameters more tightly than radial-velocity data alone.  Furthermore, planets that are discovered by direct imaging of their self-luminosity are biased to have high masses, making stability less assured.  The bias of direct imaging towards large angular separation implies that very long orbital periods will be common for such discoveries, so we foresee many stability analyses predicated on only the sky-projected positions and velocity vectors of planets.  We hope this paper proves to be a useful example of how to conduct such an analysis."
    },
    "0812/0812.3823_arXiv.txt": {
        "abstract": "We present calculations of the heliospheric solar wind charge exchange (SWCX) emission spectra and the resulting contributions of this diffuse background in the \\rosat\\ \\oqkev\\ bands. % We compare our results with the soft X-ray diffuse background (SXRB) emission detected in front of 378 identified shadowing regions during the \\rosat\\ All-Sky Survey \\citep[][]{snowden00}. This foreground component is principally attributed to the hot gas of the so-called Local Bubble (LB), an irregularly shaped cavity of $\\sim$ 50-150 pc around the Sun, which is supposed to contain $\\sim$10$^6$ K plasma. Our results suggest that the SWCX emission from the heliosphere is bright enough to account for most of the foreground emission towards the majority of low galactic latitude directions, where the LB is the least extended. On the other hand, in a large part of directions with galactic latitude above 30 degrees the heliospheric SWCX intensity is significantly smaller than the measured one. However, the SWCX R2/R1 band ratio differs slightly from the data in the galactic center direction, and more significantly in the galactic anti-centre direction where the observed ratio is the smallest. Assuming that both SWCX and hot gas emission are present and their relative contributions vary with direction, we tested a series of thermal plasma spectra for temperatures ranging from 10$^{\\,5}$ to 10$^{\\,6.5}$ K and searched for a combination of SWCX spectra and thermal emission matching the observed intensities and band ratios, while simultaneously being compatible with \\ovi\\ emission measurements. In the frame of collisional equilibrium models and for solar abundances, the range we derive for hot gas temperature and emission measure cannot reproduce the Wisconsin C/B band ratio. This implies that accounting for SWCX contamination does not remove these known disagreements between data and classical hot gas models. We emphasize the need for additional atomic data, describing consistently EUV and X-ray photon spectra of the charge-exchange emission of heavier solar wind ions. ",
        "introduction": "The diffuse soft X-ray background (SXRB), first observed in the 70's \\citep{bowyer68,williamson74,sanders77} has since been shown to be the sum of local and distant sources. Above 2 keV it is dominated by the extra-galactic background, itself a combination of unresolved point sources and warm-hot interstellar medium (WHIM) diffuse emission \\citep{hasinger93}. At lower energies it is dominated by the galactic halo \\citep{burrows91,snowden94a}, and finally below 0.3 keV it is mainly due to the unabsorbed emission from hot gas filling the so called Local Bubble (LB) \\citep{mccammon83,bloch86,snowden90a,snowden90b}, a cavity devoid of dense gas extended at high latitudes and connected to the halo \\citep{frisch83,welsh98,lallement03}. The main tools used to disentangle local and distant emission are the `shadowing' experiments, i.e. spatial variations of intensity and spectral characteristics around and towards dense, soft X-ray absorbing clouds \\citep[e.g.][]{herbstmeier95}. \\cite{snowden98,snowden00} used more than 370 \\rosat\\ shadows to produce almost full-sky mapping of the `unabsorbed' component of the emission, i.e. the LB contribution.  This was the generally accepted scenario until the discovery of X-ray emission in comets \\citep[][]{lisse96} and the identification of the emission mechanism as charge-exchange (CX) reactions between the highly charged heavy solar wind ions and the cometary neutrals \\citep[][]{cravens97}. \\cite{cox98} suggested that the CX reactions should also occur between heavy SW ions and interstellar neutrals (H and He) in interplanetary space and that the resulting X-ray emission (SWCX) should have an impact on the SXRB interpretation. \\cite{cravens00} estimated that the quiescent level of SWCX emission could be of the same order as the SXRB component attributed to the LB.%  This interplanetary, heliospheric emission is time-dependent because of the intrinsic variable nature of the solar wind. Short time-scale variations tend to be washed out by integration along the line-of-sight and the size of the emitting region \\citep{cravens01}, but longer term variations, including those related to the solar cycle, can cause more persistent changes in the heliospheric emission level. In addition, there is a contribution from the Earth's magnetosphere, due to charge-transfer with exospheric neutrals. Such emission, studied in detail by \\cite{robertson06}, reacts instantaneously to solar wind variations and magnetosphere shape variations, leading to a high variability and the occurence of high intensity peaks following solar events. The spectral characteristics of some spectacular enhancements have been recently recorded by the XMM and Suzaku satellites \\citep[][]{SCK04,henley08}. For such events both magnetospheric and heliospheric contributions may be present.  Very likely most of the sharp increases of terrestrial origin have been removed from the \\rosat\\ map along with the cleaning procedure of the Long Term Enhancements \\citep[LTEs][]{snowden94b}, as well as some heliospheric increases, especially towards the downwind side of the interstellar flow where the gravitational cone of focused helium is the most reactive region. Indeed, most points on the sky in the \\rosat\\ map were observed several times over the course of at least two days, allowing identification and removal of periods of enhanced emission. A debate is still maintained, though, about the actual level of the quasi-stationnary heliospheric contribution to the \\rosat\\ maps of unabsorbed emission. This contribution is extremely difficult to detect from time variations. On the other hand, SCWX and hot gas thermal emission have different spectral properties, i.e. the observed spectral information should help to disentangle the two processes. This is the subject of the present study. For a recent review of all types of SWCX phenomena see \\cite{bhardwaj07}.  The first estimates of the stationary heliospheric contribution \\citep{cravens00,lall04CX} were based on simplifying assumptions about the spectral characteristics of the SWCX emission. Since then the existence of the SWCX phenomenon has motivated theoretical work on exact photon yield values for the charge transfer collisions \\citep{kharchenko00,pepino04}, as well as a number of laboratory experiments devoted to the CX emission mechanism. For a recent review see \\cite{wargelin08}.  The spatial distribution of magnetospheric and heliospheric SWCX emission was modeled by \\cite{robertson03}, revealing significant variations in brightness as a function of earth location, line-of-sight direction, and activity phase. \\cite{kout06} computed similar maps for a few specific energy bands after including \\cite{pepino04}'s detailed CX emission spectra for C, N, O, and Ne ions. \\cite{lall04CX}, taking into account the specific viewing geometry of \\rosat\\, showed that the heliospheric background in the \\oqkev\\ band was nearly isotropic and could have been unnoticed in the All-Sky Survey maps, while accounting for a large portion of the signal, and possibley the major part at low galactic latitudes.  Using both the stationnary and time-dependent models \\cite{kout07} modelled four high-latitude shadowing observations and showed that in the 3/4 keV band, where the oxygen lines (\\ovii\\ triplet at 0.57 keV and \\oviii\\ line at 0.65 keV) are dominant, the SWCX emission from the heliosphere can account for all the unabsorbed, local component of the SXRB, with no need of a LB emission. In parallel, the solar wind contribution to the background and its variability have been shown to be responsible for some discrepant measurements \\citep{smith07} and for supposedly low-energy counterparts of distant objects \\citep{bregman06}.  A 10$^6$ K plasma, however, has very little emission in the 3/4 keV band and mainly emits in the \\oqkev\\ band. The \\cite{kout07} results, while not requiring any LB emission, therefore do not preclude the existence of 10$^6$ K gas. Exact calculations of the SWCX spectra and intensities below 0.3 keV are mandatory if one wants to disentangle LB hot gas diffuse emission from the SWCX background. In this paper we examine the SWCX contribution to the \\oqkev\\ spectral region, compare this contribution to observations and also to contributions from hot gas at different temperatures.  Independently of the SWCX contribution, a number of results have somewhat contradicted the interpretation of the unabsorbed soft X-ray background as the LB 10$^6$ K gas emission.  i) Data from the NASA EUVE satellite and from the dedicated CHIPS mission did not detect the EUV emission expected from surrounding 10$^6$ K gas \\citep{jelinsky95,hurwitz05}. It has been suggested that a very low metal abundance may be responsible for this non-detection, but the required depletion level corresponds to the physical state of very dense clouds, which is unlikely for 10$^6$ K, tenuous gas.  ii) The pressure of this hot gas derived from the X-ray background is far above the pressure within the local interstellar cloud and other clouds embedded in the LB \\citep{lallement98,jenkins02}.  iii) Low latitude absorption measurements of highly charged ions such as \\siiv\\, \\civ\\ and \\ovi\\ formed in conductive interfaces between the hot (10$^6$ K) gas and embedded cold:warm clouds do not seem to correspond to expectations from the models \\citep[][]{slavin02,indebetouw04}. Column densities of \\siiv\\ and \\civ\\ are too small and line-widths too narrow \\citep{welsh05}, and \\ovi\\ is detected only at the periphery of the Local Cavity, while one would also expect interfaces between the hot gas and the local clouds \\citep[][]{welsh08}.  iv) Fundamental discrepancies arise also when comparing the Wisconsin sounding rocket survey data in the B and C bands, and the \\rosat\\ All-Sky Survey data in the R1 and R2 bands. The four bands are pictured in figure \\ref{fig:spec1_4}-\\textit{upper panel}. In the low energy (0.1-0.2 keV) B band \\citep{bloch86,snowden94a}, the intensity seems to be higher than what is predicted by thermal emission models. This has been particularly well demonstrated by \\cite{bellm05} who have made a global study over the 0.1-0.3keV interval. According to this work, a best fit to all energy bands is provided by very low metallicity gas at 10$^{5.85}$ K, but inspection of their results (see their Figure 4) reveals significant discrepancies between measured and observed ratios for this best fit solution. Especially, the B/R12 band ratio favours a low temperature $\\simeq$10$^{5.8}$ K (the B band intensity is high and favours a shift of the spectrum towards low energies), while the R2/R1 ratio favours temperatures above 10$^6$ K (R2 is relatively high, favouring a shift towards high energies).  Whether or not the existence of the SWCX background can help to explain part or all these contradictions is a question that has now to be addressed. This work is a first step in this direction. In section \\ref{sec:model} we describe the the SWCX emission and spectral model we have developed and how we make use in our analysis of the Raymond \\& Smith (R-S) hot plasma model. In section \\ref{CX_IoRosat} we compute the expected SWCX emission and the contribution in \\rosat\\ R1 and R2 bands for each of the 378 shadow regions observed by \\cite{snowden00}. We compare the SWCX R1+R2 intensity with the unabsorbed component derived by the \\cite{snowden00} shadow analysis and discuss the distribution of the discrepancies between data and the SWCX model. In section \\ref{band_ratios} we compute the SWCX model R2/R1 and B/C band ratios, as well as the corresponding ratios for hot gas (R-S model) in collisional equilibrium within a large temperature range. We compare the modeled ratios with the observed band ratios during the two (Wisconsin and \\rosat ) surveys. In section \\ref{combination} we search for a combination of SWCX and hot gas emission compatible with the observed intensities and band ratios and we compare those solutions with observational constraints from \\ovi\\ and EUV background measurements. In section \\ref{discussion} we discuss the results and draw some conclusions. ",
        "conclusions": "\\label{discussion}  We have modeled the intensity and spectral characteristics of the heliospheric SWCX emission at the time of the \\rosat\\ survey and compared with the unabsorbed, local emission derived by \\cite{snowden00} in the \\oqkev\\ band. The results show that the SWCX emission can account for most of the total intensity recorded in the R1+R2 bands for most of low latitude lines-of-sight. A map of the heliospheric SWCX portion of the total signal clearly reveals the high latitude chimneys to the halo as the only regions unambiguously dominated by hot gas emission. Such a result can be interpreted as meaning that little or no hot gas exists within the galactic disk.  The spectral characteristics however reveal more complexity and preclude such a simple scenario. The SWCX band ratios disagree with the observations, especially towards the galactic anti-center and at low energies (C/B). We have thus searched for a combination of SWCX and thermal emission from hot gas in equilibrium and solar abundances able to reproduce the data. Our study shows that a combination of SWCX and thermal emission can reproduce the data in the galactic center hemisphere at low latitudes. For this solution the SWCX emission strongly dominates. The temperature of the hot gas is constrained within the interval 10$^{5.64}$-10$^{6}$. The upper limit is constrained by the lower limit on the SWCX intensity. This upper limit can be considered as firmly determined, thanks to recent observational studies above 0.3 keV that have confirmed the validity of our model \\citep[][]{kout07}. The temperatures lower than 10$^{5.64}$ are excluded by \\ovi\\ emission observations \\citep{shelton03} and interstellar ion absorption lines toward nearby stars.  On the other hand, it is difficult to fit with such a combination the Wisconsin data. In the UW (galactic center direction) a combination of hot gas and SWCX emission gives (B+C) intensities as well as C/B ratios roughly compatible with the observed values for several different temperature ranges. The main difficulty is the impossibility to account for the very low C/B ratio measured towards the galactic anti-center direction with the present input models used in our study. We note that the high B intensity is also clearly a problem for any hot gas solution, including the very high depletion hypothesis, as shown by the study of \\cite{bellm05}. The SWCX contribution, which hardens the spectra, reinforces this difficulty.  However, further investigation is required on the model's uncertainties in order to quantify their influence on the spectral hardness of SWCX emission. Further analysis of the hydrogenic ion approximation is needed, since hydrogenic ions tend to emit photons at higher energies following CX (since high-n to ground transitions are generally allowed) than the multi-electron ions considered here (because selection rules and more complicated atomic structure lead to more cascades before the final transition to ground. Therefore, it is likely that the model spectra have significantly less flux at low energies than they should. Given the steeply decreasing effective areas at lower energies (see Fig. \\ref{fig:spec1_4}), this would help explain some of the discrepancies seen in the C/B ratio. Also, the fact that no distinction was made between the neutral targets (H or He) in the \\oqkev\\ calculations, would also have an effect on the spectral hardness, since there is an effect on the electron capture level (roughly proportional to $\\sqrt{13.6\\,eV/I_n}$, where $I_n$ is the neutral target ionization potential), the capture level with He being somewhat lower than with H. The \\rosat\\ observation geometry tends to smooth out those differences, but the SWCX spectrum will tend to be softer in the DW directions (i.e., looking through or near the He cone), again helping to explain some of the C/B ratio anomaly. Finally, at the low energies of the B band, more detailed calculations of the heliospheric and magnetospheric signals must be performed, especially the low energy secondary SWCX emissions in the heliosheath and heliotail, i.e. subsequent recombinations of partially neutralized solar wind high ions."
    },
    "0812/0812.1221_arXiv.txt": {
        "abstract": "We compute $\\chi$ values for blue emission lines in active M dwarfs.  Using flux-calibrated spectra from nearby M dwarfs and spectral M dwarf templates from SDSS, we derive analytic relations that describe how the $\\chi$ values for the CaII H and K as well as the H$\\beta$, H$\\gamma$, H$\\delta$, H$\\epsilon$ and H8 Balmer emission lines vary as a function of spectral type and color.  The $\\chi$ values are useful for numerous M dwarf studies where the intrinsic luminosity of emission lines cannot be estimated due to uncertain distances and/or non-flux-calibrated spectra.  We use these results to estimate the mean properties of blue emission lines in active field M dwarfs from SDSS. ",
        "introduction": "The domination of M dwarfs in the Galaxy's stellar census makes them ideal tracers of the kinematics, structure, and evolution of the Milky Way. Many of these stars are also host to large magnetic fields that act to heat the upper atmosphere and give rise to magnetic activity (as often traced by chromospheric emission lines).  Observations of magnetic activity put important constraints on the internal structure, the relationship between magnetic field generation and rotation, atmospheric models, and the ages of M dwarfs.  Although low-mass stars are intrinsically faint, recent large surveys include an unprecedented number of M dwarfs. The Sloan Digital Sky Survey \\citep[SDSS;][]{DR6} in particular has identified $\\sim$30 million M dwarfs in the photometric database (Bochanski et al., in prep.) and over 40,000 M dwarfs with spectra \\citep{W08}.  \\citet{W04,W06} and \\citet{Boo07} used these spectra to statistically examine the magnetic activity properties (as traced by H$\\alpha$) of field M dwarfs.  Because the SDSS sample traces a large range of physical and dynamic properties, it provides an important laboratory for studying how magnetic activity changes as a function of mass, metallicity, rotation, and stellar age \\citep{W08}.  Magnetic activity is often quantified by computing the ratio of the luminosity in an emission line (typically H$\\alpha$) to the bolometric luminosity of the star \\citep[L$_{\\rm{H}\\alpha}$/L$_{bol}$;][]{Reid95,Hawley96,Burgasser02,W04}. This quantity allows for comparison of stars of different temperatures, where changes in the continua can severely affect the derived equivalent widths.  Although the activity-bolometric ratio is strictly distance independent, the challenges associated with obtaining accurate bolometric fluxes (namely good bolometric corrections for individual stars) have resulted in studies that use a color or spectral type dependent bolometric luminosity \\citep[e.g.][]{Hawley96}, and thus still require a distance to solve for the line emission luminosity.  For H$\\alpha$, the line luminosity is often calculated using the equivalent width of the emission line, and the luminosity of the continuum near H$\\alpha$; the latter can be measured directly from a flux-calibrated spectrum and a measured stellar distance.  Alternatively, the continuum luminosity can be estimated from the photometric colors \\citep{Reid95}. Because of the difficulty in calculating individual bolometric corrections for large samples of stars, the fact that many spectra are not flux-calibrated (e.g. high resolution echelle data and data taken under non-photometric conditions) and that distance estimates to most M dwarfs are still a major source of systematic uncertainty \\citep{Covey08}, an alternative approach to calculating the activity-bolometric ratio is required.  \\citet[hereafter WHW04]{WHW04} and later \\citet{W05}, introduced a novel method for calculating L$_{\\rm{H}\\alpha}$/L$_{bol}$ by removing many of the intermediate steps and requiring only a measurement of the equivalent width of the emission line (EW) and either the color or the spectral type of the star.   WHW04\\nocite{WHW04}, defined a sample of nearby stars with known distances (from trigonometric parallax), derived bolometric corrections and used flux-calibrated spectra to solve for the ratio of the H$\\alpha$ continuum region luminosity to the bolometric luminosity as a function of both spectral type and color.  This ratio is dubbed $\\chi$,  \\begin{equation} \\chi = L_{\\rm{H}\\alpha}(continuum)/L_{bol}, \\end{equation}  \\noindent and when multiplied by the EW of the H$\\alpha$ emission line yields the following L$_{\\rm{H}\\alpha}$/L$_{bol}$ ratio:  \\begin{equation} L_{\\rm{H}\\alpha}/L_{bol}=\\chi \\times EW_{\\rm{H}\\alpha} ~(\\rm{\\AA}). \\end{equation}  \\noindent The power of the $\\chi$ method (WHW04\\nocite{WHW04}) is that it does not require a flux-calibrated spectrum to obtain L$_{\\rm{H}\\alpha}$/L$_{bol}$, and it is distance independent.  Although previous studies had applied similar strategies for obtaining L$_{\\rm{H}\\alpha}$/L$_{bol}$ \\citep{Reid95,Hawley96,MB03}, they did not explicitly report recipes for obtaining ``$\\chi$'' and they relied on colors and photometric or spectroscopic parallax relations, or in the latter case, spectral models, to obtain L$_{\\rm{H}\\alpha}$(continuum). WHW04 was the first study to use flux-calibrated spectra to determine L$_{\\rm{H}\\alpha}$(continuum) and to present empirical $\\chi$ values in H$\\alpha$ for the entire M dwarf sequence.  Recently, \\citet{Reiners08} derived $\\chi$ values for H$\\alpha$ from model spectra as a function of effective temperature. Although the \\citet{Reiners08} $\\chi$ relations are similar to those in WHW04 for stars hotter than 2600 K ($\\sim$M6), they do not produce the plateau seen in the late-type empirically derived $\\chi$s from WHW04.  This discrepancy is likely due to problems with the spectral models of late-type dwarfs, which do not reproduce the correct optical colors and magnitudes of M dwarfs \\citep{Covey08}.  The difference in derived log(L$_{\\rm{H}\\alpha}$/L$_{bol}$) values using the two methods rises to 0.6 dex at a spectral type of M9.  We prefer to rely on the measured data rather than models, and adopt the WHW04 $\\chi$ values for H$\\alpha$ in our analysis.  The WHW04\\nocite{WHW04} $\\chi$s have shown great utility in computing L$_{\\rm{H}\\alpha}$/L$_{bol}$ in numerous studies \\citep[e.g.,][]{W04,Silvestri06,Boo07,Reiners08}, but H$\\alpha$ is not the only emission line produced by an active chromosphere. In fact, H$\\alpha$ may not always be the best tracer of activity as compared to other lines \\citep{Browning08, Walkowicz08}.  Traditionally, there have been several limitations in using the bluer emission lines, namely that the signal-to-noise ratio (S/N) in the blue part of an M dwarf spectrum is low because the continuum emission peaks in the near-infrared, and CCDs have been more sensitive in the red.  Several recent studies have examined magnetic activity as traced by the Ca II H and K emission lines and higher order Balmer transitions including H$\\beta$, H$\\gamma$, H$\\delta$, H$\\epsilon$ and H8 \\citep{Rauscher06, Boo07,Browning08, Walkowicz08}, but star-to-star comparisons have been limited by the lack of line specific $\\chi$ conversions to activity-bolometric luminosity ratios.  In this paper, we expand upon the WHW04\\nocite{WHW04} study, using similar methods to compute $\\chi$ values for the Ca II H and K emission lines as well as the H$\\beta$, H$\\gamma$, H$\\delta$, H$\\epsilon$ and H8 Balmer emission lines.  We describe the data and our techniques in \\S2. In \\S3 we derive $\\chi$ for each emission line as a function of both color and spectral type and use these values to investigate the mean emission line properties of active M dwarfs.  We discuss the results in \\S4. ",
        "conclusions": "The $\\chi$ values from WHW04\\nocite{WHW04} provide a distance-independent way to compare the H$\\alpha$ activity of stars with different spectral types.  We used a spectroscopic sample of nearby M dwarfs from \\citet{PH89} and the \\citet{Boo07} templates to expand the $\\chi$ values to include the H$\\beta$, H$\\gamma$, H$\\delta$, H$\\epsilon$, H8, CaII K, and CaII H emission lines.  We fit analytic relations to all of the $\\chi$ values (including H$\\alpha$) as a function of spectral type and color ($i-z$, $i-J$, and $V-I$).  The new $\\chi$ values are useful for analyzing both past and future M dwarf data.  Magnetic activity can be easily quantified if studies include EW measurements (measured using the continuum regions given in Table \\ref{tab:cont}) and either spectral types or colors.  We show that the percentage of bolometric luminosity emitted in the optical emission lines of active-field M dwarfs increases monotonically with increasing wavelength.  In addition, the emission line luminosities relative to the bolometric luminosities are constant in all measured lines for spectral types M0-M4, and then decrease at later spectral types, in agreement with the H$\\alpha$ results in \\citet{W04}.  The L$_{\\rm{CaII K}}$/L$_{bol}$ values are consistent with those measured for the H$\\alpha$ active stars from \\citet{Browning08}. Unfortunately, the H$\\alpha$ inactive templates of \\citet{Boo07} do not have a sufficient S/N at each spectral type to systematically compare the strength of the CaII K emission lines in active and inactive stars.  However, a CaII K emission line is present in the M5 inactive template. The resulting mean L$_{\\rm{CaII K}}$/L$_{bol}$ for the inactive template is an order of magnitude lower than the (H$\\alpha$) active template.   Our $\\chi$ values have already proven useful in the rotation-activity analysis of \\citet{Browning08} by allowing L$_{\\rm{CaII K}}$/L$_{bol}$ and L$_{\\rm{CaII H}}$/L$_{bol}$ to be computed from echelle spectra. Future improvements to our $\\chi$ measurements can be made by creating a larger spectroscopic sample of nearby M dwarfs with good parallaxes. The advent of synoptic telescopes like Pan-STARRS, LSST and SIM will measure parallaxes for thousands of M dwarfs, many of which will have SDSS (or other previously observed) spectra. These samples will enable tighter constraints on $\\chi$ and better inform our understanding of magnetic activity in the majority of stellar inhabitants of the Galaxy."
    },
    "0812/0812.3224_arXiv.txt": {
        "abstract": "The next generation of astrometric instruments will reach accuracies deserving new treatments. In order to get astrometric parameters achieving the precision permitted by the measurements, it will be necessary to take into account effects that were neglected until the present time. Two effects concerning the orbital elements of binary stars are considered hereafter: the former is the local perspective (LP) effect, which is due to the variation of the distance and of the orientation of the orbital plane during the observation time span; the equations describing this effect are derived for the first time. The latter effect is the light--travel time (LTT), which is also related to the orientation of the orbital plane, and which is as efficient as the preceding one. Taking these effects into account would allow to find the ascending nodes of the orbits, and lead to orbital elements more accurate than when they are ignored.  It is derived from simulations that, at a distance of 5~pc, and assuming velocities typical of Pop.I stars, the position of the right ascending node could be derived for a few simulated unresolved binaries when the astrometric measurements have errors around 1~$\\umu$as. For the resolved brown dwarf binary 2MASS J07464256 +2000321, it appears that ignoring the LP effect would result in underestimating the masses of the components by 14 per cent of the errors as soon as the astrometric errors are around 20~$\\umu$as for each measurement. However, a `degenerate LP solution', taking into account the variation of the semi-major axis when the distance is varying, should provide reliable masses when the measurement errors are larger than 1 or 2~$\\umu$as. A few binaries in the {\\it Gaia} program could deserve a degenerate LP solution, whereas a the complete LP+LTT solution could be justified for resolved binaries observed with {\\it SIM}. ",
        "introduction": "Thanks to the progresses in spatial astrometry, the errors of measurements of the positions of stars are gradually decreasing: the errors of the along--scan measurements were a few milli--arcseconds in the {\\it Hipparcos} project \\citep{hip}, but the {\\it Gaia} satellite \\citep{perryman} is expected to achieve 100 times better \\citep{Mignard}. This limit will still be pushed beyond 1~$\\umu$as when interferometry with large baseline will be feasible in space. It is expected that, with a 9-m baseline, {\\it SIM} could already get an accuracy of only 0.6~$\\umu$as for a single measurement \\citep{SIM}.  These important improvements will not only lead to more accurate results. They make also necessary to take into account effects that were never considered in the past. A first example was the correction for perspective acceleration, that was introduced for a few high-velocity single stars in the {\\it Hipparcos} catalogue. The radial velocities (RV) of these stars were used for that purpose, but \\citet{Dravins} have shown that the RV may also be derived from this effect. \\citet{Klioner03} developed a relativistic model taking into account the light deflection by solar system objects. \\citet{Klioner00} and \\citet{Mignard} outlined that some basic parameters as the trigonometric parallax, the proper motion and the RV deserve to be defined more precisely than they are at present; the definition of the RV was discussed by \\citet{linde03}. Concerning binary stars, \\citet{Angla06} have shown that the variations of the light--travel time (LTT) coming from the orbital motion may change the apparent position of the component. Since this effect is related to the radial displacement of the component, they inferred that it could be used to dissipate the ambiguity of the position of the ascending node of the orbit. Another effect which could be used for node identification is the modification of the orientation of the orbit during the time span of the astrometric program. For sake of brevity, it is called `local perspective' (LP) hereafter. In order to avoid confusion, the perspective effects concerning the barycentre of binary systems, but also single stars (i.e. perspective acceleration and change in the parallax), are called `barycentric perspective' hereafter.  In the course of a high-precision astrometric project, taking perspective and LTT effects into account could be necessary for two reasons. First of all, when the effects are relevant, the quality of the solution is improved. The astrometric and orbital parameters are then more accurate, and, even more important, the solution itself will look more `acceptable'. The decision to consider a model as adequate is usually based on the goodness-of-fit (GOF) of the solution, which may, for instance, follow the normal distribution ${\\cal N}(0,1)$ when the model is the right one. Therefore, a large GOF leads to the rejection of the model, and to the search of another one. In the course of a large astrometric survey, like {\\it Hipparcos}, objects which are not satisfactorily fitted eventually got a `stochastic' solution, which is of rather poor interest. Therefore, a binary model neglecting perspective or LTT effects could change actual binary stars in stochastic objects, losing the orbital elements. Aside from giving the right classification with accurate parameters, the second advantage of giving a precise solution is the possibility to get more informations about the object. For single stars, as mentioned above, taking into account perspective acceleration can lead to an estimation of the RV. For binaries, a consequence of the additional effects could be the complete derivation of the orientation of the orbit in space. This is investigated hereafter.  Calculations were performed in order to answer three questions: What is the smallest astrometric error still permitting to neglect LP and LTT without altering dramatically the GOF for unresolved binaries? What is the largest astrometric error permitting to derive the precise ascending node thanks to LP and LTT? What is the smallest astrometric error still permitting to neglect LP and LTT without altering the masses of resolved binary components? The basic equations of LP are derived in Section~\\ref{sec:LP}, and it is shown how to select the right ascending node thanks to this effect. A method for selecting the ascending node from LTT is presented in Section~\\ref{sec:LTT}. In Section~\\ref{sec:relat}, it is verified that the relativistic effects, which were neglected in the calculations, are much smaller than LP and LTT effects for nearby binaries with large semi--major axes. The parameters of simulated unresolved astrometric binaries were derived in Section~\\ref{sec:simulation}, and the ascending nodes were searched. Resolved binaries are treated in Section~\\ref{sec:resolved}. Section~\\ref{sec:conclusion} is the conclusion. ",
        "conclusions": "\\label{sec:conclusion}  We have study the impact of two effects which have been neglected up to now in astrometric studies of binary systems: the local perspective, and the orbital light--travel time. These effects are related to the orientation of the orbit in space, and taking them into account may lead to the position of the ascending node of the orbit, i.e. to the orientation of the spin vector. The statistical distribution of the spin orientations is a clue in the study of binary formation. It was derived and discussed in the past by \\citet{dommanget} and, more recently by \\citet{glebocki}. The global distribution of the spin looks isotropic, but Glebocki concluded from a sample of 252 orbits that subgroups seem to have asymmetric distribution. An extension of his sample would allow to re-consider this question.  We have found that neglecting LTT doesn't affect the $\\chi^2$ of the solution for unresolved binaries, since it is compensated by a bias in the orbital elements, essentially the eccentricity and the orientation of the orbit. The reason is that LTT doesn't change the shape of the apparent orbit; only the observation epochs are affected, in relation with the orbital phase. Therefore, neglecting LTT can't lead to discarding the binary model. On the other hand, taking it into account for stars with distances $d=5$~pc would significantly improve the accuracy of the orbital elements when the astrometric accuracy $\\sigma_w$ is a few $\\umu$as. For other distances, the minimum $\\sigma_w$ is varying as the inverse of $d$.  The ascending node of unresolved astrometric binary orbits may be found from the LP effect. However, for that purpose, astrometric measurements with accuracy around 1~$\\umu$as and beyond are required. When the barycentric radial velocities are known, and therefore taken into account in the correction of perspective acceleration, LP must be taken into account when the errors of the astrometric measurements are less than the 1~$\\umu$as limit. The binary star model could be erroneously discarded otherwise, since the GOF of the solution would look excessive. When the radial velocities are not available, they are derived from the barycentric perspective effects; however the corrections for LP are then necessary only when the errors of the astrometric measurements are less than 0.5~$\\umu$as. As for LTT, the results depending on LP apply to binaries at a distance $d=5$~pc, with velocities typical of population I late-type dwarfs, and with semi-major axes of 1~AU. For other parameters, they are unchanged when the astrometric errors are varying in the same proportions as the velocities and the semi-major axes, and in the inverse proportion of $d^2$.  The resolved binaries were examined through the example of 2MASS J07464256 +2000321, a brown dwarf system for which all the parameters are already known, the ascending node excepted. The LTT effect is not compensated by alteration of the orbital elements as for unresolved binaries, since it plays in opposite directions for both components. It appeared that neglecting LP and LTT would lead to erroneous masses when the errors of the astrometric measurements are around 20~$\\umu$as. However, when the astrometric error is larger than 1 or 2~$\\umu$as acceptable evaluations of the masses may be obtained from a `degenerate LP solution'. This calculation doesn't require the knowledge of the ascending node since it consists in correcting the semi-major axis for distance variations. However, when the astrometric error is around 1~$\\umu$as, the degenerate LP solution is no more acceptable, and only a calculation entirely taking into account LP and LTT may lead to masses as accurate as permitted, through the determination of the ascending node. Since these results apply to a specific binary at a distance of 12~pc and with very light components, it is highly probable that a few binaries in the forthcoming {\\it Gaia} mission will require a degenerate LP solution for getting accurate masses. On the other hand, complete LP+LTT solutions will be necessary only when measurements with errors around 1~$\\umu$as will be available for nearby binaries, possibly thanks to the {\\it SIM} project."
    },
    "0812/0812.4632_arXiv.txt": {
        "abstract": "New multiband CCD photometry is presented for the eclipsing binary GW Gem; the $RI$ light curves are the first ever compiled. Four new minimum timings have been determined. Our analysis of eclipse timings observed during the past 79 years indicates a continuous period increase at a fractional rate of $+$(1.2$\\pm$0.1)$\\times$10$^{-10}$, in excellent agreement with the value $+$1.1$\\times$10$^{-10}$ calculated from the Wilson-Devinney binary code. The new light curves display an inverse O'Connell effect increasing toward longer wavelengths. Hot and cool spot models are developed to describe these variations but we prefer a cool spot on the secondary star. Our light-curve synthesis reveals that GW Gem is in a semi-detached, but near-contact, configuration. It appears to consist of a near-main-sequence primary star with a spectral type of about A7 and an evolved early K-type secondary star that completely fills its inner Roche lobe. Mass transfer from the secondary to the primary component is responsible for the observed secular period change. ",
        "introduction": "GW Gem ($\\rm BD+27^{o}1494$, $V$=$+$10.485, $B-V$=$+$0.27) was announced by Hoffmeister (1949) to be an Algol-type variable. Subsequently, eclipse timings of the system have been reported assiduously by numerous workers, and history is now long enough to understand the binary's period behavior. Light curves have been made only by Broglia \\& Conconi (1981a, hereafter B\\&C) from photoelectric observations during the seasons of 1978 and 1979.  Unfortunately, the comparison star ($\\rm BD+28^{o}1494$) used in their observations turned out to be a $\\delta$ Scuti-type variable with a peak-to-peak amplitude of about 0.035 mag (Broglia \\& Conconi 1981b). B\\&C analyzed their $BV$ light curves, and concluded that GW Gem is a semi-detached system with the secondary star filling its inner Roche lobe. They also suggested that the period of the binary system be considered constant. Most recently, Shaw (1994) included the system in his second catalog of near-contact binaries but with an unknown subclass. In this article, we report and analyze multiband light curves and present the first detailed analysis of the $O$--$C$ diagram. ",
        "conclusions": "Based on the historical and new photometry, we have presented the first detailed period analysis and light-curve synthesis for GW Gem. The absolute stellar parameters for the system can be computed from our complete photometric solutions with the cool-spot model of Table 5 and from Harmanec's (1988) relation between spectral type and physical parameters (effective temperature, mass, and radius). The dereddened color ($B-V$)$_0$=$+$0.21 and temperature of the primary component correspond to a normal near-main-sequence star with a spectral type of about A7.  The astrophysical parameters are listed in Table 6, where the radius ($R_2$) of the secondary star results from the ratio ($r_{2}/r_{1}$=0.856) of the mean-volume radii for each component. The bolometric corrections (BCs) were obtained from the scaling between $\\log T$ and BC recalculated by Kang et al. (2007) from Flower's table. An apparent visual magnitude of $V$=+10.49 at maximum light (B\\&C), appropriate dereddening, and our computed light ratio at phase 0.75 lead to $V_1$=+10.59 and $V_2$=+13.05 for the primary and secondary stars, respectively. Using the interstellar reddening of $A_V$=0.10, the calculated value of $V_1$, and the expected value of $M_{V1}$ for the primary star, we calculated an approximate distance to the system of about 120 pc.  A comparison of the GW Gem parameters with the mass-radius, mass-luminosity, and Hertzsprung-Russell diagrams (Hilditch et al. 1988) clearly demonstrates that the primary component lies between the zero-age and terminal-age main-sequence loci, while the secondary is oversized and overluminous for its mass. In these diagrams, the locations of the two component stars fall amid those of previously-known near-contact binaries. Thus, the system is a semi-detached and FO Vir-subtype, near-contact binary consisting of a detached main-sequence primary component with a spectral type of A7 and an evolved secondary component with a spectral type of approximately K1 which fills its limiting lobe completely. Such a configuration supports the concept of mass transfer from the secondary to the primary component as indicated by our period analysis. Our results suggest that GW Gem currently is in a state of broken contact evolving from a contact configuration as predicted by thermal relaxation oscillation theory (Lucy 1976, Lucy \\& Wilson 1979). High-resolution spectroscopy will determine the astrophysical parameters and evolutionary status of the system better than is possible with photometry alone."
    },
    "0812/0812.3154_arXiv.txt": {
        "abstract": "We have developed a new method to extract halo merger rates from the Millennium Simulation. First, by removing superfluous mergers that are artifacts of the standard Friends-Of-Friends (FOF) halo identification algorithm, we find a lower merger rate compared to previous work. The reductions are more significant at lower redshifts and lower halo masses, and especially for minor mergers. Our new approach results in a better agreement with predictions from the extended Press-Schechter model. Second, we find that the FOF halo finder overestimates the halo mass by up to $50\\%$ for halos that are about to merge, which leads to an additional $\\approx20\\%$ overestimate of the merger rate. Therefore, we define halo masses by including only particles that are gravitationally bound to their FOF groups. We provide new best-fitting parameters for a global formula to account for these improvements. In addition, we extract the merger rate per {\\it progenitor} halo, as well as per descendant halo. The merger rate per progenitor halo is the quantity that should be related to observed galaxy merger fractions when they are measured via pair counting. At low mass/redshift the merger rate increases moderately with mass and steeply with redshift. At high enough mass/redshift (for the rarest halos with masses a few times the \"knee\" of the mass function) these trends break down, and the merger rate per progenitor halo decreases with mass and increases only moderately with redshift. Defining the merger rate per progenitor halo also allows us to quantify the rate at which halos are being accreted onto larger halos, in addition to the minor and major merger rates. We provide an analytic formula that converts any given merger rate per descendant halo into a merger rate per progenitor halo. Finally, we perform a direct comparison between observed merger fractions and the fraction of halos in the Millennium Simulation that have undergone a major merger during the recent dynamical friction time, and find a fair agreement, within the large uncertainties of the observations. Our new halo merger trees are available at http://www.mpe.mpg.de/ir/MillenniumMergerTrees/. ",
        "introduction": "\\label{s:intro} In the current paradigm for structure formation, the cold dark matter model \\citep{WhiteS_78a,BlumenthalG_84a,DavisM_85a,SpringelV_06a}, galaxy mergers play an important role in galaxy formation and evolution. Galaxy mergers drive gas towards central starbursts (e.g.~\\citealp{MihosJ_96a}) and supermassive black holes (e.g.~\\citealp{HernquistL_89a}), and transform galactic morphology (e.g.~\\citealp{NaabT_03a,BournaudF_05a}). Many parameters, like the gas fraction and morphology of the merging galaxies, as well as their relative orbits and orientations, affect the properties of the merger remnant. One of the most important factors is the mass ratio of the merging galaxies (e.g.~\\citealp{NaabT_06b}). Galaxy mergers of all mass ratios are frequent in a \\LCDM{ }Universe (e.g.~\\citealp{StewartK_08a}), but special importance is given to major mergers, usually considered as those with mass ratios less than $\\approx3:1$. These are thought to play a significant role in the buildup of the red sequence \\citep{ToomreA_77a,HopkinsP_07b} by transforming blue star-forming late type galaxies into red passive early type galaxies.  Many observational studies have been carried out in recent years to investigate the fraction of galaxies that show signs of major merger activity as a function of mass, luminosity and redshift. Observationally, only merger fractions can be obtained, and in order to transform them into merger rates, which can be directly compared with theoretical models, the time scale of the observed events must be estimated (e.g.~\\citealp{PattonD_00a}). Two principal approaches are used to observe galaxy merger fractions. One is pair counting (e.g.~\\citealp{PattonD_97a,LeFevreO_00a,BellE_06b,RyanR_07a,LinL_08a}), i.e.~identifying galaxies separated from one another by less that typically $\\approx20\\kpc$. This method probes the pre-merger stage and is therefore a ``progenitor galaxy'' merger fraction \\citep{DeProprisR_07a}. The second approach is identification of mergers through morphological signatures such as asymmetry and tidal tails (e.g.~\\citealp{LeFevreO_00a,ConseliceC_03a,LotzJ_08a}). This method aims at identifying mergers in their relatively late stages, and is therefore a ``descendant galaxy'' merger fraction. For each of the principal approaches different authors use different methods, as well as different selection criteria, which create systematic effects that are not easily comparable. Moreover, the intrinsic uncertainties of each method, and the effect of cosmic variance (e.g.~\\citealp{ConseliceC_08a}) contribute further to the large scatter between the obtained results.  \\citet{DeProprisR_07a} and \\citet{ConseliceC_06a} have noticed that the merger fractions obtained with the two principal methods should be carefully defined, and cannot be directly compared to each other, as they are actually different quantities. Both authors proposed (different) simple conversions between the two quantities. Nevertheless, both disregarded the fact that the mass difference between the progenitors and the descendant is up to a factor of $2$ (in equal-mass binary mergers). That may result in very large differences in the number densities of progenitors and of descendants. Since the measured merger {\\it fraction} implicitly includes the information of the {\\it number} of objects, different number densities are implicitly included in the merger rate or merger fraction per progenitor galaxy versus per descendant galaxy. \\citet{BellE_06a} and \\citet{LotzJ_08a} realised this, and have taken the different number densities into account when comparing merger fractions obtained with different methods. We show in this paper that the different definitions of the merger rate, per descendant or per progenitor halo, lead to significantly different results, particularly for major mergers, high redshift mergers, and high mass mergers.  Galaxy mergers follow their dark matter halo mergers, but the connection between the two is not straightforward. When two dark matter halos merge, the orbital angular momentum is transferred into internal degrees of freedom, while the more concentrated galaxies are at first not much affected. The galaxies lose relative angular momentum due to dynamical friction, and they start merging as well when they are more tightly bound \\citep{BarnesJ_92a}. All of the baryonic physics involved in galaxy mergers makes it difficult theoretically to quantify the galaxy merger rate reliably (P.~F.~Hopkins et al.~2009, in preparation). Nevertheless, a first step towards quantifying the galaxy merger rate would be to understand the dark matter halo merger rate, which can be studied more robustly.  \\citet{LaceyC_93a} have estimated the halo merger rate based on the extended Press-Schechter (EPS) model \\citep{PressW_74a,BondJ_91a,BowerR_91a}. \\citet{NeisteinE_08b} and \\citet{ZhangJ_08b} have recently constructed EPS-based approximations that are self-consistent, and \\citet{ZhangJ_08a} have done so for the ellipsoidal collapse model. All these models differ by factors of a few tenths up to a few. N-body simulations have only recently been used to study the halo merger rate with a large dynamical range (\\citealp{DOnghiaE_08a}; \\citealp{StewartK_08b}; \\citealp{FakhouriO_07a}, hereafter FM08). FM08 have investigated this problem based on the large Millennium Simulation \\citep{SpringelV_05a}. They have found that the dark matter halo merger rate has an almost universal form that can be separated into its dependencies on mass ratio, descendant mass and redshift. Nevertheless, the analysis of N-body simulations is also subject to uncertainties, in particular in identifying halos. FM08 present 3 ways to analyse the simulation, which differ from one another by $\\approx25\\%$, and by a more significant amount compared to the \\citet{LaceyC_93a} analytical approximation. All of these investigations have quantified the halo merger rate per {\\it descendant} halo. Recent work has also used N-body simulations to study mergers of subhalos \\citep{AnguloR_08a,WetzelA_08a}.  In this paper we extract the dark matter halo merger rate from the Millennium Simulation using a new method for identifying halos and mergers, as well as a new definition of the merger rate. In \\S\\ref{s:millennium} we review the Millennium Simulation and its post-processing. In \\S\\ref{s:new_trees} we describe how we create merger trees that are free of artificial effects not considered previously. In \\S\\ref{s:rates} we define the merger rate per descendant halo, as well as the merger rate per progenitor halo, which is related to observed galaxy merger fractions when they are measured via pair counting. In \\S\\ref{s:results} we describe our results and in \\S\\ref{s:comparison} we compare them to previous work. In \\S\\ref{s:merger_fractions} we calculate the halo merger fraction in the Millennium Simulation and compare to observations. In \\S\\ref{s:summary} we discuss our results and their relevance to observations, and summarise. ",
        "conclusions": "\\label{s:summary} We have used the Millennium Simulation to extract merger rates of dark matter halos. Our method differs from previous work in three main aspects.  First, we reject any merger between FOF groups whose descendant subhalos, at any future time, do not belong exclusively to the same FOF group. This is done by keeping such FOF groups distinct until they (if at all) irrevocably merge. Rejecting only a fraction of such events, as in previous work, leads to double counting of mergers, and to false counting of fly-by events as mergers. Therefore, our method results in a lower merger rate, especially in the minor merger regime. \\citet{LudlowA_08a} find that ejections of low mass subhalos out of their host halos typically occur in a configuration where a bound group of subhalos (which was created via past mergers) is accreted onto a large halo, and its low mass members are propelled onto high energy orbits by the multiple-body interaction. In the context of merger counting, it is not clear whether the merger of the low mass halo with the group should be counted as a merger, or merely regarded as an interaction that was interrupted by the accretion of the group onto the large halo. Our method assumes the latter for reasons described in \\S\\ref{s:new_trees}. Therefore, in a more conservative approach, our results can be regarded to as a strong lower limit for the merger rate.  Second, we define each halo's mass as the mass of all the particles gravitationally bound to it, rather than of all the particles constituting the FOF group. This definition reduces the inferred merger rate by typically $20\\%$. The motivation for this definition is our finding that the total (bound \\& unbound) mass of FOF groups artificially increases by up to $50\\%$ as they approach more massive FOF groups on their way to merge with them. This effect distorts the appropriate mass ratios of mergers, thereby changing the merger rate. A detailed discussion of this issue appears in \\S\\ref{s:mass_def}.  These two improvements make our inferred merger rate more consistent with the new predictions of \\citet{NeisteinE_08b} based on the EPS model, in the sense that the functional dependencies of the merger rate on mass ratio, halo mass and redshift are very similar. There is a constant factor of $\\approx1.5$ between the EPS merger rate and the Millennium Simulation merger rate. In \\S\\ref{s:rates_desc} we provide a simple global fitting formula (equation (\\ref{e:rate_desc})) for the merger rate per descendant halo that holds for $z\\lesssim4$ and all masses probed by the Millennium Simulation.  Third, we also extract the merger rate per {\\it progenitor} halo. This allows us to find a merger rate for the full range of mass ratios. For halos of any mass $M$, we can measure the rate at which they undergo minor mergers, major mergers and the rate at which they are being accreted as satellite halos onto more massive halos. This has significant implications for the redshift and mass dependencies of the major merger rate. We find the merger rate, in the regime where both halo masses are smaller than the knee of the mass function, to increase steeply with redshift and only slightly with mass, in a similar way to what was found in previous work. However, at high enough mass or redshift, the number of halos decreases exponentially and therefore the mass dependence of the major merger rate changes sign and starts decreasing with increasing mass. Also, the redshift dependence significantly weakens. In \\S\\ref{s:comparison_prog_desc} we provide an analytic expression for converting merger rates per descendant halo to merger rates per progenitor halo, which can be used for any theoretical merger rate.  The two different definitions of the merger rate, per progenitor halo and per descendant halo, have different physical meanings. They also correspond to the two different observational approaches towards measuring the galaxy merger rate, namely pair counting and morphological/kinematical identification. In \\S\\ref{s:merger_fractions} we discuss the importance of this distinction and its relevance to observations. Finally, we find that observed galaxy merger fractions are consistent with the halo merger fraction in the Millennium Simulation within the large observational uncertainties and uncertainties in parameters (like mass ratio or time scale) needed to compare the two. More refined comparisons still await significant improvements both in observational techniques and consistency and in the theoretical treatment of baryonic physics."
    },
    "0812/0812.1960_arXiv.txt": {
        "abstract": "{Modeling of pre-main sequence (PMS) stars through asteroseismology of PMS p-mode pulsators has only recently become possible, and spacebased photometry is one of the important sources of data for these efforts. We present precise photometry of the pulsating Herbig Ae star \\hd\\,\\,obtained in two consecutive years with the MOST (Microvariability \\& Oscilations of STars) satellite.} {Previously, only a single pulsation period was known for \\hd. The MOST photometry reveals that \\hd\\,\\,is multi-periodic. However, the unique identification of pulsation frequencies is complicated by the presence of irregular variability caused by the star's circumstellar dust disk. The two light curves obtained with MOST in 2006 and 2007 provided data of unprecedented quality to study the pulsations in \\hd\\,\\,and also to monitor the circumstellar variability.} {Frequency analysis was performed using the routine \\SigSpec\\,\\,and the results from the 2006 and 2007 campaigns were then compared to each other with the software \\cind\\,\\,to identify frequencies common to both light curves. The correlated frequencies were then submitted to an asteroseismic analysis.} {We attribute 12 frequencies to pulsation. Model fits to the three frequencies with the highest amplitudes lie well outside the uncertainty box for the star's position in the HR diagram based on published values. Some of the frequencies appear to be rotationally split modes.} {The models suggest that either (1) the published estimate of the luminosity of \\hd, based on a relation between circumstellar disk radius and stellar luminosity, is too high and/or (2) additional physics such as mass accretion may be needed in our models to accurately fit both the observed frequencies and \\hd's position in the HR diagram.} ",
        "introduction": "A star evolving from the birthline to the zero-age main sequence (ZAMS) derives most of its energy (half of which heats the star and half of which radiates away) from the release of gravitational potential energy as the star collapses.  These pre-main sequence (PMS) objects are characterized by observational features typical for their early evolutionary stage, such as emission lines, infrared and/or ultraviolet excesses, (ir)regular brightness variations, etc.  The resulting variability of PMS stars is observed on a wide range of time scales and amplitudes.  Variations on time scales of weeks with amplitudes of a magnitude and more originate from the circumstellar environment that contains gas and dust remnants of the stellar birth cloud. Accretion or chromospheric activity induces variations on time scales from several hours to days. Oscillations with periods of few hours to half an hour and amplitudes at the millimagnitude level are due to $\\delta$-Scuti-like pulsations in those PMS stars with the right combination of mass, temperature and luminosity.  There are two distinct classes of PMS objects: T Tauri stars and Herbig Ae/Be stars. The low-mass ($<$ 1$M_{\\odot}$) T Tauri stars are too cool (spectral types from late F to M) to be unstable to $\\delta$-Scuti-type p-mode pulsation. The Herbig Be stars are too massive (4 to 10 $M_{\\odot}$) for an observable PMS phase: it is believed that by the time they become first visible optically, they are already burning hydrogen in their cores. The intermediate-mass (1.5 to 4 M$_{\\odot}$) Herbig Ae stars - on the other hand - cross the instability region in the Hertzsprung Russell (HR) diagram on their way to the main sequence. They have the right combination of mass, temperature and luminosity to be pulsationally unstable.  Several pulsating PMS stars have been discovered and analyzed within the last few years.  A summary of their properties is given by Zwintz (\\cite{zwi08}). The periods of PMS pulsators are in the domain of classical $\\delta$ Scuti stars, i.e., between 20 minutes and 6 hours, and their amplitudes lie at the millimagnitude level.  Pre- and post-main sequence stars of the same mass, effective temperature and luminosity mostly differ in their inner structures, but their atmosphere properties are quite similar (Marconi \\& Palla \\cite{mar98}). PMS stars lack regions of already processed nuclear material and seem not to be affected by strong rotation gradients which can complicate their inner structures.  Asteroseismology is the only method to probe the interiors of these stars, thanks to the fact that the excited pulsational instabilities depend strongly on the stellar density profile (Suran et al. \\cite{sur01}). Since the higher-order frequency spacings for pre- and post-main sequence stars are different (Suran et al. \\cite{sur01}), the evolutionary phase of a field star may be determined from its pulsational eigenspectrum. The first successful seismic models of PMS stars have already been obtained (e.g., Guenther et al. \\cite{gue07}; Zwintz et al. \\cite{zwi07}).  In this paper we discuss MOST spacebased observations of the pulsating Herbig Ae star \\hd\\,\\,and present a first asteroseismic model for this star. ",
        "conclusions": "MOST high-precision time series photometry of the Herbig Ae star \\hd\\,\\,obtained in 2006 and 2007 was used to investigate the pulsational properties of the star. This task was complicated by the fact that \\hd\\,\\,is surrounded by a dense circumstellar dust disk which causes irregular light variations introducing high-amplitude signal at low frequencies (i.e., from 0 to 3 $d^{-1}$). We were able to disentangle the pulsational variability with amplitudes at the millimagnitude level from the irregular variations with amplitudes of up to 200 mmag in the Fourier domain.  Although numerous peaks appear to be formally significant, we identify only 12 frequencies as pulsational in origin in the data sets of both years. Therefore, only these were used for the asteroseismic analysis.  For our model fitting, we further restricted the set of frequencies by excluding frequencies below 125 $\\mu$Hz. Also, based on their proximity in frequency space, we included only the central frequency from the {f5, f9, f12} triplet, and the most significant frequency from the {f7, f8} pair, since we suspect they are rotationally split modes.  We, thus, restricted the fits to the five modes {f1, f2, f3, f7, f9}.   We are able to obtain good fits ($\\chi^2 \\le 1$) to the frequencies for $l$ = 0, 1 and 2 p-modes. Significantly, though, all the model fits lie well outside the HR diagram uncertainty box for \\hd. Even when we fit only the three most significant frequencies (f1, f2, and f3), the models with $\\chi^2 \\le 1$ lie far away from the presumed location of the star in the HR diagram. Furthermore, we note that the fundamental p-mode frequencies in the vicinity of \\hd's location in the HR diagram are higher than the lowest frequencies observed.  Higher luminosity models, though, do have lower fundamental frequencies.  The failure to find a model that fits even the three most significant modes that at the same time lies within the uncertainty box for the star's location in the HR diagram leads us to suggest three areas to investigate further:  \\noindent 1. The models are too crude and need additional modeling physics such as mass accretion during the pre-deuterium burning phases to fit the frequencies and the star's location in the HR diagram simultaneously. The presence of the circumstellar disk supports this idea but the size of the effect on p-mode frequencies needs to be determined.  \\noindent 2. The third mode in the possibly rotationally split group {f7, f8} needs to be identified to help support the hypothesis that the modes are rotationally split. This would also help mode identification in that the modes would then have to be nonradial.  \\noindent 3. The luminosity of the star needs to be better constrained. This would rule out either the model fits or the p-mode nature of the lowest frequencies observed. If the star's luminosity is correct then the lowest observed frequencies are either non-intrinsic or they are g-modes. Note that for \\hd \\,\\,no HIPPARCOS (ESA \\cite{hip97}) parallax is available and the error of the parallax from the ASCC 2.5 Catalog (Kharchenko \\cite{kha01} Kharchenko et al. \\cite{kha05}) is a factor of three larger than the parallax itself."
    },
    "0812/0812.3971_arXiv.txt": {
        "abstract": "Whereas penumbral models during the last 15 years have been successful in explaining Evershed flows and magnetic field inclination variations in terms of flux tubes, the lack of contact between these models and a convective process needed to explain the penumbral radiative heat flux has been disturbing. We report on recent observational and theoretical evidence that challenge flux tube interpretations and conclude that the origin of penumbral filamentary structure is overturning convection. ",
        "introduction": "\\label{intro} Sunspot magnetic fields and dynamics have been studied scientifically for 100 years. Despite considerable progress during the last decade, a theoretical framework that explains sunspot fine structure, dynamics, magnetic fields and energy balance in a consistent manner is only now beginning to emerge. This situation can partly be attributed to the small horizontal scales associated with sunspot fine structure and the relatively poor spatial resolution achieved with spectropolarimetric observations. In addition, realistic numerical 3D MHD simulations of sunspots have only recently become possible.  During the last few years, there has been a remarkable improvement in the quality and diversity of observational data relevant to the understanding of sunspot fine structure, dynamics, magnetic fields and energy balance. In particular, high-spatial resolution observations from the Swedish 1-m Solar Telescope (SST) and the Solar Optical Telescope (SOT) on Hinode reveal new sunspot structure and flow patterns at odds with prevailing interpretations in terms of flux tube models. In addition, theoretical arguments as well as recent 3D MHD simulations of sunspot fine structure underline problems of these interpretations and lead to the conclusion that the origin of penumbral fine structure is overturning convection.  In the present review, we describe recent progress in our understanding of penumbral fine structure and put that in context with existing models. Rather than attempting to summarize the extensive literature on penumbrae, we discuss selected key papers and attempt to describe their interconnections and to critically review conclusions drawn. We also point out connections between observed penumbral fine structure and magnetic flux concentrations {\\em outside} sunspots, such as faculae. We hope to convince the reader that the new picture emerging is one of improved consistency as regards observations and theory of sunspot penumbrae in particular, but also with respect to umbral dots, light bridges and faculae. ",
        "conclusions": "We believe that there is now strong evidence to support the conclusion that penumbral fine structure should be interpreted as the result of overturning convection, as proposed by \\citet{2006A&A...447..343S}. Evidence for this conclusion comes from recent SOT/Hinode and SST observations \\citep{2007Sci...318.1597I, 2008A&A...488L..17Z}, showing {\\em vertical} flows of the right magnitude to explain the penumbral radiative heat flux. Recent numerical 3D MHD simulations \\citep{2007ApJ...669.1390H, 2008arXiv0808.3294R} reproduce fundamental properties of observed penumbrae and confirm the convective origin of penumbral filaments. The simulations show that the nature of this convection takes place in gaps with up to 2~Mm depth and that any roll-like convection \\citep{1961ApJ...134..289D, 2004ARA&A..42..517T, 2008arXiv0808.3294R}, if present, is of small importance \\citep{2008arXiv0808.3294R}. The Evershed flow is interpreted to be identical to the horizontal flow component of this convection \\citep{2008ApJ...677L.149S}. Such horizontal flows are necessary in order to cool hot upflows by radiation.  Neither observations nor simulations lead to the conclusion that this convection is nearly field-free, as suggested \\citep{2006A&A...447..343S}. However, inferred field strengths from spectropolarimetric data are obviously limited to layers above the photosphere, whereas simulations rely on numerical diffusivities to prevent instabilities at scales corresponding to the grid separation. Other uncertainties relate to the outer parts of penumbrae where observations show nearly horizontal field and even field lines dipping down into the photosphere \\citep{2008arXiv0806.4454B}. Simulations do not show structures of this type. It appears likely that downward pumping of magnetic field by convection outside the sunspot plays a role in the outer penumbra, as proposed earlier \\citep{2002AN....323..383T}, but we disagree strongly with the conclusion that this explains the origin of the filamentary structure of the penumbra. Downward pumping by the convection {\\em inside} the penumbra also must take place and this probably explains why observations \\citep{1997Natur.389...47W} show evidence of return flux well inside the outer penumbral boundary \\citep{2008ApJ...677L.149S}.  Penumbral filaments have been successfully interpreted in terms of embedded flux tubes during a period of 15 years. While we conclude that this interpretation is misleading in terms of underlying physics, there are several reasons why this model has been so successful. We have shown that the opening of radially aligned gaps with nearly field-free convecting gas leads to a magnetic field that is much more horizontal over the gaps, giving the illusion of a flux tube \\citep{2006A&A...447..343S}. There are also other arguments for expecting nearly horizontal fields in the penumbral atmosphere: Horizontal cooling flows are most efficient near optical depth unity and if this gas is magnetized, it will aid in producing nearly horizontal magnetic fields. Also, emerging flux simulations relevant to the quiet sun suggests that convection can lead to flux expulsion in the vertical direction in addition to the horizontal direction and that this explains quiet sun horizontal magnetic fields above the photosphere. Clear evidence of such `vertical' flux expulsion is however not seen in the penumbral part of the sunspot simulated by \\citet{2008arXiv0808.3294R}.  We emphasize the connections of the moving tube model to convective processes and to the radiative cooling of such flows near the photosphere in both types of models. The similarity of the magnetic field above convecting gaps and flux tubes add to the difficulties of correctly interpreting observations and distinguishing between models.  In the embedded flux tube models, the Evershed flow is at center stage and the mechanism for heating the penumbra remains obscure. The new view of penumbral fine structure as caused by overturning convection implies that the main driver of penumbra fine structure is the energy flux below the surface and that the Evershed flow is `only' a consequence of this convection \\citep{2008ApJ...677L.149S}.  We believe that siphon flow models \\citep{1997Natur.390..485M} are of little relevance for understanding penumbrae. These are linked to the idea that there are two distinct families of field lines: those associated with dark filaments and Evershed flows and those associated with bright filaments connecting to distant magnetic regions \\citep{2004ApJ...600.1073W, 2004ARA&A..42..517T}. As far as we know, there is no observational support for this `static' picture of penumbral magnetic fields. Observations suggest life times for penumbral filaments on the order of one hour associated with flow channels opening and closing continuously \\citep{2006ApJ...646..593R}. 3D MHD Simulations \\citep{2008arXiv0808.3294R, 2007ApJ...669.1390H} confirm the transient nature of azimuthal variations in field strength and inclination. Theoretical arguments and models \\citep{2006A&A...447..343S, 2006A&A...460..605S} as well as simulations clearly lead to the conclusion that the large variations in inclination across filaments are  {\\em local} perturbations, caused by {\\em penumbral} convection and vanishing a few hundred km above the penumbral photosphere.  As regards further progress in this rapidly evolving field, we expect that even more realistic 3D MHD simulations in the near future will further improve our understanding of penumbrae, in particular as regards their outermost parts. Observed unpolarized and polarized spectra at the highest possible spatial resolution are needed. Of particular importance is such spectra giving information about the layer immediately above the photosphere. Emphasis should be given to analyzing data at +/- 90~deg from the symmetry axis of sunspots located away from disk center, as was done by \\citet{2007Sci...318.1597I}. This is in part to allow analysis of flows perpendicular to the radial direction of the filaments, but also in order to see as deep into these structures as possible. Analysis of such data need to account for the pronounced 3D nature of these filaments, caused by strong azimuthal variations in the Wilson depression, as well as strong LOS variations in the magnetic field and flow velocity. A dilemma here is that the use of inversion techniques with many nodes along the LOS raises questions of uniqueness and difficulties in comparing the results of such inversions with simulations. Existing 3D MHD simulation data allow inversion techniques to be tested with synthetic Stokes spectra from penumbral atmospheres, as done already with simulations of small-scale flux concentrations outside sunspots \\citep{2007MmSAI..78..166K, 2007ApJ...662L..31O}. The effect of assuming e.g. hydrostatic equilibrium can be evaluated quantitatively.  Presently used inversion techniques process polarized spectra pixel by pixel without constraining, for example, the magnetic field to be divergence-free. The requirement of divergence--free magnetic fields is crucial in forcing field lines to bend around convecting gaps (and flux tubes), leading to strong gradients in field strength and inclination. With spectropolarimetric observations approaching a spatial resolution of 100~km \\citep{2008arXiv0806.1638S}, which is similar to the equivalent LOS resolution achieved with inversion techniques using a small number of nodes, it is reasonable to enforce magnetic fields constrained by div({\\bf B})=0. Stray-light corrections are with most inversion techniques implemented in an ad-hoc manner pixel by pixel whereas a physical stray-light implementation would employ a point spread function that does not vary, or varies slowly, across the FOV. Micro-- and macro-turbulence parameters are used as fudge parameters to compensate spatial smearing of unresolved structures. With improved spatial resolution, modeled LOS velocity gradients should eliminate the need for such parameters in the inversions. We expect future inversion techniques to develop as `global' techniques, in the sense of fitting model parameters for a large number of connected pixels simultaneously. This will allow constraints, such as div({\\bf B})=0, and physical straylight models to be incorporated in a consistent manner to further enhance the usefulness of inversion techniques for inferring the physical state of the atmospheres above sunspots and other magnetic structures."
    },
    "0812/0812.1974_arXiv.txt": {
        "abstract": "We  present the  preliminary  results  of a  50  ks long  XMM-Newton observation of the bright Z-source GX 340+0. In this Letter we focus on the study of a broad asymmetric emission line in the Fe K$\\alpha$ energy band, whose  shape is clearly resolved and  compatible with a relativistically smeared  profile arising  from reflection on  a hot accretion  disk extending  close  to the  central accreting  neutron star. By  combining temporal and  spectral analysis, we are  able to follow  the evolution  of the  source along  its  Horizontal Branch. However, despite a significant  change in the continuum emission and luminosity,  the line profile  does not  show any  strong correlated variation.  This  broad line is produced by  recombination of highly ionized iron (Fe XXV) at an inferred inner radius close to 13 R$_g$, while the fit requires a high  value for the outer disk radius.  The inclination of the  source is extremely well constrained  at 35 deg, while the emissivity index is -2.50. ",
        "introduction": "GX 340+0  is a  bright Low-Mass X-ray  Binary (LMXB) belonging  to the class of the Z sources \\citep{hasinger89}, and its inferred luminosity is close to the Eddington limit  (2 $\\times$ 10$^{38}$ erg/s for a 1.4 M$_{\\odot}$ NS).  The source has  also a radio counterpart, from which a distance  of 11  $\\pm$ 3 kpc  has been  estimated \\citep{penninx93}; radio emission from the  counterpart is, however, highly variable, and seems to be  correlated with the X-ray flux when the  source is in its Horizontal  Branch  (HB), but  anticorrelated  on  the other  branches \\citep{oosterbroek94}; it  can also  show periods of  radio quenching, becoming extremely faint \\citep{berendsen00}.  Temporal analysis  studies have shown a  complex phenomenology, linked to the  accretion state, with  characteristics typical of  the Z-class \\citep{jonker00}. The  power spectrum  shows a low-frequency  (tens of Hz)  quasi-periodic oscillation  when  the source  resides  on its  HB (called   Horizontal  Branch   Oscillation,  HBO)   while   at  higher frequencies  the   source  shows,  at  the  same   time,  twin  kiloHz Quasi-Periodic Oscillations (kHz QPOs), whose centroid frequencies are correlated with  the HBO peak  frequency.  Although no burst  has been yet observed, the  very fast timing variability (at  more than 800 Hz) and  the  similarities with  the  other  Z-sources  indicate that  the compact object is a neutron star.  The spectral properties of the source have not been fully investigated so far;  \\citet{schulz93} studied the 2--12 keV  spectrum using EXOSAT data; the spectrum  could be well described by  a single component due to  thermal  Comptonization of  soft  photons,  emerging  from the  NS surface,  in  a hot  corona  of  moderately  optical thickness  ($\\tau \\simeq$ 5--6);  \\citet{lavagetto04}  presented  the  first  broadband  (0.1--200  keV) spectrum of  the source  using data from  a BeppoSAX  observation. The spectrum could be decomposed into  the sum of a soft thermal component of temperature  of 0.5 keV, an optically  thick Comptonized component, and an excess at energies above  20 keV that they fitted with a simple power law.   A high resolution spectrum  of the source  was studied by \\citet{ueda05},  using  a Chandra  observation;  Chandra data  clearly showed the presence of an emission line, fitted with a simple Gaussian profile, at 6.57 keV with 40 eV equivalent width.  In this Letter  we report the results of  an XMM-Newton observation of GX  340+0.   Thanks  to  the  large collecting  area  of  the  EPIC-PN instrument and  its good  spectral resolution, we  find evidence  of a strong broad emission line at $\\sim$ 6.7 keV; the line profile is very well reproduce by a smeared disk-reflected profile; we find, moreover, evidence  ot other  disk reflection  signatures supporting  the common scenario, valid  for AGNs,  galactic black hole  and NS  systems, that broad iron  fluorescence lines are  produced by reflection  of photons from a  hot corona on  the surface of  an accretion disk  that extends very close to the compact object. ",
        "conclusions": "In this Letter we report on the first unambiguous determination of the broad disk  origin of the  iron line in  the Z-source GX  340+0. Broad asymmetric iron lines  present in the energy range  6.4--6.97 keV have been  discovered first  in AGNs,  and  later, in  galactic black  hole candidates,    and   very    recently    also   in    a   NS    system \\citep{disalvo05,bhattacharyya07,cackett08}.  The  common scenario for the  origin  of  such  lines   is  disk  reflection,  where  the  line deformation primarily  arises from the high  Keplerian velocity fields of the  disk reflecting matter  and from general  relativistic effects produced by the intense gravitational  well of the compact object. Our analysis  reveals,  for  the first  time  in  a  NS accreting  at  the Eddington limit that  the iron line shape in  GX 340+0 is unambigously described by a disk reflection profile, finding also evidence of other reflection  signatures (an  absorption edge  at $\\sim$  8.8 keV  and a broad Ca XIX emission line).  During this XMM-Newton observation  we can closely follow the spectral evolution of  GX 340+0 through all  of its horizontal  branch.  At its leftmost  point, corresponding  to  the  lowest count  rate  and to  a frequency  of  the  HBO  of   $\\sim$  17  Hz,  we  derive  the  lowest temperatures both for the disk  (1.60 keV) and the hard emission (2.55 keV); as the source moves towards  the hard apex, there is a continuum rising of the temperatures and  of the X-ray flux; assuming a distance of 11$\\pm$3  kpc and an inclination  of the disk as  inferred from the diskline profile,  we derive  a change in  the source  luminosity from (1.67$\\pm$0.1) $\\times$ 10$^{38}$ erg/s (Spectrum 1) to (2.57$\\pm$0.2) $\\times$  10$^{38}$  erg/s (Spectrum  5).  The  inner  disk radius  as derived from  the \\texttt{diskbb} normalization (see Table  1) gives a value in  the 8--11 km range;  considering that the  inner disk radius can be  underestimated by a  factor of two \\citep{merloni00},  we find that  these values are  in agreement  with the  ones derived  from the \\texttt{diskline} line profile.  Despite  this remarkable continuum  change, the  profile of  the broad iron diskline does not sensibly  change.  The rest frame energy of the \\texttt{diskline} component present in  the iron region indicates that the line  is most probably  produced by He-like  iron ions at  a short distance  from  the  compact   source.  In  Spectrum  5,  because  the inclination angle and the emissivity  of the line were bad constrained by  the fit,  we  kept them  frozen to  the  best values  of the  time averaged  spectrum; we  noted, therefore,  a small  shift of  the rest frame  line  energy towards  higher  energies,  possibly indicating  a greater contribution from the H-like ions.  The equivalent width of the  line decreases from Spectrum 1 ($\\sim$ 60 eV) to  Spectrum 5 ($\\sim$ 30  eV), indicating that the  line does not respond to the soft X-ray  flux, which, on the contrary increases, but presumably to the  hard X-ray emission, above 12  keV, which falls out from the  observed spectral coverage  of the EPIC/PN  instrument.  The inclination angle of  the disk and the emissivity  index are very well constrained by  our fits,  with relative uncertainties  of only  a few percent.   Contrary  to  what  found  in  galactic  black  holes,  the emissivity  index  is  not  very  steep, indicating  a  more  extended illuminating corona  above the disk; we  are not able  to identify any significant  change  in  the   value  of  this  parameter  during  the observation,  despite a  large  increase in  the  temperatures of  the continuum components and their related fluxes.  Recently  \\citet{laurent07} proposed an  alternative scenario  for the formation  of broad  iron lines  in LMXBs,  where extensive  red wings could be formed by recoil of line photons in an optically thick medium expanding,  or  converging,   at  relativistic  velocities.   Since  a spectral line  model, adapted to  this scenario, is not  yet available for fitting  X-ray data, we are  not able to test  this scenario using our data. However,  as pointed in \\citet{pandel08} for  the case of 4U 1636-536, we note that the lack of a narrow core and the presence of a blue wing in the iron line  profile of GX 340+0 is a strong indication of the disk reflection origin.  Most  theoretical models on  the QPO  generation in  NS LMXB  agree to identify the shortest dynamical timescales, i.e. the highest frequency QPO at  $\\sim$ 1 kHz,  with the frequency  of the Keplerian  motion of matter at the inner rim of  an accretion disk.  From the HBO frequency and the  relation known to  exist between this  and the upper  kHz QPO frequency \\citep{jonker00},  we can infer the  upper kHz QPO  to be in the 550--750 Hz frequency range. For a typical 1.4 M$_{\\odot}$ NS this frequency  range corresponds  to  an inner  disk  radius of  9.8--12.1 R$_g$.   The  \\texttt{diskline}  inferred  inner radii  for  the  time selected spectra are all fully  consistent within this range, and give support to the identification of  highest QPO frequency with the inner disk  Keplerian frequency.  Our fit  results seem  also to  indicate a decrease in  the value of  the inner disk  radius as the  source moves towards higher luminosities.  The  errors associated to this parameter are, however, still too large to claim a strict correlation and longer exposures are needed to assest the significance of such a relation.  The spectra show, together with  the broad Fe K$\\alpha$ emission line, other reflection signatures, like a  broad Ca XIX resonance line and a high-energy absorbtion  edge, compatible with the Fe  XXV K-edge (8.83 keV energy  in the laboratory frame).  The  extraordinary good quality of the  EPIC-PN spectrum of this  source will allow to  test even more recent  theoretical  line  profiles  and  fully  self-consistent  disk reflection models.  We aim at expanding this analysis in a forthcoming paper."
    },
    "0812/0812.1832_arXiv.txt": {
        "abstract": "Observations of \\water~masers towards the post-AGB star and water fountain source OH~009.1--0.4 were made as part of HOPS (The \\water~southern galactic Plane Survey), with the Mopra radiotelescope. Together with followup observations using the Australia Telescope Compact Array (ATCA), we have identified \\water~maser emission over a velocity spread of nearly 400\\kms~(--109 to +289\\,\\kms). This velocity spread appears to be the largest of any known maser source in our Galaxy. High resolution observations with the ATCA indicate the maser emission is confined to a region $0\\farcs3 \\times 0\\farcs3$ and shows weak evidence for a separation of the red- and blueshifted maser spots. We are unable to determine if the water fountain is projected along the line of sight, or is inclined, but either way OH~009.1--0.4 is an interesting source, worthy of followup observations. ",
        "introduction": "Water (\\water) maser emission is found throughout our Galaxy, pinpointing unusual physical conditions towards many different astrophysical objects, including high- and low-mass star forming regions, planetary nebulae \\citep{miranda01}, Mira variables \\citep{hinkle79} and asymptotic giant branch (AGB) stars \\citep{engels96}. They are also found outside our Galaxy, such as in the centres of active galaxies \\citep{claussen84}. In order to study the occurrence of \\water~masers in our Galaxy, a large-scale untargetted survey has begun, called HOPS (The \\water~southern galactic Plane Survey; Walsh et al. {\\em in preparation}). This survey will cover 90 square degrees of the southern Galactic plane, searching for, amongst other things, \\water~masers.\\\\  Planetary nebulae (PN) are the end-phase of evolution for intermediate-mass ($\\sim$~1-8 M$_\\odot$) Main Sequence stars. These stars spend a significant fraction of their lives ($\\sim$~10$^6$~yr) on the Asymptotic Giant Branch (AGB), a phase characterised by copious mass-loss which results in thick, dusty, molecular envelopes (see e.g. Herwig 2005 for a review of AGB evolution) with expansion velocities of order ~10-20~km~s$^{-1}$. In the O-rich case, the envelopes of these AGB stars can display emission from the masing species of SiO, OH and H$_2$O. The velocity profiles of these species (e.g. Deacon et al. 2004, 2007) are usually consistent with them having arisen in largely spherical shells around their host star (e.g. Cohen 1989); additionally, the velocity extent of the OH typically lies outside of and encloses that of the H$_2$O. Interferometric observations of these species reveal spatiokinematic distributions of maser spots that largely corroborate the suggested overall sphericity, sometimes with small deviations within the inner molecular envelopes (e.g. Bains et al. 2003). As the objects evolve past the tip of the AGB, it is surmised that a fast wind emanates from the exposed stellar core \\citep{kwok78} and gradually the maser species are extinguished due to the loss of coherence length in the disrupted shell (see e.g. van Winckel 2003 for a review of post-AGB evolution). It is this fast wind that is thought to amplify any existing asymmetries in the slower AGB envelope, leading to the morphological evolution of the object from largely spherical to the array of shapes displayed by PN (e.g. Balick \\& Frank 2002).\\\\  Recent imaging of post-AGB dust shells (e.g. Gledhill 2005) in scattered light has revealed a preponderance of axisymmetric structures, suggesting the window for turning on the shaping mechanism can be narrowed to the late AGB/early post-AGB phase. Further, the slowly-growing subclass of `water fountain' sources (first discovered and aptly named by Likkel \\& Morris 1988), with about 10 members \\citep{imai07}, can potentially elucidate the mystery of the shaping mechanism. Water fountains are bipolar post-AGB stars that display highly collimated and high velocity expanding regions (typically $v_{\\rm exp}$~$<$~100~km~s$^{-1}$) of \\water~maser emission. The spectral profile of the \\water~maser emission displays a larger velocity range of emission compared to that of OH, where present; maps of the sources reveal that the OH emission traces the remnant molecular shell while the \\water~emission is shock-excited in collimated bipolar jets \\citep{imai02}. Observations seem to confirm the theory that these high velocity jets are likely driven by magneto-hydrodynamical processes \\citep{vlemmings06}. However, whether the observed strong magnetic fields originate from single stars or through binary interactions is still unclear. Water fountains may therefore represent an extremely short-lived phase of post-AGB evolution that all such objects pass through, where much of the shaping takes place via the sculpting caused by the high velocity jets. Alternatively, they may be a special subset of stars, perhaps the more massive ones that evolve quicker and have sufficiently thick envelopes to shield clouds of \\water~vapour and preserve them to later be shocked into masing by the fast wind (Suarez et al. 2007). Either way, their discovery represents an exciting new addition to the study of the late stages of stellar evolution.\\\\  In this Letter, we report on the detection of a \\water~maser towards the young post-AGB star known as OH~009.1--0.4 and b292 \\citep{sevenster97}, which shows maser emission over the widest velocity range known in our Galaxy, confirming it as a new addition to the short list of known water fountains. The post-AGB star is associated with the IRAS source 18043--2116. It is unusual in that it exhibits 1665\\,MHz and 1612\\,MHz OH maser emission, but not at 1667\\,MHz \\citep{deacon04}. Furthermore, it is the only known post-AGB star to show 1720\\,MHz OH maser emission, which is thought to arise in the region where the AGB and starting planetary nebula winds collide \\citep{sevenster01}. Previous observations of \\water~maser emission towards OH~009.1--0.4 \\citep{deacon07} showed emission covering 210\\,\\kms, although the authors noted that there could be other maser features that were outside the observed velocity range.\\\\ ",
        "conclusions": "We have observed the water fountain OH~009.1--0.4 with both the Mopra and ATCA radiotelescopes. We find water maser spots covering an unprecedented velocity range of 398\\,\\kms~which we believe is the largest range of velocities for any known maser site in our Galaxy. Using the ATCA interferometer, we have partially resolved the water maser emission, which covers a projected area of $1900 \\times 1900$\\,AU. From our current observations, it is not clear if the water fountain is aligned along the line of sight, or if it is inclined. If it is aligned along the line of sight, then this orientation offers a unique perspective on water fountains with the jet pointed towards us. If it is inclined to the line of sight, then the three dimensional velocities will be even larger than measured here. Either way, OH~009.1--0.4 is certainly a unique and interesting water fountain, worthy of further investigation.\\\\  We suggest that followup observations OH~009.1--0.4 at higher spatial resolution, such as with VLBI, will be able to tell us more about the nature of this region.\\\\"
    },
    "0812/0812.2693_arXiv.txt": {
        "abstract": "{ Supersymmetric models with $t-b-\\tau$ Yukawa unification at $M_{GUT}$ qualitatively predict a sparticle mass spectrum including first and second generation scalars at the 3--15 TeV scale, third generation scalars at the (few) TeV scale and gluinos in the sub-TeV range. The neutralino relic density in these models typically turns out to lie far above the measured dark matter abundance, prompting the suggestion that instead dark matter is composed of an axion/axino mixture. We explore the axion and thermal and non-thermal axino dark matter abundance in Yukawa-unified SUSY models. We find in this scenario that {\\it i}).\\ rather large values of Peccei-Quinn symmetry breaking scale $f_a\\sim 10^{12}$ GeV are favored and {\\it ii}).\\ rather large values of GUT scale scalar masses $\\sim 10-15$ TeV allow for the re-heat temperature $T_R$ of the universe to be $T_R\\agt 10^6$ GeV. This allows in turn a solution to the gravitino/Big Bang Nucleosynthesis problem while also allowing for baryogenesis via non-thermal leptogenesis. The large scalar masses for Yukawa-unified models are also favored by data on $b\\to s\\gamma$ and $B_s\\to \\mu^+\\mu^-$ decay. Testable consequences from this scenario include a variety of robust LHC signatures, a possible axion detection at axion search experiments, but null results from direct and indirect WIMP search experiments. } ",
        "introduction": "\\label{sec:intro}  The celebrated unification of gauge couplings at scale $M_{GUT}\\simeq 2\\times 10^{16}$ GeV under minimal supersymmetric standard model (MSSM) renormalization group evolution\\cite{gaugeunif} strongly suggests that nature is described by some sort of supersymmetric grand unified theory (SUSY GUT) model at very high energy scales. While the GUT gauge group $SU(5)$\\cite{su5} has many compelling and beautiful properties, the gauge group $SO(10)$ has an even greater appeal\\cite{so10}, and also some indirect experimental support in terms of how well see-saw neutrino mass fits into the general scheme\\cite{seesaw}. While both $SU(5)$ and $SO(10)$ SUSY GUT theories admit gauge coupling unification, $SO(10)$ theories also yield {\\it matter unification}, in that all superfields of a single Standard Model (SM) generation-- plus a SM gauge singlet superfield $\\hat{N}^c$ containing a right-hand neutrino state-- fit neatly into the 16-dimensional spinorial representation $\\hat{\\psi}_{16}$ of $SO(10)$.  In the simplest $SO(10)$ SUSY GUT models, the MSSM Higgs superfields-- $\\hat{H}_u$ and $\\hat{H}_d$-- both live in the fundamental representation $\\hat{\\phi}_{10}$. In these models, the superpotential has the form \\be \\hat{f}\\ni f\\hat{\\psi}_{16}\\hat{\\psi}_{16}\\hat{\\phi}_{10} +\\cdots \\ee where the dots represent model-dependent terms which, for instance, might include further Higgs fields responsible for the GUT gauge symmetry breaking. The coupling $f$ represents the {\\it unified} Yukawa coupling of each generation: thus, just as $SU(5)$ models often predict $f_b-f_\\tau$ unification, simple $SO(10)$ SUSY GUT models predict the more restrictive $t-b-\\tau$ Yukawa coupling unification at the GUT scale\\cite{old}.  It is probably fair to say that at this moment no compelling $SO(10)$ SUSY GUT model yet exists. Models based in four spacetime dimensions require large, unwieldy Higgs representations to break the $SO(10)$ GUT symmetry. Newer models formulated in extra spacetime dimensions are able to do away with the large Higgs reps and break the GUT symmetry via compactification of the extra dimensions\\cite{exdimguts}. In our approach here, we will adopt a pragmatic view, assuming that the MSSM (or MSSM plus gauge singlets) is the correct effective theory describing physics between the weak and GUT scales, and we will explore some of the consequences of requiring the three third generation Yukawa couplings to have a high degree of unification at $M_{GUT}$, as suggested by simple $SO(10)$ SUSY GUT models.  In our calculation to check third generation Yukawa coupling unification, we begin with the measured gauge couplings and third generation fermion masses, and use renormalization group methods to evolve the coupled gauge and Yukawa couplings up to the GUT scale. The calculation ends up being sensitive to the entire SUSY particle mass spectrum through weak scale threshold corrections involved in transitioning between the SM and MSSM effective field theories\\cite{hrs}. The parameter space of the model is given by \\be m_{16},\\ m_{10},\\ M_D^2,\\ m_{1/2},\\ A_0,\\ \\tan\\beta, \\ {\\rm sign}(\\mu ) \\label{eq:pspace} \\ee where $m_{16}$ is the GUT scale mass of all matter scalars, $m_{10}$ is the GUT scale mass of Higgs scalars, $M_D$ is their D-term value, $m_{1/2}$ is the (unified) GUT scale gaugino mass, $A_0$ is the unified GUT scale SSB trilinear term, $\\tan\\beta\\equiv v_u/v_d$ is the weak scale ratio of Higgs field vevs, and $\\mu$ is the superpotential Higgs bilinear term, whose magnitude--but not sign--is determined by the scalar potential minimization conditions.  In practice, the two Higgs field soft breaking terms-- $m_{H_u}^2$ and $m_{H_d}^2$-- cannot be degenerate at $M_{GUT}$ and still allow for an appropriate radiative breakdown of electroweak symmetry (REWSB). Effectively, $m_{H_u}^2$ must be less than $m_{H_d}^2$ at $M_{GUT}$ in order to give $m_{H_u}^2$ a head start in running towards negative values at $M_{weak}$. We parametrize the Higgs splitting as $m_{H_{u,d}}^2=m_{10}^2\\mp 2M_D^2$ in accord with nomenclature for $D$-term splitting to scalar masses when a gauge symmetry undergoes a breaking which reduces the rank of the gauge group. A $D$-term splitting should apply to matter scalar SSB terms as well; in practice, better Yukawa unification is found when the splitting is only applied to the Higgs SSB terms. Such a GUT scale Higgs mass splitting might arise via GUT scale threshold corrections\\cite{bdr2}.  In previous work, the above parameter space was scanned over (via random scans\\cite{bf,abbbft} and also by more efficient Markov Chain Monte Carlo (MCMC) scans\\cite{bkss}) to search for Yukawa unified solutions using the Isasugra subprogram of \\Isajet\\cite{isajet} for sparticle mass computations. The quantity \\be R=\\frac{max(f_t,f_b,f_\\tau )}{min(f_t,f_b,f_\\tau )}\\ \\ \\ ({\\rm evaluated\\ at\\ Q=M_{GUT}}), \\label{eq:R} \\ee was examined, where solutions with $R\\simeq 1$ gave apparent Yukawa coupling unification. For superpotential Higgs mass parameter $\\mu >0$ (as favored by $(g-2)_\\mu$ measurements), Yukawa unified solutions with $R\\sim 1$ were found but only for {\\it special} choices of GUT scale boundary conditions\\cite{bf,bdr1,bdr2,abbbft,drrr,bkss}: \\be m_{16}\\sim 3-15~{\\rm TeV},~ A_0\\sim -2m_{16}\\,,~ m_{10}\\sim 1.2 m_{16}\\,,~ m_{1/2}\\ll m_{16}\\,,~ \\tan\\beta \\sim 50\\,. \\label{eq:HSpar} \\ee Models with this sort of boundary conditions were derived even earlier in the context of inverted scalar mass hierarchy models (IMH) which attempt to reconcile suppression of flavor-changing and $CP$-violating processes with naturalness via multi-TeV first/second generation and sub-TeV scale third generation scalars\\cite{bfpz}. The Yukawa-unified spectral solutions were thus found in Refs.~\\cite{abbbft,bkss} to occur with the above peculiar choice of boundary conditions as long as $m_{16}$ was in the multi-TeV regime.   Based on the above work\\cite{abbbft,bkss}, the sparticle mass spectra from Yukawa-unified SUSY models are characterized qualitatively by the following conditions: \\begin{itemize} \\item first and second generation scalars have masses in the $3-15$ TeV regime, \\item third generation scalars, $\\mu$ and $m_A$ have masses % in the TeV to few TeV regime (owing to the inverted scalar mass hierarchy), \\item gauginos, including the gluino, have sub-TeV masses, \\item the lightest neutralino $\\tz_1$ is nearly pure bino with mass typically $m_{\\tz_1}\\sim 50-80$ GeV. \\end{itemize} The presence of a bino-like $\\tz_1$ along with multi-TeV scalars gives rise to a neutralino cold dark matter (CDM) relic abundance that is typically in the range $\\Omega_{\\tz_1}h^2\\sim 10-10^4$, {\\it i.e.} far above\\cite{auto} the WMAP measured\\cite{wmap5} value \\be \\Omega_{CDM}h^2=0.110\\pm 0.006\\ \\ \\ \\ \\ \\ee by several orders of magnitude.  One solution to the CDM problem in Yukawa-unified SUSY models occurs for $m_{16}\\sim 3$~TeV, where one can have $m_{\\tz_1}\\simeq m_h/2$. In this case, the $\\tz_1$ would pair-annihilate in the early universe through the light Higgs $h$ resonance at a sufficient rate to obtain the desired relic density\\cite{bkss}, and could hence be the stable lightest SUSY particle (LSP). However, since $m_{16}$ is relatively low, Yukawa coupling unification only occurs at the $R\\sim 1.09$ level.\\footnote{Another possibility, having $m_A$ light enough that $\\tz_1\\tz_1$ can annihilate through the $A$ resonance, appears to be excluded because these cases violate limits on $BF(B_s\\to\\mu^+\\mu^- )$ decay\\cite{bkss}.}  Another very compelling way out of the Yukawa-unified dark matter abundance problem occurs if one makes the additional assumption that the strong $CP$ problem is solved by the Peccei-Quinn mechanism\\cite{pq}, which leads to the presence of a light pseudoscalar particle, the axion~$a$~\\cite{ww,axreview}. Since we are working in a supersymmetric theory, the axion occurs as part of an axion supermultiplet\\cite{nillesraby}, which contains not only the axion, but a spin-0 saxion (with mass of order the weak scale), and a spin-$1\\over 2$ {\\it axino} $\\ta$. The axino is $R$-parity odd. While the saxion is expected to have a mass of order the SUSY breaking scale, the axino mass is very model dependent,  and can lie anywhere in the keV-GeV range\\cite{axmass,ckkr}. The axino then can serve as the  LSP instead of the lightest neutralino\\cite{wilczek,steff_rev}.  In the case of an axino LSP, the supposed neutralino relic abundance is greatly reduced since $\\tz_1\\to \\ta\\gamma$ decay can occur with a lifetime in the range of $\\sim 10^{-5}-10^{1}$ sec (depending on parameters). This decay time is sufficiently short that late-time neutralino decay to axino in the early universe should not upset successful predictions of Big Bang Nucleosynthesis (BBN)\\cite{ckkr}. The neutralino abundance then gets reduced by the ratio $m_{\\ta}/m_{\\tz_1}$ which can be in the range $10^{-1}-10^{-4}$. The $\\ta$ coming from $\\tz_1$ decay would actually constitute {\\it warm} dark matter as long as $m_{\\ta}\\alt 1$ GeV\\cite{jlm}. However, axinos can also be produced thermally in the early universe, and these would consititute CDM as long as $m_{\\ta}\\agt 100$ keV\\cite{ckkr,steffen}. Thus, in the axino LSP case, we could have a mixed dark matter (DM) scenario\\cite{julien} with \\bi \\item thermally produced cold or warm axino DM, \\item an admixture of warm axino DM arising from $\\tz_1\\to\\ta\\gamma$ decays and \\item a possibly large presence of cold {\\it axion} DM. \\ei  The axino LSP scenario turns out to be even more compelling cosmologically than just as a means to reconcile the dark matter relic abundance with Yukawa-unified models. In this class of Yukawa-unified solutions with $m_{16}$ in the multi-TeV range, we expect from supergravity theory that scalar SSB terms should be directly related to the gravitino mass $m_{3/2}$, and so we also expect the gravitino $\\tG$ to lie in the multi-TeV range. The cosmological gravitino problem-- wherein gravitinos produced thermally in the early universe suffer a late-time decay, thus destroying the successful predictions of BBN-- can be avoided. For $m_{3/2}\\alt 5$ TeV, the re-heat temperature $T_R$ must be $T_R\\alt 10^5$ GeV\\cite{kohri}, thus creating tension with most viable mechanisms for baryogenesis\\cite{buchm}. However, for $m_{3/2}\\agt 5$ TeV, the re-heat bound is much higher: $T_R\\alt 10^8-10^9$ GeV. This range of $T_R$ is too low for thermal leptogenesis (which requires $T_R\\agt 10^{10}$ GeV)\\cite{buchm}, but is exactly what is needed for baryogenesis via {\\it non-thermal} leptogenesis\\cite{ntlepto}, wherein the heavy right-hand neutrino states are not produced thermally, but rather via inflaton decay. It was pointed out in Ref. \\cite{bs} that this is also the exact range needed to generate a dominantly {\\it cold} axino DM universe. Thus, the whole scenario fits together to offer a consistent cosmological picture of BBN, non-thermal leptogenesis and CDM composed of axions and/or axinos, and one solves the strong $CP$ problem to boot\\cite{bs}!  In this paper, we make a detailed study of the cosmological consequences of Yukawa-unified SUSY GUT models with an axino LSP. The initial study proposed in Ref.~\\cite{bs} assumed only a negligible component of axion dark matter. In this study, we now fold in the axion contribution to the dark matter abundance. In addition, we have updated the value of $m_t$ in our calculations, and refined the Yukawa coupling 1-loop beta function threshold effects in \\Isajet . We also incorporate in this study the results of our Markov Chain Monte Carlo (MCMC) approach to finding Yukawa-unified solutions. Whereas the Yukawa-unified solutions found via a random scan in Ref.~\\cite{bs} had $m_{16}\\sim 15-20$ TeV, the more efficient MCMC scans used here are able to find many solutions for $m_{16}$ values as low as $3-10$~TeV.  The paper is organized as follows. In Sec.~\\ref{sec:mass}, we update our sparticle mass predictions for Yukawa-unified SUSY models from \\Isajet\\ using 1. an improved beta-function threshold decoupling, 2. an updated value of the top mass $m_t=172.6$ GeV, and 3. Markov Chain Monte Carlo (MCMC) scans of the parameter space. We also exhibit plots of the $b\\to s\\gamma$ and $B_s\\to\\mu^+\\mu^-$ branching fractions versus $m_{16}$ and find these favor $m_{16}$ values $\\agt 10$ TeV. Moreover, we perform MCMC scans for Yukawa-unified solutions with an alternative spectrum generator, \\Softsusy, and compare these results with those gained from \\Isajet. In Sec.~\\ref{sec:dm}, we review elements of axion and axino dark matter cosmology, including plots of the axion relic abundance, axino relic abundance and neutralino lifetime. In Sec.~\\ref{sec:axdm}, we discuss the gravitino/BBN problem, non-thermal leptogenesis via inflaton decay, and mixed axion/axino dark matter. Our main findings are located in Sec.~\\ref{ssec:results}. Here, we calculate all three components of axion/axino dark matter in Yukawa-unified models for $m_{16}$ values of 5, 8, 10 and 15 TeV. We explore scenarios wherein axions constitute either most, or hardly any, of the dark matter abundance. Our results are presented in the $m_{\\ta}$ vs. $T_R$ plane. We find that models with a rather large value of the Peccei-Quinn breaking scale $f_a\\sim 10^{12}$ GeV are favored, as well as models with $m_{16}\\sim m_{3/2}$ on the high side: $\\sim 10-15$ TeV. In these cases, the re-heat temperature of the universe can range above $10^6$ GeV, allowing for a solution to the BBN gravitino problem as well as allowing for non-thermal leptogenesis. Our conclusions are presented in Sec.~\\ref{sec:conclude}. In an Appendix, we list updated Yukawa-unified benchmark points from \\Isajet\\ 7.79. ",
        "conclusions": "\\label{sec:conclude}  One vestige of supersymmetric $SO(10)$ grand unified theories may be that the third generation $t-b-\\tau$ Yukawa couplings unify, in addition to gauge couplings and matter multiplets. Assuming the MSSM is the low energy effective theory at energy scales $Q<M_{GUT}$, we are able to find parameter space solutions that yield a superparticle mass spectrum with Yukawa coupling unification good to 5\\% or better, using the \\Isajet\\,7.79 and \\Softsusy\\,2.0.18 programs. The sparticle mass spectrum that qualitatively emerges is that first and second generation scalars lie in the multi-TeV regime, third generation scalars, $\\mu$ and $m_A$ lie in the few TeV range, and gauginos lie in the sub-TeV range. The neutralino relic density turns out to be $10-10^4$ times the measured dark matter density, prompting the suggestion that in this case, the axino is a better LSP candidate, so that the dark matter of the universe would be composed of an axion/axino mix.  Yukawa-unified SUSY models should thus give rise to {\\it three} components to the dark matter density: axion dark matter produced via vacuum mis-alignment at the QCD phase transition, non-thermally produced axinos from $\\tz_1\\to\\ta\\gamma$ decay (likely warm dark matter if $m_{\\ta}\\alt 1$ GeV), and thermally produced axinos which are likely cold dark matter unless $m_{\\ta}\\alt 100$ keV. We compute the abundance of all three components of dark matter in Yukawa-unified models for four different scenarios containing either dominant axion or dominant axino cold dark matter. The relative abundances depend on $\\Omega_{\\tz_1}h^2$, $m_{\\tz_1}$, $m_{\\ta}$, $f_a/N$ and $T_R$. It is found that for all solutions with $m_{16}\\sim m_{3/2}\\agt 5$ TeV, the value of $T_R$ needed is below limits calculated from BBN constraints, thus solving the gravitino BBN problem.  We also find that it is very difficult for models with PQ symmetry breaking scale $f_a$ lower than $\\sim 2\\times 10^{11}$ GeV (case C1) to generate dominantly cold dark matter and a sufficiently large $T_R$ to be consistent with at least non-thermal leptogenesis, for any allowed value of $m_{16}>5$ TeV. However, if the PQ scale is large enough ($\\agt 4\\times 10^{11}$ GeV), then we can generate a universe with axions as the dominant component of CDM (cases C2 and C3), and only a smaller component is composed of possibly warm axinos with $m_{\\ta}\\alt 10^{-3}$ GeV. These solutions work out if $m_{16}\\agt 10$ TeV, so that the higher side of the $m_{16}\\sim 3-15$ TeV range is preferred. This is also in accord with results presented in Sec. \\ref{ssec:bsg}, where it is found that $BF(b\\to s\\gamma )$ and $BF(B_s\\to\\mu^+\\mu^- )$ decay also prefer $m_{16}\\agt 10$ TeV. Finally, we show a case C4 with accidentally small axion CDM component, but which gives cold thermally produced axino dark matter, as long as $m_{\\ta}\\agt 100$ keV. This scenario allows for solutions with $m_{16}$ as low as 5  TeV.  As far as tests go of the Yukawa-unified SUSY models with mixed axion/axino dark matter, it is clear from previous work that very large signal rates from gluino pair production and subsequent cascade decays should be visible at LHC relatively soon after start-up\\cite{so10lhc}. The $\\tz_1$ produced in cascade decay events will still yield events with $\\eslt$, since the $\\tz_1\\to \\ta\\gamma$ decay occurs far outside the detector. In cases C2 and C3, there is a good chance for direct detection of relic axion dark matter at experiments such as ADMX\\cite{admx}. However, direct and indirect searches for WIMP dark matter would likely turn up null results.  \\appendix"
    },
    "0812/0812.4189_arXiv.txt": {
        "abstract": "We present timing and spectral analysis of RXTE-PCA observations of the accretion powered pulsar 4U 1907+09 between June 2007 and August 2008. 4U 1907+09 had been in a spin-down episode with a spin-down rate of $-3.54\\times10^{-14}$ Hz s$^{-1}$ before 1999. From RXTE observations after March 2001, the source showed a $\\sim 60$\\% decrease in spin-down magnitude and INTEGRAL observations after March 2003 showed that source started to spin-up. We found that the source recently entered a new spin-down episode with a spin-down rate of $-3.59 \\times 10^{-14}$ Hz s$^{-1}$. This spin-down rate is pretty close to the previous long term spin-down rate of the source measured before 1999. From the spectral analysis, we showed that Hydrogen column density varies with the orbital phase.  {\\bf{Keywords:}} X-rays: binaries, pulsars:individual:4U 1907+09, stars:neutron, accretion, accretion discs ",
        "introduction": "4U 1907+09 is a High Mass X-ray Binary (HMXB) system discovered by the Uhuru survey (Giacconi et al. 1971). The system consists of an X-ray pulsar accreting mass from a blue supergiant companion star. An orbital period of 8.38 days was found from Ariel V observations (Marshall \\& Ricketts 1980). Makishima et al. (1984) found the spin period value of 437.5 s using Tenma observations. Eccentricity of the orbit of the system was found to be $\\sim$0.16 from EXOSAT observations (Cook \\& Page 1987). From RXTE observations, In't Zand et al. (1998) found revised orbital parameters of the system. The evidence of two phase-locked flares at the folded orbital profile was reported (Marshall \\& Ricketts, 1980; In't Zand et al. 1998). The presence of the two flares per orbit leaded to the suggestion of an equatorial disk like envelope around a Be type stellar companion (Cook \\& Page 1987). Although  the variation in H$\\alpha$ emission in the optical observations strengthened the suggestion of a Be star (Iye 1986), Cox et al. (2005) suggested the possibility of an OB supergiant companion previously mentioned by Schwartz et al. (1980) and Van Kerkwijk et al. (1989). Cox et al. (2005) identified  the companion as an O8-09 Ia supergiant with a distance of $\\sim$5 kpc.  4U 1907+09 had been steadily spinning down for more than 15 years with an average rate of $\\dot \\nu=-3.54\\times 10^{-14}$ Hz s$^{-1}$ (Cook \\& Page 1987; in't Zand et al. 1998; Baykal et al. 2001; Mukerjee et al. 2001). Afterwards, RXTE observations in 2001 showed $\\sim$60\\% decrease in the magnitude of the spin down rate (Baykal et al. 2006). It has recently been reported that the pulse period evolution has shown a torque reversal with $2.58\\times 10^{-14}$ Hz s$^{-1}$, confirming a spin up episode after March 2003 (Fritz et al. 2006).  Spectral studies (Schwartz et al. 1980; Marshall \\& Ricketts 1980; Makishima et al. 1984; Cook \\& Page 1987; Chitnis et al. 1993; in't Zand et al. 1997; Coburn et al. 2002) on 4U 1907+09 indicate a highly variable hydrogen column density over the binary orbit, between $\\sim1\\times10^{22}$ cm$^{-2}$ and $\\sim9\\times10^{22}$ cm$^{-2}$. The power law photon index is between $\\sim$0.8 and $\\sim$1.5 with a high energy cutoff around $\\sim$13keV. Extensive energy spectra generated from Ginga (Mihara 1995; Makishima et al. 1999) and BeppoSAX (Cusumano et al. 1998) observations exhibit absorption features at $\\sim$19keV and $\\sim$39keV referring to the fundamental and the second harmonic of a cyclotron absorption feature corresponding to a surface magnetic field strength $2.1\\times10^{12}$ G.  In this paper, we present timing and spectral analysis of RXTE monitoring observations of 4U 1907+09 between  MJD 54280 and MJD 54690. First, we provide a brief description of observation and analysis procedure in Section 2. The pulse timing analysis, indicating a new spin-down episode, is reported in Section 3. In Section 4, we present the results of X-ray spectral analysis. ",
        "conclusions": "After more than 15 years long spin-down episode with an almost constant spin-down rate of  $\\dot \\nu=-3.54(2)\\times 10^{-14}$ Hz s$^{-1}$ (Baykal et al.2001), RXTE observations between March 2001 and March 2002 showed a $\\sim 0.6$ factor of decrease in spin down rate compared to long-term spin down rate of the source. From the observations of INTEGRAL between March 2003 and September 2005, Fritz et al. (2006) found that source was spinning-up with a rate of $2.58\\times 10^{-14}$ Hz s$^{-1}$. We found, in this work, that the source began to spin-down again with a spin-down rate of $-3.59(2) \\times 10^{-14}$ Hz s$^{-1}$ which is pretty close to the previous long term spin-down rate of the source.  According to Ghosh\\& Lamb (1979) model, net torque exerted on the neutron star by the accreting matter from a prograde accretion disk can be expressed as  \\begin{equation} N=I\\dot{\\Omega}=n(\\omega_s)\\dot{M}l, \\end{equation}  where I is the moment of inertia of the neutron star, $\\dot{\\Omega}$ is the spin frequency derivative, $n(\\omega_s)$ is the dimensionless torque, $\\dot{M}$ is mass accretion rate and $l$ is specific angular momentum of the accreting matter. Here, $\\omega_s$ is the fastness parameter defined as the ratio of neutron star's spin frequency to the Keplerian frequency at the magnetospheric radius, which is almost equal to the accretion disk's inner radius.  Dimensionless torque, $n(\\omega_s)$ in Equation 2, is the ratio of the total (magnetic and material torque) to the material torque. It is, for a slowly rotating neutron star ($\\omega_s<0.35-0.95$), expected to be positive and of the order of unity leading to spin up of the star, whereas it is expected to be negative for a fast rotating neutron star ($\\omega_s>>0.35-0.95$) leading to spin down of the star (Ghosh\\&Lamb 1979; Wang 1995; Li\\&Wang 1996, 1999) . Since the dimensionless torque depends on the fastness parameter (and thus the magnetospheric radius), it decreases and may even be negative with an increase in magnetic field or decrease in mass accretion rate. Decrease in mass accretion rate leads to a decrease in X-ray flux and due to the accretion geometry changes, X-ray spectral variations may also be evident.  For 4U 1907+09, observed torque reversal cannot simply be explained by Ghosh\\&Lamb (1979) model since spin-up episode observed by INTEGRAL and spin-down episodes with different spin-down rates have not shown significant X-ray flux and X-ray spectral variations (Baykal et al. 2001; Fritz et al. 2006; Baykal et al. 2006). Using this model, it is also difficult to explain similar spin rate magnitudes observed in spin-down and spin-up episodes of 4U 1907+09. Absolute value of the spin rate has always been found to be  $\\sim 2-3\\times 10^{-14}$ Hz s$^{-1}$.  In case of accretion via stellar wind, simulations showed that transitions from spin-down to spin-up is possible even when there is not significant mass accretion rate change (Anzer et al. 1987;Taam\\&Fryxell 1988a, 1988b, 1989; Blondin et al. 1990, Murray et al. 1998). However, torque reversal timescale of 4U 1907+09 is of the order of years which is much greater than the typical timescale of torque reversals in the simulations, found to be only of the order of hours. Moreover, 4U 1907+09 has shown transient QPOs (quasi-periodic oscillations) which is a sign of the presence of an accretion disk (in't Zand et al. 1998a; Mukerjee et al. 2001; Baykal et al. 2006).  It may also be possible that negative torques (corresponding to spin-down episodes) may come from a retrograde Keplerian accretion disk (Nelson et al. 1997). Spin-down torques may also be the result of an advection dominated sub-Keplerian disk for which the fastness parameter should be higher than that of the rest of the Keplerian disk causing a net spin-down (Yi et al. 1997) or the warping of the disk so that the inner disk is tilted by more than 90 degrees (van Kerkwijk et al. 1998). Torque and X-ray luminosity correlation is expected from these models which is not found for 4U 1907+09.  A recent model of Perna et al. (2006) explains torque changes in accreting neutron stars without a need for retrograde disks. In this model, the scenario in which the neutron star's magnetic field is tilted with respect to the axis of rotation of the neutron star is considered. For an accretion disk located in the neutron star's equatorial plane, the magnetic field strength depends on the azimuthal angle and thus the location of magnetospheric radius is variable. This can lead to regions in the disk where the propeller effect is locally at work, while accretion from other regions is possible. Perna et al. (2006) shows that this accretion geometry may cause cyclic torque reversal episodes without a need for significant mass accretion rate variations.  4U 1907+09 has an eccentric orbit and shows occasional X-ray flux dips (in't Zand et al. 1997). These are indications of the transient accretion disk formation around the neutron star. On the other hand, model of Perna et al. (2006) is proposed for persistent prograde accretion disks like Ghosh\\&Lamb (1979) model. To understand torque reversal of 4U 1907+09, transient nature of the accretion should also be taken into account.  From Figure 4, orbital dependence of $n_H$ is evident, on the other hand there is no significant variation of power law index. $n_H$ increases just after the periastron due to the accretion flow at the periastron passage and it remains at high values until the apastron. Roberts et al. (2001) also found similar $n_H$ dependence on orbital phase and interpreted this $n_H$ excess as the neutron star passing through an equatorially enhanced matter envelope in an inclined orbit.  {\\bf"
    },
    "0812/0812.1848_arXiv.txt": {
        "abstract": "We use numerical simulations of turbulent, multiphase, self-gravitating gas orbiting in the disks of model galaxies to study the relationships among pressure, the vertical distribution of gas, and the relative proportions of dense and diffuse gas.  A common assumption is that the interstellar medium (ISM) is in vertical hydrostatic equilibrium.  We show that the disk height and mean midplane pressure in our multiphase, turbulent simulations are indeed consistent with effective hydrostatic equilibrium, provided that the turbulent contribution to the vertical velocity dispersion and the gas self-gravity are included. Although vertical hydrostatic equilibrium gives a good estimate for the mean midplane pressure $\\langle P\\rangle_{\\rm midplane}$, this does not represent the pressure experienced by most of the ISM.  Mass-weighted mean pressures $\\langle P \\rangle_{\\rho}$ are typically an order of magnitude higher than $\\langle P\\rangle_{\\rm midplane}$ because self-gravity concentrates gas and increases the pressure in individual clouds without raising the ambient pressure.  We also investigate the ratio $R_{\\rm mol}=M_{\\rm H_2}/M_{\\rm HI}$ for our hydrodynamic simulations. \\citet{2006ApJ...650..933B} showed that $R_{\\rm mol}$ is proportional to the estimated midplane pressure in a number of systems.  We find that for model series in which the epicyclic frequency $\\kappa$ and gas surface density $\\Sigma$ vary together as $\\kappa \\propto \\Sigma$, we recover the empirical relation.  For other model series in which $\\kappa$ and $\\Sigma$ are varied independently, the midplane pressure (or $\\Sigma$) and $R_{\\rm mol}$ are not well correlated.  We conclude that the molecular fraction -- and hence the star formation rate -- of a galactic disk inherently depends on its rotational state, not just the local values of $\\Sigma$ and the stellar density $\\rho_{\\ast}$.  The empirical result $R_{\\rm mol} \\propto \\langle P\\rangle_{\\rm midplane}$ implies that the three ``environmental parameters'' $\\kappa$, $\\Sigma$, and $\\rho_{\\ast}$ are interdependent in real galaxies, presumably as a consequence of evolution: real galaxies trend toward states with Toomre $Q$ parameter near unity. Finally, we note that $R_{\\rm mol}$ in static comparison models far exceeds both the values in our turbulent hydrodynamic simulations and observed values of $R_{\\rm mol}$, when $\\Sigma > 10 \\msun \\ \\pc^{-2}$, indicating that incorporation of turbulence is crucial to obtaining a realistic molecular fraction in numerical models of the ISM. ",
        "introduction": "All phases of the interstellar medium (ISM) are turbulent, and this turbulence has many effects.  In the astrophysical literature, turbulence is often treated as yielding a simple addition to the thermal pressure, $P_{\\rm total}= \\rho(c_s^2 + v^2_{\\rm turb})$, where $v^2_{\\rm turb}$ is the dispersion in the (one-dimensional) turbulent velocity, and $c_s^2=P/\\rho= f\\kB T/\\mu$ for gas with a total number density $f n$ and mass density $\\mu n$.  This approach is often adopted when analyzing the stratification of interstellar gas clouds and the ISM as a whole, with the combined pressure gradients taken to balance the gravitational force per unit volume such that hydrostatic equilibrium is maintained by the total pressure.  The turbulent pressure is believed to be especially important in the cold components of the ISM, for which observed linewidths far exceed the values of $c_s$ inferred from excitation of atomic and molecular lines.  Models of effective hydrostatic equilibrium in the vertical direction, usually assuming the turbulent and thermal velocity dispersions are constants independent of height, are often applied to observations of the large-scale Galactic ISM, and to observations of the ISM in external galaxies (e.g. \\citealt{1991ApJ...382..182L,1994ApJ...433..687M,1995ApJ...448..138M, 1997A&A...326..554C,2000MNRAS.311..361O, 2002A&A...394...89N,2004ApJ...608..189D, 2004ApJ...612L..29B,2006ApJ...650..933B, 2008AstL...34..152K}).  For example, \\cite{2002A&A...394...89N} showed that the observed atomic and molecular disk thicknesses in the Milky Way can be fit well by assuming effective hydrostatic equilibrium, and accounting for both the gas self-gravity and the external gravitational potential of stars and dark matter.  \\cite{2004ApJ...612L..29B} and \\cite{2006ApJ...650..933B} (hereafter BR06) used a simplified approach to hydrostatic equilibrium in order to estimate the midplane gas pressure in a sample of disk galaxies, adopting a single velocity dispersion for the gas, treating the gravitational potential as dominated by the stars, and assuming the stellar disk's scale height is independent of radius.  \\cite{2008AstL...34..152K} extended the analysis of BR06 but instead of adopting a constant scale height for the stellar disk, they assumed that the velocity dispersion for the stars is consistent with a state of marginal gravitational instability (with Toomre parameter $Q_{\\ast}=1.5$) for the corresponding stellar surface density.  They then assumed hydrostatic equilibrium for all (gaseous and stellar) components separately, and computed the self-consistent midplane pressure, finding differences of order $30-40\\%$ from the simplified BR06 approach.  Although widely adopted, the effective hydrostatic equilibrium model for the large-scale ISM has not, to our knowledge, been explicitly verified using actual turbulent flows.  One of the goals of this work is to test this formulation systematically, using the solutions of time-dependent numerical hydrodynamic simulations of turbulent, multiphase gas.  In addition to providing support against gravity, pressure also affects the phase balance in the ISM.  For a static system at a given mean density $\\bar n$, changing the pressure alters the proportions of mass divided between dense clouds and diffuse intercloud medium; e.g. for cold and warm components in pressure equilibrium, the mass ratio of cold to warm gas is $M_{\\rm cold}/M_{\\rm warm}=[\\bar n/n_{\\rm warm} - 1]/[1 - \\bar n/n_{\\rm cold}]= [\\bar n k T_{\\rm warm} - P]/[P - \\bar n k T_{\\rm cold}]$. The mean density itself, however, depends on pressure through the condition of vertical hydrostatic equilibrium. Turbulent pressure, as it affects the response to external and self-gravity, can be expected to change both the mean density and the mass fractions of dense and diffuse gas.  Here, we investigate these effects quantitatively.  The fraction of ISM mass in dense gas is important from the point of view of galactic evolution, since this component is the immediate precursor to star formation.  A recent observational study of external disk galaxies by BR06 identified a linear relationship between the mean ratio of molecular-to-atomic mass, $R_{\\rm mol}$, and an estimate for the total midplane pressure $\\propto \\sqrt{\\rho_{\\ast}} \\Sigma$, where $\\rho_{\\ast}$ is the stellar volume density and $\\Sigma$ is the total gaseous surface density.  BR06 propose that the molecular fractions in widely-varying types of galaxies -- and hence their respective star formation efficiencies -- are therefore determined essentially by a single parameter, the midplane pressure.  To investigate this proposal, we use multiphase turbulence simulations in which we independently vary the input galactic ``environmental'' parameters. The observational study of BR06 focused on the dependence of $R_{\\rm mol}$ on $\\rho_{\\ast}$ and $\\Sigma$, but another important -- and independent -- environmental parameter is the angular rotation rate $\\Omega$ (and the associated epicyclic frequency $\\kappa^2= R^{-3}d (\\Omega^2 R^4)/dR$).  Using our data sets from turbulent simulations, we compare the pressure estimate of BR06 to the true value of the pressure, and also test how $R_{\\rm mol}$ relates to the mean pressure measured in two different ways.  We note that a number of recent numerical studies have investigated the formation of ISM structures with internal densities reaching those similar to Giant Molecular Clouds (GMCs).  Some studies \\citep[e.g.][]{2002ApJ...564L..97K,2005ApJ...633L.113H,2006ApJ...648.1052H,2007A&A...465..431H,2007A&A...465..445H,2008A&A...486L..43H,2007ApJ...657..870V} have focused on how this may occur as a consequence of the collision of large-scale high-velocity flows that shock and cool, becoming turbulent at the same time.  Other studies have focused on the ability of self-gravitating instabilities to induce converging flows over sufficiently large scales that massive, high-column density structures similar to observed GMCs are created \\citep[e.g.][]{2001ApJ...559...70K,2007ApJ...660.1232K, 2005ApJ...620L..19L,2006ApJ...639..879L}; these models include the galactic shear and rotation that are important on these large scales, and in some cases also include magnetic effects \\citep[e.g.][]{2002ApJ...581.1080K,2003ApJ...599.1157K}.  As spiral arms are observed to be strongly associated with high molecular fractions and star formation, some studies have focused on the interaction between large-scale spiral shocks and self-gravity in inducing GMC formation \\citep[e.g.][]{2002ApJ...570..132K,2006ApJ...646..213K,2008MNRAS.tmp.1249D}. The details of conversion from diffuse to dense gas by cooling downstream from spiral shock fronts has also recently been studied in the absence of self-gravity \\citep[e.g.][]{2008MNRAS.389.1097D,2008ApJ...681.1148K}.  Taken together, these and other recent studies have shown that significant quantities of dense gas form naturally as a result of large-scale ISM dynamical processes.  Of course, dense gas in the ISM is also returned to the diffuse phases by the energetic inputs from star formation.  In the present work, by incorporating feedback, we are able to evolve our models until a quasi-steady state is reached.  This enables an analysis of the correlations among statistical properties of the system, in terms of their influence on the fraction of dense gas when the system has reach a quasi-steady state of cloud formation and destruction.   This paper is organized as follows: In \\S 2 we briefly summarize our numerical methods.  The specification of model parameters and the results of statistical analysis in comparison to the vertical-equilibrium approximation are presented in \\S 3.  In \\S 4, we discuss the molecular fraction and investigate how it relates to the ISM pressure in our models.  We summerize our results and discuss implications for ISM structure and evolution in \\S 5. ",
        "conclusions": "We have used numerical simulations of turbulent, multiphase, self-gravitating gas orbiting in the disks of model galaxies to study the relationships among pressure, the vertical distribution of gas, and the relative proportions of dense and diffuse gas.  In particular, we compare the results on vertical stratification obtained from space-time averages of fully-dynamic -- and often turbulence-dominated -- systems with simple estimates based on single-component effective hydrostatic equilibria.  We also investigate how vertical-equilibrium estimates for the pressure compare with measured mean values of the pressure in our models.  Empirical studies by BR06 have identified a linear relation between the molecular-to-atomic mass ratio $R_{\\rm mol}$ and a midplane ISM pressure estimate, $P_{\\rm BR} \\propto \\Sigma \\sqrt{ \\rho_{\\ast}}$.  We study the origin and implications of this relation by testing the correlations among $R_{\\rm mol}$, $P_{\\rm BR}$, and the directly-measured midplane and mean pressures in our models.   Our chief conclusions, and their implications, are as follows:   1. The average disk scale height is well represented by estimates that assume hydrostatic equilibrium and an effective total pressure based on the total (thermal + turbulent) vertical velocity dispersion (see Fig. \\ref{fig:Have} and eq. \\ref{Hest_eq}).  Thus, provided that gas surface densities, vertical velocity dispersions, and stellar density can be measured, an accurate estimate for the disk thickness can be obtained.  Hydrostatic equilibrium with an effective turbulent pressure is commonly assumed in both Galactic and extragalactic observational studies (e.g. \\citealt{1991ApJ...382..182L,1994ApJ...433..687M,1995ApJ...448..138M, 1997A&A...326..554C,2000MNRAS.311..361O, 2002A&A...394...89N,2004ApJ...608..189D, 2004ApJ...612L..29B,2006ApJ...650..933B, 2008AstL...34..152K}), but to our knowledge the relations that are generally adopted have not previously been tested with direct numerical simulations.  Our hydrodynamic studies demonstrate that for determining the scale height $H$, the effective hydrostatic equilibrium assumption is indeed sufficient, even when turbulent support far exceeds thermal support (and provided that magnetic effects are sub-dominant; see below).  Thus, measured disk thicknesses in edge-on disk galaxies could in principle be used to determine the unobservable vertical velocity dispersion, and measured line-of-sight velocity dispersions in face-on galaxies could be used to determine the unobservable disk thickness.   The basic reason the hydrostatic formula can be used to obtain an accurate measure of $H$ is that what is really being equated is the total vertical momentum flux $\\rho (\\kB T/\\mu + v_z^2)$ averaged over the midplane, and the total vertical weight of the ISM, $\\int dz\\, \\rho\\, g_z \\sim \\rho 4 \\pi G (\\rho + \\rho_{\\ast}) H^2$, averaged over the horizontal direction. Provided that the time-averaged value of the momentum per unit volume in the midplane does not change, momentum conservation including gravitational source terms demands that the difference between vertical momentum flux and vertical weight must be zero, independent of details of the dynamics.  The formula $H^2\\approx \\sigma_z^2/[4\\pi G (\\rho_{\\ast} + \\Sigma/H\\sqrt{2\\pi})]$ is therefore fundamentally an expression of momentum conservation.   2. Mass-weighted mean pressures $\\langle P \\rangle_{\\rho}$ in our hydrodynamic models significantly differ from the mean midplane pressure $\\langle P\\rangle_{\\rm midplane}$, while these quantities are quite similar to each other in our static comparison models. Typically, the hydrodynamic models yield values of $\\langle P \\rangle_{\\rho}$ an order of magnitude larger than $\\langle P\\rangle_{\\rm midplane}$. The difference can be attributed to self-gravitating condensation, which makes concentrated clouds with high internal pressure rather than a horizontally-uniform gas distribution with more moderate pressure.  Simple estimates of the pressure based on vertical hydrostatic equilibrium fall between mass-weighted and midplane values, with the formula used by BR06 (see our eq. \\ref{PBR_eq}) comparable to the geometric mean $P_{\\rm BR} \\sim \\sqrt{\\langle P \\rangle_{\\rho} \\langle P\\rangle_{\\rm midplane}}$. A single-component estimate for the midplane thermal pressure that accounts for self-gravity and the mean thermal and turbulent velocity dispersions (see eq. \\ref{eq:Pnewth}) follows the measured midplane pressure fairly closely, especially at high $\\Sigma$.  Thus, if turbulent and thermal vertical velocity dispersions can be measured directly (for face-on galaxies), a good estimate of the midplane total or thermal pressure can be computed via equation (\\ref{eq:Pnew}) or (\\ref{eq:Pnewth}).  For an edge-on system in which the scale height is measured, the midplane total or thermal pressure can be estimated as $\\Sigma/(H \\sqrt{2\\pi})\\times \\langle \\sigma_z^2\\rangle \\ {\\rm or}\\ \\langle c_s^2\\rangle$.  Midplane pressure estimates based on large-scale observables that assume hydrostatic equilibrium can be quite accurate, but this depends on an accurate measure of the vertical velocity dispersion or vertical thickness.  Even if the velocity dispersion is not known, the {\\it relative} midplane pressures of different regions within a galaxy (or from one galaxy to another) can be obtained using the hydrostatic formulae, provided the variation in the (unknown) velocity dispersion within the observational sample is small compared to the variation in the stellar volume and gaseous surface densities.  Midplane pressure estimates made in this way should not, however, be treated as a proxy for the pressure in the typical mass element, $\\langle P \\rangle_{\\rho}$, which can be much larger than the pressure in the typical volume element.       3. Based on calculations of molecular abundance as a function of hydrogen volume density $n$ and column density $N$ combined with the resolution and measured virial ratios in our simulations, we adopt a working definition of gas at $n\\ge 100 \\cm^{-3}$ as ``molecular'' and $n<100 \\cm^{-3}$ as ``atomic''.  We then investigate the ratio $R_{\\rm mol}=M_{\\rm H_2}/M_{\\rm HI}$ for all our models.  We find that Series Q and R, which have rotation rate $\\Omega \\propto \\Sigma$, show correlations between $R_{\\rm mol}$ and $P_{\\rm BR}$ (or $R_{\\rm mol}$ and $\\Sigma$) that are similar to the empirical result reported by BR06, $R_{\\rm mol}\\propto P_{\\rm BR} \\propto \\Sigma\\sqrt{\\rho_{\\ast}}$. On the other hand, Series K and S, in which $\\Sigma$ and $\\Omega$ do not vary together, depart from the empirical relation $R_{\\rm mol} \\propto P_{\\rm BR}$.  We conclude that (i) the molecule fraction inherently depends on the rotational state of a galactic disk, not just on the local values of the stellar volume and gaseous surface densities $\\rho_{\\ast}$ and $\\Sigma$; and (ii) the empirical relation $R_{\\rm mol} \\propto P_{\\rm BR}$ identified by BR06 implies that the third ``environmental parameter,'' the epicyclic frequency $\\kappa = \\sqrt{2} \\Omega$ (assuming a flat rotation curve), is not independent of $\\rho_{\\ast}$ and $\\Sigma$ in real galaxies.  This dependence can be accomplished by evolution: for example, disk galaxies may convert gas into stars until they reach a state in which the Toomre parameter $\\propto \\kappa/\\Sigma$ approaches a critical value.  4. We have tested the correlation between $R_{\\rm mol}$ and the measured pressures in our models, $\\langle P \\rangle_{\\rm midplane}$ and $\\langle P \\rangle_\\rho$, and find a good correlation in all Series only for the latter.  The correlation between $\\langle P \\rangle_\\rho$ and $R_{\\rm mol}$ is potentially useful as a way to estimate the typical internal pressure within gravitationally-bound regions when only the total molecular-to-atomic mass is easily accessible, as for low-resolution observations.  This internal pressure is important in the small-scale aspects of star formation such as determining the IMF \\citep{2007ARA&A..45..565M}, as well as in molecular chemistry.  The lack of correlation between $R_{\\rm mol}$ and $\\langle P \\rangle_{\\rm midplane}$ in Series K implies that the molecular content cannot in general be predicted solely from $\\Sigma$ and $\\rho_{\\ast}$ (i.e. without knowledge of $\\kappa$), as noted above. This reflects the fact that the formation of self-gravitating clouds is regulated not just by gravitational processes and pressure, but also by angular momentum.   5.  For our non-turbulent comparison models, we find that $R_{\\rm mol}$ far exceeds observed values.  This indicates that turbulence is essential to setting the observed phase balance in the ISM. Recent theoretical investigations of the origin of Kennicutt-Schmidt laws have focused on the dependence of star formation rates on the molecular, rather than total, gas surface density (e.g. \\citealt {2007arXiv0711.1361N,2008ApJ...680.1083R}).  Since turbulence is crucial in determining the abundance of dense gas, in simulations that aim to compute this abundance realistically it is necessary to incorporate the feedback effects that drive turbulence, and to run on a fine enough mesh (or with sufficient SPH particles) that the turbulence is well resolved. While technically challenging in global disk models, local models may offer a more immediate route to this goal.   {\\it Caveats --} The models analyzed in this paper are subject to a number of limitations, which could potentially affect some of our conclusions. The chief limitations of the simulations are that (i) they are two-dimensional, representing cuts in the $R-z$ plane, rather than three-dimensional; (ii) we have adopted a very simple model to implement turbulent driving as a star formation feedback effect from HII regions, and we have not included other drivers of turbulence such as supernovae, spiral shocks, and shear instabilities; (iii) we have not included magnetic fields (or cosmic rays). We intend to pursue these extensions in future work.  Inclusion of magnetic fields and altered turbulent driving would certainly affect the specific quantitative findings for $H_{\\rm ave}$, $\\langle P \\rangle_{\\rm midplane}$, $\\langle P \\rangle_\\rho$, and $R_{\\rm mol}$ in our models.  We believe, however, that the results we have emphasized regarding physical {\\it relationships} are robust.  In particular, with appropriate modifications to include magnetic stresses, the time-averaged vertical momentum flux through the midplane must still equal the time-averaged vertical weight if the mean vertical momentum is conserved.  This can be used to predict the total midplane pressure (including the magnetic pressure) and $H$ given the values of $\\Sigma$, $\\rho_{\\ast}$, and the thermal, turbulent, and Alfv\\'en velocities. Thus, we anticipate that inclusion of magnetic fields and alternate turbulence sources would not fundamentally alter the conclusion that reasonable estimates of scale heights can be made using observable quantities even in highly-dynamic systems.    Further, we expect that our conclusions regarding the presence or absence of correlations between $R_{\\rm mol}$ and $\\langle P \\rangle_\\rho$ or $\\langle P \\rangle_{\\rm midplane}$ would continue to hold in models that include additional turbulence sources and magnetic fields, although the details of correlations might change.  Namely, angular momentum inherently must be important in permitting or preventing formation of dense, self-gravitating clouds.  Our present models account for angular momentum effects, and show that $R_{\\rm mol}$ does not in general have a one-to-one relationship with $\\langle P \\rangle_{\\rm midplane}$ or $\\Sigma \\sqrt{\\rho_{\\ast}}$; we expect this result would carry over into any model that incorporates sheared background rotation of the galactic disk.  Thus, if a one-to-one relationship between $R_{\\rm mol}$ and $\\Sigma \\sqrt{\\rho_{\\ast}}$ indeed exists empirically, it implies that $\\Sigma$, $\\rho_{\\ast}$, and $\\kappa$ are not all independent quantities in real galaxies.    We are grateful to the referee for a number of comments that have helped improve our presentation.  Numerical computations used in this project were carried out on the OIT High Performance Computing Cluster, and the CTC cluster in the Department of Astronomy, at the University of Maryland.  This work was supported by grant NNG05GG43G from NASA."
    },
    "0812/0812.1079_arXiv.txt": {
        "abstract": " ",
        "introduction": "Inflation is an important concept in cosmology, and is supported by recent satellite-based measurements. The present authors have studied the effect of the length of inflation and pre-inflation on the power spectrum and the angular power spectrum [1], and it has been shown that the suppression of the spectrum at $l=2$ as indicated by Wilkinson Microwave Anisotropy Probe (WMAP) data [2] may be explained to a certain extent by the finite length of inflation for inflation of 50-60 $e$-folds [3]. The length of inflation is generally believed to be very long, but the physically interesting period of inflation is that in the last 60 $e$-folds. Although a correct inflationary potential based on superstring or supergravity theory has yet to be established, it could be assumed that the potential term in this latter period of inflation can be expressed in a simple form, such as a small-field or large-field model, or a hybrid model. In the present paper, the time-dependent behavior of the inflaton field during the last 100 $e$-folds in inflation is investigated using a simple inflation model. Specifically, numerical calculations are performed using a chaotic inflation model ($\\phi^2$ model), and the effects of the time dependence of the inflaton field on cosmological and inflationary parameters such as slow-roll parameters, the Hubble parameter, and the spectral index are investigated. The $k$-dependence of the power spectrum of the curvature perturbation has been investigated using Taylor expansion or the numerical method [4-7]. Here, using the derived parameters, i.e., the slow-rolled parameters, Hubble parameter and  the inflaton field  which have the time dependence, the power spectrum is calculated with the numerical method and the Bessel approximation, and it is compared with the power spectra derived from the familiar methods. ",
        "conclusions": ""
    },
    "0812/0812.2882_arXiv.txt": {
        "abstract": "Infrared Dark Clouds (IRDCs) are cold, dense regions of giant molecular clouds that are opaque at wavelengths $\\sim 10\\:{\\rm \\mu m}$ or more and thus appear dark against the diffuse Galactic background emission. They are thought to be the progenitors of massive stars and star clusters.  We use 8~$\\rm \\mu m$ imaging data from {\\it Spitzer} GLIMPSE to make extinction maps of 10 IRDCs, selected to be relatively nearby and massive. The extinction mapping technique requires construction of a model of the Galactic IR background intensity behind the cloud, which is achieved by correcting for foreground emission and then interpolating from the surrounding regions. The correction for foreground emission can be quite large, up to $\\sim 50\\%$ for clouds at $\\sim 5\\:{\\rm kpc}$ distance, thus restricting the utility of this technique to relatively nearby clouds. We investigate three methods for the interpolation, finding systematic differences at about the 10\\% level, which, for fiducial dust models, corresponds to a mass surface density $\\Sigma = 0.013\\:{\\rm g\\:cm^{-2}}$, above which we conclude this extinction mapping technique attains validity. We examine the probability distribution function of $\\Sigma$ in IRDCs. From a qualitative comparison with numerical simulations of astrophysical turbulence, many clouds appear to have relatively narrow distributions suggesting relatively low ($<5$) Mach numbers and/or dynamically strong magnetic fields. Given cloud kinematic distances, we derive cloud masses. Rathborne, Jackson \\& Simon identified cores within the clouds and measured their masses via mm dust emission. For 43 cores, we compare these mass estimates with those derived from our extinction mapping, finding good agreement: typically factors of $\\lesssim 2$ difference for individual cores and an average systematic offset of $\\lesssim 10\\%$ for the adopted fiducial assumptions of each method. We find tentative evidence for a systematic variation of these mass ratios as a function of core density, which is consistent with models of ice mantle formation on dust grains and subsequent grain growth by coagulation, and/or with a temperature decrease in the densest cores. ",
        "introduction": "\\label{S:intro}  Large fractions of stars form in clusters from the densest {\\it clumps} within giant molecular clouds (GMCs) (Lada \\& Lada 2003). These regions are also responsible for the birth of essentially all massive stars (de Wit et al. 2005). It is possible that our own solar system formed in such a region near a massive star (Hester et al. 2004), so the process of massive star and star cluster formation may be directly involved in our own origins. Understanding the formation of star clusters is also important as a foundation for understanding global galactic star formation rates (Kennicutt 1998) and thus the evolution of galaxies.  In spite of this importance, there are many gaps in our knowledge of how massive stars and star clusters are formed. For massive star formation there is a basic debate about whether the process is simply a scaled-up version of low-mass star formation from gas cores (Shu, Adams, Lizano 1987), albeit requiring the high pressures found in the centers of star-forming clumps (McKee \\& Tan 2003), or whether a qualitatively different mechanism is involved, such as stellar collisions (Bonnell, Bate, Zinnecker 1998) or competitive Bondi-Hoyle accretion (Bonnell et al. 2001; Bonnell, Vine, \\& Bate 2004). A prediction of the scenario involving formation from cores is the presence of massive, gravitationally-bound starless cores as initial conditions.  For star cluster formation there is a debate about whether the {\\it protocluster} (or {\\it star-forming clump}) is in quasi-virial equilibrium (Tan, Krumholz, \\& McKee 2006) or is undergoing rapid global collapse (Elmegreen 2000, 2007; Hartmann \\& Burkert 2007). This corresponds to a debate about the timescale of star cluster formation: does it take many or just a few free-fall times?  To help resolve these issues we require knowledge about the initial conditions of massive star and star cluster formation. The star-forming clumps that are the sites of these processes have typically been identified from the radiative emission and activity of their young stars: e.g. radio emission from ultracompact \\ion{H}{2} regions created by massive stars or maser emission from hot, dense gas near massive protostars. Unfortunately by the time this activity signposts the region, it has typically also erased much of the memory of the initial conditions from the system. This is especially true of protostellar outflows, which have a mechanical power large enough to significantly stir the gas of the star-forming clump and halt any large scale gravitational collapse (Nakamura \\& Li 2007).  The initial conditions for massive star and star cluster formation are expected to be dense, cold cores and clumps of gas. In some formation models there may be of order hundreds of solar masses of gas compressed to within a few tenths of a parsec. For a spherical cloud the mean mass surface density, $\\Sigma$, in this case is \\beq \\label{Sigma} \\Sigma \\equiv \\frac{M_c}{\\pi R_c^2} = 0.665 \\frac{M_c}{100\\sm}\\left(\\frac{R_c}{0.1{\\rm pc}}\\right)^{-2}\\:{\\rm g\\:cm^{-2}}. \\eeq For reference, $\\Sigma=1\\:{\\rm g\\:cm^{-2}}$ corresponds to $4800\\:M_\\odot\\:{\\rm pc^{-2}}$, $N_{\\rm H}=4.3\\times 10^{23}\\:{\\rm cm^{-2}}$ and $A_V\\simeq 230$~mag, for local diffuse ISM dust properties (e.g. Draine 2003). The extension of the extinction law into the mid-infrared (MIR) is somewhat controversial and uncertain (Lutz et al. 1996; Draine 2003; Indebetouw et al. 2005; Rom\\'an-Z\\'u\\~niga et al. 2007), but nevertheless, such column densities are expected to correspond to several magnitudes of extinction at 8~$\\rm \\mu m$.  Such Infrared Dark Clouds (IRDCs) have been identified in images of the Galactic plane from the {\\it Infrared Space Observatory} (ISO) (P\\'erault et al. 1996), the {\\it Midcourse Space Experiment} (MSX) (Egan et al. 1998). Simon et al. (2006a) identified about 10,000 potential IRDCs from intensity contrast features in the MSX survey. Simon et al. (2006b) identified the $^{13}$CO emission from about 300 of the darkest of these in the Galactic Ring Survey (GRS) (Jackson et al. 2006), thus deriving their kinematic distances. Rathborne, Jackson, \\& Simon (2006) surveyed the 1.2~mm dust continuum emission from 38 of these IRDCs, identifying 140 mm-emission cores within these clouds.  The temperatures in IRDCs are measured to be $\\lesssim 20$~K (Carey et al. 1998). High deuteration fractions have been reported by Pillai et al. (2007). Under these conditions one expects high depletion of volatiles onto ice mantles of dust grains (Dalgarno \\& Lepp 1984). Larger molecular line surveys of IRDCs have been carried out by, for example, Ragan et al. (2006) and Sakai et al. (2008).  In this paper we present a method of extinction mapping of IRDCs using the diffuse Galactic IR emission observed by {\\it Spitzer} as the background source. This method complements that based on measuring the extinction to individual stars (e.g. Rom\\'an-Z\\'u\\~niga et al. 2007), by providing a measurement of the extinction at the location of almost every pixel in the cloud image, including at very high column densities, thus allowing a detailed investigation of the cloud structure. ",
        "conclusions": "We have presented a new method of diffuse mid-IR ($\\rm 8 \\mu$m) extinction mapping of infrared dark clouds to probe the initial conditions for massive star and star cluster formation. The technique can naturally probe to higher mass surface densities than NIR extinction mapping using background stars and provides an estimate of the mass surface density at scales that are only limited by the resolution of the mid-IR image (e.g. a couple of arcseconds for {\\it Spitzer} IRAC). However, the method involves a spatial interpolation and smoothing of the diffuse background, which inevitably introduces additional uncertainties. In particular it is a relative mass surface density of the cloud compared to the region used to ascertain the background that is derived. Also, small scale background fluctuations are not captured and local regions containing mid-IR bright sources cannot be treated. One of the largest uncertainties for more distant clouds is the correction for foreground mid-IR emission from hot dust, and so we have generally restricted the analysis to relatively nearby clouds.  We have examined three different methods of estimating the diffuse background behind IRDCs. The Large-scale Median Filter (LMF) method is the simplest by not requiring information about the cloud size or location, but then must smooth over relatively large scales. The Small-scale Median Filter (SMF) achieves higher spatial resolution of the background given defined cloud boundaries, and we generally regard this method as being more accurate where such information is available. Finally, for very thin, filamentary clouds, more detailed background fitting can be attempted by considering perpendicular strips across the filament. Comparing these methods on a sample of 10~IRDCs we find systematic differences due to background fitting at about the 10\\% level. This means that the mid-IR extinction mapping technique becomes unreliable for $\\Sigma\\lesssim 0.013\\:{\\rm g\\:cm^{-2}}$ (given our adopted fiducial 8~$\\rm \\mu$m [{\\it Spitzer} IRAC Band 4 and diffuse Galactic background spectrum weighted] opacity of $7.5\\:{\\rm cm^{2}\\:g^{-1}}$).  We have then used this method to measure the mass surface densities of the 10 IRDCs. The probability distribution functions of $\\Sigma$ for these clouds show a range of shapes, that is qualitatively consistent with that expected in numerical simulations of turbulence, though the relatively narrow distributions of the observed clouds indicate that the Mach numbers are probably $\\lesssim 5$ and/or that magnetic fields are dynamically important. A more quantitative comparison of the observed clouds with numerical models, and a search for evolutionary trends associated with the development of self-gravitating structures is planned for a future study.  Given IRDC kinematic distances, we then estimate masses of the clouds and, more accurately, of their dense cores. Comparing to mass estimates from mm dust continuum emission (Rathborne et al. 2006) we find very good agreement: a $\\lesssim 10\\%$ systematic offset for the population of 43 studied cores, and a dispersion of less than a factor of two in the distribution.  Finally, we examine the ratio of masses estimated by mm dust emission to those found from mid-IR dust extinction as a function of core density, tentatively finding a trend which can be explained as being due to opacity changes due to ice mantle formation and grain growth or by a temperature decrease of about 5~K in the densest cores. Future studies of the mid-IR extinction law in these cores and molecular line studies of their chemistry can help to test the reality of this result and distinguish between these explanations."
    },
    "0812/0812.5081.txt": {
        "abstract": "In the paper the closed Friedmann--Robertson--Walker model with quantization in the presence of the positive cosmological constant and radiation is studied. For analysis of tunneling probability for birth of an asymptotically deSitter, inflationary Universe as a function of the radiation energy a new definition of a ``free'' wave propagating inside strong fields is proposed. On such a basis, tunneling boundary condition is corrected, penetrability and reflection concerning to the barrier are calculated in fully quantum stationary approach. For the first time non-zero interference between the incident and reflected waves has been taken into account which turns out to play important role inside cosmological potentials and could be explained by non-locality of barriers in quantum mechanics. Inside whole region of energy of radiation the tunneling probability for the birth of the inflationary Universe is found to be close to its value obtained in semiclassical approach. The reflection from the barrier is determined for the first time (which is essentially differs on 1 at the energy of radiation close to the barrier height). The proposed method could be easily generalized on the cosmological models with the barriers of arbitrary shape, that has been demonstrated for the FRW-model with included Chaplygin gas. Result is stable for variations of the studied barriers, accuracy are found to be 11--18 digits for all coefficients and energies below the barrier height. ",
        "introduction": "}  In order to understand what really happens in the formation of the Universe, many people came to the point of view that a quantum consideration of this process is deeper. The first papers with the quantum approach for the description of Universe formation and its initial expansion may be~\\cite{DeWitt.1967,Wheeler.1968}, and shortly afterwards many other papers appeared in this field, pointing to a rapid development of the quantum approach in cosmology (for example, see the first some papers \\cite{Vilenkin.1982.PLB,Hartle.1983.PRD,Linde.1984.LNC,Zeldovich.1984.LNC,Rubakov.1984.PLB,Vilenkin.1984.PRD,Vilenkin.1986.PRD,Atkatz.1984.PRD} and some discussions in Refs.~\\cite{Vilenkin.1994.PRD,Rubakov.1999} with references therein).  Today, among all variety of models one can select two approaches which are the most prevailing: these are the Feynman formalism of path integrals in multidimensional spacetime, developed by the Cambridge group and other researchers, called the \\emph{``Hartle--Hawking method''} (for example, see Ref.~\\cite{Hartle.1983.PRD}), and a method based on direct consideration of tunneling in 4-dimensional Euclidian spacetime, called the \\emph{``Vilenkin method''} (for example, see Refs.~\\cite{Vilenkin.1982.PLB,Vilenkin.1984.PRD,Vilenkin.1986.PRD,Vilenkin.1994.PRD}). Here, according to Ref.~\\cite{Vilenkin.1995}, in the quantum approach we have the following picture of the Universe creation: a closed Universe with a small size is formed from ``nothing'' (vacuum), where by the word ``nothing'' one refers to a quantum state without classical space and time. A wave function is used for a probabilistic description of the creation of the Universe and such a process is connected with transition of a wave through an effective barrier. Determination of penetrability of this barrier is a key point in estimation of duration of the formation of the Universe, and dynamics of its expansion in the first stage.  However, in majority of these models (with the exception of some exactly solvable models) tunneling is practically studied in details in the semiclassical approximation mainly (for example, see Refs.~\\cite{Vilenkin.1994.PRD,Rubakov.1999}). An attractive side of such an approach is its simplicity in the construction of decreasing and increasing partial solutions for the wave function in the tunneling region, an outgoing wave function in the external region, and a possibility to define and to estimate in a simply enough way a penetrability of the barrier, which can be used for obtaining the duration of the nucleation of the Universe. A \\emph{tunneling boundary condition} \\cite{Vilenkin.1995,Vilenkin.1994.PRD} could seems to be the most natural and clear, where the wave function should represent an outgoing wave only in the enough large value of the scale factor $a$. However, whether is such a wave really free in the asymptotic region? If to draw attention on increasing of a modulus of the potential with increasing of the scale factor $a$ and increasing of a gradient of such a potential, used with opposite sign and having a sense of force, acting ``through the barrier'' on this wave, then one come to a serious contradiction: \\emph{influence of the potential on this wave increases strongly at increasing of $a$}! Now a new question has been appeared: what should the wave represent in a general in the cosmological problem? This problem connects with another (new) else more general one in general quantum physics --- a real importance \\emph{to define a ``free'' wave inside \\underline{strong} fields} and we need in mathematical stable tools for its study and work. It becomes unclear whether a connection between exact solutions for the wave function at turning point and ``free'' wave defined in the asymptotic region is corrected. If in frameworks of semiclassical approach (up to the second order) it is corrected then any deformations of tail of the potential after the barrier cannot change the penetrability absolutely, while force mentioned above and acting on the asymptotic ``free'' wave could be increased up to infinity! Answer on such misunderstanding could be found in non-locality of definition of the penetrability in quantum mechanics, which is reduced to minimum in the semiclassical approach (i.~e. this is so called ``error'' of the cosmological semiclassical approach). Now, if to refuse application of the semiclassical approximation, then we come to a necessity to define such a wave as maximal correctly and accurately as a possible. The problem of the correct definition of the wave in cosmology is reinforced else more, if to calculate incident and reflected waves in the internal region. \\emph{Even with the known exact solution for the wave function there is uncertainty in determination of these waves!} But, namely, the standard definition of the coefficients of penetrability and reflection is based on them. In particular, I have not found papers yet where the coefficient of reflection is defined and estimated in this problem (which differs essentially from unity at the energy of radiation close to height of the barrier and, therefore, such a characteristics could be interesting from a physical point of view). Note that the semiclassical approximation prejudices a possibility of its definition at all \\cite{Landau.v3.1989}.  Thus, in order to estimate probability of the formation of the Universe as accurately as possible, we need in the fully quantum definition of the wave. Note that the non-semiclassical penetrability of the barrier in the cosmological problems has not been studied in details and, therefore, a development of fully quantum methods for its estimation is a perspective task. I would like to note undoubtedly perspective direction of researches~\\cite{AcacioDeBarros.2007.PRD}, where the penetrability was estimated on the basis of tunneling of wave packet through the barrier. However, a stationary boundary condition has uncertainty that could lead to different results in calculations of the penetrability. The stationary approach could allow to clarify this deeply. It is able to give stale solutions for the wave function (and results in Ref.~\\cite{Maydanyuk.2008.EPJC} have confirmed this at zero energy of radiation), uses the standard definition of the coefficients of the penetrability and reflection, is more accurate in their estimation.  Aims of this paper are: (1) to define the wave in the quantum cosmological problem; (2) to construct the fully quantum (non-semiclassical) stationary method of determination of the coefficients of penetrability of the barriers and reflection from them on the basis of such a definition of the wave; (3) to estimate how much the semiclassical approach is differed in the estimation of the penetrability from the fully quantum one. In order to achieve this, at first it needs to construct tools for stable calculation of partial solutions of the wave function. In order to resolve the questions pointed out above, we shall restrict ourselves by a simple cosmological model, where the potential has a barrier and internal above-barrier region. The paper is organized so. After a short description of closed Friedmann--Robertson--Walker (FRW) model with quantization in a presence of a positive cosmological constant and radiation in Sec.~2, in Sec.~3 a new formalism for calculation of two linear independent partial solutions for the wave function of the Universe for the scale factor inside the region $0 \\le a \\le 100$ and the energy of radiation from zero up to the barrier height is presented. In Sec.~4 a fully quantum definition of the wave is formulated (for the first time), a new approach for determination of the incident, reflected and transmitted waves relatively the barrier is constructed on such basis, the boundary condition is corrected. In Sec.~5 a fully quantum stationary method of determination of coefficients of penetrability and reflection relatively the barrier with analysis of uniqueness of solution is presented, where at first time non-zero interference between the incident and reflected waves is analyzed and for its estimation the coefficient of mixing is introduced. In such an approach the penetrability and reflection for the barrier for the studied cosmological model with $A=36$, $B=12\\,\\Lambda$ parameters at $\\Lambda=0.01$ are calculated with estimation of accuracy and in comparison with the semiclassical results and with results of non-stationary quantum approach \\cite{AcacioDeBarros.2007.PRD}. In Sec.~6 the penetrability is estimated for the FRW-model with included Chaplygin gas which has internal hole before the barrier. Conclusions finalize the paper.  % *******************************************************************************************************************     % ******************************************************************************************************************* ",
        "conclusions": "}  In the paper the closed Friedmann--Robertson--Walker model with quantization in the presence of a positive cosmological constant and radiation is studied. I have solved it numerically and have determined the tunneling probability for the birth of an asymptotically deSitter, inflationary Universe as a function of the radiation energy. Note the following. % \\begin{enumerate} \\item A formalism for calculation of two linear independent partial solutions for the wave function of the Universe for the scale factor inside the region $0 \\le a \\le 100$ and the energy of radiation from zero up to the barrier height has been constructed.  \\item A fully quantum definition of the wave which propagates inside strong field and interact minimally with them, has been formulated for the first time, and approach for its stable determination has been constructed.  \\item A new stationary approach for determination of the incident, reflected and transmitted waves relatively a barrier has been constructed, the tunneling boundary condition has been corrected.  \\item A quantum stationary method of determination of coefficients of penetrability and reflection relatively the barrier with analysis of uniqueness of solution has been developed, where for the first time non-zero interference between the incident and reflected waves has been taken into account and for its estimation the coefficient of mixing has been introduced.  \\item A criterion of estimation of accuracy of the determination of these coefficients has been proposed on the basis of check of eq.~(\\ref{eq.4.2.3}). \\end{enumerate} % In such a quantum approach the penetrability of the barrier for the studied quantum cosmological model with $A=36$, $B=12\\,\\Lambda$ parameters at $\\Lambda=0.01$ has been estimated with a comparison with results of other known methods. Note the following. % \\begin{itemize} \\item According to the calculations, inside whole region of energy of radiation the tunneling probability for the birth of an asymptotically deSitter, inflationary Universe is very close to its value, obtained in the semiclassical approach by eqs.~(\\ref{eq.6.1}) and (\\ref{eq.6.2}), but essentially differs on the results obtained before by the quantum non-stationary approach in Ref.~\\cite{AcacioDeBarros.2007.PRD} (see Tabl.~1 and 2 in Appendix).  \\item The coefficient of reflection from the barrier in the internal region has been determined at first time. According to calculations, it is differed essentially on 1 at the energy of radiation close enough to the barrier height (see Tabl.~2 in Appendix).  \\item The modulus of the coefficient of mixing is less $10^{-19}$ for all energies, that points out that \\emph{there is no \\underline{interference} between the found incident and reflected waves close to the internal turning point.}  \\item On the basis of the calculated coefficients I reconstruct a property (\\ref{eq.4.2.3}) with accuracy of the first 11--18 digits (see Tabl.~2 in Appendix) inside the whole studied region of the energy of radiation. \\end{itemize} % One can assume that the method proposed can be easily generalized on the other cosmological models with the barriers with arbitrary complicated shapes. However, one can suppose that a visible change can be found in difference between the penetrabilities in the fully quantum and semiclassical approaches after inclusion of new components of density into the model (for example, after use of the component of Chaplygin gas in Refs.~\\cite{Chaplygin.1904,Kamenshchik.2001.PLB,Bilic.2002.PLB,Bento.2002.PRD}). Of course, it can be interesting to compare the estimations of the coefficients of penetrability and reflection with results which could be obtained on the basis of the \\emph{Improved WKB approach} on the basis of Refs.~\\cite{Casadio.2005.PRD.D71,Casadio.2005.PRD.D72,Luzzi.2006.PhD}.  % *******************************************************************************************************************     % ******************************************************************************************************************* \\appendix"
    }
}